{
    "0609/astro-ph0609025_arXiv.txt": {
        "abstract": "{We quantitatively investigate how collisional avalanches may develop in debris discs as the result of the initial breakup of a planetesimal or comet-like object, triggering a collisional chain reaction due to outward escaping small dust grains. We use a specifically developed numerical code that follows both the spatial distribution of the dust grains and the evolution of their size-frequency distribution due to collisions. We investigate how strongly avalanche propagation depends on different parameters  (e.g., amount of dust released in the initial breakup, collisional properties of dust grains, and their  distribution in the disc). Our simulations show that avalanches evolve on timescales of $\\sim\\!1000$\\,years, propagating outwards following a spiral-like pattern, and that their amplitude exponentially depends on the number density of dust grains in the system. { We estimate a probability for witnessing an avalanche event as a function of disc densities, for a gas-free case around an A-type star, and find that features created by avalanche propagation can lead to observable asymmetries for  dusty systems with a $\\beta$\\,Pictoris-like dust content or higher.  Characteristic observable features include: (i) a brightness asymmetry of the two sides for a disc viewed edge-on, and (ii) a one-armed open spiral or a lumpy structure  in the case of  face-on orientation. A possible system in which avalanche-induced structures might have been observed is the edge-on seen debris disc around HD\\,32297, which displays a strong luminosity difference between its two sides. } } ",
        "introduction": "Direct imaging of circumstellar discs \\citep[e.g.,][]{Heap00, ClampinKrist03, Liu04, Schneider05} have provided resolved disc morphologies for several systems (e.g.,  \\bpic, HD\\,141569A, HD\\,100546, HD\\,32297) and have shown that dust distribution is not always smooth and axisymmetric. Warps,  spirals, and other types of asymmetries are commonly observed \\citep[e.g.,][ for the \\bpic~system]{Kalas95}. These morphological features can provide hints on important ongoing processes in the discs and improve our understanding of the evolution of circumstellar discs and of planetary formation.   The usual explanation proposed for most of these asymmetries is the perturbing influence of an embedded planet. As an example, the warp in the \\bpic~disc has been  interpreted as induced by a jovian planet on an inclined orbit \\citep{Mouillet97, Au01}. Likewise, for annulus-like discs with sharp inner or outer edges, the most commonly proposed explanation is  truncation or gap opening due to planets or bound stellar companions \\citep[e.g.,][]{Au04}, although alternative mechanisms such as gas drag on dust grains within a gas disc of limited extent have also been proposed \\citep{TA01}. For spiral structures, authors have also been speculating on gravitational instabilities \\citep{Fuka04}, as well as on a bound stellar companion \\citep{Au04}.   The catastrophic breakup of one single large object releasing a substantial amount of dust fragments could be an alternative explanation for some observed asymmetries.  \\cite{WyattDent02} have examined how such collisionally produced bright dust clumps could be observed in Fomalhaut's debris disc. Likewise, such clumps have been proposed by \\citet{Telesco05} as a possible explanation for mid-infrared brightness asymmetries in the central \\bp~disc, but only based on preliminary order of magnitude estimates. More recently, the detailed study of \\citet{KenyonBromley05} investigated the possibility of detecting catastrophic two-body collisions in debris discs and found that such a detection would require the breakup of 100-1000\\,km objects. The common point between these different studies is that they focus on global luminosity changes due to the debris cloud directly produced by the shattering events themselves.  In the present paper, we re-examine the consequences of isolated shattering impacts from a different perspective, i.e., by considering the collisional evolution of the produced dust cloud $after$ its release by the shattering event. The main goal here is to study one possibly very efficient process, first proposed by \\citet{Art97}, but never quantitatively studied so far, i.e., the so-called collisional avalanche mechanism. The basic principle of this process is simple. After a localized disruptive event, such as the collisional breakup of a large cometary or planetesimal-like object, a fraction of the dust then produced is driven out by radiation pressure on highly eccentric or even unbound orbits. These grains moving away from the star with significant radial velocities can breakup or microcrater other particles farther out in the disc, creating in turn even more small particles propagating outwards and colliding with other grains. Should this collisional chain reaction be efficient enough, then a significant increase in the number of dust grains could be achieved. In this case, the consequences of a single shattering event, in terms of induced dust production, could strongly exceed that of the sole initially released dust population. The outward propagation of the dusty grains could then induce observable asymmetric features in the disc, even if the initially released dust cloud is undetectable.  The goal of this work is to perform the  first quantitative study of the avalanche process and investigate the morphology of avalanches in debris discs, under the assumption that dust dynamics is not controlled by gas \\citep{Lagrange00}. For this purpose we have created a numerical code, described in Sect.~\\ref{sec:model}, that enables us to simulate the coupled evolution of dynamics  and size-frequency distribution of dusty grains. The results of our simulations, which explore the effect of several parameters (total mass and radial distribution of  dust in the disc, mass and size distribution of the planetesimal debris, physical properties of the grains and the prescription   for collisional outcome for grain-grain collisions) are presented in Sect.~\\ref{sec:res}. In Sect.~\\ref{sec:obs} we examine under which  conditions avalanche-induced features might become observable. We end with a discussion of the probability of witnessing an avalanche (Sect.~\\ref{sec:discussion}) and finally a summary (Sect.~\\ref{sec:sum}). ",
        "conclusions": "\\label{sec:discussion}  \\subsection{Probability of witnessing an avalanche event} \\label{sec:Pobs}   The  numerical investigation of the previous sections has shown that collisional avalanches are a powerful and efficient mechanism that naturally develops in debris discs after the breakup of a large planetesimal. However, in our nominal case of a $\\beta$\\,Pic-like system and $M_0=10^{20}$\\,g of dust initially released, the asymmetric features produced by the avalanche probably remain too weak to be observable in scattered light (Sect.~\\ref{sec:obs_nom}). This result should, however, be taken with great care since our parameter exploration has shown that avalanche strength strongly depends on several critical and often poorly constrained parameters. The first set of parameters is those linked to the initial breakup event. Here we obtain the intuitive result that higher amounts of initially released dust leads to more powerful avalanches (see Sect.~\\ref{sec:mo}), with the avalanche strength scaling linearly with $M_0$. This is not unique to the avalanche mechanism: \\citet{KenyonBromley05} find a similar dependence when only considering the signature of the cloud of primary debris produced immediately after the planetesimal breakup. What distinguishes our results from studies in which only dust released at impact is considered is that avalanches strongly depend on the number density of  dust in the disc. Section~\\ref{sec:fp} has indeed shown that the global optical depth of the dust disc $\\tau_{\\|}$ is the parameter avalanche development depends most on, the dependence being close to an exponential. We have seen that other parameters, mostly related to the way the physical response of grains to collisions is modeled, might also lead to observable events when stretched to the extreme values that were numerically explored here. This is in particular the case for $Q_*$, for which very low $\\simeq 10^{6}$\\,erg/g values might lead to powerful avalanches.  We shall however leave these ``technical'' parameters aside to focus on the 2 parameters directly related to the system's properties themselves, i.e., the optical depth, both $\\tau_{\\bot}$ and $\\tau_{\\|}$, and the initial amount of dust released $M_0$, and  derive   an order-of-magnitude estimate for the  probability of witnessing avalanche events as a function of these parameters. From the results of Sect.~\\ref{sec:res}, the $L_\\mathrm{av}/L_\\mathrm{d}$ criterion for observability might be written  \\begin{equation} \\left( \\frac{\\tau_{\\bot}}{\\tau_\\mathrm{\\bot,nom}} \\right)^{-1} \\frac{F_\\mathrm{max}}{F_\\mathrm{max(nom)}}\\, \\frac{M_0}{10^{20}\\,\\mathrm{g}} \\ga 100, \\label{visi0} \\end{equation} which is equivalent to saying that the luminosity ratio between avalanche and field grains should be at least 100 times higher than in the nominal case (for which $L_\\mathrm{av}/L_\\mathrm{d} \\sim 10^{-2}$). Section~\\ref{sec:mc} has shown that $F_\\mathrm{max}$ is independent of $M_0$, so that in our approximation $F_\\mathrm{max}$ is only a function of  $\\tau_{\\|}$, and this  $\\tau_{\\|}$ dependence is given by Eq.~\\ref{eq:Ftau}. Thus, Eq.~\\ref{visi0} reduces to  \\begin{equation} \\exp \\left[ 5.3 \\left( \\frac{\\tau_{\\|}}{\\tau_\\mathrm{\\|,nom}} \\right)^{0.55}\\right] \\, \\frac{M_0}{10^{20}\\,\\mathrm{g}} ~ \\frac{\\tau_{\\bot,nom}}{\\tau_\\mathrm{\\bot}}\\,  \\ga  2\\times10^{4}, \\label{visi1} \\end{equation} which gives a direct link between a given disc density ($\\tau_{\\bot}$ and $\\tau_{\\|}$) and the { minimum mass of released dust} able to produce a visible avalanche in such a disc (the denser the disc, the smaller the corresponding $M_0$ value). The other important issue affecting witnessing probabilities is  the duration of an avalanche. Our simulations show that the typical lifetime  of an avalanche-induced pattern is  $t_\\mathrm{av} \\sim 10^3$\\,yrs. With this value and Eq.~\\ref{visi1}, one can estimate the probability $P_\\mathrm{obs}$ of witnessing an observable avalanche event in a given disc:  \\begin{equation} P_\\mathrm{obs} = \\frac{t_\\mathrm{av}}{t_\\mathrm{shatt(M_0,\\tau_{\\bot})}}, \\label{probv0} \\end{equation} where $t_\\mathrm{shatt(M_0,\\tau)}$ is the average time between 2 shatterings producing $M_0$ of dust in a disc of average optical depth $\\tau_\\bot$, with $M_0$,  $\\tau_{\\|}$ and  $\\tau_\\bot$ satisfying  Eq.~\\ref{visi1}. { As suggested in Sect.~\\ref{sec:plb}, we consider that the object releasing $M_0$ of dust has a mass $M_\\mathrm{PB}\\simeq 10M_0$.} { From unpublished results of the \\citet{Th03} simulations of collisional rates and outcomes in the inner \\bpic~disc, we determine that the approximate timescale for the shattering of a $M_\\mathrm{PB}=10M_0$ object to occur in the innermost $<50$\\,AU (the typical location for the initial shattering events considered in our simulations) of a \\bpic~like system is $t_\\mathrm{shatt}\\simeq 150[(10M_0)/10^{21}$\\,g$]^{1.25}$\\,yrs. Since, for systems with similar spatial distributions, the frequency of collisional events is proportional to the square of a system's total mass, we get the empirical relation:  \\begin{equation} t_\\mathrm{shatt(M_0,\\tau_\\bot)} \\simeq \\,150\\, \\left( \\frac{\\tau_\\bot}{\\tau_\\mathrm{\\bot,nom}} \\right)^{-2} \\left( \\frac{M_0}{10^{20}\\, \\mbox{g}} \\right)^{1.25} \\,\\mbox{yrs,} \\label{eq:tshat} \\end{equation} where we implicitly assume that the system's spatial distribution is the same as in the nominal case, so that the ratio between two systems' total masses is equal to the ratio $\\tau_\\bot / \\tau_\\mathrm{\\bot,nom}$ anywhere in the disc. This equation should of course be regarded as giving a $very$ rough estimate, since $t_\\mathrm{shatt(M_0,\\tau_\\bot)}$ depends on many poorly constrained parameters, like the number density of planetesimals and their average kinetic energy at impact. Equation~\\ref{eq:tshat}, however, gives the global trend of the way $t_\\mathrm{shatt(M_0,\\tau_\\bot)}$ increases with $M_0$. Taking the lowest $M_0$ value satisfying Eq.~\\ref{visi1} and plugging it into Eq.~\\ref{eq:tshat}, we get, from Eq.~\\ref{probv0}:  \\begin{equation} P_\\mathrm{obs}  \\approx 3 \\times 10^{-5} \\left( \\frac{\\tau_\\bot}{\\tau_\\mathrm{\\bot,nom}} \\right)^{0.75} \\exp \\left(6.6 \\left( \\frac{\\tau_\\|}{\\tau_\\mathrm{\\|,nom}} \\right)^{0.55}\\right). \\label{probv2} \\end{equation} Equation \\ref{probv2} indicates that $P_\\mathrm{obs}\\simeq 0.03$ for the nominal case field particle disc, which means that we have about a $3\\%$ chance of witnessing the avalanche caused by the breakup of a $M_\\mathrm{PB}=10M_0 \\sim\\!\\! 10^{23}$\\,g object ($M_0=10^{22}$\\,g being the smallest released dust mass able to trigger a visible avalanche for such a disc, as given by Eq.~\\ref{visi1}). } This makes it a rather unlikely event, although it cannot be completely ruled out. Nevertheless, slightly denser discs (i.e., higher $\\tau_{\\bot}$ and $\\tau_{\\|}$) can easily raise  $P_\\mathrm{obs}$ up to 1. As a matter of fact, the dependence on $\\tau_{\\|}$ is so sharp that $P_\\mathrm{obs}=1$ is obtained for $\\tau_{\\|}\\simeq 2.1\\tau_{\\|,nom} \\simeq  0.046$. We thus see that a $\\beta$\\,Pic-like system is below, but not too far from the limit for which chances of witnessing an avalanche are high, especially when considering the uncertainties regarding avalanche strength due to its dependence on several poorly constrained parameters related to the collision-outcome prescription (also keeping in mind that higher $\\tau_{\\|}$ values could alternatively be achieved for a thinner disc of the same dust mass (see Sect.~\\ref{sec:dd})). { Moreover, our $L_\\mathrm{av}/L_\\mathrm{d}>1$ criterion for observability is probably too conservative, and avalanche-induced patterns might be detectable for lower luminosity excess values. Taking, for example, $L_\\mathrm{av}/L_\\mathrm{d}>0.1$ would raise the detection probability to $\\simeq 45\\%$ for a $\\beta$\\,Pic-like system}.     {  \\subsection{Avalanches in observed systems, perspectives}  We defer a detailed application of our model to specific circumstellar discs to a future study. However, the present results can already give a good idea of the typical profile for a ``good'' avalanche-system candidate  Our numerical exploration has shown that structures that are the most likely to be associated with avalanche-events have two-sided asymmetry for discs viewed edge-on and open spiral patterns for discs viewed pole-on or at intermediate inclinations. An additional requirement is that these discs should be dust-rich systems, with a dustiness at least equal to, and preferably higher than that of $\\beta$-Pic. Note also that our model makes an additional prediction, i.e., that avalanche affected regions should consist of grains significantly smaller than the \"field\" particles in the rest of the disc. If the blow-out radius of grains is of the order of the wavelength of the observed light, then this should translate into color differences between avalanche (bluer) and non-avalanche (redder) regions.  In this respect, one good edge-on candidate might be the recently discovered HD\\,32297 system, which exhibits a strong two-sided asymmetry. As reported by \\citet{Schneider05} and \\citet{Kalas05}, this system is a $\\beta$\\,Pic analog with its SW side significantly brighter than the NE one within $\\simeq$ 100\\,AU \\citep{Schneider05} and possibly outside 500\\,AU \\citep{Kalas05}. Such a two-sided asymmetry would be compatible with the ones obtained in our simulations (as shown for example in the bottom panel of Fig.~\\ref{fig:ass_const}). Furthermore, \\citet{Kalas05} also reported a color asymmetry between the two sides, with the brighter one (SW) being significantly bluer. This seems to indicate that this side is made of smaller, possibly submicron grains \\citep{Kalas05}. As previously discussed, this is what should be expected for an avalanche-affected region. However, an alternative scenario, like the collision with a clump of interstellar medium proposed by \\citep{Kalas05}, might also explain the HD\\,32287 disc structure. Future imaging and spectroscopic observations are probably needed before reaching any definitive conclusions.  Among all head-on observed systems, the one displaying the most avalanche-like structure is without doubt HD\\,141569 \\citep{ClampinKrist03}, with its pronounced spiral pattern. Furthermore, the disc's mass, significantly higher than $\\beta$-Pictoris, makes it a perfect candidate in terms of witnessing probabilities. Of course, avalanche is not the only possible scenario here, and several alternative explanations, like an eccentric bound planet or stellar companion, or a stellar flyby have already been proposed \\citep[e.g.,][]{Au04, Wyatt05, Ardila05}. One should, however, be aware that this system is strictly speaking not a \"standard\" debris disc as defined by \\citet{Lagrange00} and as considered in the present study. Indeed, several studies seem to suggest the presence of large amounts of primordial gas \\citep{Zuckerman95,Ardila05}.  Gas drag effects have been left out of the present study on purpose, mainly because, in the strict sense of the term, debris discs are systems where dust dynamics is not dominated by gas friction \\citep{Lagrange00}. Moreover, the correct description of dust-gas coupling adds several additional free parameters (gas density and temperature distributions, etc.) and requires a full 2-D or 3-D treatment of gas by far exceeding the scope of the present paper. However, the issue of avalanches in a gaseous medium might be a crucial one for those systems that are most favorable for avalanches, i.e., discs more dusty than $\\beta$-Pic, a system which is already at the upper end of debris-discs in terms of dustiness \\citep[e.g.,][]{Spangler01}. Such more massive systems should fall into a loosely defined category of \"transition\" discs between T-Tauri or Herbig Ae protoplanetary systems and \"proper\" debris discs \\citep[see for example Sect.~4 of][]{Dutrey04}. For such systems (of which HD141569 is a typical example), which are younger than the more evolved debris discs, risks (or chances) of encountering large amounts of remaining gas are high. This crucial issue will be the subject of a forthcoming paper.}"
    },
    "0609/astro-ph0609814_arXiv.txt": {
        "abstract": "The first three years of observation of the {\\it Wilkinson Microwave Anisotropy Probe} (WMAP) have provided the most precise data on the anisotropies of the cosmic microwave background (CMB) to date. We investigate the impact of these results and their  combination with data from other astrophysical probes on cosmological models with a dynamical dark energy component. By considering a wide range of such models, we find that the constraints on dynamical dark energy are significantly improved compared to the first year data. ",
        "introduction": "Observations of type Ia supernovae (SNe Ia) \\cite{Astier:2005qq,Riess:2004nr}, structure formation (LSS) \\cite{Percival:2001hw,Tegmark:2003ud} and the cosmic microwave background (CMB) \\cite{Spergel:2006hy,Readhead:2004gy,Goldstein:2002gf} all agree on an accelerated expansion of our Universe. This rather unexpected phenomenon can be explained by modifying 4-D gravity \\cite{Dvali:2000hr,Bekenstein:2004ne} or adding a new component to the total energy momentum tensor. The simplest such component is a cosmological constant. It fits all current observations flawlessly and has a simple interpretation in terms of a vacuum energy. Yet, its observed value is 120 orders of magnitude off from the naive estimate $\\Lambda \\sim \\mplank^4$, where $\\mplank$ is the reduced Plank mass. The coincidence between this minute dark energy contribution and the observed energy density of matter is rather puzzling. If not given by chance via some sort of anthropic principle, it necessitates a mechanism that explains this coincidence. An immediate possibility is a coupling \\cite{Wetterich:1994bg,Amendola:1999er,Amendola:2006qi} between (dark) matter and dark energy (though there might be problems due to quantum effects \\cite{Doran:2002bc}).  Another solution is an attractor behavior \\cite{Wetterich:fm,Ratra:1987rm,Caldwell:1997ii} of dark energy that leads to an almost constant ratio between the fractional energy density $\\od(z)$ of dark energy and the species otherwise dominating the expansion, i.e. photons and neutrinos during radiation domination and matter during matter domination. Coincidentally, such an attractor behavior corresponds to a scalar field with exponential potential that arises in string theories and when solving the  cosmological constant problem from the point of view of dilatation symmetry \\cite{Wetterich:fm}. The non-vanishing $\\od(z)$ at higher redshifts alleviates the problem of  explaining the coincidence of matter and dark energy today $\\om^0 \\approx \\od^0$. Instead of fine tuning $\\Lambda$ to many orders of magnitude, the tuning needed is of the order of $10^{-3}$. However, the tuning needed for such \\emph{early dark energy} cosmologies increases the less dark energy there is at earlier times. A detection of early dark energy, on the other hand, would give crucial hints to fundamental laws of nature. The aim of this study therefore is to investigate the implications of the three year data of WMAP on dynamical dark energy models in general and their respective fractions of early dark energy.  In view of the theoretical uncertainties many different techniques have been employed in the analysis of the dark energy, ranging from atttempts to reconstruct the potential of a scalar field dark energy (e.g. Ref. \\cite{Sahlen:2006dn}) to the principal component approach of Ref. \\cite{Huterer:2002hy}.  We consider the redshift dependence of the fractional dark energy $\\od(z)$ as a free function to be ``measured'' by observation. We investigate in this note various parameterizations and an interpolated model. The possible coupling between dark energy and dark matter is neglected in this study. ",
        "conclusions": "The three year data of WMAP was used to estimate cosmological parameters for a wide range of dynamical dark energy models. We have shown that this new data in combination with large scale structure data  constrain the average amount of dark energy during the time of structure formation to $\\osf \\lesssim 4\\%$ for a fair sample of  dark energy models from the literature. We have also constructed a parameterization of $\\od(a)$ which linearly interpolates between the dark energy fraction at several redshift bins. This allows for a considerably higher fraction of $\\osf$. The analysis also shows that the values of the standard cosmological parameters for the dynamical dark energy models are well compatible with the values found by WMAP for \\lcdm. The effect of dynamical dark energy on $\\sigma_8$ and on the formation of nonlinear structure offer promising routes for further constraints on the time evolution of dark energy or a possible falsification of the \\lcdm\\ model."
    },
    "0609/astro-ph0609739_arXiv.txt": {
        "abstract": "Traditional thermal evolution models of giant planets employ arbitrary initial conditions selected more for computational expediency than physical accuracy.  Since the initial conditions are eventually forgotten by the evolving planet, this approach is valid for mature planets, if not young ones.  To explore the evolution at young ages of jovian mass planets we have employed model planets created by one implementation of the core accretion mechanism as initial conditions for evolutionary calculations.  The luminosities and early cooling rates of young planets are highly sensitive to their internal entropies, which depend on the formation mechanism and are highly model dependent.  As a result of the accretion shock through which most of the planetary mass is processed, we find lower initial internal entropies than commonly assumed in published evolution tracks. Consequently  young jovian planets are smaller, cooler, and several to 100 times less luminous than predicted by earlier models.  Furthermore the time interval during which the young jupiters are fainter than expected depends on the mass of planet.  Jupiter mass planets ($1\\,\\rm M_J$) align with the conventional model luminosity in as little at 20 million years, but $10\\,\\rm M_J$ planets can take up to 1 billion years to match commonly cited luminosities, given our implementation of  the core accretion mechanism.  If our assumptions, especially including our treatment of the accretion shock, are correct and if extrasolar jovian planets indeed form with low entropy, then young jovian planets are substantially fainter at young ages than currently believed.  Furthermore early evolution tracks should be regarded as uncertain for much longer than the commonly quoted $10^6$ years.  These results have important consequences both for detection strategies and for assigning masses to young jovian planets based on observed luminosities. ",
        "introduction": "In the past decade, a number of nearby star associations have been recognized as being quite young, less than $10\\,\\rm Myr$ old (e.g. IC 348, TW Hydrae, MBM12, $\\eta$ Cha \\citep{Lada95, Webb99, Luhman04}).  Such associations are likely well stocked with recently formed, presumably bright giant planets that should in principle be easy prey for a variety of planet detection technologies.  Planning for the hunt, however, requires knowledge of the expected luminosity of young giant planets as a function of time since their formation, particularly at young ages when they are presumably easy game.  While models of   the luminosity evolution of giant planets have a long pedigree \\citep[e.g.,][]{Grossman72, Graboske75}, the early work focused on the evolution of the solar system giants, attempting to explain their current luminosity at an age of 4.5 Gyr. Since planets lose memory of their initial conditions over time, initial conditions were selected more for computational convenience than for accuracy.  Many improvements have subsequently been made to the models, particularly in the characterization of jovian atmospheres at various effective temperatures, although essentially the same initial conditions are still employed \\citep{Hubbard80a, Burrows97, Chabrier00b}.  The standard evolution model begins with a hydrogen-helium sphere having a large radius, high internal entropy, and large effective temperature.  Such an object is not necessarily one that would be the result of any particular planet formation model.  This model planet is allowed to radiate and cool over time.  Since there have been no detections yet of young planets with measured masses, the applicability of this initial condition is untested, although there are data from more massive objects.  A pair of eclipsing, young ($\\sim 1\\,\\rm Myr$) brown dwarfs with known dynamical masses (57 and 36 Jupiter masses ($\\rm M_J$)) indeed  have radii exceeding five times that of Jupiter \\citep{Stassun06}, confirming that young massive brown dwarfs, at least, are in fact large and hot.  But giant planets that formed in a disk around a primary star, or even isolated planet-mass objects, may have experienced very different initial conditions.  Evolution tracks computed in the usual way--even though they do not necessarily reflect any particular planet formation theory--have been used to evaluate detection strategies for true giant planets orbiting solar type stars \\citep[e.g.,][]{Burrows05}, and they have been used to characterize isolated, very low mass brown dwarfs. \\citet{Chauvin04}, for example, reported on the detection of a faint companion to the M8 brown dwarf \\object{2MASSWJ1207334-393254} (hereafter 2MASS1207) in TW Hydrae with an estimated age of 8 Myr.  Using its observed luminosity and applying published evolution tracks, they estimated a mass of just $5\\,\\rm M_J$ for the mass of the companion.  The early modelers certainly did not foresee that direct detections of putative young planets would be compared against the models at exceptionally young ages, at times when the model planet may not yet have forgotten its hot start. \\citet{stevenson82} wrote that evolution calculations ``...cannot be expected to provide accurate information on the first $10^5$--$10^8$ years of evolution because of the artificiality of an initially adiabatic, homologously contracting{\\footnote{See \\citet{Stahler88} for a discussion of homologous contraction.}} state.''  More recently, \\citet{Baraffe02} examined the uncertainties in evolution tracks of brown dwarfs at young ages and cautioned about the applicability of evolution models at ages less than a few million years, on the lower end of Stevenson's uncertainty range. \\citet{Wuchterl05} has also  expressed concern that standard evolution models do not capture the early evolution correctly.  Given the clear imperatives to interpret observations of young, low mass objects and to plan for future direct detections of giant planets formed in orbit about solar type stars, there is a need to connect models of giant planet formation to giant planet evolution.  Our goals here are both to help fill the void in physically plausible models of extrasolar giant planets (EGPs) at young ages and to better quantify the age beyond which the evolution models are robust and applicable.  We aim to understand whether or not the current generation of evolution models can reliably predict the luminosity of giant planets at young ages and, if not, then define the age beyond which current models are reliable. Instead of using an arbitrary starting condition, we employ planets formed by one implementation of the core accretion model.  In this scenario gas giant planets form by rapid accretion of gas onto a solid core that grew by accretion of planetesimals in the nebula.  This mechanism is one of two competing scenarios of gas giant formation, the other being the gas instability model by which giants form from a local disk instability \\citep{Boss98} that creates a self-gravitating clump of gas.  The planet resulting from such a clump could also be used as the starting point of an evolutionary calculation (see \\citet{Bodenheimer74, Bodenheimer76, Bodenheimer80}), but we choose here to focus solely on the core accretion mechanism as it currently seems the more promising mechanism for explaining the formation of the giant planets (see \\citet{Lissauer06} for a review).  For specificity, we rely upon the implementation by \\citet{Hubickyj05} of the core accretion mechanism.   By necessity, their work  makes a host of assumptions that ultimately affect the properties of newly born giant planets.  As we will demonstrate, following the end of accretion this model predicts that giant planets are substantially fainter than standard evolution models.  While we believe that this conclusion is secure, we stress that the precise numerical value of the post-accretion luminosity depends upon the particular assumptions employed by \\citet{Hubickyj05}.  Therefore, we first briefly review this model and highlight the assumptions upon which the work rests in Section 2.  We describe our method of evolving these model planets over time in Section 3 and compare our results with standard giant planet evolution models. We find in Section 4 that the initial conditions influence subsequent planetary evolution for longer than generally appreciated and that planets formed by the  \\citet{Hubickyj05} recipe for core accretion are substantially fainter than standard models have previously predicted.  We end by cautioning those who wish to rely upon evolution models to characterize detected young giant planets that may have grown by the core accretion mechanism would be wise to be judicious in their estimation of the model-dependent uncertainties. ",
        "conclusions": "We have computed the first giant planet evolution models that couple planetary thermal evolution to the predicted core mass and thermal structure of a core accretion planet formation model.   \\citet{Baraffe06} investigated the evolution of planets with core sizes and heavy element abundances derived from the core accretion models of \\citet{Alibert05}.  However, \\citet{Baraffe06} did not attempt to match the thermal structure (and hence, temperature, entropy, and density) at the interface between planetary formation and subsequent evolution.  Our implementation of the core accretion model processes most of the planetary mass through an accretion shock in which the accreting gas loses most of its internal entropy.  As a result, our young giant planets are cooler, smaller, fainter, and take longer to evolve than the standard hot-start model giant planets.  We note, however, that our accretion model  does not resolve the radiative transfer within the shock, but rather uses the  shock boundary conditions of \\citet{stahler80}.  A more complete or detailed treatment of accretion at the surface of the planet could very well result in different initial conditions, including possibly a warmer, larger, brighter and more conventional young planets. Specifically, a three-dimensional hydrodynamic simulation of gas accretion by giant planets, allowing for material accreted through a circumplanetary disk and shock radiation, would provide more rigorous post-formation models for subsequent evolution calculations.  Until such models are available, our approach--which likely provides a lower limit to the post accretion luminosity--demonstrates that plausible initial conditions can lead to very different early evolution tracks for giant planets than the `hot start' models that are commonly relied upon.  For example, at $10^7$ years--a time greater than the age of the TW Hydrae association--our $10\\,\\rm M_J$ core accreted planet is more than a factor of 100, or 5 magnitudes, fainter than the equivalent hot start planet.   The luminosity difference falls with decreasing mass, so that our model luminosity for a $1\\,\\rm M_J$ planet is comparable to the standard case.  Thus the thermal luminosity of young, massive giant planets, which have been assumed to be easy targets for coronagraphy, may be substantially less than previously assumed. If this result is correct, then searches for the thermal emission from young, several Jupiter mass planets must be far more sensitive than previously anticipated in order to detect these relatively faint, young planets.  Thermal infrared planet searches by the Large Binocular Telescope, the James Webb Space Telescope, and other planned telescopes would all be impacted, although efforts to detect planets in reflected light would not.  This conclusion holds true even to ages as great as that of the Pleiades for the most massive planets considered here. Ironically the least massive, most intrinsically faint planets (1 to $2\\,\\rm M_J$) match their hot-start luminosity tracks by just a few tens of millions of years, or less and are much less underluminous before that time.   Direct detections of young giant planets with dynamically measured masses will test this conclusion.  Since the numerical values for luminosity that we derive depend primarily upon our treatment of energy radiated from the accretion shock,  these results should not  be viewed as specific predictions of the core accretion model.  Rather our point is that core accretion naturally leads to gas accretion through a shock, which may result in low entropy planets.  The viability of giant planet formation via core accretion depends on physical processes happening earlier in the accretionary process (at smaller masses) than those processes that we have shown to be crucial for the luminosity of young planets of Jupiter's mass and larger.  We note in passing that the faintness predicted for young Jupiter-mass planets compounds with non-equilibrium chemistry to make detection of young giant planets at M band particularly challenging.  \\citet{Marley96}, after the discovery of \\object{Gl229B}, suggested that a substantial M-band 4 - 5 $\\mu$m flux peak should be a universal feature of giant planets and brown dwarfs.  In addition to the intrinsic emergent flux, this spectral range has looked promising for planet detection due to the favorable planet/star flux ratio \\citep[e.g.][]{Burrows05}. However, it has been known since the 1970s \\citep[see][]{Prinn77} that Jupiter's 5 $\\mu$m flux is suppressed by absorption by CO present in amounts exceeding that predicted by equilibrium chemistry.  This same effect has now been observed in brown dwarf M-band photometry \\citep{Golim04, Leggett06}, as anticipated by \\citet{Fegley96}.  Excess CO leads to strong absorption at 4.5 $\\mu$m, leading to diminished flux in M-band \\citep{Saumon03}.  This effect further suppresses the M band fluxes of young planets below the existing models.  Taken together, fainter young planets and reduced M-band flux may well reduce the catch from what had seemed a promising fishing hole for direct planet detection.  We also conclude that the predicted evolution of giant planets objects, regardless of formation mechanism, is far more sensitive to the precise  conditions at the termination of accretion than has been previously recognized.  Most workers have assumed that evolving model planets `forget' their initial conditions within $10^6$ \\citep{Baraffe03b} to $10^8$ years \\citep{stevenson82} of the first time step.  While 1 to 2 Jupiter mass planets do have a short ($\\sim10^7\\,\\rm yr$) memories, we have shown that more massive planets remember their initial thermal state far longer.  The evolution time scale for young, hot planets depends exponentially on their initial entropy.  Until the initial thermal state of young, low mass objects--even isolated brown dwarfs--is known with more certainty, the early evolution tracks must be regarded with some skepticism. Any effort to assign a mass to a very young putative giant planet must consider the uncertainties in these early evolutionary tracks."
    },
    "0609/astro-ph0609213_arXiv.txt": {
        "abstract": "Roughly $80\\%$ of Ultraluminous Infrared Galaxies (ULIRGs) show blue shifted absorption in the resonance lines of neutral sodium, indicating that cool winds are common in such objects, as shown by Rupke et al and by Martin. The neutral sodium (NaI) columns indicated by these absorption lines are $\\sim10^{13}-3\\times10^{14}\\cm^{-2}$, while the bolometric luminosity varies by a factor of only four. We show that the gas in ULIRG outflows is likely to be in photoionization equilibrium. The very small ULIRG sample of Goldader et al. demonstrates that the ratio of ultraviolet flux to far infrared flux varies by a factor $\\sim100$ from object to object. While the Goldader sample does not overlap with those of Rupke et al. and Martin, we show that such a large variation in ultraviolet flux will produce a similar variation in the column of neutral sodium for a fixed mass flux and density. However, if the cold gas is in pressure equilibrium with a hot outflow with a mass loss rate similar to the star formation rate, the range of ionization state is significantly smaller. Measurements of the UV flux for objects in the Martin and Rupke et al. catalogs will definitively determine if photoionization effects are responsible for the wide variation seen in the sodium columns. If they are, a determination of the gas density and mass loss rate in the cool winds will follow, with attendant improvements in our understanding of wind driving mechanisms and of the effects of galaxies on their surroundings. ",
        "introduction": "Recent observations of the sodium resonance line (NaD) doublet in Ultra-luminous Infra-red Galaxies (ULIRGS) by Rupke and coworkers \\cite{2002ApJ...570..588R,2005ApJ...631L..37R,2005ApJS..160...87R,2005ApJS..160..115R} and by Martin (2005; 2006) show that cool outflows from ULIRGS are common. Martin (2005) finds that 15 of 18 ULIRGS ($83\\%$) possess such flows, while Rupke et al. (2005b) find a detection rate of $80\\pm7\\%$ for 30 local ULIRGs. These results indicate that the outflow emerges in most directions. The absorption troughs, which typically extend over $\\sim300\\kms$, are not black, which at first blush might suggest that the outflows are optically thin in the NaD line. However, the doublet line ratio, which is equal to 2:1 in optically thin gas, rarely has that ratio in Martin's sample, while Rupke et al. (2005a) find optical depths for the weaker of the doublet lines that range from $0.06$ to $7$, with an average of $1.5$. This finding shows that some sight lines to ULIRG galaxies are optically thick in the NaD line, that the optically thick outflow covers only about $25-30\\%$ of the (optically emitting) galaxy, but that this optically thick component is seen toward $\\sim80\\%$ of ULIRGS. Given the patchy nature of the interstellar medium in most galaxies, the last finding is not entirely surprising.  However, another implication of the observations is very surprising. As Figure \\ref{Fig:column} shows, the mean column averaged over both samples (Rupke et al. 2005a; Martin 2006) is $13.7$ in the log, with a standard deviation of $0.66$ dex, or from $10^{13}\\cm^{-2}$ to $3\\times10^{14}\\cm^{-2}$. The actual Na I columns vary from $7\\times10^{12}\\cm^{-2}$ to $5\\times10^{14}\\cm^{-2}$ in the sample of \\cite{2005ApJS..160...87R}. Since the luminosity, and star formation rate of the galaxies in both samples varies by only a factor of about 4 (the standard deviation is a factor of 2), one might expect the $N_H$ column and hence the $N_{NaI}$ column toward the galaxies to vary by a similar factor; in simple galactic wind models the mass outflow rate increases with increasing star formation rate.  The Na I column can be used to find the mass loss rate from each galaxy, in principle. For example, in a smooth flow, the density is related to the mass loss rate by \\be \\label{eq:mdot}% \\dot M_w = \\Omega r^2\\mu(r) n(r) v(r), \\ee % where $\\mu(r)$ is the mean molecular weight, $n(r)$ is the number density of hydrogen, $v(r)$ is the velocity of the flow at a distance $r$ from the center of the galaxy, and $\\Omega\\le4\\pi$ is the global wind cover factor in steradians. For simplicity we employ a spherical model, although ULIRGs appear to form a substantial fraction of their stars in a kiloparsec (or smaller) scale disk \\cite{1991ApJ...378...65C}.  The hydrogen column through such a wind is \\be % N_H = {\\dot M_w\\over \\Omega}\\int_{r_0}^\\infty {dr\\over r^2 \\mu(r) v(r)}. \\ee % A very rough estimate is found by taking $\\mu(r)=m_p$ and $v(r)=v_\\infty$, where the latter is the terminal velocity of the outflow, which in the spirit of approximation adopted here is taken to be the maximum observed blue shift: \\be \\label{eq:column}% N_H\\approx{\\dot M_w\\over\\Omega m_p v_\\infty r_0}, \\ee % where $r_0$ is the radius at which the absorption line forms. Solving for the mass loss rate in terms (as far as possible) of observed quantities, \\be % \\dot M_w \\approx \\Omega m_p r_0 v_\\infty N_{NaI} \\left({N_{Na}\\over   N_{NaI}}\\right) \\left({N_H\\over N_{Na}}\\right)d_{Na}, \\ee % where $N_{NaI}$ is the (observed) column of gas phase neutral sodium, $N_{Na}$ the total gas phase sodium column, $N_H$ is the total hydrogen column, and $d_{Na}$ is the ratio of total to gas phase Na; it is larger than unity since much of the Na is locked up in dust grains. Neither of the two terms in parentheses nor $d_{Na}$ are measured, nor is $r_0$ well constrained. For solar abundances $N_H/N_{Na}=4.9\\times10^5$ \\cite{AG}, while \\cite{1996ARA&A..34..279S} find that the gas phase abundance of Na in the cool gas of the diffuse cloud toward $\\zeta$ Ophiuci is a factor of $d_{Na}=8.9$ smaller than the solar abundance of Na.  There is evidence of dust in ULIRG outflows; first, NaD is a resonance line, so that it scatters continuum photons rather than destroying them. Since the winds are nearly spherical (they are seen toward 80\\% of ULIRGS), any photons removed from our line of sight by the NaD transition should appear along some other line of sight as redshifted emission; since we see no such emission, some other mechanism is removing the scattered photons, absorption by dust being the number one suspect. Second, there is a correlation seen between reddening and NaD absorption toward luminous infrared galaxies, again suggesting dust in the outflow \\cite{1995ApJS...98..171V}; \\cite{2000ApJS..129..493H}. In the rest of this paper we will adopt solar metalicities and the Milky Way ISM $d_{Na}$ just quoted. Using these values, the typical observed column $NaI=10^{14}\\cm^{-2}$ would correspond to a hydrogen column $N_H\\approx10^{21}\\cm^{-2}$ if somewhat less than half the gas-phase sodium were neutral.  Any local heating, such as might be produced by shocks, will lead to lower depletion levels, but by less than the factor of ten corresponding to all the grains being destroyed, given the evidence for dust cited above.  The force responsible for expelling the wind is uncertain. One candidate is ram pressure from outflowing supernova-heated gas, which will entrain cold gas from the interstellar medium of the ULIRG. A variant on this is ram pressure from hot gas produced by an accreting massive black hole (an active galactic nucleus, or AGN). The Compton temperature of both Seyfert and Quasar nuclei is around $10^7$K, similar to that from supernova heating. Momentum driving from supernovae can be significant, if most supernovae explode while surrounded by cold gas. Another candidate for momentum driving is radiation pressure on dust grains embedded in the cold gas. The radiation could arise either from stars or from a central black hole. The radiation pressure from a black hole can exceed that from a starburst. Whether such a centrally driven outflow can produce an outflow that will be detected from $80\\%$ of the sky is far from certain.  In fact it seems likely that both energy and momentum driving operate. ULIRGS are enveloped in x-ray emitting gas, Arp 220 being an excellent example \\cite{2003ApJ...591..154M,2002ApJ...581..974C}; the presence of hot gas is a necessary (but not sufficient) condition for energy driving of a cool outflow.  As we have just noted, the winds are most likely dusty. The wind optical depth to optical or UV emission is likely substantial, while the galaxies themselves are optically thick even to the far infrared. The outflow may well be multiphase; the questions to be addressed revolve around which phase carries more mass, momentum, energy, and metals into the surrounding intergalactic medium.  To answer such questions, we would like to measure the mass loss rate in cold gas. If the NaI columns are taken at face value, and the ratio $N_H/N_{NaI}$ is assumed to be constant, the mass loss rates in cold gas vary from ULIRG to ULIRG by a factor of $\\sim50$. Is this really true?  There are numerous possible reasons for the NaI column seen in galactic outflows to vary from object to object. The mass loss rate will vary from object to object, so that the total column of hydrogen, and hence Na, will vary, even at a fixed metallicity. The metallicity of the outflowing gas may vary from object to object. Both these properties are likely to depend on the star formation rate (and history); the star formation rate, at least, varies over only a small range for the samples we are discussing. A more likely cause of variation in NaI column is a variation in either or both the depletion on to dust grains and the ionization state of Na. Neutral sodium is a rather delicate atom, easily ionized by ultraviolet radiation. The wind may be irradiated either by starlight from the host galaxy, or by shocks in the wind.  In this paper we argue that the outflow is in photoionization equilibrium, except possibly for the most UV dim systems, and then show that if Martin and Rupke et al.'s galaxies have spectral energy distributions like those of other nearby ULIRGs that have been observed in the ultraviolet, then it may be that the mass loss rates are more nearly equal than the observed NaD columns imply; unless the neutral gas is very dense, $n\\sim 10^5\\cm^{-3}$, the ionization fraction, or the ratio of neutral to total sodium, $(NaI/Na)$, will vary by a large factor from object to object. This follows from the observed variation of $\\sim100$ in the ratio of UV to FIR flux seen in nearby ULIRGs (Goldader et al. 2002), and simple photoionization calculations, as we show below. ",
        "conclusions": "The mass loss rates of ultraluminous infrared galaxies, as determined by observations of the NaD lines, are believed to be substantial---of order one fifth of the star formation rate.  This estimate assumes that the ionization fraction of sodium is $\\sim0.1$.  The gross physical properties of ULIRGs, including their stellar masses, gas masses, sizes, and star formation rates vary by factors of only a few, so at first glance the assumption of fixed NaI ionization fraction appears reasonable. However, one property of ULIRGS that appears, based only on a few galaxies, to vary by a large amount is the ultraviolet flux. We have shown that cool ULIRG winds are likely to be in photoionization equilibrium. We used a simple argument to show that the large object-to-object UV flux variations seen in the small number of galaxies that have been observed in the ultraviolet would be enough to produce similarly large variations in NaI column density, if the density of the cool gas is $n\\lesssim10^5(500\\pc/r)^2\\cm^{-3}$. More complete photoionization calculations using CLOUDY produce the same magnitude of variation in NaI/Na with variation in UV flux. They also show, for these moderate densities, that the MgII column should not vary appreciably with changes in UV flux.  If, however, the cold gas is in pressure equilibrium with a hot wind having a mass loss rate similar to the star formation rate of a typical ULIRG, the cool gas will have a density $n\\gtrsim 10^6(500\\pc/r)^2)\\cm^{-3}$. In that case the ionization parameter is $U_{NaI}<10^{-3}$, and there will be little variation in NaI column from object to object. As a check, the (so far unobserved) MgII columns are predicted to show variations that are anti-correlated with the UV/FIR ratio.  We conclude that the observed factor of $\\sim100$ variation in the UV/FIR flux density ratios in ULIRGs with similar bolometric luminosities is consistent with the factor of $\\sim40$ variations in the column density of Na I outflows seen by Martin (2006) and \\cite{2005ApJS..160...87R}; the different UV fluxes may well be the cause of the variation seen in the NaI columns, if the cool gas has a density $n\\lesssim 10^5\\,\\cm^{-3}$. A tight correlation between Na I column and UV flux would indicate that the cool gas is not accelerated by a hot wind.  The origin of the object-to-object variation in the UV/FIR ratio is not clear. The known correlation between reddening and NaI column in lower luminosity galaxies suggests that dust in the outflow may play a role. However, the NaI optical depth toward luminous infrared galaxies in the sample of Rupke et al. (2005b) actually exceeds that toward their (more luminous) ULIRGs. It is unlikely that the UV/FIR ratio is smaller in the less luminous objects. Thus dust absorption in the outflow is probably not the sole reason for the low UV fluxes in ULIRGS.  Whether the galaxies studied by Martin (2005) and by Rupke et al. (2005a,2005b) actually have large object-to-object variations in their near-UV fluxes is not known. We note that the results of Goldader et al. (2002) show that near IR or even optical flux measurements are not a good proxy for the UV flux; see the data for VV114, IRAS 15250+3609, and Mrk 273 in Figure \\ref{Fig:UV}. This demonstrates that optical observations will not be sufficient to answer this question. Observations by GALEX or by the ACS on the Hubble Space Telescope could answer this question definitively."
    },
    "0609/astro-ph0609449_arXiv.txt": {
        "abstract": "We report on the detection in Sloan Digital Sky Survey data of at least three, roughly parallel components in a $65\\arcdeg$-long, stellar stream complex previously identified with the Anticenter or Monoceros Ring.  The three-stream complex varies in width from $4\\arcdeg$ to $6\\arcdeg$ along its length and appears to be made up of two or more narrow substreams as well as a broader, diffuse component. The width and complexity of the stream indicate that the progenitor was likely a dwarf galaxy of significant size and mass. The stream is 8.9 kpc distant and is oriented almost perpendicularly to our line of sight. The visible portion of the stream does not pass near any known dwarf galaxies and a preliminary orbit does not point to any viable progenitor candidates. Orbits for the narrower substreams can be modeled with velocity offsets from the broad component of $\\approx 8$ km s$^{-1}$. We suggest that the broad component is likely to be the remains of a dwarf galaxy, while the narrower streams constitute the remnants of dynamically distinct components which may have included a native population of globular clusters. While the color of the main sequence turn-off is not unlike that for the Monoceros Ring, neither the visible stream nor any reasonable projection of its orbit passes through Monoceros or Canis Major, and we conclude that this stream is probably unrelated to the overdensities found in these regions. ",
        "introduction": "The value of large scale digital sky surveys to studies of Galactic structure and the Local Group has become abundantly clear in recent years, and particularly in the last few months.  In addition to the large scale features attributed to past galaxy accretion events \\citep{yann03,maje2003,roch04}, Sloan Digital Sky Survey (SDSS) data were used to detect the remarkably strong tidal tails of Palomar 5 \\citep{oden2001,rock2002, oden2003, grill2006b}, NGC 5466 \\citep{belo2006a, grill2006a}, and streams due to extant or extinct globular clusters \\citep{grill2006c} and dwarf galaxies \\citep{belo2006b, grill2006d}. \\citet{will2005}, \\citet{zuck2006}, and \\citet{belo2006c} recently used SDSS data to discover several new dwarf satellites of the Milky Way.  In this paper we continue our analysis of the SDSS database to search for extended structures in the Galactic halo. We briefly describe our analysis in Section \\ref{analysis}. We discuss the structure of a stream complex, attributed to the Monoceros Ring by \\citet{belo2006b} in Section \\ref{discussion}, estimate distances in Section \\ref{distance}, and put constraints on the orbit in Section \\ref{orbit}. ",
        "conclusions": "}  Apparent in Figure 1 is a broad and complex stream running from north to south across the field. Portions of this stream are visible in \\citet{newberg02} and \\citet{belo2006b}. The narrow, curved stream running towards the northeast is part of the 63$\\arcdeg$-long globular cluster stream found by \\citet{grill2006c}, hereafter referred to as GD-1.  The broad, east-west stream just above the main gap in the data is the Sagittarius stream discussed by \\citet{belo2006b}, though somewhat muted by our filtering due to its greater distance.  The stream complex of interest extends from (R.A., decl.) = ($126.4\\arcdeg, -0.7\\arcdeg$) to (R.A., decl.) = ($133.9\\arcdeg, 64.2\\arcdeg$), and runs in a $65\\arcdeg$, nearly great circle path from Ursa Major in the north to Hydra in the south. The stream is truncated at both the southern and northern ends by the limits of the available data. We note that there is also an apparent concentration of stars at (R.A., decl.) = ($134\\arcdeg, 3.4\\arcdeg$), surrounded by faint, banded substructure roughly parallel to that of the main western stream.  From the reddening map of \\citet{schleg98}, the maximum values of $E(B-V)$ are $\\approx 0.2$ (near the northern tip), with typical values near 0.03 along the remainder of the stream.  There are diminutions here and there in the stream that could be attributed to regions of higher reddening, but there are no long features with a north-south orientation which could be held to account for either the appearance of the stream as a whole, or for the different components within it.  Sampling at several representative points, the main stream complex appears to be about $5\\arcdeg$ wide on average. This is significantly broader than the globular cluster streams found by \\citet{oden2003, grill2006a, grill2006b} and \\citet{grill2006c}. $5\\arcdeg$ corresponds to about 800 pc at our estimated distance to the stream (see below), which is much larger than the tidal diameters of globular clusters. We conclude that the progenitor was considerably more extended than a globular cluster and was most likely a dwarf galaxy.  The stream is clearly not just a broad swath of stars. Rather, it appears to be made up of a $\\sim 2\\arcdeg$ wide broad component running roughly along the center of the stream complex, and at least two narrower, $1\\arcdeg$-wide streams to the east and west of the broad component. The fine structure in the stream complex is illustrated in Figure 2, where we have made several east-west slices across the complex at various declinations. Comparing Figures 1 and 2, there are indications of additional structure within the broad component, and still other, more tenuous parallel streams to the east and west of the three major components.  Integrating the background subtracted, weighted star counts along the stream over a width of $\\approx 5\\arcdeg$ we find that the total number of stars in the discernible stream down to $g = 22.5$ is $9200 \\pm 1500$. The mean surface density of stars in the southern portion of the stream ($-1\\arcdeg < \\delta < 14\\arcdeg$) is about 67 stars deg$^{-2}$, which is roughly twice the 32 stars deg$^{-2}$ (corrected for $cos(\\delta)$) found in the northern section ($39\\arcdeg < \\delta < 52\\arcdeg$). The highest local surface densities exceed 200 stars deg$^{-2}$. If we assume a globular cluster-like luminosity function in the stream, then we can use the color transformation equations of \\citet{smith2002}, the deep M 4 luminosity function of \\citet{richer2002}, and the mass-luminosity relation of \\citet{baraffe97} to extrapolate to fainter magnitudes. Integrating over $3 < M_V < 17$, we estimate a total number of stars in the visible portion of the stream of $33,000 \\pm 5400$, a total luminosity of $5500 \\pm 900 L_\\odot$, and a total estimated mass of $9300 \\pm 1500 M_\\odot$.   \\subsection{Color-Magnitude Distribution and Distance to the Stream \\label{distance}}   In Figure 3 we show $g, g - i$ color-magnitude distributions for the stream stars, extracted by generating a Hess diagrams of stars lying along $2\\arcdeg$-wide regions covering the western and eastern halves of the stream complex and subtracting a similar field star distribution sampled over $\\approx 500$ deg$^2$ to the east and west of the stream complex.  Despite the somewhat limited statistics, a clear signature of the turn-off and main sequence is evident in the stream population.  Moreover, the turn-off regions of the distributions match the dereddened, shifted main sequence locus of M 13 fairly well. The turn-offs in the eastern and western halves of the complex lie at dereddened $g - i = 0.27$ and $0.3$, respectively. A similar diagram in $g - r$ yields dereddened turn-off colors of $g - r = 0.23$ and 0.25, respectively, with the difference most likely a due to $\\sim 0.02$ mag measurement uncertainty. These estimates lie between values of 0.26 and 0.22 measured by \\citet{newberg02} for the Monoceros Ring and the Sagittarius stream, respectively. They are also much bluer than \\citet{newberg02}'s estimate of dereddened $g - r = 0.40$ for thick disk stars.  Based solely on turn-off color, we can therefore rule out association of the stream complex with thick disk stars. However, within the uncertainties, the turn-off colors are consistent with those of either the Monoceros Ring or the Sagittarius stream.   Varying the magnitude shift applied to M 13's main sequence locus from -1.0 to +3.0 mag, we measured the foreground-subtracted, mean surface density of stream stars in the regions $-1\\arcdeg < \\delta < 9\\arcdeg$, $17\\arcdeg < \\delta < 39\\arcdeg$, and $39\\arcdeg < \\delta < 63\\arcdeg$.  To avoid potential problems related to a difference in age between M 13 and the stream stars, we used only the portion of the filter with $19.5 < g < 22.5$, where the bright cutoff is 0.8 mag below M 13's main sequence turn-off. Though this reduces the stream contrast somewhat, it provides sufficient integrated signal-to-noise to enable main sequence fitting.  Fitting Gaussians to the mean surface densities as a function of magnitude shift (e.g. \\citet{grill2006c}, we find that the highest contrasts occur for shifts of +0.31, +0.29, and +0.37 mag for the eastern half of the stream complex over the declination ranges given above, respectively. For the western half, we find that the filter response peaks at shifts of +0.34, +0.25, and +0.38 mag, respectively. The magnitude shifts for the two halves are highly consistent with one another, and we conclude that there is no significant distance offset from one side of the complex to the other. Adopting a distance to M 13 of 7.7 kpc \\citep{harris96} we find an average heliocentric distance of $d = 8.9 \\pm 0.2$ kpc. The stream is roughly perpendicular to our line of sight and slightly curved about the Galactic center as expected. Our distance estimate is in excellent agreement with the $\\approx 9$ kpc found by \\citet{ibata03} in their WFS-0801 field, which is situated on the western edge of the stream complex at $\\alpha = 120.5, \\delta = 40.3$.   \\subsection{Constraints on the Orbit \\label{orbit}}  The visible portion of the stream complex spans the Galactic anticenter direction and, projecting a great circle path, is inclined by $35\\arcdeg$ to the Galactic plane. Though we are currently limited by a lack of velocity information, for a given model of the Galactic potential the progenitor's orbit is actually fairly well constrained by the observed distance and orientation of the stream. Using the Galactic model of \\citet{allen91} (which includes a disk, bulge, and spherical halo, and which \\citet{grill2006a, grill2006b} and \\citet{grill2006c} found to work reasonably well for NGC 5466, Pal 5, and GD-1), we use a least squares method to fit both the orientation on the sky and the distance measurements in Section \\ref{distance}. In addition to a number of normal points lying along the central component of the stream, we chose as a velocity fiducial point a position at the northern end of the stream at (R.A., decl) = (125.463\\arcdeg, 51.492\\arcdeg).  If we allow the proper motions to be free ranging and uninteresting parameters, the model which best fits the data predicts $v_{LSR} = -18 \\pm 10$ km s$^{-1}$ at the fiducial point, where the uncertainty corresponds to the 95\\% confidence interval. A projection of this orbit is labeled C in Figure 4. We note that the uncertainty is primarily determined by the large lever arm over which it has been possible to measure relative distances. The 95\\% range in $v_{LSR}$ in turn predicts a range in perigalactic and apogalactic radii of $6.6 < R_p < 6.9$ kpc and $16.8 < R_a < 17.3$ kpc. Of course, these ranges do not take into account uncertainties in the absolute distance of the stream (which depends on the uncertainty in M 13's distance) or of the validity of \\citet{allen91}'s Galactic model.  Given the very similar distances estimated for the eastern and western portions of the stream, it is highly unlikely that the complex could be a superposition of multiple wraps around the Galaxy. The two narrower streams (E and W in Figure 4) can be reasonably well modeled by 0.18 $mas$ yr$^{-1}$ offsets in east-west proper motion. At the distance of the stream complex this amounts to $\\approx 8$ km s$^{-1}$. The entire stream complex is therefore likely to be the remains of a dwarf galaxy which contained distinct components spanning a range of binding energies. The narrower streams might, for example, be the remnants of the parent galaxy's globular cluster population. Piecing together the original structure and evolution of the stream's progenitor will require detailed N-body modeling.  Integrating orbits for parameter sets spanning the range above, we find that, with the exception of the Sagittarius dE, there are no known dwarf galaxies within $5\\arcdeg$ of the projected orbit.  The Sagittarius dE lies $3.1\\arcdeg$ from the projection of the best-fit orbit, but the orbital planes of the the Sagittarius dE and the new stream are clearly distinct (Figure 4); we attribute this apparent proximity to the expected confluence of orbit projections in the direction of the Galactic center and not to any physical association between them.  The orientation of the stream on the sky puts fairly strict limits on the plane of the orbit. The visible portion of the stream passes through Lynx, Cancer, and Hydra. We find no reasonable combination of parameters that would place the southern projection of the stream in either Monoceros or Canis Major. Nor does the visible extent of the stream complex fit either the prograde or retrograde models of the Monoceros stream computed by \\citet{penarrubia05}. Thus, even while the turn-off colors appear to be similar, we conclude on orbital grounds that this stream complex is unlikely to be related to either the Monoceros stream or the Canis Major overdensity."
    },
    "0609/astro-ph0609163_arXiv.txt": {
        "abstract": " ",
        "introduction": "A large body of observational data provides evidence for the existence of dark matter, which manifests itself only through its gravitational influence on stellar systems of various scales (from individual galaxies to clusters of galaxies and superclusters), determining to a great extent their dynamics and structure. The controversy surrounding the following two main questions related to the dark matter is still continuing: What is its nature and to what extent does it exceed in mass the luminous matter? The first question is far from being resolved. As regards the second question, on the scales of individual galaxies, it is formulated as follows: Beginning from which regions (inner or outer) does the dark halo mass prevail over the luminous mass?   There are several observational constraints on the mass and extent of the dark halos in galaxies; these constraints do not depend on what the nature of the dark matter is. In general, they give an estimate of the lower limit for the dark halo mass within a fixed radius. For spiral galaxies, this estimate primarily follows from the analysis of the contributions from various components of the system to its rotation curve for the so-called maximum disk model (a classical example of the separation of the contributions from the disk and the dark halo to the rotation curve of a galaxy, NGC 3198; van Albada et al. 1985). In this model, the dark halo mass within the optical radius of a galaxy generally does not exceed the stellar disk mass. When specific systems are investigated, a joint interpretation of the data on the rotation curves and radial velocity dispersion profiles, along with considerations regarding the marginal stability of stellar disks (see, e.g., Bottema and Gerritsen 1997; Khoperskov et al. 2001; Zasov et al. 2001; Khoperskov and Tyurina 2003)\\footnote{In these papers, the results of numerical $N$-body simulations with a variable disk mass were used to explain the observed stellar velocity dispersion for a number of spiral galaxies.}, often raises significantly the lower limit for the dark halo mass within a fixed radius (occasionally, interpretation of the data on the stellar velocity dispersion based on numerical simulations increases the dark halo mass by almost an order of magnitude compared to its value given by the maximum disk model\\footnote{The galaxy NGC 891 may be cited as an example (Khoperskov et al. 2001).}).  The upper limit for the dark halo mass is much more difficult to constrain. Where the region within the optical radius of a galaxy is involved, stability-related considerations are invoked. For example, the presence of a regular spiral pattern related to the propagation of density waves in galaxies requires a gravitationally ``active'' stellar disk, which reduces the contribution from the spherical component to the total gravitational field of the spiral galaxy and to the total mass (Athanassoula et al. 1987). Another important test for the presence of a massive halo is the length of the tidal tails in interacting spiral galaxies. In interacting systems, we quite often observe extended tidal features stretching to distances as large as 40--100 kpc. Based on numerical simulations of interacting disk galaxies with dark halos that correspond to models obtained in cosmological calculations (Navarro et al. 1999), Dubinski et al. (1999) showed that the calculated parameters could be reconciled with the parameters of observed systems only for two types of models: models with an extended moderate-mass halo and a rotation curve that is determined almost entirely by the stellar disk within the optical radius and models with a compact low-mass halo that makes the main contribution to the rotation curve in the inner disk. Reconciling the rotation curves of specific interacting systems (e.g., NGC 4038/39 and NGC 7252) and the sizes of their tidal features with the model parameters leads one to conclude that the stellar disks in such systems dominate in mass within the optical radius, which gives a strong upper limit on the dark mass. A similar conclusion was also drawn from the simulations of the famous Mice galaxy, NGC 4676 (Sotnikova and Reshetnikov 1998).  Comparison of the various estimates for the dark mass in galaxies shows that the upper and lower limits do not overlap (see, e.g., McGaugh and Block 1998). This gives reason to revise a number of tests, in particular, the test related to the thickness of the stellar disks in spiral galaxies suggested by Zasov et al. (1991, 2002). These authors argue that the thickness of a stable stellar disk at a given halo mass is limited below. If the disk thickness is smaller than a certain value, then bending instability will lead to increasing the vertical velocity dispersion and to a thickening of the system. The growth of bending perturbations is stabilized by a massive dark halo: the more massive the halo is, the thinner the stable disk can be. As we showed previously (Sotnikova and Rodionov 2005), not only a massive dark halo has a stabilizing effect on the growth of bending perturbations, leaving the stellar disk fairly thin. A low-mass compact bulge also produces such an effect. In this case, the lower limits on the dark mass can be mild and consistent with the upper limits. In this paper, we analyze the relationship between the stellar disk thickness and the dark halo mass that is derived both from theoretical considerations and from numerical simulations of the dynamical evolution of thin disks in the presence of spherical components with different density profiles and different masses. The derived relationship predicts moderate dark halo masses even for the thinnest galaxies. ",
        "conclusions": ""
    },
    "0609/astro-ph0609480_arXiv.txt": {
        "abstract": "{A study of the structural and scaling properties of the temperature distribution of the hot, X-ray emitting intra-cluster medium of galaxy clusters, and its dependence on dynamical state, can give insights into the physical processes governing the formation and evolution of structure.} {Accurate temperature measurements are a pre-requisite for a precise knowledge of the thermodynamic properties of the intra-cluster medium.} {We analyse the X-ray temperature profiles from XMM-Newton observations of 15 nearby ($z<0.2$) clusters, drawn from a statistically representative sample. The clusters cover a temperature range from 2.5 keV to 8.5 keV, and present a variety of X-ray morphologies. We derive accurate projected temperature profiles to $\\sim 0.5\\,\\rv$, and compare structural properties (outer slope, presence of cooling core) with a quantitative measure of the X-ray morphology as expressed by power ratios. We also compare the results to recent cosmological numerical simulations.} {Once the temperature profiles are scaled by an average cluster temperature (excluding the central region) and the estimated virial radius, the profiles generally decline in the region $0.1\\,\\rv \\lesssim R \\lesssim 0.5\\,\\rv$. The central regions show the largest scatter, attributable mostly to the presence of cool core clusters. There is good agreement with numerical simulations outside the core regions. We find no obvious correlations between power ratio and outer profile slope. There may however be a weak trend with the existence of a cool core, in the sense that clusters with a central temperature decrement appear to be slightly more regular. } {The present results lend further evidence to indicate that clusters are a regular population, at least outside the core region. } ",
        "introduction": "The temperature and density are the key measurable characteristics of the hot, X-ray emitting intracluster medium (ICM). The determination of important derived properties such as entropy, pressure, and, under the assumption of hydrostatic equilibrium, the total mass, is dependent on accurate estimation of these quantities. Because of limited photon statistics\\footnote{Also the need for an azimuthally symmetric approximation for purposes of deprojection.} it is usual to measure the density and temperature in terms of radial profiles. However, while the density of the ICM is relatively easy to measure from the surface brightness profile of a given cluster, the temperature determination requires sufficient photon statistics to build, and fit, a spectrum. Thus ICM temperature profiles are typically determined with considerably less spatial resolution than density profiles.  The measurement of radial temperature profiles is further complicated by the density squared ($n_e^2$) dependence of the X-ray emission. The steep drop of the X-ray surface brightness with distance from the centre, combined with the background from cosmic, solar and instrumental sources, makes accurate measurement of the temperature distribution at large distances from the centre a technically challenging task.  The earliest temperature profiles were measured with {\\it Einstein}, {\\it EXOSAT}, {\\it Spacelab-2\\/} and {\\it GINGA\\/} only for the nearest, brightest clusters (e.g., \\citealt*{fab1,fab2,hughes,eyles,koy}). The low, stable background of {\\it ROSAT} made possible spatially resolved spectroscopy of poor clusters \\citep*[e.g.][]{david}; however, limited spectral resolution and bandwidth made such measurements difficult for hotter clusters \\citep*[e.g.][]{hbn,bh94,hb95}. {\\it ASCA\\/} and {\\it BeppoSAX\\/} had sufficient high-energy sensitivity to accurately measure the temperatures of hot clusters. However both of these satellites suffered from significant PSF blurring, which, in the case of {\\it ASCA},  was exacerbated by a significant energy dependence.  As a result, at the end of the {\\it ASCA/BeppoSAX\\/} era, the exact shape of cluster temperature profiles was still under vigorous debate \\citep*{mark98,irwin99,whi00,ib00,fin01,dm02}.  {\\it Chandra\\/} and \\xmm do not suffer from major PSF problems. The on-axis {\\it Chandra\\/} PSF is negligible, while the \\xmm PSF becomes an issue only for clusters with very centrally peaked core emission; in addition, neither is energy-dependent. Recent observations of moderately large samples consisting primarily of nearby cooling core clusters with \\xmm \\citep{piff} and {\\it Chandra\\/} \\citep{vikh05} have largely validated the original {\\it ASCA\\/} results of Markevitch et al., which suggested that temperature profiles declined from the centre to the outer regions. However, other {\\it Chandra} and \\xmm observations have found flatter profiles \\citep{allen01,kaa,app}. As of the time of writing, no systematic attempt has been made, with either \\xmm or {\\it Chandra}, to look at the temperature profiles of a representative sample of nearby clusters\\footnote{Some work has been done on medium-distant clusters, see \\citep{zhang,kotov}}. Although other projects on representative samples are in progress \\citep[e.g.,][]{reiprich}, they are not expected to be able to map the temperature distribution out to large radius.  In this paper we deal with observations of 15 clusters from a statistically representative sample observed with \\xmm. We describe in detail the data reduction and background subtraction, and compare our results with previous work and with those from cosmological hydrodynamical simulations. We also make a preliminary investigation of correlations with quantitative morphological measures. We present only projected temperature profiles in this paper -- such profiles are direct observables and do not depend on complicated PSF and deprojection algorithms. We will deal with correction of the profiles in forthcoming papers which make use of observations of the full sample. All results are given assuming a $\\Lambda$CDM cosmology with $\\Omega_m=0.3$ and $\\Omega_\\Lambda=0.7$ and $H_0 = 70$~km~s$^{-1}$~Mpc$^{-1}$. Unless otherwise stated, errors are given at the $68$ per cent confidence level. ",
        "conclusions": "We have used \\xmm observations of 15 clusters drawn from a statistically representative, luminosity-selected sample to investigate the behaviour of the temperature profiles. The clusters range from morphologically relaxed looking objects with strong central surface brightness peaks (e.g., RXC\\,J1044\\,-0704), to diffuse structures with significant amounts of surrounding substructure (e.g., RXC\\,J1516\\,+0056), and constitute a representative subsample. We find that, once scaled appropriately, the temperature profiles are similar in the radial range from $0.1\\,\\rv$ to $0.5\\,\\rv$, declining steadily from the central regions to the outer boundary of the measurements with a relative dispersion of $\\sim 10$ per cent out to $0.5\\,\\rv$. The region interior to $0.1\\,\\rv$ is the region of greatest scatter in the scaled profiles: the relative scatter of $\\sim 25$ per cent is likely a lower limit. A preliminary comparison with numerical simulations shows relatively good agreement outside $\\sim 0.1\\,\\rv$, with all of the measured temperature profiles falling within the scatter of the simulated profiles.  Calculating power ratios for the sample, we investigate whether there are correlations between the power ratio measured in an aperture corresponding to $R_{500}$ and the shape of the temperature profile. We characterise the temperature profile shape in two ways: with the ratio $T (0.5\\,\\rv)/T (0.2\\,\\rv)$, a measure of the outer slope, and with the ratio $T_c/\\langle T \\rangle$, a measure of the central temperature drop. There is no obvious correlation of outer slope with power ratio; neither is there a correlation of central temperature dip with $P_2/P_0$ or $P_3/P_0$. There is evidence for a weak correlation of the central temperature dip with $P_4/P_0$. The analysis thus suggests that the outer slope of the temperature profile is not particularly sensitive to the morpho-dynamical state, although there may be some correlation with the existence of a central temperature drop. Further investigation with power ratios evaluated in other apertures, for the entire sample, should be undertaken before definitive conclusions can be drawm.  The overall conclusion from this work on a statistically representative sample indicates that clusters are a relatively regular population, at least outside the cool core regions, with the caveat that comparisons between samples or with simulations is limited by the available temperature profile resolution. The observed similarity in density \\citep{neuarn,croston} and temperature profiles \\citep[][this work]{mark98,dm02,piff,vikh05} indicates both similarity in the underlying gravitational mass distribution (such as has already been seen in the X-ray mass profiles of morphologically relaxed clusters, e.g., \\citealt*{point}), and similarity in the entropy of the ICM (such as has been seen by e.g., \\citealt*{psf,pap}). In this case a single integrated temperature, excluding the core region, should be a good proxy for the total mass. The observed regularity thus has important implications for the use of clusters as cosmological probes.  In future papers, we will reinvestigate the trends with the full sample, make maps of quantities such as temperature, entropy and pressure, and estimate the mass, baryon fraction and entropy in the clusters. A more extensive comparison with numerical simulations will also be undertaken."
    },
    "0609/astro-ph0609355_arXiv.txt": {
        "abstract": "{The BH mass (and the related Eddington ratio, $l$ = $L_{bol}/L_{edd}$) in broad line AGN is usually evaluated by combining estimates (often indirect) of the BLR radius and of the $FWHM$ of the broad lines, under the assumption that the BLR clouds are in Keplerian motion around the BH. Such an evaluation depends on the geometry of the BLR.  There are two major options for the BLR configuration: spherically symmetric or ``flattened''. In the latter case the inclination to the line of sight becomes a relevant parameter. This paper is devoted to evaluate the bias on the estimate of the Eddington ratio when a spherical geometry is assumed (more generally when inclination effects are ignored), while the actual configuration is ``flattened'', as some evidence suggests. This is done as a function of luminosity and redshift, on the basis of recent results which show the existence of a correlation between the fraction of obscured AGN and these two parameters up to at least z=2.5 (date at larger redshifts being insufficient.) The assumed BLR velocity field is akin to the ``generalized thick disk'' proposed by Collin et al. (2006). Assuming an isotropic orientation in the sky, the mean value of the bias is calculated as a function of luminosity and redshift. It is demonstrated that, on average, the Eddington ratio obtained assuming a spherical geometry is underestimated for high luminosities, and overestimated for low luminosities. This bias converges for all luminosities at $z$ about 2.7, while nothing can be said on this bias at larger redshifts due to the lack of data. The effects of the bias, averaged over the luminosity function of broad line AGN, have been calculated. The results imply that the bias associated with the a-sphericity of the BLR make even worse the discrepancy between the observations and the predictions of evolutionary models.   \\keywords {Galaxies: active - galaxies: nuclei - quasars: general} } ",
        "introduction": "The accretion rate and the black hole mass are the two fundamental parameters in our understanding of the Active Galactic Nuclei (AGN) phenomenon. Measurements of these two quantities are, unfortunately, not devoid of significant uncertainties.  The accretion rate $\\dot m$ is derived from the bolometric luminosity, $L_{bol}$, under assumptions on the efficiency $\\eta$ for the conversion of gravitational energy, $L_{bol}$ = $\\eta$$\\dot m$$c^2$. For non-rotating black holes (BH), $\\eta$ = 0.057 is generally adopted, assuming that effective (for the observer of the electromagnetic radiation) conversion takes place down to the marginally stable circular orbit at three times the Schwarzschild radius, 3$R_s$; for rotating BH, $\\eta$ can reach the maximum value of 0.42. $L_{bol}$ is generally obtained from the luminosity observed in a given band, multiplied by a factor based on the Spectral Energy Distribution (SED) attributed to the specific class the AGN belongs to. This procedure is regarded to be rather safe, but in fact there are still uncertainties on the luminosity and/or redshift dependence of the bolometric correction.  For the BH mass, two ``direct'' methods have been followed. The first, applicable only to AGN, is the reverberation mapping (RM) method. This method is based on the principle (Blandford \\& McKee 1982) that the delay in the response of the lines from the Broad Line Region (BLR) to variations of the continuum is a measure of the size of this region, $R_{BLR}$. Assuming that the line widths are due to motions governed by the BH, the combination of $R_{BLR}$ and a velocity derived from the line profiles yields a ``Keplerian'' estimate of the BH mass. The other method is based on the fairly strict correlation between the mass of the BH and properties of the stellar bulge of the host galaxy. This method, having proved very reliable for a rather large sample of galaxies (Gebhardt et al. 2000a, Ferrarese \\& Merritt 2000, Tremaine et al. 2002), represents a benchmark for the previous method when it can be applied to AGN (Gebhardt et al. 2000b, Ferrarese et al. 2001, Onken et al. 2004). The agreement found, although far from perfect (e. g. Collin et al. 2006), has encouraged the extension to many more AGN (especially the distant and more luminous ones), for which both the RM and the bulge methods can hardly be applied, of a ``secondary'' method. The latter is based on the estimate of $R_{BLR}$ through an empirical correlation between this quantity and the luminosity (see Sect. 2) which has emerged from the RM measurements.  The BH mass can be univocally converted into the Eddington luminosity, $L_{edd}$. Thus a quantity $l$ can be defined: $l = L_{bol}/L_{edd}$, which, altough it tells us nothing precise about the accretion rate, is of high interest because $L_{edd}$ is a very significant physical limit. This quantity is often referred to as the Eddington ratio.  Several papers have been devoted to explore the behaviour of $l$, in particular as a function of $L$ and the cosmological epoch, namely the redshift $z$. Among the uncertainties and the selection effects which may plague the results, the present paper is devoted to point out and evaluate a particular bias, linked to the possibility that the spatial distribution of the BLR clouds is far from spherical, a situation supported by various lines of evidence. The evaluation is based on a recent result on the fraction of AGN which are photoelectrically absorbed in the X-rays (column density $N_H$ $>$ $10^{22}$ H atoms/cm$^2$, and Compton thin), which can be summarized as follows. Calling $\\xi$ the ratio of the absorbed ones to the total, it turns out that $\\xi$ is a function of $L_X$ (hence of $L_{bol}$) (Ueda et al. 2003, La Franca et al. 2005) as well as of $z$ (La Franca et al. 2005). Qualitatively speaking, in the local Universe this fraction decreases with increasing luminosity; as the redshift grows, the anticorrelation remains but it becomes progressively shallower. If this behaviour is associated with a luminosity and redshift dependence of the opening angle of the absorbing matter, within which the BLR can be observed, it should introduce a bias on the estimate of the BH mass of broad line AGN (AGN 1 for short), when this is performed using the RM method and its ``secondary'' extrapolation. To this effect it is important to stress (see Fiore et al. 2003, Perola et al. 2004) that the value  $N_H$ = $10^{22}$ H atoms/cm$^2$ works as a good (the exceptions are a minority) discriminant between AGN which are optically classified as type 1 and type 2.  The plan of the paper is as follows. In Sect. 2 the results on $l$ from the literature are summarized. In Sect. 3 two lines of evidence, again from the literature, which favour a non-spherical distribution of the BLR clouds are briefly described. In Sect. 4 the bias associated with the a-sphericity of the BLR, and its dependence on $L_{bol}$ and $z$, as it can be predicted on the basis of the abovementioned finding, is quantified. A discussion follows in Sect. 6. ",
        "conclusions": "The curves in Fig. \\ref {qmedioZLbol} show that the mean value of the bias, $<$q$>$, is typically greater than 1 and can be as large as about 2. Furthermore, this bias increases with the luminosity,  converging for all luminosities towards a value of order unity at $z$ larger than 2. As noted previously, the estimate of the opening angle $i_0$, hence of  $<$q$>$, ceases to be reliable in this regime of the cosmic time, because the lack of data prevented La Franca et al. (2005) from evaluating the actual evolution of $\\xi$ further back in time. The ``saturation'' effect adopted in $\\xi$ corresponds to an angle $\\simeq 40^{\\circ}$,  such that, according to eq. (\\ref {qmedio}), $<$q$>$ is $\\simeq 1$ for $a$=0.3, or $\\simeq 0.7$ for $a$=0.1. Since the weight of the solid angle within this value of $i_0$ is relatively modest, the mean value of q will result significantly lower than unity only for $i_0$ much lower than $i_{\\ast}$. If, for instance, $\\xi$ were to reach 90\\% (that is $i_0$ = 26$^{\\circ}$) at values of $z$ much greater than 2, than  $<$q$>$ would be 0.32 ($a$=0.1) or 0.56 ($a$=0.3).  We also calculated the rms of the $q$ distribution and found values of 0.9, 0.6 and 0.2 for $a$=0.1 and $i_0$=90$^{\\circ}$, 60$^{\\circ}$, 30$^{\\circ}$, respectively (the corresponding values of $<$q$>$ are 2.0, 1.3 and 0.4). Similar values are found for $a$=0.3. The expected scatter in masses is thus reassuringly smaller than that found in the Black Hole mass (estimated from the reverberation mapping) vs. bulge dispersion velocity relationship (Onken et al. 2004).  Turning to the parameter $l$, this quantity is {\\it underestimated} when $<$q$>$ $>$ 1, if $M_{BH}$ is ``measured'' assuming a spherical distribution of the BLR, the other way round if $<$q$>$ were less than unity.  A possible way to illustrate the effects of this bias consists in calculating its value averaged over the entire Luminosity Function. La Franca et al. (2005) give both the Luminosity Function (LF) and the $\\xi$ which best fit the data (fit \\#4 in their Table 2), after taking into account the selection effects (in the X-ray and optical bands) for the samples used. The LF behaviour with cosmic time follows a Luminosity Dependent Density Evolution. One can therefore combine the LF and $\\xi$ to obtain, for a given $z$, the average bias, $<$q$>$$_{L}$, over the entire luminosity range:  \\begin{equation} \\label{qmedioL} <q>_{L}=\\frac{\\int_{\\log\\!L_{X1}}^{\\log\\!L_{X2}}\\frac{\\mathrm{d}\\Phi_{1}(L_{x},z)}{\\mathrm{d}\\!\\log\\!L_{X}}\\, \\mathrm{d}\\!\\log\\!L_{X}\\int_{0}^{i_{0}(L_{X},z)}q\\, \\sin i\\, \\mathrm{d}i} {\\int_{\\log\\!L_{X1}}^{\\log\\!L_{X2}}\\frac{\\mathrm{d}\\Phi_{1}(L_{x},z)}{\\mathrm{d}\\!\\log\\!L_{X}}\\, \\mathrm{d}\\!\\log\\!L_{X}\\int_{0}^{i_{0}(L_{X},z)}\\sin i\\, \\mathrm{d}i}, \\end{equation} where $\\Phi_{1}$ is the luminosity function of unabsorbed AGN (which depends on $\\xi$) and the parameter $a$ is included in the function $q$ (eq. (\\ref {qi})). This result should, by construction, be free from selection effects. For a choice of $\\log\\!L_{X1}$ = 42 and $\\log\\!L_{X2}$ = 48, $<$q$>$$_{L}$  is illustrated in the left panel of  Fig. \\ref{qmedioL} (the corresponding values of $\\log\\!L_{bol}$, obtained applying the luminosity dependent bolometric correction of Marconi et al. (2004), are: $\\log\\!L_{bol1}$ = 43.030 and $\\log\\!L_{bol2}$ = 50.937).   \\begin{figure*}[ht] \\hbox{ \\includegraphics[width=7 cm]{Lamastra06fig4a.ps} \\hspace{0.02cm} \\includegraphics[width=7 cm]{Lamastra06fig4b.ps} } \\caption{Left: mean value of q, $<$q$>$$_{L}$, between 42$<\\log\\!L_{x}<$48 as a function of z, for a=0.1 (solid line) and a=0.3 (dotted line). Right: mean value of q, $<$q$>$$_{L}$, between 41$<\\log\\!L_{x}<$48 (dashed line) and  42$<\\log\\!L_{x}<$48 (solid line) as a function of z, for a=0.1.} \\label{qmedioL} \\end{figure*}      To illustrate the effect arising from the inclusion of the numerous AGN in the decade around  $\\log\\!L_{X}$ = 41 ($\\log\\!L_{bol}$ = 41.893), in the right panel of Fig. \\ref {qmedioL} (and only for the case $a$=0.1) $<$q$>$$_{L}$ is compared with the result shown in the left panel of the same figure. The effect is that, the higher is the luminosity, the steeper is the dependence of the mean bias $<$q$>$$_{L}$ on $z$.    A direct application of our results to the samples used by the various authors, quoted in Sect. 2, is not straightforward, because this would require proper evaluation of their specific observational selection effects, a hard task which goes beyond the aims of this paper. However, some general remarks can be made. The results shown in Fig. \\ref {qmedioL} imply that, if $l$ were intrinsically constant with $z$, and if its ``mean'' value were calculated with the spherical approximation (or, more generally, with a constant value of $\\kappa$ in eq. (\\ref {VblrVfwhm})) then one should observe an increase of $l$ with $z$. Since the observational results indicate instead a constant value, the bias discussed in this paper implies that the actual value of $l$ \\textit{decreases} with $z$.  In semianalytical models which link the evolution of the galaxies in the hierarchical clustering scenario with the quasar evolution (e.g. Menci et al. 2003, 2004), the black hole accretion is triggered by galaxies encounters. In this scenario, at high $z$ the protogalaxies grow rapidly by hierarchical merging, meanwhile much cold gas is imported and also destabilized, so that the black holes are fueled at their Eddington rates. At lower $z$ the accretion rate of cold gas onto the central black hole diminish due to the combined effects of the decrease of the galaxies merging and encounter rates and the decrease of the amount of galactic cold gas, which was already converted into stars or accreted onto the black hole. This model predicts an average Eddington ratio dropping from $l$ $\\simeq$ 1 at $z$ $\\simeq$ 2.5 to $l$ $\\simeq$0.01 at $z$ $\\simeq$0 (Menci et al. 2003). Our results imply that the bias associated with the a-sphericity of the BLR make even worse the discrepancy between the observations and the predictions of the models."
    },
    "0609/astro-ph0609433_arXiv.txt": {
        "abstract": "We present spectra of 59 nearby stars candidates, M dwarfs and white dwarfs, previously identified using high proper motion catalogues and the DENIS database. We review the existing spectral classification schemes and spectroscopic parallax calibrations in the near-infrared $J$-band and derive spectral types and distances of the nearby candidates. 42 stars have spectroscopic distances smaller than 25~pc, three of them being white dwarfs. Two targets lie within 10~pc, one M8 star at 10.0~pc (APMPM J0103-3738), and one M4 star at 8.3~pc (LP 225-57). One star, LHS 73, is found to be among the few subdwarfs lying within 20 pc. Furthermore, together with LHS 72, it probably belongs to the closest pair of subdwarfs we know. ",
        "introduction": "Recent discoveries of cool objects, such as M stars or brown dwarfs, closer than 5 parsecs show that even the immediate solar neighbourhood sample is still incomplete \\citep{delfosse2001,scholz2003,teegarden2003,hambly2004}.  \\citet{henry1997} estimated that about 130 systems over 359 (36\\%) are missing within 10 pc. The missing fraction is even larger within 25 pc (63\\%) with a deficit of about 3500 systems over the 5500 expected ones \\citep{henry2002}. Statistical comparisons from the local sample in the northern hemisphere led to a less pessimistic result, the current 10 pc sample being $\\sim$75\\% complete \\citep{reid2003a}.  New surveys in the near infrared such as DENIS \\citep{epchtein1997} and 2MASS \\citep{cutri2003} provide unprecedented data for a systematic search for low luminosity cool dwarfs. The use of these data together with high proper motion catalogues is a powerful tool for discovering our neighbours. Hundreds of stars closer than 25 parsecs have been discovered this way \\citep[e.g.][]{phanbao2001,phanbao2003,reid2002a,reid2002b,reyle2002,reid2003b,reid2004,hambly2004,reyle2004,lodieu2005,scholz2005}.  As spectroscopy provides much more information than photometry alone, spectroscopic observations were also carried out to identify and classify nearby stars. Recent studies dealing with sample sizes larger than 10 objects are cited in Table~1. All together, they revealed 490 stars in the 25 pc sample and 25 within 10 pc. However, we note that the samples investigated by the different authors do have many stars in common so that the total number of newly discovered neighbours is much smaller.  \\begin{table} \\label{tab1} \\caption{Nearby stars identified by spectroscopy } \\begin{tabular}{lccc} \\hline Reference \t&Sample &\\multicolumn{2}{c}{Number of stars within} \\\\ & size                        &25 pc\t&10 pc \\\\ \\hline \\citet{cruz2002}\t  \t&70\t\t\t&28\t\t& \\\\ \\citet{henry2002} \t&34\t\t\t&6\t\t&2 \\\\ \\citet{reid2003b}  \t&357\t\t&127\t&9 \\\\ \\citet{lodieu2005}\t&71\t\t\t&25\t\t&3 \\\\ \\citet{crifo2005}\t\t&39\t\t\t&31\t\t&1 \\\\ \\citet{scholz2005}\t&322\t\t&226\t&8 \\\\ \\citet{phanbao2006}\t&45\t\t\t&5\t\t& \\\\ This study\t\t\t&59\t\t\t&42\t\t&2 \\\\ \\hline\\\\ \\end{tabular} \\end{table}   In previous papers \\citep{reyle2002,reyle2004}, we have determined the photometric distances of high proper motion stars that we cross-identified with the DENIS survey. We have reported the discovery of 115 nearby candidates, probably lying within 25 pc, the limit of the Catalogue of Nearby Stars \\citep[CNS3]{gliese1991}. We selected the closest candidates for a spectroscopic follow-up performed at La Silla Observatory (Chile).  Our sample of nearby candidates is described in \\S~\\ref{sample}. \\S~\\ref{spectro} describes the spectroscopic observations. In \\S~\\ref{classification} we review the spectroscopic classification schemes and compute the spectral type of the stars. Computation of distances is detailed in \\S~\\ref{distances}. In \\S~\\ref{halpha} we discuss the chromospheric activity of the stars in our sample and the conclusions are given in \\S~\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion} We obtained spectral types and spectroscopic distances for 59 high proper motion stars. These stars are nearby candidates, as seen from their photometric distances. 42 stars have indeed spectroscopic distances below the 25~pc limit of the Catalogue of Nearby Stars. These new neighbours are few in regards to  the large number of missing stellar systems within 25 pc ($\\sim$ 2000) according to \\citet{henry2002}. However, this study, joined to similar ones, allows to increase, step by step, the completeness of the solar neighbourhood sample and to cross-check the new nearby candidates in different samples (see e.g. Table 1). Nevertheless, it is clear that further efforts are needed in order to reach a much higher completeness level.  The neighbours recovered from our study are mainly M dwarfs. Three are white dwarfs, cooler than 6200~K. One star, APMPM J0541-5349, appeared to be a binary consisting of two M2 stars with an angular separation of 3.33\\arcsec, at a distance of 53.4 $\\pm$ 4.1 pc from the Sun. We computed the radial velocity of the subdwarf candidate LHS 73 and found a space motion compatible with that of an object belonging to an old population. With our determined distance of 18.7 pc, LHS 73 is among the few subdwarfs within 20 pc. Furthermore, LHS 73 has a companion, LHS 72. LHS 72/LHS 73 is probably the closest known pair of subdwarfs. Two stars lie within 10~pc: APMPM J0103-3738 is a M8 star at distance d = 10.0 $\\pm$ 0.8 pc, and LP 225-57 a M4 star at d = 8.3 $\\pm$ 2.8 pc.  The new sources that we identified are worthy of more detailed observations, including high-resolution imaging to search for low-mass companions, and trigonometric parallax measurements. We plan more follow-up for the nearest objects and hope to place the most promising targets on new or existing parallax programs such as that of the Cerro Tololo Inter-American Observatory Parallax Investigation \\citep[CTIOPI]{jao2005}."
    },
    "0609/astro-ph0609119_arXiv.txt": {
        "abstract": "We present high sensitivity sub-arcsecond resolution images of the Herbig Ae star AB Aurigae at 11.6 and 18.5 $\\mu$m taken with Michelle on Gemini North. Bright extended dust emission close to the star is resolved at both wavelengths, with quadratically subtracted FWHM of 17$\\pm$4 AU at 11.6 $\\mu$m and 22$\\pm$5 AU at 18.5 $\\mu$m. Additional, fainter emission is detected out to a radius of 280 AU at 11.6 $\\mu$m and 350 AU at 18.5 $\\mu$m down to the sensitivity limit of the observations. The latter value is identical to the measured size of the millimeter-continuum disk, but much smaller than the CO disk. Assuming moderately absorbing material, we find that larger particles ($\\sim 1$ $\\mu$m) dominate the mid-IR emission in the inner ($<$ 100 AU) regions of the disk, and smaller particles ($<$ 0.3 $\\mu$m) dominate in the outer regions of the disk.  A model of a nearly face-on passive flared disk with an inner rim accounts well for our observations. ",
        "introduction": "Located only 144 pc away \\citep{anc98}, the 2 Myr old A0 star AB Aurigae is the brightest (V=7.06) of the original sample of Herbig stars \\citep{her60}, which are intermediate mass (2 to 8 M$\\sun$) pre-main sequence stars. Consequently, it is not only the best-studied Herbig object, but it has become an important touchstone for our understanding of the class. The spectral energy distribution (SED) of this source shows emission in excess of the photosphere throughout the infrared region indicative of circumstellar (CS) dust. Different models, among which are a highly inclined passive flared disk with an inner rim \\citep{dul01}; a flat thick disk surrounded by a halo \\citep{vin03}; and a halo alone \\citep{eli04}, have been used to explain the spatial distribution of this dust. Despite their differences, all of these models reproduce reasonably well the observed SED, which indicates the need for high-resolution imaging to provide additional critical constraints.  Spatial observations at various wavelengths imply that the CS dust in the AB Aur system lies in a disk and some type of more extended structure. An inhomogeneous envelope extending to 1300 AU is apparent in optical scattered light \\citep{gra99}, while closer to the star, in optical and near-IR scattered light, one sees what appears to be a disk with quasi-spiral structure, a radius of 580 AU, and an inclination of 30$^\\circ$ (face-on = 0$^\\circ$), assuming flat geometry \\citep{gra99, fuk04}.  CO observations reveal a complex disk with an inner hole of about 70 AU extending out to 1000 AU and possibly having non-Keplerian motions, while 1.4 mm continuum observations indicate a disk with an inner radius of 110 AU and an outer radius of 350 AU \\citep [Pi\\'{e}tu et al. 2005; see also][]{man97, cor05}. Near-IR interferometric studies resolve the inner 0.7 AU region of the disk \\citep{mil01}.  However, previous mid-IR studies present somewhat contradictory results.  Marsh et al. (1995) report extended structure at 17.9 $\\mu$m with a semi-major axis of 80$\\pm$20 AU and an inclination of 75$^\\circ$.  Chen \\& Jura (2003), using Keck, do not confirm that detection of extended structure, and at 20.5 $\\mu$m using deconvolved images with a resolution of 0.$\\arcsec$6, Pantin et al. (2005) report an elliptical ring-like structure at an average distance of 280 AU from the star. In addition, Liu et al. (2004) resolve the inner disk interferometrically at 10.3 $\\mu$m determining a size of 27$\\pm$3 AU and an inclination of 45$^\\circ$. There is also recent evidence that AB Aur could be the brighter component of a binary system, with a companion separation most likely between 1 and 3 arcseconds \\citep{bai06}.  In this paper we present deep mid-infrared images of AB Aur obtained at Gemini North.  We have resolved the emission close to the star at 11.6 and 18.5 $\\mu$m, and we find an additional extended component that appears to be roughly circularly symmetric.  We show how these observations of the thermal emission from dust in the AB Aur system help establish a more coherent picture of the dust geometry consistent with most observations at other wavelengths. ",
        "conclusions": "Our mid-IR images reveal two different emission components in AB Aur.  The central stronger emission is resolved, with quadratically deconvolved FWHM sizes of 17$\\pm$4 AU and 22$\\pm$5 AU at 11.6 $\\mu$m and 18.5 $\\mu$m, respectively.  We also detect fainter extended emission out to a radius of 280 AU at 11.6 $\\mu$m and 350 AU at 18.5 $\\mu$m.  Emission is slightly elongated at 12 $\\mu$m, indicating a disk inclination angle in the range 29$^\\circ \\pm 11^\\circ$ and a PA of 80$^\\circ \\pm 11^\\circ$.  The morphology at 18 $\\mu$m is consistent with an inclination angle of 12$^\\circ \\pm 12^\\circ$. However, within the uncertainties, inclination angles at 12 and 18 $\\mu$m are the same.  Assuming moderately absorbing material, we derive average radii of the mid-IR emitting dust in the system and find that larger particles ($a \\sim 1$ $\\mu$m) dominate the mid-IR emission in the inner ($<$ 100 AU) regions of the disk, and smaller particles ($a <$ 0.3 $\\mu$m) dominate in the outer regions of the disk.  Our results are reasonably well accounted for by a model of a passive flared disk with an inner rim.  The presence of a more spherical component fails to account for the particle size segregation derived from our observations."
    },
    "0609/astro-ph0609605_arXiv.txt": {
        "abstract": "{Observations of X-Ray sources harbouring a black hole and an accretion disc show the presence of at least two spectral components. One component is black-body radiation from an optically thick standard accretion disc. The other is produced in a optically thin corona and usually shows a powerlaw behaviour. Electron-proton (ep) bremsstrahlung is one of the contributing radiation mechanisms in the corona. Soft photons from the optically thick disc can Compton cool the electrons in the corona and therefore lead to a two-temperature plasma, where electrons and ions have different temperatures.} {We qualitatively discuss effects on ep-bremsstrahlung in the presence of such a two-temperature plasma.} { We use the classical dipole approximation allowing for non-relativistic electrons and protons and apply quantum corrections through high-precision Gaunt factors. } {In the two-temperature case ($T_\\textrm{e}< T_\\textrm{p}$) the protons cause a significant fraction of the ep-bremsstrahlung if their speed is high compared to the electrons. We give accurate values for ep-bremsstrahlung including quantum-mechanical corrections in the non-relativistic limit and give some approximations in the relativistic limit.} {The formulae presented in this paper can be used in models of black hole accretion discs where an optically thin corona can comprise a two-temperature plasma. This work could be extended to include the fully relativistic case if required.} ",
        "introduction": "Models of black hole accretion discs are widely used to explain spectral characteristics of X-Ray observations of such objects. These models usually consist of an optically thin and hot corona and a cool, optically thick standard accretion disc. The optically thin corona consists of a two-temperature plasma. Electrons cool by bremsstrahlung, mainly by electron-proton bremsstrahlung. The electron-proton system has got a dipole moment while the electron-electron and proton-proton system can only emit radiation through the (much weaker) quadrupole moment.  \\citet{1975ApJ...199L.153E} presented a model for Cyg X-1 incorporating a two-temperature plasma in the inner parts of the accretion flow. This kind of model has been subsequently developed \\citep[e.g.][]{1976ApJ...204..187S,1995MNRAS.277...70Z} and applied to different objects \\citep{1991ApJ...380L..51H,1993ApJ...413..507H} and different geometries \\citep{1991ApJ...380...84W}.   The calculation of the rate of bremsstrahlung coming from a hot gas has received much attention over a long period of time \\citep[e.g.][]{1962RvMP...34..507B,1970RvMP...42..237B,1980ApJ...238.1026G,1981ApJ...243..677G}. Today the bremsstrahlung rate is known with a very high accuracy. All these calculations however assume electron and proton temperature to be equal.  In the following we limit ourselves to a purely classical and non-relativistic treatment of electron-proton bremsstrahlung. We do this in order to highlight the elementary physical processes involved. Since we apply Gaunt factors to account for quantum mechanical effects, our results are strictly accurate only in the non-relativistic case. An expansion to the fully relativistic case is feasible but beyond the scope of this paper.   In this contribution we first highlight the physical mechanism in Sect.~\\ref{sect:mechanism} and then present a recalculation of the electron-proton Bremsstrahlung rate for a two-temperature plasma in the non-relativistic limit in Sect.~\\ref{sect:calc}. While we limit ourselves to a electron-proton plasma, the formulae can easily be modified to consider mixtures of ions with different mass and charge. We give a discussion of the results and our conclusions in Sect.~\\ref{sect:conclusions}. ",
        "conclusions": "\\label{sect:conclusions}  We have recalculated electron-proton bremsstrahlung for a two-temperature plasma. Owing to the increasing importance of the proton speed relative to the electron speed for $T_\\textrm{e}< T_\\textrm{p}$, the protons can contribute significantly to the electron-proton bremsstrahlung. They dominate the electron-proton Bremsstrahlung losses for electron temperatures lower than $m_\\textrm{e}/m_\\textrm{p}\\approx 1/1836$ times the proton temperature.  While our results are strictly valid only in the non-relativistic regime, we give a crude extrapolation for the relativistic case. The work presented here should only be considered as exploratory work to outline the influence of a two-temperature plasma on the electron-proton bremsstrahlung. If calculated for both the non-relativistic and relativistic regime, one needs to account for the relativistic kinematics and consider the \\citet{betheheitler} cross section, adjusted for the effect presented here. While it is certainly needed and desirable, such a treatment is beyond the scope of this paper. The apparent dominance of electron-electron bremsstrahlung for electron temperatures in excess of $10^9$ K, however, does not make it that necessary.  The mechanism presented here, seen in the context of Coulomb collisions between electrons and protons, is already well known in plasma physics. The NRL Plasma Formulary \\citep{NRL_FORMULARY_06}, a standard reference in plasma physics for more than 25 years, gives a formula for the Coulomb collision rate which depends only on the electron temperature for $T_\\textrm{e}>m_\\textrm{e}/m_\\textrm{p} T_\\textrm{p}$. This is the classical \\citet{1962pfig.book.....S} result. For higher proton temperatures, the electron temperature dependence weakens and the collision rate only depends on the proton temperature owing to their larger speed compared to the electrons. This effect is seen as well in the work by \\citet{1983MNRAS.202..467S} who derives the Coulomb collision rate in a fully relativistic framework. Similar to our extrapolation for the electron-proton bremsstrahlung rate, they find a change in the exponent of the temperature dependence of $+1/2$, i.e. from $T^{-3/2}$ to $T^{-1}$.  The importance of the effect presented in this paper needs to be examined by comparing the characteristic timescales for Coulomb collisions to equilibrate electron and proton temperature and the Bremsstrahlung timescale. While it may not be of that strong influence for the energetics of the two-temperature plasma, it could have an observable signature due to the higher high-energy cut-off for the Bremsstrahlung created by the kinetic energy of the protons, if it is not hidden behind some more important emission mechanisms at these energies."
    },
    "0609/astro-ph0609575_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "We discussed the expected neutrino flux from $\\gamma$-transparent accelerators of cosmic rays, and emphasized the case of young SNR. We showed that for RX J1713.7-3946 (the best known SNR in TeV sky, thanks to H.E.S.S.) the expectations are stable: $\\sim 5$ events per km$^2$ per year in an ideal detector. The median neutrino energy is $3$ TeV. Since the detected $\\mu$ are softer than the $\\mu$ in the production point (that in turn is softer than the impinging $\\nu$) several events will fall in an energy region where the atmospheric background and the role of imperfect detection efficiency are important: see the lectures of Lipari and Lucarelli. Thus, we believe that it would be desirable to have a detailed discussion of the characteristics of a detector that aims to see a $\\nu$ signal from RX J1713.7-3946.  \\vskip2mm  Sometimes soon H.E.S.S.\\ should tell us more on another intense VHE $\\gamma$-ray source, RX J0852.0-4652 (Vela Jr). This SNR has $F_\\gamma=6.5$ million $\\gamma$-rays $E^{-2.1}/(\\rm TeV\\ km^2\\ yr)$ below $10$ TeV. If the exponential cutoff is at $E_{\\rm \\gamma cut}=50$ $(150)$ TeV (and again, if these $\\gamma$ rays are of hadronic origin) we get $N_\\mu + N_{\\bar\\mu}=10$ $(14)/(\\rm km^2 yr)$ in an ideal detector, with a significantly higher energy~\\cite{vis06}.     \\subsection*"
    },
    "0609/astro-ph0609096_arXiv.txt": {
        "abstract": "We investigate the dynamical interaction between a galactic disk and surrounding numerous dark subhalos as expected for a galaxy-sized halo in the cold dark matter (CDM) models. Our particular interest is to what extent accretion events of subhalos into a disk are allowed in light of the observed thinness of a disk. Several models of subhalos are considered in terms of their internal density distribution, mass function, and spatial and velocity distributions. Based on a series of N-body simulations, we find that the disk thickening quantified by the change of its scale height, $\\Delta z_d$, depends strongly on the individual mass of an interacting subhalo $M_\\mathrm{sub}$. This is described by the relation, $\\Delta z_d/R_d \\simeq 8 \\sum_{j=1}^N (M_{\\mathrm{sub},j}/M_d)^2$, where $R_d$ is a disk scale length, $M_d$ is a disk mass, and $N$ is the total number of accretion events of subhalos inside a disk region ($\\le 3R_d$). Using this relation, we find that an observed thin disk has not ever interacted with subhalos with the total mass of more than 15~\\% disk mass. Also, a less massive disk with smaller circular velocity $V_c$ is more affected by subhalos than a disk with larger $V_c$, in agreement with the observation. Further implications of our results for the origin of a thick disk component are also discussed. ",
        "introduction": "The cold dark matter (CDM) paradigm has become a standard framework for understanding the structure formation in the Universe. According to this theoretical paradigm, the growing process of self-gravitating structures is hierarchical in the sense that small dark matter halos virialize first, and aggregate successively into larger and larger objects. This clustering process of dark matter halos is successful for explaining a wide variety of observations including the large-scale distribution of galaxies.  In this CDM scenario, N-body simulations are an important tool in order to investigate the non-linear growth of cosmological structures. Early N-body simulations based on the CDM models suffered from the so-called {\\it over-merging} problem, i.e., substructures are disrupted very quickly within dense environments (\\cite{summers1995}). However, recent high-resolution N-body simulations have revealed the presence of hundreds of dark matter substructures (subhalos) which survive in not only cluster scales but also galactic scales (\\cite{moore1999}; \\cite{klypin1999}). This large number of subhalos in a galaxy-sized halo is in contrast to only about a dozen satellite galaxies in the Galaxy, which confronts so-called ``the Missing Satellite Problem''. Several authors have argued that this apparent discrepancy could be resolved by considering some suppressing process for star formation, such as gas heating by an intergalactic ionizing background or energy feedback from evolving stars. In whatever models relying on the suppression of galaxy formation, a typical galaxy-sized halo should contain numerous dark subhalos.  Then, there is a possibility that a large amount of subhalos interact frequently with a stellar disk embedded in the center of a halo, so that the disk would be dynamically heated and thickened. On the other hand, an observed galactic disk is rather thin: the scale height (or half thickness) is only about $\\sim 250$ pc in the Galaxy. Likewise, recent observations of external disk galaxies (\\cite{kregel2002}) suggest that the observed scale height of a disk, $z_d$, is confined to some limiting value relative to the scale length of a disk, $R_d$, i.e, $z_d / R_d < 0.2$.  This observed thinness of a disk provides important limits on the disk heating due to infalling satellites. T\\'{o}th \\& Ostriker (1992) analytically evaluated this effect and concluded that an observed disk like that of the Galaxy within the solar radius should have interacted with satellites with no more than 4~\\% of the present disk mass within the last 5 Gyr. Subsequent numerical simulations of an interaction between a disk and a single satellite (e.g. \\cite{velazquez1999}) showed that their analytical estimation for the disk heating was somewhat too high because an actual interaction process is highly non-linear and more complicated than simplified analytical representation. Interactions with many subhalos would be much more complicated and thus require a more detailed analysis.  \\citet{Font01} have conducted numerical simulation of interaction between a disk and numerous subhalos based on the CDM models. They concluded that the effect of subhalos on a disk is rather small, and therefore subhalos do not conflict with the presence of a thin disk since their orbit seldom take them near the disk. However, it is worth noting that in their simulation the initial scale height of a disk (700 pc) is already thick compared with the observed one in the Galaxy ($\\sim250$ pc), thereby leading to possibly the underestimation of the disk heating effect. Their simulation is also limited to only one realization of subhalos; it is yet unclear whether the derived weak effect of subhalos on a disk is general or not. \\citet{ardi2003} have investigated more details in this disk heating by subhalos. They found that a more massive subhalo is more effective to heat the disk than a less massive one. However, in their calculation subhalos are represented by rigid bodies which never lose their mass irrespective of tidal effects of a host galaxy, so that the disk heating is overestimated. Also, the applicability of their result to an actual disk, especially, to what extent accretion events of subhalos into a disk are allowed remains unclear.  Our aim in this paper is thus to set more useful limits on the dynamical interaction between numerous subhalos and a galactic disk. For this purpose, we conduct a series of numerical simulations, in which a self-gravitating disk is embedded in a dark halo containing many subhalos. In this work we set an initially thin disk with the scale height of 250 pc, in contrast to previous numerical studies starting from the scale height of $\\sim 700$ pc much larger than the observed one (\\cite{velazquez1999}; \\cite{Font01}). Several models for the system of subhalos in a host halo are taken into account in terms of their mass function, spatial distribution, and velocity distribution. We also consider two different models for the internal density distribution of subhalos: point-mass and extended-mass models. In the latter model, subhalos are affected by a tidal field of a host galaxy so that they lose their mass in the course of their orbital motions. Based on our simulations, we investigate the dependence of the disk heating on the model parameters and apply our analysis to understanding an observed thin disk in the context of the disk heating by subhalos.  This paper is organized as follows. In \\S~\\ref{sec:model} we describe our galaxy model which is composed of halo, bulge, and disk components. The models of subhalos are also described in this section. In \\S~\\ref{sec:result} we present the results of our numerical simulations. In \\S~\\ref{sec:discussion} we analyze our results and present our prediction for the relation between the disk heating by subhalos and an observed thin or thick disk. Finally, in \\S~\\ref{sec:conclusion} we present our conclusions. ",
        "conclusions": "\\label{sec:conclusion}  We summarize our conclusions as follows:  \\begin{itemize} \\item The dynamical effects of subhalos on a disk are represented by the relation between the change of the disk scale height $\\Delta z_d$ (measured at the disk edge $R=3R_d$) and individual masses of subhalos $M_{\\mathrm{sub}}$, i.e., $\\Delta z_d/R_d \\simeq 8 \\sum_{j=1}^N (M_{\\mathrm{sub},j}/M_d)^2$, where $R_d$ is a disk scale length, $M_d$ is a disk mass, and $N$ is the total number of accretion events of subhalos inside a disk region ($\\le 3R_d$). \\item If subhalos with the total mass of more than 15~\\% disk mass interact with a disk, then the disk thickness is made larger than the observed range. \\item A less massive disk with smaller circular velocity $V_c$ is found to be more affected by subhalos than a disk with larger $V_c$, which is in agreement with the observed properties of a thin disk. \\item Stars in a significantly thickened disk by subhalos appear to be well mixed and show a vertical gradient in their rotation velocity, being similar to the observed properties of the thick disk in the Galaxy. \\end{itemize}  We note that the relation (\\ref{eq:relation2}) we have obtained here is universal and thus useful for the applications to any relevant issues, including the dynamics of an evolving stellar disk at the center of a growing dark halo. Such detailed studies of a galactic disk in comparison with recently increasing datasets of a remote disk galaxy will be of great importance and is left to future work.  \\bigskip The numerical computations reported here were carried out on GRAPE systems (project ID: g05b05) kindly made available by the Astronomical Data Analysis Center (ADAC) at the National Astronomical Observatory of Japan (NAOJ). This work has been supported in part by a Grant-in-Aid for Scientific Research (15540241, 17540210) from the Japanese Ministry of Education, Culture, Sports, Science and Technology."
    },
    "0609/astro-ph0609743_arXiv.txt": {
        "abstract": "The Baikal Neutrino Telescope has been operating in its NT200 configuration since April, 1998. The telescope has been upgraded in April, 2005 to the 10\\,Mton scale detector NT200+. It's main physics goal is the detection of signals from high energy neutrino cascades. NT200+ reaches a 3-year sensitivity of $2 \\times 10^{-7}$cm$^{-2}$s$^{-1}$sr$^{-1}$GeV for an all-flavor diffuse cosmic $E^{-2}$ neutrino flux for energies 10$^2$~TeV~$\\div$~10$^5$~TeV.  Desgin and sensitivity of NT200+ are described. NT200+ is forming the basic building block of a future km3-scale (Gigaton-Volume) Baikal Telescope. Research and development work on that next stage detector has started. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609269_arXiv.txt": {
        "abstract": "We obtained 98 $R$-band and 18 $B$, $r'$, $i'$ images of the optical afterglow of GRB 060526 ($z=3.21$) with the MDM 1.3m, 2.4m, and the PROMPT telescopes in Cerro Tololo over the 5 nights following the burst trigger. Combining these data with other optical observations reported in GCN and the \\swift-XRT observations, we compare the optical and X-ray afterglow light curves of GRB 060526. Both the optical and X-ray afterglow light curves show rich features, such as flares and breaks. The densely sampled optical observations provide very good coverage at $T>10^4$~sec. We observed a break at $2.4\\times10^5$~sec in the optical afterglow light curve. Compared with the X-ray afterglow light curve, the break is consistent with an achromatic break supporting the beaming models of GRBs. However, the pre-break and post-break temporal decay slopes are difficult to explain in simple afterglow models. We estimated a jet angle of $\\theta_j \\sim 7^{\\circ}$ and a prompt emission size of $R_{prompt} \\sim 2\\times10^{14}$~cm. In addition, we detected several optical flares with amplitudes of $\\Delta m \\sim 0.2$, 0.6, and 0.2~mag.  The X-ray afterglows detected by \\swift\\ have shown complicated decay patterns.  Recently, many well-sampled optical afterglows also show decays with flares and multiple breaks. GRB~060526 provides an additional case of such a complex, well observed optical afterglow. The accumulated well-sampled afterglows indicate that most of the optical afterglows are complex. ",
        "introduction": "In the fireball model \\citep{meszaros97,sari98}, the afterglow emission of gamma-ray bursts (GRBs) are thought to be synchrotron emission in the external shocks. After the launch of \\swift\\ \\citep{gehrels04}, with its rapid localization of GRBs and the dedicated on-board XRT instrument, the afterglow models can be tested extensively with the regularly obtained XRT light curves. More than half of the \\swift-XRT light curves show complicated decay patterns with multiple breaks and giant X-ray flares \\citep{burrows05,nousek06,obrien06}.  New ingredients were added to the models to interpret these features \\citep[e.g.,][]{zhang06,nousek06,panaitescu06a}.  Compared with the large number of \\swift-XRT afterglows, only a few bursts have good optical afterglow coverage, which limits the multi-wavelength study of GRB afterglows.  Moreover, a large fraction of the well-studied optical afterglows also show complicated behaviors (e.g., Guidorzi et al. 2005; Blustin et al. 2006; Rykoff et al. 2006; Stanek et al. 2006), challenging the simple, smooth-decay afterglow models.  Another important aspect of the models is that GRBs are thought to be collimated in jets, based on the achromatic breaks observed in many optical GRB afterglow light curves \\citep[e.g.,][]{stanek99}. Different jet models have been proposed for GRBs either under a uniform jet model \\citep[e.g.,][]{rhoads99,frail01,granot02} or a structured jet model \\citep[e.g.,][]{lipunov01,zhang02,rossi02,lloyd-ronning04,zhang04,dai05}. Some of these jet models can also unify the closely related phenomena of X-ray flashes with GRBs \\citep{yamazaki03,zhang04,lamb05,dai05}. Since the distinct signature that a jet imposes on GRB afterglow light curves (an achromatic break) is simply a geometric effect, it is important to test the wavelength independence across the broadest possible wavelength range, for example between optical and X-ray afterglow light curves.  To date, the lack of wavelength dependence has only been confirmed across different optical bands.  Recently, GRB~050525A \\citep{blustin06}, GRB~050801 \\citep{rykoff06a}, and GRB~060206 \\citep{stanek06} show possible achromatic breaks across optical and X-ray light curves.  However, in GRB~050801, the break is interpreted as energy injection or the onset of the afterglow, and in GRB~060206, it is debated whether the break is achromatic \\citep{monfardini06}.  As many XRT light curves show multiple breaks, it is not obvious which of them should be associated with the optical break and which of them interpreted as the jet break. Recently, \\citet{panaitescu06b,fan06} showed that some of the X-ray breaks (1--4 hours after the burst trigger) are chromatic from X-rays to optical bands. However, as the X-ray light curves have several breaks, it is possible that the achromatic jet break occurs at some later time. In addition, as many optical afterglows also show rich features such as flares and multiple breaks (e.g., Stanek et al. 2006), fits to poorly sampled optical light curves may not be reliable.  In this paper, we report the optical follow-up of GRB 060526 with the MDM 1.3m, 2.4m telescopes and the PROMPT at Cerro Tololo, and our detection of an achromatic break across the optical and X-ray bands. GRB 060526 was detected by the BAT on board \\swift\\ at 16:28:30 UT on May 26, 2006 (Campana et al. 2006). The XRT and UVOT rapidly localized the burst location.  The burst was followed up with ground-based telescopes by several groups.  In particular, \\citet{berger06} reported the burst redshift of $z=3.21$. We organize the paper as follows.  First, we describe the data reduction in \\S\\ref{sec:data}.  In \\S\\ref{sec:evo}, we describe the evolution of the GRB afterglow and perform a comparison between the optical and X-ray light curves.  Finally, we discuss our results in \\S\\ref{sec:dis}. ",
        "conclusions": "} We present well-sampled optical and X-ray afterglow light curves of GRB~060526. As discussed in the previous section, the evolution of the afterglow is complicated with multiple breaks and flares both in the optical and the X-ray bands. The combination of flares and incomplete data sampling present severe challenges to measuring the temporal decay slopes of the afterglow, even for a well-sampled burst such as GRB~060526. Below, we proceed by assuming that our analysis results are not significantly affected by these factors.  We detected a possible achromatic jet break in the optical and X-ray afterglow light curves. Before this late-time break ($T\\sim2.4\\times10^5$~sec), the afterglow is consistent with many \\swift\\ afterglows \\citep{nousek06}. The X-ray light curve started with a steep decay ($\\alpha_1 = 2.3\\pm0.2$), which is interpreted as the tail of the prompt emission due to the ``curvature effect'' \\citep{kumar00}.  The spectral and temporal indices are constrained as $\\alpha=\\beta+2$, which is consistent with the X-ray spectral index of $\\beta_1 = 0.55\\pm0.15$.  The X-ray light curve then entered into a shallow decay stage with $\\alpha_2 \\sim 0.5$, which we interpret as energy injection.  Then it decays into a normal afterglow stage with $\\alpha_3 = 1.14\\pm0.02$ as constrained from the optical observations. In addition, the X-ray light curve shows two huge flares which are commonly seen in \\swift\\ X-ray afterglows and are attributed to late time central engine activities.  After the achromatic break, the afterglow enters a very steep stage with $\\alpha_4=3.4\\pm0.2$. The X-ray and optical-to-X-ray spectral indices after $5000$~sec are consistent with $\\sim 1$, suggesting that the optical and X-ray bands are on the same power-law segment of the spectral energy distribution. The optical spectral index is marginally consistent with the X-ray index, although the error-bar is large.  The achromatic break observed in optical afterglows is traditionally interpreted as the jet break. As mentioned in the introduction, the achromatic break is not yet confirmed across optical and X-ray light curves.  Here, we present a case in GRB~060526 where such an achromatic break is observed across the optical and X-ray afterglows, supporting the beaming model of the GRBs.  However, the afterglow decay slopes before and after the break are hard to reconcile with simple afterglow-jet models.  The late time decay slope after the jet break should follow $\\alpha = p$ \\citep{sari99}, where $p$ is power-law index for the electrons $N(\\gamma) \\propto \\gamma^{-p}$.  The post-jet slope, $3.4\\pm0.2$, is too steep for a pre-jet slope of $1.14\\pm0.02$ under any combination of either constant or wind medium and relative positions between $\\nu$, $\\nu_m$, and $\\nu_c$.  Since the achromatic break is most easily explained by a jet, it is possible that more complicated afterglow models are needed \\citep[e.g.,][]{panaitescu06b} with non-standard micro-physical parameters. Another possibility is that energy injection or flares continued contributing significant flux and significantly affected the temporal decay slope. We estimated the jet angle (half opening angle for uniform jets or observer's viewing angle for structured jets) using $t_j \\simeq 6.2(E_{52}/n)^{1/3}(\\theta_j/0.1)^{8/3}$~hr \\citep{sari99} and obtained $\\theta_j \\sim 7^{\\circ}$ assuming ambient density $n=1~$cm$^{-3}$. We further estimated the size of the $\\gamma$-ray prompt emission by combining the measured jet angle and the X-ray tail emission detected before 225~sec using $t_{tail} = (1+z)(R_{prompt}/c)(\\theta_j^2/2)$ \\citep{zhang06} and obtained $R_{prompt}\\sim2\\times10^{14}$~cm.  We also observed multiple optical flares in the light curve of GRB~060526.  Optical flares or re-brightenings have been observed in both pre-\\swift\\ \\citep[e.g., GRB~970508, GRB~021004, and GRB~030329,][]{galama98,lazzati02,bersier03,mirabal03,lipkin04} and \\swift\\ bursts \\citep[e.g., GRB~050525A, GRB~050820A, GRB~060206, GRB~060210, GRB~060605, GRB~060607, and GRB~061007,][]{blustin06,cenko06,stanek06,schaefer06,nysewander06,bersier06}. The fraction of bursts with optical flares seemed small.  However, recently many well-sampled bursts show complex optical decay behaviors.  The accumulating observations argue that it is possible that most of optical afterglows are complex and the appearance of simplicity was a consequence of poor sampling. There are several interpretations for the optical flares, such as models of density fluctuations, ``patchy shell'', ``refreshed shock'', and late central engine activities \\citep[e.g.,][]{jakobsson04,ioka05,gorosabel06}. We estimated the quantities $\\Delta t/t = 0.23\\pm0.08$, $0.75\\pm0.17$, and $0.22\\pm0.11$ and $\\Delta F_{\\nu}/F_{\\nu} \\sim 0.2$, 0.7, and 0.2, respectively, for the three flares detected in GRB~060526. The $\\Delta t/t$ values are small which do not favor the models of patchy shell and refreshed shock, since the model predictions are $\\Delta t/t > 1$ and $>1/4$ for these two models \\citep{ioka05}. The properties of the flares barely satisfy Ioka et al.'s constraint, $\\Delta F_{\\nu}/F_{\\nu} < 1.6 \\Delta t/t$, under the density fluctuation model. Recently, \\citet{nakar06} also modeled the effects of density fluctuations on the afterglow light curves and found that they cannot produce the sharp features observed in many bursts. Another possibility is that the flares (or breaks) indicate the onset of the afterglow \\citep{rykoff06a,stanek06} scaled with the isotropic energy as $T_{onset} \\propto E_{iso}^{1/3}$ \\citep{sari97}. The onset time also depends on the density of the ambient medium and the initial Lorentz factor that are more difficult to measure. We might expect a correlation between $E_{iso}$ and $T_{onset}$ for a large sample of bursts, or if the densities and Lorentz factors for the bursts only spread in a narrow range. We tested this hypothesis by plotting the two properties for bursts with optical flares in Fig.~\\ref{fig:et}, and did not detect positive correlation between the two properties. However, we notice the difference between flares that occur before the optical afterglow has decayed and those that occur afterward. The flares in GRB~050820A, GRB~060210, GRB~060605, GRB~060607, and GRB~061007 possibly belong to the category of flares that occur before the afterglow has faded, and they roughly follow the scaling between isotropic energy and flare time. However, a larger sample is needed to fully test the model. In any case, the flares in GRB~060526 are unlikely to be associated with the onset of the afterglow. It is possible that these flares are from late central engine activities, which can have arbitrary variabilities.  However, we are open to other theoretical models which can be tested extensively with our well-sampled light curve."
    },
    "0609/astro-ph0609082_arXiv.txt": {
        "abstract": "We follow up on our (Radhakrishnan \\& Deshpande, 2001: RD01) radically different interpretation of the observed structures and morphologies in the x-ray observations of the nebulae around young pulsars (PWNe). In our general model for PWNe (RD01), originally motivated by the Chandra observations of the Vela X-ray nebula, the bright arcs, the jet-like feature and the diffuse components in such nebulae can be explained together in detail, wherein the arcs are understood as traces of the particle beams from the two magnetic poles at the shock front. We consider this as important evidence for collimated particle beams from pulsars' magnetic poles. In this paper, we discuss the variability in the features in the Vela X-ray nebula observed by Pavlov et al. (2003), and assess the relevance and implication of our model to the observations on the Crab and other remnants. Our basic picture after incorporating the signatures of free precession of the central compact object can readily account for the variability and significant asymmetries, including the bent jet-like features, in the observed morphologies. The implications of these findings are discussed. ",
        "introduction": "} The superb capabilities of the Chandra telescope have over the last 6 years revealed several spectacular images in X-rays of the nebulae surrounding pulsars, highlighting and resolving the various spatial structures richly loaded with information about many aspects. While most, if not all, of these images bear a remarkable commonality that is emphasized by their overall symmetric morphology about what is readily identified as the direction of the projected rotation axis, the forms and proportions of the components appear to differ significantly.  An overwhelming majority of the observers and theorists interpreting these observations seem to suggest and endorse the following basic picture. The jet-like features nearly along the symmetry axis, bisecting the arcs and the diffuse glow spread about them, are identified with collimated outflows of relativistic particles along the spin axis of the central compact remnant, a pulsar. The two arc-like features lie along circular rings highlighting shocks in which the energy of an outflowing equatorial wind is dissipated to become the source of synchrotron emission for the compact nebula and the incompleteness of the rings is attributed to preferential Doppler boosting of the emission in the forward direction. The two rings, if apparent, straddle the equator symmetrically, and the deficit of emission exactly in the equatorial plane is related to the fact that this is where the direction of a toroidally wrapped magnetic field changes sign i.e. the field may vanish there.  However, Radhakrishnan \\& Deshpande (2001, hereafter RD01) have suggested an alternative interpretation that differs from this mainstream model {\\it in practically every aspect}, particularly those regarding the arcs and the jet-like features. The bright arcs, the jet-like feature and the diffuse components (for example, in the Vela X-ray nebula) are explained together in detail by our ``Rotating Vector model\" in which the arcs are understood as traces of the particle beams from the two magnetic poles at the shock front. We consider this as important evidence for collimated particle beams from pulsars, a point that we find necessary to reemphasize.  In this paper, we follow up on the RD01 model to address the new clues provided by the variability of various features, such as those in the Vela X-ray nebula observed by Pavlov et al. (2001, 2003), and assess the relevance and implication of our model to the observations on other remnants, such as the Crab. In the next section, we begin with a brief summary of our model (RD01) for the x-ray nebulae around pulsars. In section 3, we revisit the Vela nebula story, now incorporating signatures of free precession of the central compact object, to explain the observed variability. We try to model, in section 4, the observed morphology of the X-ray nebula surrounding the Crab pulsar, and the significant asymmetries, including the bent jet-like features. The implications of these findings are discussed in the last section.   \\noindent{ ",
        "conclusions": "}  The cavity shapes suggested by our modeling of the observed x-ray nebulae might appear to differ significantly in the two cases (namely, the Vela and the Crab) discussed here. The part of the cavity profile sampled by the particle beams from the Crab pulsar appears to be significantly shallower than that in the case of the Vela pulsar. Given the age and the proper motion of the Vela pulsar, the relative dimensions of its x-ray arcs appear to favour an hour-glass shaped cavity, rather than an ellipsoid. The visibility of the jet-like feature apparently extending well beyond the diffuse component is not inconsistent with both of these components sharing a common origin, i.e. the latitudinal spread of a small fraction of the otherwise collimated particle flow, since the apparent relative brightness of the two components could differ significantly for intrinsic reasons, as well as those dictated by the viewing geometry.  In our picture, the observed extents of the jet-like features, after due accounting of the projection effect, provide lower limits for the dimension of the cavity in the latitudinal direction, which appears to exceed the equatorial dimension, at least in the Vela case. The apparent asymmetry in the `jet' and the `counter-jet' extents in the Vela nebula is consistent with our expectation of the associated radiation being symmetric with respect to the direction of the magnetic axis, modified by the viewing geometry. If the jet-like emission were to be due to any physical flow of matter at relativistic speeds along the rotation axis (e.g. as estimated by Pavlov et al., 2003), the counter-jet should have been more prominent, both in its extent and its brightness, given that it would be pointing closer to our sightline. This is definitely not what is observed, and in fact, this aspect was noted as a cause for concern by Pavlov at al. (2003). Further, such a physical jet flow cannot escape its dissipation and consequent termination at the relevant region of the cavity wall where a cap-like emission feature should have been apparent, but is not observed. In their `physical jet' picture, the jet instabilities leading to bending etc. are expected to be due to and along the proper motion of the star. The Crab case is clearly inconsistent with this expectation, implying at the least that the apparent bending has little to do with the motion of the star. These inconsistencies and the other aspects discussed by RD01 argue strongly against a physical jet along the rotation axis.  Apparent bends in the jet-like features in general, and particularly those with anti-symmetry about the equatorial plane, are interpreted in our model as due to possible free-precession. Its successful application to the Crab and the Vela pulsars, differing in the viewing geometry and rotation history, is not a chance coincidence. It would not surprise us therefore, if the bent jet-like feature in the x-ray nebula around pulsar B1509-58 (Gaensler et al., 2002) would have a similar interpretation. The variability of these and other features is then a natural consequence of such a rotational history, as illustrated by the multi-epoch observations of the Vela x-ray nebula. The variability seen in the Crab nebula at optical wavelengths would be another illustration of the same effect, though on much longer time-scales. If our assumed precession period of 14 years for the Crab pulsar is correct, a significant change in the orientation of the jet-like feature should be expected in the coming years. As commented earlier, the yet unexplored interplay between the glitches and possible precession would determine how the rotation/precession history evolves with time. Hence it is not at all clear whether free-precession, if any, would be unaffected through the glitch episodes, and even if it does, whether any clear signature of free-precession of the star would be apparent from the radio pulsar timing residuals obtained after fitting for period glitches and the associated recoveries commonly observed in young pulsars. It remains to be seen if any significant variation in the shape and intensity of the radio pulses from the Crab pulsar reveals a signature consistent with the rotational history of the star suggested by the above mentioned observations at high energies.  There have been suggestions of free-precession in a few pulsars, but they are based on apparent changes in pulse shapes and intensities (e.g., DM96; Shabanova et al., 2001), or in pulse arrival times (Stairs et al., 2000). We wish to point out that a direct and independent way of probing any changes in the orientation of the rotation axis of the star, as in precession, is through monitoring of the polarization position angle (PA) sweep across the radio pulses, and measuring systematic changes, if any, particularly in the PA sweep-rate (Hari Dass \\& Radhakrishnan, 1975). We have indeed begun recently such a monitoring of the Vela pulsar using the Giant Meter-wave Radio Telescope (GMRT) in India.  To summarize, we find compelling evidence for collimated particle beams from pulsars, and for free-precession of the Vela and the Crab pulsars, based on the Chandra observations of the respective x-ray nebulae and their apparent systematic variability. Quantitative estimates of the intensities \\& the locations of the different elements of the nebulae and their variability, when compared in the framework of our simple interpretation, should pave the way to important clues on the properties of the cavities created by the pulsars, and the energy spectra associated with the collimated particle beams."
    },
    "0609/astro-ph0609561_arXiv.txt": {
        "abstract": "Young, rapidly rotating neutron stars could accelerate protons to energies of $\\sim 1$ PeV close to the stellar surface, which scatter with x-rays from the stellar surface through the $\\Delta$ resonance and produce pions. The pions subsequently decay to produce muon neutrinos. We find that the energy spectrum of muon neutrinos consists of a sharp rise at $\\sim 50$ TeV, corresponding to the onset of the resonance, above which the flux drops as $\\epsilon_\\nu^{-2}$ up to an upper-energy cut-off that is determined by either kinematics or by the maximum energy to which protons are accelerated. We predict event rates as high as 10-50 km$^{-2}$ yr$^{-1}$ from relatively young, close neutron stars. Such fluxes would be detectable by IceCube. ",
        "introduction": "A new window in high-energy astronomy is opening as existing neutrino detectors are improved and new ones are developed. Since neutrinos produced in astrophysical systems are unimpeded by interstellar matter on their way to Earth, detections will provide a new way to study the highest energy phenomena in the Universe. Astrophysical neutrinos are expected to arise in many environments in which neutrinos are produced by the decay of pions created through hadronic interactions ($pp$) or photomeson production ($p\\gamma$). Neutrinos may be produced by cosmic accelerators, like those in supernova remnants \\cite{pro98}, active galactic nuclei \\cite{lm00}, micro-quasars \\cite{dist02} and gamma-ray bursts \\cite{wb97,dai01}. To detect these neutrinos, several projects are underway to develop large-scale neutrino detectors under water or ice. AMANDA-II \\cite{ah04}, in the South Pole, and Baikal \\cite{ay06} are the two neutrino telescopes currently running, whereas ANTARES \\cite{ca03} and NESTOR \\cite{tz03} are under construction in the Mediterranean Sea. Those telescopes belong to a first generation with instrumented volume smaller than 0.02 $\\rm km^3$. The IceCube detector \\cite{ha06}, with a volume of about 1 $\\rm km^3$, is under construction on the same site of AMANDA. The NEMO project \\cite{mig06} is in its starting phase, and will be a cubic kilometer size detector located at Capo Passero, Southern Italy.  As neutrino astronomy comes of age, it is important to get some idea of what the sources might look like to aid in their detection. Recently, we proposed \\cite{lb05,lb06} that young ($t_{age}\\lesssim 10^5 \\rm yr$) and rapidly-rotating neutron stars could be  intense neutrino sources. Here we summarize this work. ",
        "conclusions": "Large-area neutrino detectors use the Earth as a medium for conversion of a muon neutrino to a muon, which then produces \\v{C}erenkov light in the detector. The conversion probability in the Earth is $P_{\\nu\\mu \\rightarrow \\mu} \\simeq 1.3 \\times 10^{-6} (\\epsilon_{\\nu}/\\rm 1\\,\\, TeV)$, where $\\epsilon_{\\nu}$ is the energy of the incident muon neutrino \\cite{ghs95}. The muon event rate is \\begin{equation} \\frac{dN}{dAdt} = \\int d\\epsilon_\\nu \\frac{d\\phi_\\nu}{d\\epsilon_\\nu} P_{\\nu_\\mu \\rightarrow \\mu}. \\label{rate} \\end{equation} Estimated conversion probabilities for the Crab, Vela and 6 other pulsars are given in Table 1 (column 8) for a characteristic acceleration length $L=0.1$. The final column gives the estimated event rates. Complete consideration of the kinematics gives event rates are a factor of $\\sim 10-30$ lower than estimated in our previous paper \\cite{lb05}.  Neutrinos are produced at relatively high rates only if the protons are accelerated through the resonance close to the star. In this case, we obtain integrated count rates of several to $\\sim 50$ km$^{-2}$ yr$^{-1}$.  Such count rates should be easily detected by IceCube, and possibly by AMANDA-II or ANTARES with integration times of about a decade (IceCube is planned to have replaced AMANDA-II by then) for depletion factors of $f_d\\simeq 1/2$. While the characteristics of the spectrum presented here are robust, we caution that the events rates we obtain are very rough upper limits, subject to many uncertainties. For example, we have assumed that the neutrinos are beamed into the same solid angle as the radio beam, which might not be a correct assumption. The radio beam is thought to be produced at about $10R$ \\cite{cor78}. In our model, the pions are produced much closer to the star. They then propagate to $\\sim 1000R$ before decaying to neutrinos. At this distance from the star, the field is unlikely to be dipolar, and it is difficult to say anything definite about the distribution of pion velocities in this region. If the neutrinos form a beam, the beam may be more or less collimated than the radio beam. If it is more collimated, the neutrino event rates could be higher than estimated here.  Results of 807 d of data from AMANDA-II are now available \\cite{gr05}. AMANDA-II has detected 10 events (over a background of 5.4) from the direction of the Crab pulsar, with energies higher than 10 GeV. This result, though intriguing, is not statistically significant; IceCube will be able to confirm or refute this result. While it would be more exciting to see neutrinos from pulsars, the accumulation of null results over the next decade would be interesting as well; it would probably mean that photomeson production is ineffective or non-existent in the neutron star magnetosphere, thus providing a bound on the accelerating potential that exists near the neutron star surface."
    },
    "0609/astro-ph0609611_arXiv.txt": {
        "abstract": "The high-latitude Galactic H~I cloud toward the extragalactic radio source 3C~225 is characterized by very narrow 21~cm emission and absorption indicative of a very low H~I spin temperature of about 20~K. Through high-resolution optical spectroscopy, we report the detection of strong, very narrow Na~I absorption corresponding to this cloud toward a number of nearby stars.  Assuming that the turbulent H~I and Na~I motions are similar, we derive a cloud temperature of 20$^{+6}_{-8}$~K (in complete agreement with the 21~cm results) and a line-of-sight turbulent velocity of 0.37~$\\pm$~0.08~~km~s$^{-1}$ from a comparison of the H~I and Na~I absorption linewidths.  We also place a firm upper limit of 45~pc on the distance of the cloud, which situates it well inside the Local Bubble in this direction and makes it the nearest-known cold diffuse cloud discovered to date. ",
        "introduction": "Recent H~I 21~cm absorption surveys have revealed a significant population of cold (T~$<$~40~K) diffuse clouds in the Galactic interstellar medium (ISM) \\citep{hei03,gib05,kav05}.  Two of the most remarkable examples of such clouds were actually discovered long ago by \\citet{ver69} through their very narrow 21~cm emission.  He found two closely-spaced, degree-sized patches of H~I gas at $l\\sim226\\arcdeg,b\\sim+44\\arcdeg$ (``Cloud A'') and $l\\sim236\\arcdeg,b\\sim+46\\arcdeg$ (``Cloud B'') each with line widths indicative of kinetic temperatures below 30~K.  \\citet{kna72} subsequently mapped the 21~cm emission of these clouds at higher resolution and derived respective H~I spin temperatures of $\\approx$24~K and $\\approx$17~K for Cloud~A and Cloud~B.  In another approach, \\citet{cro85} measured spin temperatures of $\\approx$20~K amd $\\approx$14~K from the Cloud~A and Cloud~B 21~cm absorption toward the extragalactic radio sources 3C~225 and 3C~237.  As part of their Millenium Arecibo 21~cm Absorption-Line Survey, \\citet{hei03} revisited these two cold clouds in the context of new observations of 3C~225 and 3C~237.  Utilizing the Leiden-Dwingeloo 21~cm sky survey data \\citep{har97}, they found that Cloud~A and Cloud~B are the predominant components of a narrow ($\\approx$2$\\arcdeg$), broken ribbon of cold H~I gas stretching over 20$\\arcdeg$ across the constellation Leo.  Their observations of 3C~225a,b (6.3$\\arcmin$~separation) yield spin temperatures of 22~K and 17~K and H~I column densities of $1.7\\times10^{19}$~cm$^{-2}$ and $3.2\\times10^{19}$~cm$^{-2}$ for the intervening Cloud~A gas. Assuming thermal pressure equilibrium under these conditions leads to an extremely thin, sheetlike geometry for the cloud \\citep{hei03}.  A more detailed physical assessment including the specific aspect ratio, actual size, and mass of the cloud is limited by the unknown cloud distance. In this {\\it Letter}, we present the first optical absorption-line study of the cold Leo clouds and conclusively show that Cloud~A is located well inside the Local Bubble. ",
        "conclusions": "Figure~3 illustrates the strength of the Na~I absorption at v$_{LSR}$~$\\approx$~3~km~s$^{-1}$ toward our program stars as a function of sky position relative to the 21~cm emission contour of the cold H~I gas \\citep{hei03} and the location of 3C~225 and 3C~237. It is clear from Figure~3 that the Na~I absorption is an excellent tracer of the H~I emission for this cold gas.  In the case of comparisons with the 21~cm absorption observed toward 3C~225a,b, four stars with matching-velocity Na~I absorption (HD~84194, HD~84182, HD~83683, and HD~83509) are located within 1$\\arcdeg$ of this radio source.  The mean linewidth of this Na~I absorption (and the standard error of the mean) is b~$=$~0.54~$\\pm$~0.10~km~s$^{-1}$ (0.90~$\\pm$~0.17~km~s$^{-1}$~FWHM) as compared with the 1.3~km~s$^{-1}$~FWHM H~I linewidth measured toward both 3C~225a and 3C~225b by \\citet{hei03}.  Assuming that the turbulent H~I and Na~I motions are similar, a comparison of the H~I and Na~I linewidths through the expression $b^2=(2kT/m)+2v_t^2$ yields a cloud kinetic temperature of 20$^{+6}_{-8}$~K and a one-dimensional rms turbulent velocity (v$_t$) of 0.37~$\\pm$~0.08~km~s$^{-1}$.  Thus, the Na~I measurements are completely consistent with the remarkably cold temperatures previously derived for Cloud~A from the H~I 21~cm spin excitation toward 3C~225.  The Na~I observations also provide a stringent constraint on the distance of Cloud~A.  The two closest program stars within the Cloud~A H~I contour are at {\\it Hipparcos} distances \\citep{per97} of 42.1$^{+1.8}_{-1.6}$~pc (HD~83683) and 42.8$^{+1.9}_{-1.7}$~pc (HD~85091) and they both exhibit strong Cloud~A Na~I absorption. Another nearby star at a distance of 41.5$^{+1.7}_{-1.6}$~pc (HD~83808) just outside the Cloud~A contour exhibits very weak Na~I absorption at the appropriate velocity.  The only program star closer than these objects is at a distance of 39.2$^{+1.5}_{-1.4}$~pc (HD~80218) and it lies just outside the H~I contour of the lowest-longitude cold Leo cloud.  Consequently, the lack of Na~I absorption toward HD~80218 is not conclusive with respect to the distance of this cloud.  Although it is likely that the three cold Leo clouds illustrated in Figure~3 are at the same distance based on their similar velocities, narrow linewidths, and close position on the sky, the Na~I measurements are not definitive in this regard.  Nevertheless, the Na~I absorption observed toward HD~83683, HD~85091, and HD~83808 does allow us to place a firm upper limit of 45~pc on the distance of Cloud~A.  With this limit, Cloud~A is not only one of the coldest diffuse clouds discovered to date, it also becomes the nearest-known cold interstellar cloud to the Sun.  The proximity of Cloud~A is especially interesting in that it places this cold cloud well inside the Local Bubble of hot ($\\sim$10$^6$~K), tenuous (n$_H$~$\\sim$~0.01~cm$^{-3}$) gas surrounding the Sun out to distances of $\\sim$100~pc or more \\citep{cox87}.  The Local Bubble is believed to have originated through multiple supernovae explosions over the past $\\sim$5 to $\\sim$15~million years \\citep{smi01,bre06}. Utilizing Na~I absorption measurements toward $\\approx$1000 early-type (earlier than A5) stars with {\\it Hipparcos} distances, \\citet{lal03} have mapped the neutral gas ``edge'' of the Local Bubble across the sky.  In the direction of the cold Leo clouds, they place this boundary at a distance of about 100~pc.  Since our star sample in the Leo field of Figure~3 is much larger (33) than that (2) of Lallement et al.\\ (due in part to our inclusion of objects with later spectral types), we can explore this boundary in greater detail.  Among the 17~program stars with d~$>$~157~pc, nine exhibit interstellar Na~I absorption at LSR velocities (ranging from $-$11.5 to $-$1.6~km~s$^{-1}$) distinct from those of the cold Leo clouds whereas none of the 16 stars with d~$<$~157~pc display such absorption.  As illustrated in Figure~1, such distance distinctions are also apparent when the Na~I spectra are compared in localized sky regions.  Allowing for patchiness in the more distant Na~I absorption and the larger stellar distance uncertainties past $\\approx$100~pc, the nearest neutral gas clouds beyond Cloud~A are at a distance between 100 and 150~pc.  Thus, Cloud~A is indeed closer to the Sun than it is to the neutral gas boundary of the Local Bubble.  Within the Local Bubble, there are known to be a number of warm, partially-ionized clouds which collectively have an average temperature of 6700~K and a mean thermal pressure (P/k) of 2300~K~cm$^{-3}$ \\citep{red04}.  Although the thermal pressures of these clouds appear to be below that of the surrounding hot gas in the Local Bubble, they provide a benchmark in considering the physical characteristics of Cloud~A.  Assuming that Cloud~A is in thermal pressure equilibrium at 2300~K~cm$^{-3}$, the 20~K temperature and mean N(H~I) of $2.5\\times10^{19}$~cm$^{-2}$ toward 3C~225 yield a cloud thickness of $2.2\\times10^{17}$~cm or 0.07~pc (this equilibrium thickness would be $\\approx$4~times less at the hot gas thermal pressure). If Cloud~A is at a distance of 40~pc, its $2\\arcdeg\\times7\\arcdeg$ H~I angular extent would then correspond to a physical size of $1.4\\times4.9$~pc, a very thin length-to-width-to-thickness aspect ratio of 70:20:1, and an H~I mass of about 1.4~solar masses.  At the other extreme, if we assume that Cloud~A is at a distance ($\\approx$2~pc) where its projected width is equal to its equilibrium thickness of 0.07~pc, the cloud size would be $0.07\\times0.25$~pc with an H~I mass of about 3.6~Jupiter masses. In any case, given our measured v$_t$ of 0.37~km~s$^{-1}$, the line-of-sight turbulent crossing time over the equilibrium thickness of 0.07~pc is $1.8\\times10^5$~years.  In other words, it appears that Cloud~A is a much more transient phenomenon than the Local Bubble and thus, is likely to have originated somehow in this seemingly hostile environment to cold clouds.  Recently, several theoretical studies have proposed the formation of cold clouds with filamentary or sheet-like geometry through converging flows of warm interstellar gas \\citep{aud05,vaz06}. In particular, Vazquez-Semadeni et al.\\ find that they can come close to reproducing the \\citet{hei03} description of Cloud~A through transonic compression in colliding warm gas streams. They acknowledge that their timescale to produce this cold structure is rather long ($\\sim$1~Myr) but may be consistent with the observations if Cloud~A is overpressured and v$_t$ reflects the gas flow rather than turbulent motions.  Such models are promising but will need to consider other factors such as magnetic fields and certainly thermal conduction in order to provide a full explanation of the nature and evolution of Cloud~A inside the hot Local Bubble.  For example, \\citet{sla06} has shown that the evaporative and condensation time scales for such cold clouds are measured in millions of years even if immersed in a hot surrounding medium.  It is also possible that the hot gas in the Local Bubble may be less pervasive than generally thought, especially given recent results suggesting significant foreground contamination by heliospheric soft X-ray emission \\citep{lal04}.  In any case, the size of Cloud~A on the sky and its isolated nature make it ideal for further optical/UV absorption-line and H~I 21~cm emission-line studies of its distance, small-scale structure, velocity fields, and other physical characteristics such as dust depletion, electron density, and radiation field. A better understanding of the origin and survival of this nearby 20~K cloud inside the Local Bubble will almost certainly have a bearing on interpreting the continuing discoveries of such cold diffuse clouds elsewhere in the Galactic ISM."
    },
    "0609/astro-ph0609427_arXiv.txt": {
        "abstract": "In the last few years, radio detection of cosmic ray air showers has experienced a true renaissance, becoming manifest in a number of new experiments and simulation efforts. In particular, the LOPES project has successfully implemented modern interferometric methods to measure the radio emission from extensive air showers. LOPES has confirmed that the emission is coherent and of geomagnetic origin, as expected by the geosynchrotron mechanism, and has demonstrated that a large scale application of the radio technique has great potential to complement current measurements of ultra-high energy cosmic rays. We describe the current status, most recent results and open questions regarding radio detection of cosmic rays and give an overview of ongoing research and development for an application of the radio technique in the framework of the Pierre Auger Observatory. \\vspace{1pc} ",
        "introduction": "About 40 years ago, Jelley et al. \\cite{Jelley1965} measured pulsed radio emission originating from extensive air showers (EAS) for the first time. As the radio technique proved to be too difficult to handle with the technical limitations of the 1960s and 1970s, however, the interest in radio detection of cosmic rays diminished completely within the following decade. It was only recently, now having powerful digital technology at our disposal, that the concept of radio detection of cosmic rays experienced its renaissance \\cite{FalckeGorham2003}. By now a number of new projects dedicated to the measurement of radio emission from EAS, most prominently the LOPES project \\cite{HornefferArena2005,Horneffer2006} and the CODALEMA project \\cite{Codalema}, have been established.  The radio technique for measuring cosmic rays has a number of merits in its own right: it can watch for EAS with nearly 100\\% duty cycle, even in populated areas, and it measures a bolometric signal that is only very slightly attenuated in the atmosphere, thus allowing the observation of highly inclined showers. But naturally, the most interesting application is to combine the technique with other detection methods, in particular ground-based particle detector arrays and air fluorescence measurements. Each of these techniques yields different observables, and a combination of the methods allows so-called ``hybrid detection'' of cosmic rays, yielding much more information than the individual techniques alone.  In this article, we review the goals, the status and the results so far gathered within the LOPES project followed by an outlook on the application of the radio technique on large scales for the measurement of ultra-high energy cosmic rays. ",
        "conclusions": "In the last few years, radio detection of cosmic ray air showers has once again become a very active field of research. LOPES, as one of the projects studying radio emission from EAS with modern digital technology, has made important contributions. Radio emission from EAS is proven to be coherent at 40--80~MHz, and the field strength correlation with geomagnetic angle strongly indicates that the radiation is of dominantly geomagnetic origin. Aspects such as the absolute field strength, the detailed angular correlations and the scale radius of the lateral dependence will soon be analysed with the new LOPES30 data. Our next-generation Monte Carlo code calculating the radio emission based on highly realistic CORSIKA-based air showers will also allow many new studies with unprecedented detail. Unlike 40 years ago, the study of radio emission today is making good progress.  A possible application of the radio technique on large scales for hybrid detection of ultra-high energy cosmic rays is now being studied in the framework of the Pierre Auger Observatory. These activities are still in a relatively early phase, but with the experience gained in the LOPES and CODALEMA experiments, the challenges to unlock the great potential of the radio technique on large scales can be tackled.     \\vspace{0.4cm}  {\\small\\noindent Acknowledgements: LOPES was supported by the German Federal Ministry of Education and Research (Verbundforschung Astroteilchenphysik) and is part of the research programme of the Stichting voor Fundamenteel Onderzoek der Materie (FOM), which is financially supported by the Nederlandse Organisatie voor Wetenschappelijk Onderzoek (NWO). The KASCADE-Grande experiment is supported by the German Federal Ministry of Education and Research, the MIUR of Italy, the Polish Ministry of Science and Higher Education and the Romanian National Academy for Science, Research and Technology.}"
    },
    "0609/astro-ph0609388_arXiv.txt": {
        "abstract": "{ We study the distribution of  fermionic dark matter at the center of galaxies using  NFW, Moore and isothermal density profiles and show that dark matter becomes degenerate for  particle masses of a few {\\rm keV} and for  distances less than a few parsec from the center of our galaxy. A compact degenerate core forms after galaxy merging and boosts the growth of supermassive black holes at the center of galaxies. To explain the galactic center black hole of mass of $\\sim 3.5 \\times 10^{6}M_{\\odot}$ and a supermassive black hole of $\\sim 3 \\times 10^{9}M_{\\odot}$ at a redshift of 6.41 in SDSS quasars, we require a degenerate core of mass between $3 \\times 10^{3} M_{\\odot}$ and $3.5 \\times 10^{6}M_{\\odot}$. This constrains the mass of the dark matter particle between   $0.6 \\ {\\rm keV}$ and $82~{\\rm keV}$. The lower limit on the dark matter mass is improved to ~{\\rm 7~keV} if exact solutions of Poisson's equation are used in the isothermal power law case. We argue that the constrained particle could be the long sought dark matter of the Universe that is interpreted here as a sterile neutrino. ",
        "introduction": "Precision observations of the cosmic microwave background and of large scale structure confirm the picture in which 96\\% of the matter density of the Universe is made of dark energy and dark matter (DM), which could be revealed only by their gravitational interaction (Spergel  et al. \\cite{spergel03}). The nature of these two forms of matter is still unknown. Many candidates have been proposed for DM. These include Cold Dark Matter (CDM) particles of masses heavier than $1 {\\rm GeV/c^{2}}$ (Bertone et al. \\cite{bertone05}) and Warm Dark Matter (WDM) particles such as sterile neutrinos (Dodelson \\& Widrow \\cite{dw94}). Recently, there has been a renewed interest in sterile neutrinos as candidates for DM as they could be a natural extension of the minimal  standard model (MSM) of electroweak interactions (Weinberg \\cite{weinberg67}; Glashow \\cite{glashow61}). One of these sterile neutrinos could be the DM, while the other two could help explain baryogenesis (Asaka \\& Shaposhnikov \\cite{asaka05}; Shaposhnikov \\cite{shapo06}).  Numerical cosmological N-body simulations suggest DM density profiles which follow  $\\rho \\sim r^{-\\gamma}$ law, with $\\gamma \\approx 3$ in the outer parts of the halos and $1 \\stackrel {\\textstyle <}{\\sim}   \\gamma  \\stackrel {\\textstyle <}{\\sim} 2$ inside  a few kpc (Navarro et al. \\cite{nfw97}; Moore et al. \\cite{moore98}, \\cite{moore99}; Klypin  et al. \\cite{klypin02} and Power et al. \\cite{power03}). Although the Milky Way is well studied in the inner 3 - 10 kpc region, little is known about the DM distribution on smaller scales, i.e. $r \\stackrel {\\textstyle <}{\\sim}0.1 \\ {\\rm pc}$ where the gravitational potential of baryons dominates over DM. In addition, there is a mounting evidence that  a black hole of mass $M \\sim 3.5 \\ 10^{6}M_{\\odot}$ dominates  the mass distribution in the inner one parsec of the Galaxy (Sch\\\"odel et al. \\cite{schodel03}; Ghez et al. \\cite{ghez05}). Thus the investigation of the  distribution of DM around the black hole in the Galactic center is important as it could give some hints on the nature of the  DM particle (Gnedin \\& Primack \\cite{gnedin04};  Bertone \\& Merritt \\cite{bm05}). Recently, we  have established that the  inner DM density profile scales as $\\rho \\sim r^{-3/2}$ under the assumption  that the black hole feeds from degenerate fermionic DM (Munyaneza \\& Biermann \\cite{mb05}). It was then shown that a DM particle of order of $m_{s} \\sim 10$ keV could explain the growth of supermassive black holes of $10^{6}$ to $10^{9}M_{\\odot}$ from stellar seed black holes. It is interesting to note that a  sterile neutrino mass in overlapping mass range,  with a small mixing angle with the active neutrinos, has  also been suggested by Kusenko (\\cite{kusenko04}) to explain the high velocities of pulsars at birth in supernovae. Moreover, a DM particle mass in the keV range  can be constrained from  X-ray background studies (Drees \\& Wright \\cite{dw00};  Abazajian et al.  \\cite{abaza01} and Dolgov \\& Hansen \\cite{dh02}). Here, it is worth mentioning that  an upper limit of  a sterile neutrino mass of $6.3 \\ {\\rm keV}$ from X-ray background seems to be in conflict with the lower limit of $14 \\ {\\rm keV}$ from SDSS Lyman-alpha forest (Abazajian \\& Koushiappas \\cite{ak06}). However, the non-thermal phase distribution of the DM particles has the potential to modify the Lyman alpha limits, while leaving the X-ray limits intact, so eliminating the contradiction with this additional degree of freedom. Depending on the specific model for the production of the DM particles, their initial phase space distribution is possibly largely sub-thermal. Recently, Biermann \\& Kusenko (\\cite{bk06}) established that the decay of a such sterile neutrino could speed up the formation of molecular hydrogen and boost the early star formation and reionisation in agreement with the  WMAP 3-year results (Spergel et al.  \\cite{spergel06}).  The purpose  of this Letter is to study the constraints on the DM particle mass in the central region of the galaxy. Given that there is still uncertainty about the DM profile around the central black hole, we will assume standard Navarro-Frenk-White (NFW), Moore or isothermal gas sphere profiles and investigate what happens if the  DM particles become degenerate and form a degenerate dark matter star also  called  a fermion ball (Munyaneza \\& Viollier \\cite{mv02}; Munyaneza \\& Biermann \\cite{bm05}). We  then get the lower limits on the DM particle mass by assuming that the mass of the fermion ball  cannot be more than the mass of  $3.5 \\times 10^{6}M_{\\odot}$ for the Galactic center black hole. An upper limit to the DM mass will be established from the condition that the mass of the fermion ball should be greater than  about $3 \\times 10^{3}M_{\\odot}$. The choice of this mass comes from the growth mechanism arguments to form supermassive black holes of mass of $10^{9}M_{\\odot}$ in SDSS quasars (Munyaneza \\& Biermann \\cite{mb05}). Moreover, a mass of $3 \\times 10^{3} M_{\\odot}$ is the upper limit of masses for which there is no black hole in galaxies such as M33 (Gebhardt et al. \\cite{geb01}, \\& Barth et al. \\cite{barth05}). Thus the Galactic center provides us with a  fertile testing ground for DM theories and cosmological structure evolution. ",
        "conclusions": "We have determined the constraints on the DM mass under the assumption that the NFW, Moore or isothermal density profiles  become degenerate near the center. Assuming that the mass of the degenerate fermion ball is between about $3 \\times 10^{3}M_{\\odot}$  and $3.5\\times 10^{6}M_{\\odot}$, we have found that the mass of the DM particle should have a lower limit of about 7 keV in the isothermal power law case. The  above limit on the DM mass is obtained from a simple model which does not take into account baryonic matter and a proper model for DM with anisotropic distribution in phase space (Evans \\& An \\cite{evans06};  Hansen \\& Stadel \\cite{hs06}) and baryonic matter would modify these limits. The obtained mass range is in full agreement with the Tremaine-Gunn lower bound of about 0.5 {\\rm keV }on the mass of any fermionic DM when applied to dwarf spheroidal galaxies (Tremaine \\& Gunn \\cite{tg79}). The constraints obtained  on DM particle mass  overlap with the upper limit on sterile neutrino masses of about $8 \\ {\\rm keV}$ to obey the X-ray emission constraints from the Virgo cluster observations (Abazajian \\cite{abaza05}). In addition,  a sterile neutrino mass in the mass range of 2-20 keVs  was derived to explain the high velocities up to 1000 km/s experienced by pulsars at birth  in supernovae  (Kusenko \\cite{kusenko04}; Barkovitch et al. \\cite{barko04} and Fuller et al.  \\cite{fuller03}). Moreover, the obtained constraints agree well with the  lower bound of $m_{s} \\approx 2 \\  {\\rm keV}$ on the mass of sterile neutrinos obtained from the analysis of the Lyman- $\\alpha$ forest data (Hansen et al. \\cite{hansen02}; Viel et al. \\cite{viel05}, and Abazajian \\cite{abaza05}). Recently, Biermann \\& Kusenko (\\cite{bk06}) have shown that the decay of such a sterile neutrino could help initiate star formation in the early Universe and the detection of such an X-ray line would confirm whether such a sterile neutrino exists or not and this could be done by observing the X-ray decay line using XMM and Chandra satellites (Boyarski et al. \\cite{boya06}, Riemer-S$\\o$rensen et al. \\cite{riemer06}).  To summarize, we find that NFW, Moore and isothermal  density profiles become degenerate for fermion DM  particle of ${\\rm keV}$ masses at a size of the order  parsec. Assuming that supermassive black holes grow from degenerate fermion cores of masses of a few $10^{3} M_{\\odot}$ to $\\sim 10^{6}M_{\\odot}$ at galactic centers, we have found that the  sterile neutrino mass should be in the range from 0.6~{\\rm keV} to about 82~{\\rm keV}, with an improved lower limit of about 7 keV in the isothermal power law density profile. As this range of sterile neutrino masses overlaps with other results on sterile neutrino masses discussed in this Letter, this leaves the sterile neutrino to be an excellent candidate for DM. The formation and growth details of degenerate fermion balls is beyond the scope of this Letter and will be discussed in another paper."
    },
    "0609/astro-ph0609207_arXiv.txt": {
        "abstract": "We briefly discuss the past, present, and future state of astronomical science with laser guide star adaptive optics (LGS AO).  We present a tabulation of refereed science papers from LGS AO, amounting to a total of 23 publications as of May 2006.  The first decade of LGS AO science (1995--2004) was marked by modest science productivity ($\\approx$1~paper/year), as LGS systems were being implemented and commissioned. The last two years have seen explosive science growth ($\\approx$1~paper/month), largely due to the new LGS system on the Keck~II 10-meter telescope, and point to an exciting new era for high angular resolution science.  To illustrate the achievable on-sky performance, we present an extensive collection of Keck LGS performance measurements from the first year of our brown dwarf near-IR imaging survey.  We summarize the current strengths and weaknesses of LGS compared to {\\sl Hubble Space Telescope}, offer a list of desired improvements, and look forward to a bright future for LGS given its wide-scale implementation on large ground-based telescopes. ",
        "introduction": "\\label{sect:intro}  %  Astronomers have envisioned using laser guide star adaptive optics (LGS AO) to achieve diffraction-limited observations from ground-based telescopes for over two decades\\cite{1985A&A...152L..29F, 1987Natur.328..229T, 1994OSAJ...11..263H}.  (See Refs.~\\citenum{2004aoa..book.....R} and~\\citenum{1998aoat.conf.....H} for a historical review of LGS AO development.)  The realization of these visions has been an arduous effort, but we are now entering a new epoch as LGS systems are commissioned on the largest ground-based telescopes. The scientific promise of near diffraction-limited imaging and spectroscopy from the ground over most of the sky is finally being realized.  At this key juncture, the purpose of this paper is to briefly review published LGS science to date; to provide a snapshot of the science that is being done with LGS AO; and to look ahead to a future path where LGS AO is a ubiquitous tool for observational astronomy.   \\begin{table}[t] \\vskip -0.1in \\caption{Refereed science papers from LGS AO as of May 2006, listed by publication date. Some titles have been truncated and some title words abbreviated (``LGS'' and ``AO''). ``Field'' gives the area of study: ``SS'' = solar system, ``Gal'' = galactic, ``Xgal'' = extragalactic.  ``$N_{obj}$'' indicates the number of science targets/fields observed with LGS.} \\label{tab:bib} \\small \\begin{center} \\vskip -0.2in \\begin{tabular}{c|p{1.2in}|p{2.75in}|p{0.33in}|c|c|c} \\hline \\hline \\rule[-1ex]{0pt}{3.5ex} \\# & Authors, Journal & Title & Facilty & Field & $\\lambda\\lambda$ & $N_{obj}$  \\\\ \\hline \\rule[-1ex]{0pt}{3.5ex} 1 & McCullough \\etal\\ 1995, ApJ   & Photoevaporating Stellar Envelopes Observed with Rayleigh Beacon Adaptive Optics & SOR & Gal & H$\\alpha$ & 1 \\\\ \\hline \\rule[-1ex]{0pt}{3.5ex} 2 & Christou \\etal\\ 1995, ApJ     & Rayleigh Beacon AO Imaging of ADS 9731: Measurements of the Isoplanatic Field of View & SOR & Gal & $IJH$ & 1 \\\\ \\hline \\rule[-1ex]{0pt}{3.5ex} 3 & Drummond \\etal\\ 1995, ApJ     & Full AO Images of ADS 9731 and $\\mu$ Cassiopeiae: Orbits and Masses & SOR & Gal & $IJH$ & 1 \\\\ \\hline \\rule[-1ex]{0pt}{3.5ex} 4 & ten Brummelaar \\etal\\ 1996, AJ & Differential Binary Star Photometry Using the AO System at Starfire Optical Range & SOR & Gal & $ri$ & 10 \\\\ \\hline \\rule[-1ex]{0pt}{3.5ex} 5 & Glenar \\etal\\ 1997, PASP      & Multispectral Imagery of Jupiter and Saturn Using AO and Acousto-Optic Tuning & SOR & SS & 0.7--1.0 \\micron & 2 \\\\ \\hline \\rule[-1ex]{0pt}{3.5ex} 6 & Koresko \\etal\\ 1997, ApJ      & A Multiresolution Infrared Imaging Study of LkH$\\alpha$ 198 & SOR & Gal & $H$ & 1 \\\\ \\hline \\rule[-1ex]{0pt}{3.5ex} 7 & Drummond \\etal\\ 1998, Icarus  & Full AO Images of Asteroids Ceres and Vesta: Rotational Poles \\& Triaxial Ellipsoid Dimensions & SOR & SS & $i$ & 2 \\\\ \\hline \\rule[-1ex]{0pt}{3.5ex} 8 & Barnaby \\etal\\ 2000, AJ       & Measurements of Binary Stars with the Starfire Optical Range AO System & SOR & Gal & $ri$ & 1 \\\\ \\hline \\rule[-1ex]{0pt}{3.5ex} 9 & Hackenberg \\etal\\ 2000, AA    & Near-Infrared AO Observations of Galaxy Clusters: Abell 262 at z=0.0157 ... & Calar Alto & Xgal & $K$ & 1 \\\\ \\hline \\rule[-1ex]{0pt}{3.5ex} 10 & Perrin \\etal\\ 2004, Science   & Laser Guide Star AO Imaging Polarimetry of Herbig Ae/Be Stars & Lick & Gal & $JHK_S$ & 2 \\\\ \\hline \\rule[-1ex]{0pt}{3.5ex} 11 & Melbourne \\etal\\ 2005, ApJL   & Merging Galaxies in GOODS-S: First Extragalactic Results from Keck Laser AO & Keck & Xgal & $K^\\prime$ & 1 \\\\ \\hline \\rule[-1ex]{0pt}{3.5ex} 12 & Gal-Yam \\etal\\ 2005, ApJL     & A High Angular Resolution Search for the Progenitor of the Type Ic  Supernova 2004gt & Keck & Xgal & $K_S$ & 1 \\\\ \\hline \\rule[-1ex]{0pt}{3.5ex} 13 & Brown \\etal\\ 2005, ApJL       & Keck LGS AO Discovery and Characterization of a Satellite to Large Kuiper Belt Object 2003 EL61 & Keck & SS & $K^\\prime$ & 1 \\\\ \\hline \\rule[-1ex]{0pt}{3.5ex} 14 & Muno \\etal\\ 2005, ApJ         & A Remarkable Low-Mass X-ray Binary within 0.1 pc of the Galactic Center & Keck & Gal & $K^{\\prime}L^{\\prime}$ & 1 \\\\ \\hline \\rule[-1ex]{0pt}{3.5ex} 15 & Liu \\etal\\ 2005, ApJ          & Kelu-1 is a Binary L Dwarf: First Brown Dwarf Science from Laser Guide Star AO & Keck & Gal & $JHK^\\prime$ & 1 \\\\ \\hline \\rule[-1ex]{0pt}{3.5ex} 16 & Cohen \\etal\\ 2005, ApJL       & To Be or Not to Be: Very Young Globular Clusters in M31 & Keck & Xgal & $K^\\prime$ & 6 \\\\ \\hline \\rule[-1ex]{0pt}{3.5ex} 17 & Ghez \\etal\\ 2005, ApJ         & The First LGS AO Observations of the Galactic Center: Sgr A*'s Infrared Color ... & Keck & Gal & $K^{\\prime}L^{\\prime}$ & 1 \\\\ \\hline \\rule[-1ex]{0pt}{3.5ex} 18 & Marchis \\etal\\ 2006, Nature   & A Low Density of 0.8 g cm$^{-3}$ for the Trojan Binary Asteroid 617 Patroclus & Keck & SS & $HK^\\prime$ & 1\\\\ \\hline \\rule[-1ex]{0pt}{3.5ex} 19 & Brown \\etal\\ 2006, ApJ        & Satellites of the Largest Kuiper Belt Objects & Keck & SS & $K^\\prime$ & 4 \\\\ \\hline \\rule[-1ex]{0pt}{3.5ex} 20 & Krabbe \\etal\\ 2006, ApJL      & Diffraction Limited Imaging Spectroscopy of the SgrA* Region Using OSIRIS & Keck & Gal & 2.0--2.4 \\micron & 1 \\\\ \\hline \\rule[-1ex]{0pt}{3.5ex} 21 & Gelino \\etal\\ 2006, PASP      & Evidence of Orbital Motion in Binary Brown Dwarf Kelu-1AB & Keck & Gal & $HK^\\prime$ & 1\\\\ \\hline \\rule[-1ex]{0pt}{3.5ex} 22 & Liu \\etal\\ 2006, ApJ          & SDSS J1534+1615AB: A Novel T Dwarf Binary Found with Keck LGS AO and the Role of Binarity in the L/T Transition & Keck & Gal & $JHK^\\prime$ & 1 \\\\ \\hline \\rule[-1ex]{0pt}{3.5ex} 23 & Sheehy \\etal\\ 2006, ApJ       & Constraining the AO PSF in Crowded Fields: Measuring Photometric Aperture Corrections & Keck & Xgal & $H$ & 1 \\\\ \\hline \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": "More than two decades since its conception, LGS AO is now entering a new phase in its growth, opening a new era in high angular resolution science.  The first decade of astronomical science was marked by modest science productivity, as these LGS AO systems were commissioned and optimized ($\\approx$1~science paper/year).  Benefiting from the fruits of this effort, the LGS AO system on the Keck~II 10-meter telescope has had a highly successful first 1.5 years of science ($\\approx$1~paper/month).  Several LGS systems are planned to come online in the next two years, including those at the Gemini-North (2006), VLT (2006), Palomar 5-meter (2006), MMT (2006), Subaru (2007), Gemini-South (2007), and the Keck~I (2008) telescopes.  Given the size and quality of the science communities associated with these new AO systems, we can look forward to even more significant growth in LGS science soon ($\\approx$1~paper/week?).  We conclude with a non-comprehensive wishlist for future LGS developments, in order to highlight ongoing efforts and to review some outstanding challenges:  \\begin{itemize}  \\item {\\em Field of view:} Single LGS systems produce a corrected field of view the size of the isoplanatic angle.  In fact, nearly all LGS science thus far has been restricted to objects that span $\\lesssim$10\\arcsec\\ across (and most typically $\\lesssim2$\\arcsec). The upcoming Gemini-South multi-conjugate AO system is notable, as it will be the first LGS AO system to correct much larger fields of view ($\\approx$1--2\\arcmin).  \\item {\\em Observing efficiency:} Nighttime LGS AO operations currently require many more personnel compared to regular seeing-limited observations.  (See also paper 6270-12 by LeMignant \\etal\\ in this Proceedings.)  Greater automation will reduce this burden on telescope staff and also should lead to greater efficiency, as observing procedures are streamlined.  For ``survey-style'' science programs, in some sense the science return is proportional to the number of targets observed; therefore more efficient observing is a significant benefit.  \\item {\\em Robust real-time performance predictions:} LGS image quality varies depending on nightly weather conditions, laser performance, AO performance, sodium layer density and structure, and tiptilt star properties.  To maximize the science return, it would be desireable to be able to robustly predict LGS performance for any given target on any given night.  This capability would allow observations to be tailored to achieve the desired science goal.  Queue-scheduled observing is a key element here, but also better real-time understanding of seeing conditions, LGS performance and the conditions in the sodium layer are need.   \\item {\\em Near-IR tiptilt sensors:} While much of the sky is available for LGS, the most obscured regions (\\eg, star-forming regions) are not due to the lack of optically visible tiptilt stars.  Tiptilt sensors working at near-IR wavelengths would open the door to studying the youngest stages of star and planet formation.  \\item {\\em High quality catalogs for tiptilt stars:} Our brown dwarf imaging survey finds about 1 in 10 tiptilt stars with $R\\lesssim18$~mag from the USNO-B1.0 catalog\\cite{2003AJ....125..984M} are unsuitable, either because they turn out to be faint galaxies or they turn out to be much fainter than the reported magnitudes.  The Pan-STARRS project\\cite{2002SPIE.4836..154K} will provide precise, multi-band photometry over the entire sky visible from Hawaii, with the initial PS-1 telescope beginning operations later this year. Combined with the SkyMapper Telescope\\cite{2005AAS...206.1509S} in the southern hemisphere, high quality all-sky catalogs should be available in a few years for robust selection of tiptilt stars.  \\item {\\em Improved PSF stability and characterization:} The time-variability of the LGS PSF will remain a significant concern for the foreseeable future.  Post-processing software techniques can provide some immediate assistance, \\eg, deconvolution and/or PSF modeling techniques.\\cite{1993ApJ...415..862J, 2000A&AS..147..335D, 2006astro.ph..4551S} Development of algorithms to use AO telemetry data to estimate the real-time PSF would be a valuable capability for LGS.\\cite{2005PASP..117..847S} On a somewhat longer timescale, instruments tailored to handle the challenges of LGS AO imaging, such as dual-channel imaging systems, can circumvent this problem for some types of science programs.  Finally, next-generation LGS AO systems should produce higher Strehl imaging, leading to more stable and well-behaved PSFs.  \\end{itemize}  While the technology is far from mature, LGS AO is entering a phase of rapid growth.  There is little doubt that it will quickly become a key capability for a very broad range of astrophysics, spanning the nearest solar system bodies to the highest redshift galaxies and the entire universe in-between."
    },
    "0609/astro-ph0609800_arXiv.txt": {
        "abstract": "Observations of two of the formaldehyde (H$_2$CO) masers (A and D) in Sgr~B2 using the VLBA+Y27 (resolution $\\approx$ 0\\farcs01) and the VLA (resolution $\\approx 9\\arcsec$) are presented.  The VLBA observations show compact sources ($\\la 10$ milliarcseconds, $\\la 80$~AU) with brightness temperatures $>10^8$~K.  The maser sources are partially resolved in the VLBA observations.  The flux densities in the VLBA observations are about 1/2 those of the VLA; and, the linewidths are about 2/3 of the VLA values. The applicability of a core-halo model for the emission distribution is demonstrated.  Comparison with earlier H$_2$CO absorption observations and with ammonia (NH$_3$) observations suggests that H$_2$CO masers form in shocked gas.  Comparison of the integrated flux densities in current VLA observations with those in previous observations indicates that (1) most of the masers have varied in the past 20 years, and (2) intensity variations are typically less than a factor of two compared to the 20-year mean.  No significant linear or circular polarization is detected with either instrument. ",
        "introduction": "The nature of Galactic formaldehyde (H$_2$CO) masers is a growing mystery. While hundreds of Galactic OH, H$_2$O, and CH$_3$OH masers are known, only five Galactic star-forming regions have associated H$_2$CO maser emission. To date, this emission is seen only in the $1_{10}\\rightarrow 1_{11}$ transition at 6 cm wavelength.  Shortly after the discovery of the first H$_2$CO maser in NGC~7538 (Downes \\& Wilson 1974; Forster et al.\\ 1980), a radiative pumping model was proposed (Boland and de Jong 1981).  The H$_2$CO masers discovered subsequently did not meet the conditions required for this mechanism (Gardner et al.\\ 1986; Mehringer, Goss, \\& Palmer 1994, hereafter MGP94; Hoffman et al.\\ 2003; hereafter H03).  Thus, twenty-five years after the discovery of the first H$_2$CO maser, these sources remain rare and the excitation mechanism remains unknown.  Sgr~B2, the northernmost component of the extended Sgr~B radio source, is located within a few hundred pc of the Galactic center (Reid et al.\\ 1988). (The distance to the Galactic center is assumed to be 8.5~kpc in this paper.)  Sgr~B2 is comprised of three main star-forming complexes designated north (N), middle or main (M), and south (S), and many smaller H{\\sc ii} regions.  The H$_2$CO masers occur throughout Sgr~B2, shown in Figure~\\ref{fig1}.  The heating mechanisms and complex chemistry of the region are subjects of ongoing study (e.g., Gaume \\& Claussen 1990; Goicoechea et al.\\ 2004).  Sgr~B2 contains nine individual H$_2$CO maser regions, several of which have multiple velocity components.  All of the masers are unresolved at 1\\arcsec\\ angular resolution, except for maser C which MGP94 suggest consists of several masers blended within the beam.  These regions are near H{\\sc ii} regions distributed over the $\\sim$3.6 arcmin$^2$ complex (MGP94).  Whiteoak and Gardner (1983) and MGP94 designated the maser regions with letters (Fig.\\ \\ref{fig1}).  The H$_2$CO masers are observed over the velocity range $+40$~\\kms$\\lesssim v_{\\rm LSR} \\lesssim +80$~\\kms, while other species, such as H$_2$O masers, are observed over a larger range $-30$~\\kms$\\lesssim v_{\\rm LSR} \\lesssim +120$~\\kms\\ (Kobayashi et al.\\ 1989; McGrath, De Pree, and Goss 2004).  Of the nine maser regions observed in Sgr~B2 by MGP94, the G maser was shown to be time variable, at least quadrupling in intensity over 10~yr.  (Similarly, H03 found one NGC~7538 feature to triple in intensity over $\\approx 10$~yr.)  As initially noted by Whiteoak and Gardner (1983), all of the Sgr~B2 H$_2$CO masers lie close to OH, H$_2$O, CH$_3$OH, and NH$_3$ masers.  For most of the masers in MGP94, the separation to an OH maser was less than 0.05 pc.  Recent successes in search techniques for new masers (Araya et al.\\ 2004, 2005, 2006a) and in high-resolution observational techniques (H03) promise to provide empirical constraints for the development of a realistic model for the Galactic H$_2$CO maser emission.  The necessary steps in compiling an empirical picture of the H$_2$CO emission in Sgr~B2 are (1) detailed imaging of the masers in order to quantify the intrinsic properties of the emission (e.g., brightness temperature), (2) assessment of intensity variability in the masers, and (3) precise astrometry for elucidating spatial relationships between the H$_2$CO masers and more common masers (OH, H$_2$O, CH$_3$OH).  In this paper, we present new observations of the H$_2$CO masers in Sgr~B2 using the the Very Long Baseline Array (VLBA) and Very Large Array (VLA) of the NRAO\\footnote{The VLA and VLBA are components of the National Radio Astronomy Observatory (NRAO), a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.}. ",
        "conclusions": "We present VLBA+Y27 images of the Sgr~B2 A and D H$_2$CO masers.  The measured sizes ($\\la 80$~AU) and brightness temperatures ($>10^8$~K) are comparable to those found in other VLBI studies of H$_2$CO masers. However, about half of the flux density from these regions is resolved out with the VLBA data.  A comparison between VLA and VLBA observations shows that the missing flux density exhibits a broader linewidth than the emission from the compact VLBI source.  We demonstrate quantitatively the applicability of a core-halo model for these masers.  We also present new VLA observations of the H$_2$CO masers in Sgr~B2 with improved velocity resolution.  We have detected variability in several of the masers.  H$_2$O and CH$_3$OH masers discovered near the H$_2$CO masers may indicate associations among the species, suggestive of a related origin.  The association of H$_2$CO masers in Sgr~B2 with hot NH$_3$ shells proposed by Mart\\'{\\i}n-Pintado et al. (1999), together with the arguments for shock excitation of the maser region in NGC~7538 in H03, provide an encouraging stepping stone toward a solution of the problem of the excitation of these poorly understood masers."
    },
    "0609/gr-qc0609076_arXiv.txt": {
        "abstract": "By the assumption that the thermodynamics second law is valid, we study the possibility of $\\omega=-1$ crossing in an interacting holographic dark energy model. Depending on the choice of the horizon and the interaction, the transition from quintessence to phantom regime and subsequently from phantom to quintessence phase may be possible. The second transition avoids the big rip singularity. We compute the dark energy density at transition time and show that by choosing appropriate parameters we can alleviate the coincidence problem. ",
        "introduction": "Recent observations suggest that the universe is undergoing an accelerated expansion \\cite{acc}. This acceleration may be explained by the assumption that $70\\%$ of the universe is filled by a perfect fluid with negative pressure, dubbed dark energy. Some present data seem to favor an evolving dark energy, corresponding to an equation of state (EOS) parameter less than $\\omega=-1$ at present epoch from $\\omega>-1$ in the near past \\cite{Bo}. Many candidates for dark energy has been proposed such as the cosmological constant \\cite{dark}: A constant quantum vacuum energy density which fills the space homogeneously, corresponding to a fluid with a constant EOS parameter $\\omega=-1$; dynamical fields with a suitably chosen potential to make the vacuum energy vary with time \\cite{field}, and so on. Recently, using holographic principle, a new candidate for dark energy which is independent of any specific field has been suggested \\cite{Li},\\cite{hol}. Based on the holographic principle (which relates the number of degrees of freedom of a physical system to the area of its boundary), in order to allow the formation of black holes in local quantum field theory, Cohen et al \\cite{cohen} proposed a relationship between UV and IR cutoff. This yields an upper bound on the zero-point energy density, which by a suitable choice of the infrared cutoff, can be viewed as the holographic dark energy density. In \\cite{Li}, three candidates for the infrared cutoff was proposed: the Hubble radius, the particle horizon and the future event horizon. There was shown that among these options only the future event horizon may be identified with the desired infrared cutoff. To study the coincidence problem and also to have other choices for the infrared cutoff, e.g. the Hubble radius, interaction between dark matter and dark energy \\cite{inter} may be considered in the holographic dark energy model. As we have mentioned, based on astrophysical data, we may take into account the possibility of $\\omega=-1$ (phantom divide line) crossing. Therefore dark energy models which can describe phantom divide line crossing, has been also studied vastly in the literature \\cite{divide}. The phantom like behavior of interacting holographic dark energy was studied in \\cite{abd}, where it was claimed that by selecting appropriate interaction parameters the transition from the dark energy EOS parameter $\\omega_D>-1$ to $\\omega_D<-1$ is possible. Despite this, in \\cite{antabd} it was shown that the dark energy effective EOS parameter cannot cross $\\omega_D^{eff.}=-1$.  In this paper we consider interacting holographic dark energy model and study the ability of the model to describe the transition from quintessence to phantom regime and vice versa. After preliminaries in section two, where we introduce the interacting  holographic dark energy model and some of its general properties used in the subsequent sections, in section three we study the possibility of crossing $\\omega=-1$. In section four we derive necessary conditions for existence of two transitions in our model. The first transition is from quintessence to phantom phase and the second the transition from phantom to quintessence regime. The importance of the second transition lies on the fact that it avoids the big rip singularity. We discuss also the behavior of Hubble parameter and dark energy density at transitions times.  We use $\\hbar=G=c=k_B=1$ units throughout the paper. ",
        "conclusions": "In this paper we considered the holographic dark energy model with a general interaction between dark matter and dark energy, see (\\ref{3}). We took the infrared cutoff as a linear combination of the future an particle horizon see (\\ref{9}). We derived an expression for EOS parameter of the universe, $\\omega$, in terms of the ratio of dark energy density and total energy density of the flat FRW space-time, $\\Omega_D$, see (\\ref{20}). Using (\\ref{20}) and the expression obtained for time derivative of $\\Omega_D$ (\\ref{21}), and by assumption that the thermodynamics second law is still valid, we studied the possibility (and necessary conditions) for quintessence to phantom phase transition and vice versa with a differentiable Hubble parameter. Using some theorems about solutions of cubic equation satisfied by $\\Omega_D^{1/2}$ see (\\ref{25}), we showed that such transitions occur provided we appropriately choose the parameters of the system see(\\ref{47}) and (\\ref{50})."
    },
    "0609/astro-ph0609031_arXiv.txt": {
        "abstract": "Although observations point to the neutrality and lack of currents on large scales in the universe, many mechanisms are known that can generate charges or currents during the early universe.  We examine the question of survivability of relic charges and currents in a realistic model of the universe.  We show that the dynamics of cosmological perturbations drive the universe to become electrically neutral and current-free to a high degree of accuracy on all scales, regardless of initial conditions.  We find that charges are efficiently driven away in a time small compared to the Hubble time for temperatures $100 \\, \\mathrm{GeV} \\gtrsim T \\gtrsim 1 \\, \\mathrm{eV}$, while the same is true for currents at all temperatures $T \\gtrsim 1 \\, \\mathrm{eV}$.  The forced neutrality relaxes constraints on the generation of an electric charge in the early universe, while the forced erasure of currents disfavors many mechanisms for the early origins of large-scale magnetic fields. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609540_arXiv.txt": {
        "abstract": "Brighter Type~Ia supernovae (SNe~Ia) have broader, more slowly declining \\Bband\\ light curves than dimmer SNe~Ia.  We study the physical origin of this width-luminosity relation (WLR) using detailed radiative transfer calculations of Chandrasekhar mass SN~Ia models. We find that the luminosity dependence of the diffusion time (emphasized in previous studies) is in fact of secondary relevance in understanding the model WLR.  Instead, the essential physics involves the luminosity dependence of the \\emph{spectroscopic/color} evolution of SNe~Ia.  Following maximum-light, the SN colors are increasingly affected by the development of numerous Fe~II/Co~II lines which blanket the \\Bband\\ and, at the same time, increase the emissivity at longer wavelengths. Because dimmer SNe~Ia are generally cooler, they experience an earlier onset of Fe~III to Fe~II recombination in the iron-group rich layers of ejecta, resulting in a more rapid evolution of the SN colors to the red.  The faster \\Bband\\ decline rate of dimmer SNe~Ia thus reflects their faster ionization evolution. ",
        "introduction": "Although normal Type~Ia supernovae (SNe~Ia) are generally considered a homogeneous class, they nevertheless show substantial ($\\sim 1$~mag) variations in luminosity at peak.  The use of SNe~Ia for cosmology measurements thus relies on empirical calibration techniques.  Most common among these is the width-luminosity relation \\citep{Phillips_1999}.  Bright SNe~Ia generally have broad \\Bband\\ light curves (LCs) which decline slowly after peak. Dimmer SNe~Ia have narrower, more quickly declining LCs.  Understanding the physical origin of the width-luminosity relation (WLR) is a primary goal of the theory and modeling of SNe~Ia. Unfortunately, a full theoretical description of the SN~Ia explosion mechanism is still lacking.  Normal SNe~Ia are widely believed to be the thermonuclear disruption of carbon-oxygen white dwarfs near the Chandrasekhar limit, however a number of uncertainties remain regarding the structure of the progenitor, the precise ignition conditions, and the physics of the turbulent nuclear combustion that unbinds the star.  Fortunately, these uncertainties need not be fully resolved in order to study the WLR. In SN~Ia explosions, hydrodynamic and nuclear burning processes last only $\\sim 1$~minute.  The subsequent luminosity is powered entirely by the decay of radioactive elements synthesized in the explosion, in particular \\Nifs\\ in the decay chain $\\Nifs \\rightarrow\\ \\Cofs \\rightarrow\\ \\Fefs$.  The LCs of SNe~Ia are therefore fully determined by the composition and density structure of the material ejected in the explosion.  Naturally, the mass of \\Nifs\\ produced (\\Mni) is the primary determinate of the peak brightness of the event.  Observations indicate that for normal objects, \\Mni\\ spans the range $0.4-0.9$~\\Msun, with a typical value $\\Mni \\approx 0.6~\\Msun$.  The challenge thus falls to radiative transfer theory to explain why SNe~Ia with larger \\Mni\\ also have broader \\Bband\\ LCs.  Most previous transfer studies have emphasized diffusion arguments of one sort or another \\citep{Hoeflich_WLR, Hoeflich_99by, Pinto-Eastman_WLR, Mazzali_WLR}.  In SNe~Ia, the ejecta remain optically thick for the first several months after explosion.  The width of the bolometric LC is related to the time scale for photons to escape the ejecta by diffusion.  In principle, the WLR could be explained if brighter SNe~Ia have higher effective opacities and hence a longer diffusion time.  Different authors have invoked different physical arguments to motivate this sort of opacity dependence (see \\S\\ref{Sec:Others}).  In this paper, we stress that the WLR does not in fact hinge on the bolometric diffusion time, but is instead principally a \\emph{broadband} phenomenon.   In particular, the \\Bband\\ LC depends sensitively on the rate at which the SN colors evolve progressively redward following maximum light.  Dimmer SNe~Ia exhibit a faster color evolution than brighter SNe~Ia; this turns out to be a primary reason for their relatively faster \\Bband\\ decline rates.  A relevant theoretical explanation of the WLR should therefore focus on the effect \\Mni\\ has on the \\emph{spectroscopic/color evolution} of SNe~Ia, rather than on the overall diffusion time-scale.  In what follows, we use realistic time-dependent multi-group radiative transfer calculations to explain the physical origin of the WLR in a set of Chandrasekhar-mass SN~Ia models. We demonstrate that the faster color evolution of dimmer models reflects their faster \\emph{ionization evolution}.  Following maximum-light, the SN colors are increasingly affected by the development of numerous Fe~II/Co~II lines which blanket the \\Bband\\ and, at the same time, increase the emissivity at longer wavelengths.  Because dimmer SNe~Ia are generally cooler, they experience an earlier onset of Fe~III to Fe~II recombination in the iron-group rich layers of ejecta, resulting in a more rapid evolution of the SN colors to the red.  This explains their faster \\Bband\\ decline rate. ",
        "conclusions": ""
    },
    "0609/astro-ph0609295_arXiv.txt": {
        "abstract": "During the past decade $\\eta$ Car has brightened markedly, possibly indicating a change of state.  Here we summarize photometry gathered by the Hubble Space Telescope as part of the HST Treasury Project on this object.  Our data include STIS/CCD acquisition images, ACS/HRC images in four filters, and synthetic photometry in flux-calibrated STIS spectra.  The HST's spatial resolution allows us to examine the central star separate from the bright circumstellar ejecta.  {\\it Its apparent brightness continued to increase briskly during 2002--06, especially after the mid-2003 spectroscopic event.\\/} If this trend continues, the central star will soon become brighter than its ejecta, quite different from the state that existed only a few years ago.  One precedent may be the rapid change observed in 1938--1953.  We conjecture that the star's mass-loss rate has been decreasing throughout the past century. ",
        "introduction": "%  Eta Carinae's photometric record is unparalleled among well-studied objects, especially since it has been near or exceeded the classical Eddington limit during the past two or three centuries.  From 1700 to 1800 it gradually brightened from 4th to 2nd magnitude, and then experienced its famous Great Eruption or ``supernova impostor event'' beginning about 1837.  For twenty years it was one of the brightest stars in the sky, rapidly fluctuating between magnitudes 1.5 and 0.0, briefly attaining $V \\approx -1.0$. After 1858 it faded below 7th magnitude, presumably enshrouded in the nascent Homunculus nebula.  Subsequent behavior, however, has been more complex than one might have expected. A mysterious secondary eruption occurred in 1887--1900; then the apparent brightness leveled off around $m_{pg} \\approx 8$ for about 40 years, followed by a rapid increase in 1938--53;  after that it brightened at a fairly constant rate for another 40-year interval, and most recently the rate accelerated in the 1990's. Some, but not all, of the secular brightening can be attributed to decreasing obscuration as the Homunculus nebula expands.  However, in truth this is more complex than it appears. The star's physical structure has been changing in a decidedly non-trivial way which is, at best, only dimly understood.  For historical and observational details see  \\citet{dh97,frew05,dvauc52, oconnell56,feinstein67,feinstein74,aavso,kdetal99a,kdetal99b, vangenderen99, sterken99,phot1}.  Spectroscopic changes have occurred along with the brightness variations.  The 5.5-year spectroscopic/photometric cycle \\citep{gaviola53,zanella84,whitelock94,whitelock04,damineli96,phot1} is not apparent in data obtained before the 1940's \\citep{feast01,rmh05}.  Brief ``spectroscopic events'' marking the cycle are most likely mass-ejection or wind-disturbance episodes, probably regulated by a companion star \\citep{zanella84,kd99,smith03,heii}. At visual wavelengths, the associated ephemeral brightness changes represent mainly emission lines in the stellar wind, while the longer-term secular brightening trend involves the continuum \\citep{phot1,martinconf}.  \\citet{rmh05}, \\citet{halpha05}, and \\citet{kd05} have speculated that the four obvious disruptions in the photometric record -- c.\\ 1843, 1893, 1948, and 2000 -- might indicate a quasi-periodicity of the order of 50 years.\\footnote{ So far as we know, this idea was first voiced by Humphreys at two meetings in 2002, but it did not appear in the published proceedings} In any case the star has not yet recovered from its Great Eruption seen 160 years ago.  The Hubble Space Telescope (HST) Treasury Program for $\\eta$ Car was planned specifically to study the 2003.5 spectroscopic event.  We employed the Space Telescope Imaging Spectrograph (STIS) and Advanced Camera for Surveys (ACS), following earlier STIS observations that began in 1998 \\citep{etatp04}.   Fortuitously, the STIS data almost coincide with a rapid secular brightening which began shortly before 1998 (see Section 4 below).  Those and the ACS images are of unique photometric value for at least two reasons: \\begin{enumerate} \\item{At visual wavelengths, normal ground-based observations have been dominated by the surrounding Homunculus ejecta-nebula, which, until recently, appeared much brighter than the central star and which has structure at all size scales from 0.1 to 8 arcseconds. {\\it So far, only the HST has provided well-defined measurements of just the central star.\\/}\\footnote{ At least this is true for visual and UV wavelengths.  The near-infrared photometry reported by \\citet{whitelock94} and \\citet{whitelock04} may be strongly dominated by the central star. Those observations probably represent free-free emission in the wind at larger radii than the visual wavelength data.  They show both the spectroscopic events and the brightening trend better than other ground-based measurements.} The Homunculus is primarily a reflection nebula, but the Homunculus/star brightness ratio has changed substantially. During 1998-99, for instance, the star nearly tripled in apparent brightness while ground-based observations showed only about a 0.3-magnitude brightening of Homunculus plus star \\citep{kdetal99b}.  This rather mysterious development is known from HST/STIS and HST/ACS data. }  \\item{Numerous strong emission lines perturb the results for standard photometric systems.  H$\\alpha$ and H$\\beta$ emission, for example, have equivalent widths of about 800 and 180 \\mbox{\\AA} respectively in spectra of $\\eta$ Car.  Broad-band $U$, $B$, $R$, and $I$ magnitudes, and most medium-band systems as well, are therefore poorly defined for this object.  Photometry around 5500 \\mbox{\\AA}, e.g. broad-band $V$, is relatively free of strong emission lines, but transformations from instrumental magnitudes to a standard system require the other filters \\citep{kdetal99b,sterken01,vangenderen03}. This difficulty is somewhat lessened for HST observations restricted to the central star, whose spectrum has fewer emission lines than the bright ejecta;  and some of the HST/ACS filters are fairly well-adapted to the case.  At any rate the STIS and ACS data appear to be stable and internally consistent.  Detector and filter systems used in most ground-based work, on the other hand, require fluctuating instrumental and atmospheric corrections, and do not give any major advantage for this object. } \\end{enumerate} In this paper we present the complete set of photometric data gathered for the $\\eta$ Car HST Treasury Project.  The central star has brightened, especially in the UV, since the HST results described by \\citet{phot1}. We report three types of later measurements: \\begin{enumerate} \\item{The star's brightness in acquisition images made with STIS before that instrument's failure in early 2004.  These images represent a broad, non-standard wavelength range from 6500 to 9000 {\\AA}.} \\item{The star's brightness in ACS/HRC images made in four filters (F220W, F250W, F330W, \\& F550M) from October 2002 to the present.} \\item{Synthetic photometry in flux-calibrated STIS spectra.} \\end{enumerate}  After presenting  the data below, we briefly discuss the observed trends and what they bode for the near future of $\\eta$ Car.  Our principal reason for reporting these data now is that the Treasury Program observations have been completed;  future HST/ACS observations are possible but not assured.  These last Treasury Program observations are essential for demonstrating the secular trend in brightness (see Section \\ref{trend}), that these brightness changes suggest a fundamental change in the state of the star (see Section \\ref{discussion}) and what the near future may hold for $\\eta$ Carinae (\\ref{predictions}). ",
        "conclusions": "}  %  The observed brightening of $\\eta$ Car is not easy to explain. It cannot signify a major increase in the star's luminosity, because that would exceed the Eddington limit, producing a giant eruption.  It cannot be a standard LBV-like eruption;  in that case the energy distribution should have shifted to longer wavelengths, the Balmer emission lines should have decreased, and the spectrum should have begun to resemble an A- or F-type supergiant \\citep{lbv94}. In fact, qualitatively the star's spectrum has changed little in the past decade, and it has become bluer, not redder.\\footnote{ The change in color is modest, however, too small to confidently quote here.  Dust near $\\eta$ Car has long been known to have an abnormally small reddening/extinction ratio, see \\citet{dh97}, \\citet{fos95}, and refs.\\ cited therein.}  The most obvious remaining explanation involves a change in the circumstellar extinction,  which, in turn, probably requires a subtle change in the stellar wind.  Mere ``clearing of the dust'' -- i.e., motion of a localized concentration of dusty ejecta -- cannot occur fast enough \\citep{kdetal99b}.  Therefore one must consider either destruction of dust grains, or a decrease in the formation of new dust, or both; and, if these account for the observations, why should they happen now?  \\subsection{Dust Near the Star}  The hypothetical decreasing extinction probably occurs within 2000 AU ($\\sim$ 1{\\arcsec}) of the star, and preferably closer, because: \\begin{enumerate} \\item{In various observations between 1980 and 1995, the star did not appear as bright as expected relative to the Weigelt blobs;  the discrepancy was a factor of the order of 10, based on simple theoretical arguments \\citep{dh86,fos95}.  Evidently, then, our line of sight to the star had substantially larger extinction at visual wavelengths, even though its projected separation from the blobs was less than 0.3{\\arcsec}.  The required extra extinction was of the order of 3 magnitudes.  Since then the star has brightened far more than the Weigelt objects have; therefore, if this involves localized extinction, its size scale must be a fraction of an arcsec, only a few hundred AU. } \\item{No known process seems likely to destroy dust more than 2000 AU from the star in a timescale of only a few years.} \\item{Ground-based photometry and HST images have shown only a modest, fraction-of-a-magnitude increase in the brightness of the large-scale Homunculus lobes during the past decade \\citep{laplata,phot1}. } \\end{enumerate}  Dust grains should condense in $\\eta$ Car's wind at a distance of 200--600 AU, 2 to 10 years after the material has been ejected.\\footnote{ Here, lacking a specific dust-formation model for the unusual case of $\\eta$ Car, we suppose that appreciable grain condensation begins in the outward flowing material at the location where the equilibrium grain temperature is around 1000 K.  This is a fairly conventional assumption and the precise choice of temperature has little effect on our reasoning. The quoted time-after-ejection assumes typical ejecta speeds of 200--700 km s$^{-1}$.}  Since newly-formed dust moves outward in a timescale of several years, the circumstellar extinction seen at any time depends partly on the current dust formation rate.  This, in turn, depends on local wind density, radiation density, etc., and newly formed hot grains ($T_d > 800$ K) are susceptible to destruction.  The dust column density can thus be sensitive to small changes in the stellar parameters. Moreover, the wind is latitude-dependent and our line of sight is close to the critical latitude where wind parameters can vary rapidly \\citep{smith03}.  All these factors appear suitable for the proposed explanation.  On the other hand, near-infrared observations imply that extinction within $r < 2000$ AU  has been quite small along most paths outward from the star.  In Fig.\\ 3 of \\citet{cox95}, for instance, the 2--6 $\\mu$m flux indicates the high end of the dust temperature distribution.  Modeling this in a conventional way, we find that less than 5\\% of the total luminosity was absorbed and re-emitted by inner dust with $T_d > 500$ K during the years 1973 to 1990 when those observations were made.\\footnote{ The measured flux was approximately a power law $f_{\\nu} \\sim {\\nu}^{-3.7}$ at wavelengths around 4 ${\\mu}$m. Assuming a typical emission efficiency dependence $Q_{\\nu} \\sim {\\nu}$, the observed spectral slope can be explained by a grain temperature distribution $dN/dT \\sim T^{-8.7}$.  The result noted in the text is obtained by normalizing this to match the observed flux around 4 or 5 $\\mu$m and then integrating the total emitted flux at all wavelengths due to grains above 500 K.}  Therefore, {\\it our line of sight must be abnormal in order to have a large amount of extinction near the star.\\/}  In principle one might view this as an argument against our proposed scenario, but no plausible alternative has been suggested to explain the apparent faintness of the central star before 1999 and its ratio to the Weigelt blobs \\citep{speckle86,dh86,speckle88,fos95,kdetal99b},  The spatial distribution of dust is probably quite inhomogeneous near the star.  The Homunculus lobes have a conspicuously ``granular'' appearance;  the equatorial ejecta are clumpy, including the Weigelt knots; and stars near and above the Eddington limit tend to produce clumpy outflows \\citep{shaviv05}.  Consequently the radiative transfer problem includes macroscopic effects which have not yet been modeled.  If the grain albedo is sufficient, light may escape mainly by scattering through interstices between condensations.  In that case, high-extinction lines of sight may be fairly common in the inner region even though most of the light escapes along other paths, not necessarily radial.  Incidentally, the near-infrared photometric trends reported by \\citet{whitelock94,whitelock04} are not straightforward to interpret.  The fairly-constant 3.5 $\\mu$m flux, for instance, represents a complicated mixture of dust formation parameters and does not necessarily indicate a constant amount of dust; see comments by \\citet{kdetal99b}.  \\subsection{The Role of the Stellar Parameters \\label{stellarparams}}  If the observed brightening represents a decrease in circumstellar extinction, the likeliest reason for this to occur is through some change in the star -- no one has yet proposed a suitable alternative.  The most relevant stellar parameters are the radius,  current luminosity, and surface rotation rate, which together determine the wind's velocity, density, and latitude structure.  All of these may still be changing today, 160 years after the Great Eruption;  thermal and rotational equilibrium in particular are likely to be poor assumptions for the star's internal structure \\citep{smith03,kd05}.  As a working hypothesis to explain $\\eta$ Car's photometric and spectroscopic record in the past 100 years, let us tentatively suppose that {\\it the mass-loss rate is gradually decreasing,\\/} while the surface rotation rate may be increasing.  Historical considerations include: \\begin{enumerate} \\item{High-excitation \\ion{He}{1} emission, now observed at most times, was consistently absent before 1920 \\citep{feast01} and probably before 1940 \\citep{rmh05}.  If a hot companion star is present as most authors suppose, then the most obvious way to hide or suppress its helium ionization is to immerse the entire system in an extremely dense wind -- i.e., the primary star's mass-loss rate must have been larger then. This idea is far from straightforward \\citep{kd99}, but so far as we know it is the only qualitative explanation yet proposed.  Informally, based on Zanstra-style arguments (i.e., assessing the volume emission measure $n_{He} n_e V$ needed to absorb all the photons above 25 eV), we estimate that a rate of the order of 10 times the present value, i.e. $\\sim 10^{-2}$ $M_\\odot$ y$^{-1}$, would have been required early in the twentieth century in order to suppress the helium recombination emission. } \\item{Twenty years ago the amount of fresh dust, indicated by the near-infrared flux, appeared consistent with a mass-loss rate somewhat above $10^{-3}$ $M_{\\odot}$ yr$^{-1}$ \\citep{iue86}.  This absorbed only a small fraction of the luminosity (Section 5.1 above), but the substantially higher mass-loss rate suspected for earlier times would have produced enough hot inner dust to absorb a non-negligible fraction. } \\item{The brightness observed between 1900 and 1940 is rather mysterious.  Judging from its mass and present-day optical thickness, around 1920 the Homunculus (then only half as large as it is today) should have had at least 5 magnitudes of visual-wavelength extinction;  in a simple model the object should have been fainter than 10th magnitude instead of $m_{pg} \\approx 8$ as was observed. No doubt the inhomogeneities mentioned earlier played a role, but no model has been calculated. Moreover, why did the brightness remain fairly constant even though the Homunculus expanded by about 70\\% in 1900--1940? This interesting problem has received practically no theoretical attention.} \\item{\\ion{He}{1} emission first appeared, and $\\eta$ Car's brightness suddenly increased, between 1938 and 1953 as we mentioned in Section 1.  This might conceivably be explained by a decrease in the wind density;  but \\citet{kd05} and \\citet{halpha05} have conjectured that 1940--1950 may have been the time when rotation became fast enough to produce latitude structure in the wind.  If so, a higher-excitation, lower-density zone then developed at low latitudes \\citep{smith03}. } \\end{enumerate}  The above points inspire two hypotheses that may explain the rapid brightening trend shown in Fig.\\ \\ref{plot1}. First, if the mass-loss rate has been decreasing, this tends to reduce the column density of recently-formed dust along our line of sight.  Meanwhile (or alternatively), perhaps the wind's latitude structure is continuing to evolve so that its dense zone is now moving out of the line of sight.  HST data suggest that our line of sight has been fairly close to the critical boundary latitude separating the two phases \\citep{smith03}.  A small increase in surface rotation rate, or some other parameter change, might conceivably move the dense zone to higher latitudes, decreasing the amount of dust that forms along our line of sight.  This idea is appealing because it suggests a way in which the effective extinction may be very sensitive to the stellar parameters.  This problem obviously requires detailed models far beyond the scope of this paper, combining the star's changing structure, its wind, dust formation, and possibly dust destruction.  \\subsection{Concerning the 5.5-year Cycle}  Figs.\\ \\ref{plot1} and \\ref{plot2} reveal no major surprises about the 2003.5 spectroscopic event, but several comments are worthwhile.  First, the sharp drop in UV brightness (filters F220W and F250W) is qualitatively understood and does not involve circumstellar dust.  During both the 1998 and the 2003 events, STIS data showed very heavy ultraviolet blanketing by ionized metal lines;  indeed the star became quite dark at some wavelengths between 2000 and 3000 {\\AA} \\citep{ironcurtain}. We further note that just before the spectroscopic event, a slight increase occurred at wavelengths below 4000 {\\AA} (filters F220W, F250W, F330W), but not at visual and far-red wavelengths (F550M and F25ND3).  Ground-based visual-wavelength and near-IR photometry showed a qualitatively similar effect \\citep{vangenderen03,whitelock04,laplata}.  The brightening is particularly prominent in J, H, and K which are dominated by free-free emission \\citep{whitelock04}. The ACS F550M and STIS F25ND3 data primarily measure the continuum brightness, while the other HST filters are heavily influenced by strong emission or absorption lines.    At about the same time \\ion{He}{1} emission in the central star also went through a similar increase in brightness \\citep{martinconf}.  The minor pre-event brightening thus appears to represent an increase in some emission features implying a temporary increase in ionizing flux.  The primary star may provide additional UV photons or the hypothetical hot companion star may excite the primary wind more than usual at that time (just before periastron), but no quantitative model has been attempted.  Figs.\\ \\ref{plot1} and \\ref{plot2} contain interesting hints about the timescale for the star's post-event recovery.  Four months after the 2003.5 event, for instance, the 2--10 keV X-ray flux had increased almost to a normal level \\citep{xray}. The HST/ACS F220W and F250W brightnesses, however, were still quite low at that time, and they required about eight months to recover.  This timescale must be explained in any valid model for the spectroscopic events.  \\citet{halpha05} noted serious differences between STIS spectra of the 1998.0 and 2003.5 events, and interpreted them as evidence for a rapid secular physical change in $\\eta$ Car. \\citet{damineli99} had earlier found that \\ion{He}{1} emission became progressively weaker after each of the last few spectroscopic events.  These clues are obviously pertinent to our comments in Section \\ref{stellarparams} above.  Fluctuations {\\it between\\/} spectroscopic events have received little attention in the past.   For instance, Fig.\\ \\ref{plot1} shows a brief 0.2-magnitude brightening at 2001.3; measured by the STIS in both imaging and spectroscopic mode.  It was correlated with the behavior of a strange unidentified emission line near 6307 {\\AA}, and with other subtle changes in the spectrum \\citep{uemit}.  This is interesting because mid-cycle events have not been predicted in any of the competing scenarios for the 5.5-year cycle.  Perhaps the effects seen in 2001 indicate the level of basic, LBV-like fluctuations in $\\eta$ Car.  \\subsection{Eta Carinae in the Near Future \\label{predictions}}  The appearance of $\\eta$ Car and the Homunculus nebula has changed dramatically.  Twenty years ago the entire object could have been described as ``a bright, compact nebula having an indistinct eighth-magnitude central core;''   but a few years in the future, if recent trends continue, it will be seen instead as ``a fifth- or even fourth-magnitude star accompanied by some visible nebulosity.''  Meanwhile the color is gradually becoming bluer.  This overall development has long been expected \\citep{kd87}, but now appears to be occurring 20 years ahead of schedule.   If it signals an irregularity in the star's recovery from the Great Eruption, then this may be a highly unusual clue to the highly abnormal internal structure.  Unsteady diffusion of either the thermal or the rotational parameters would be significant for stellar astrophysics in general.  There are several practical implications for future observations of this object.  Valid ground-based spectroscopy of the star (strictly speaking its wind) is becoming feasible for the first time, as its increased brightness overwhelms the emission-line contamination by inner ejecta.  Unfortunately this implies that the inner ejecta -- particularly the mysterious Weigelt knots -- are becoming difficult to observe.  In fact, since the HST/STIS is no longer available, they are now practically impossible to observe.  When some new high-spatial-resolution spectrograph becomes available in the future, the inner ejecta will probably be much fainter than the star.  The expected future of the larger-scale Homunculus nebula is also interesting.  At present it is essentially a reflection nebula.  However, based on the presence of high-excitation emission lines such as [\\ion{Ne}{3}] close to the star, the system almost certainly contains a source of hydrogen-ionizing photons with energies above 13.6, and helium-ionizing photons above 25 eV.  (See, e.g., \\citet{zanella84};  most recent authors assume that this source is a hot companion star.)  Eventually, when circumstellar extinction has decreased sufficiently due to expansion and other effects, the UV source will begin to photoionize the Homunculus.  This is especially true if the primary stellar wind is weakening as we conjectured above. First the inner ``Little Homunculus'' will become a bright compact \\ion{H}{2} region, and then the bipolar Homunculus lobes.  The time when that will occur is not obvious, but it may be within the next few decades if current trends continue."
    },
    "0609/astro-ph0609776_arXiv.txt": {
        "abstract": "{} {\\hete\\ is the seventh known  X-ray transient accreting millisecond pulsar and has been in outburst for more than one year. We compared the data on \\hete\\ with other similar objects and made an attempt at deriving constraints on the physical processes responsible for a spectral formation.} {The broad-band spectrum of the persistent emission in the 2--300 keV energy band and the timing properties were studied using simultaneous {{\\it INTEGRAL} } and publicly available {{\\it RXTE}} data obtained in October 2005. The properties  of the X-ray bursts observed from \\hete\\ were also investigated.} {The spectrum is well described by a two-component model consisting of a blackbody-like soft X-ray emission at 0.8 keV temperature and  a thermal Comptonized spectrum with electron temperature of 30 keV and Thomson optical depth $\\tau_{\\rm T} \\sim 2$ for the slab geometry. The source is detected by {{\\it INTEGRAL}} up to 200 keV at a luminosity of $5\\times10^{36}$ erg s$^{-1}$ (assuming a distance of 5 kpc) in the 0.1--200 keV energy band. We have also detected one type I X-ray burst which shows  photospheric radius expansion. The burst occurred at an inferred persistent emission level of $\\sim$ 3--4\\% of the Eddington luminosity. Using  data for all X-ray bursts observed to date from \\hete, the burst recurrence time is estimated to be about 2 days. No pulsations  have been detected either in the {{\\it RXTE}} or in the {{\\it INTEGRAL}} data which puts interesting constraints on theories of  magnetic field evolution in neutron star low-mass X-ray binaries. } {} ",
        "introduction": "\\label{sec:intro}  The detection of  X-ray millisecond pulsation in persistent emission from low-mass X-ray binaries (LMXBs) remained elusive for many years until the discovery of the first accreting millisecond pulsar (MSP) by \\citet{wvdk98}. Since that time, a total of seven accreting MSP transients  have been detected. They are weakly magnetized ($\\sim 10^{8}-10^9$ G) neutron stars (NS) with spin frequencies in the 180--600 Hz range and orbital periods between 40 min and 5 hr \\citep[see reviews by][]{w05,p06}. Their companion stars have been found to be either highly   evolved  white or brown dwarfs. For the first time,  the predicted decrease of the NS spin period during accretion was measured  in the accreting MSP IGR~J00291+5934 \\citep{mfb05,burd06}. This provided a strong confirmation of the theory of `recycled' pulsars in which old neutron stars in LMXBs become millisecond radio pulsars through spin-up by transfer of angular momentum by the accreting material.  MSP energy spectra can be well  described by a two-component model consisting of a soft black body (or multi-color  blackbody) and a hard power-law like tail. The soft thermal component could be associated with radiation from the accretion disc and/or the heated NS surface around the shock \\citep[see e.g.][]{gp05,p06}. The hard emission is likely to be produced by thermal Comptonization in the hot accretion shock on the NS surface \\citep{gdb02,pg03} with seed photons coming from the stellar surface. The observed hard spectra are similar to the spectra observed from atoll sources in their hard, low-luminosity state \\citep{b00}.  \\hete\\ was discovered during a bright X-ray burst by the {\\it High Energy Transient Explorer 2} ({\\it HETE-2}) on 14 June 2005 \\citep{vanderspek05}. Followup observations with the {\\it Rossi X-ray Timing Explorer} ({\\it RXTE}) identified the source as the seventh X-ray accreting millisecond pulsar, with a pulse frequency of 377.3 Hz, an orbital period of 83 min, and most likely a 0.016--0.07 M$_{\\odot}$ brown dwarf companion \\citep{kaaret06}. The detected burst was consistent with a type I X-ray burst with  photospheric radius expansion. Assuming that the bolometric burst peak luminosity during  photospheric radius expansion saturated at the Eddington limit, \\citet{kawai05} estimated the distance to the source to be $\\sim5$ kpc assuming helium burst burning and canonical NS values.  An optical counterpart candidate had an $R$-band magnitude of 18.02 and a broad HeII emission line spectrum \\citep[][]{fox05,steeghs05a}. A similar line was previously observed in  IGR~J00291+5934 \\citep{rjs04,ffc04}. The optical counterpart of the X-ray source is located at coordinates $\\alpha_{\\rm J2000} = 19^{\\rm h}00^{\\rm m}08\\fs65$ and $\\delta_{\\rm J2000} = -24{\\degr}55\\arcmin13\\farcs7$ with an uncertainty of $0\\farcs2$. Near-infrared observations detected the optical candidate at a constant  magnitude of J=17.6. No radio counterpart consistent with the \\hete\\ coordinates was  detected by the VLA \\citep[][]{steeghs05b,rupen05}.  One year after the discovery, \\hete\\ is still active (see Fig. \\ref{fig:asm}). Compared to  other accreting MSPs with outburst periods of a few days to a month, \\hete\\ shows evidence of being a ``quasi-persistent'' X-ray source. This source also has  other properties atypical for accreting MSPs. During the first $\\sim30$ days after the discovery, \\hete\\ showed significant flux emission variability in the fractional rms amplitude \\citep{galloway06b}. On July 8, 2005 (MJD 53559) during the flux brightening, the observed pulse frequency decreased by $\\Delta\\nu/\\nu\\sim6\\times10^{-7}$, and in the subsequent observations, the pulsations were suppressed  \\citep{kaaret06}.  In this paper we report the  {\\it INTEGRAL} observations of \\hete\\ obtained simultaneously with {\\it RXTE}. We study the broad-band spectral and timing properties of the source. The X-ray burst properties are also investigated.   \\begin{figure} \\centerline{\\epsfig{file=6457fig1.eps,width=8.0cm} } \\caption{{\\it RXTE}/ASM light curve for \\hete\\ averaged over 1-day intervals from May 5, 2005 (53500 MJD). The count rate has been converted into flux using 1 Crab Unit for 75 cts s$^{-1}$ \\citep{l96}. The arrows indicate the times of the detected X-ray bursts \\citep[][and this work]{vanderspek05,barbier05,galloway05}. } \\label{fig:asm} \\end{figure} ",
        "conclusions": "\\label{sec:conclusions}  We have found that the source spectrum is similar to other accreting X-ray millisecond pulsars having a high plasma temperature around 30 keV and a Thomson optical depth $\\sim2$. This source differs from other MSP in requiring thermal soft X-ray emission with nearly double the temperature. From our spectral fits we infer that this emission is not likely be produced in a multi-temperature accretion disk but  more likely arises from thermal emission at the NS surface. One might expect that the lack of pulsations could be due to a particularly high optical depth, but our spectral fits rule out this possibility.   The reason for the lack of coherent pulsations in the persistent emission from LMXBs is still an open question. Different explanations have been put forward to explain this phenomenon, including models which invoke gravitational lensing, electron scattering, or weak surface magnetic fields due to magnetic screening \\citep[][and references therein]{w88,brainerd87,titarchuk02,cumming01}. The high accretion rate inferred for \\hete\\ relative to the other known MSP transients suggests that we may be observing the evolution of the NS's magnetic field due to magnetic screening in this source \\citep{cumming01}. The  transition of \\hete\\ from an X-ray millisecond pulsar to a persistent LMXB could indicate that there is a population of suppressed  X-ray millisecond pulsars among the non-pulsating LMXBs. Detailed observations of this source at the epoch of pulsation suppression can help to solving the long-standing issue of missing pulsations in persistent LMXB emission."
    },
    "0609/astro-ph0609406_arXiv.txt": {
        "abstract": "This paper reports on a near-infrared survey of early-type galaxies designed to provide information on bar strengths, bulges, disks, and bar parameters in a statistically well-defined sample of S0-Sa galaxies. Early-type galaxies have the advantage that their bars are relatively free of the effects of dust, star formation, and spiral structure that complicate bar studies in later type galaxies. We describe the survey and present results on detailed analysis of the relative Fourier intensity amplitudes of bars in 26 early-type galaxies. We also evaluate the {\\it symmetry assumption} of these amplitudes with radius, used recently for bar-spiral separation in later-type galaxies.  The results show a wide variety of radial Fourier profiles of bars, ranging from simple symmetric profiles that can be represented in terms of a single gaussian component, to both symmetric and asymmetric profiles that can be represented by two overlapping gaussian components. More complicated profiles than these are also found, often due to multiple bar-like features including extended ovals or lenses. Based on the gravitational bar torque indicator $Q_b$, double-gaussian bars are stronger on average than single-gaussian bars, at least for our small sample. We show that published numerical simulations where the bar transfers a large amount of angular momentum to the halo can account for many of the observed profiles. The range of possibilities encountered in models seems well-represented in the observed systems. ",
        "introduction": "S0 galaxies were introduced to the Hubble sequence by Hubble (1936) as a means of bridging the apparently catastrophic gap between E7 and Sa galaxies. After real examples were discovered (see Sandage 1961), the hallmark of the class became a disk shape (definitely ``later than'' E6 or E7) and an absence of spiral arms or star formation. SB0 galaxies were originally classified as SBa by Hubble (1926) even though they also lacked arms. This inconsistency was corrected in Sandage (1961).  Barred S0 galaxies are extremely interesting because they help to take some of the mystery out of S0s in general: a bar is usually a disk feature that is closely related to spiral structure in many galaxies (e.g., Kormendy and Norman 1979). In addition, ring features are directly related to bars in spirals, and SB0s may show vestiges of similar rings. One could therefore ask whether bar properties in S0s might provide any clues as to how S0s and spirals might be related in an evolutionary sense.  We are interested in the distribution of bar strengths in S0 galaxies, a property of very early-type galaxies that has not yet been tapped for what it might tell us about the evolutionary history of S0 galaxies. Early-type galaxies have a remarkable array of bar morphologies whose Fourier and other properties are worth characterizing in more detail. The advantage we have for examining these issues is that early-type galaxies show their bars largely unaffected by dust, star formation, and spiral structure. We can look for subtle structural differences between different bars and isolate possible different bar types.  We are also interested in comparing early-type barred galaxies with the models of Athanassoula (2003, 2005), who demonstrated the importance of angular momentum transfer to the halo as a means of producing strong bars. The Fourier profile information we provide here is ideally suited to comparison with $n$-body models, and for evaluating the symmetry assumption used for bar-spiral separation (Buta, Block, \\& Knapen 2003).  The Near-InfraRed S0 Survey (NIRS0S) is an attempt to obtain a statistically well-defined database of images of S0s from which the properties of S0 bars may be fairly compared to those of spirals. The Ohio State University Bright Galaxy Survey (OSUBGS, Eskridge et al. 2002) provides a valuable dataset for studying the properties of spiral galaxies, and has been fully tapped for bar strength studies by Block et al. (2002), Buta, Laurikainen, \\& Salo (2004=BLS04), Laurikainen, Salo, \\& Buta (2004a), Laurikainen et al. (2004b), and Buta et al. (2005). The questions we address with the NIRS0S sample are: (1) how strong do S0 bars get compared to spiral bars? (2) how does the distribution of bar strengths in S0 galaxies compare with that for spirals? (3) what characterizes the morphology of bars in S0 galaxies? and (4) what are the near-IR luminosity ratios and profile characteristics of bulges in S0 galaxies?  In describing early-type disk galaxies, we will use specific terminology. An ``ansae\" bar is one showing bright enhancements near the ends (Sandage 1961). A ``regular\" bar does not show such enhancements. Ansae bars are preferentially found in early-type galaxies while regular bars are preferentially found in later types, although no statistical study has quantified the difference. A lens is a feature showing a shallow brightness gradient interior to a sharp edge (Kormendy 1979). An oval is a broad elongation in the light distribution; it differs from a conventional bar in lacking higher-order Fourier terms. If intrinsically elongated, a lens can also be an oval. Ovals are discussed further by Kormendy and Kennicutt (2004).  In this paper, we focus on a set of 26 galaxies from the NIRS0S having ovals and/or bars, and investigate radial profiles of relative Fourier intensity amplitudes. We seek to examine the diversity in early-type bars according to the symmetry of such profiles.  In section 2, we describe first the rationale for the survey and the sample selection criteria. In section 3, we describe the observations made with different instruments and detectors. The study of the relative Fourier intensity amplitudes is presented in section 5. A discussion of the results is presented in section 6. ",
        "conclusions": "Interpretation of the results in this paper requires a theory that accounts not only for bar strength and pattern speed, but also the varied shapes of bars. Studies by Athanassoula (2005 and references therein) show that angular momentum exchange is at the heart of all of these issues. Critical to the properties of bars is how much angular momentum is transferred to the halo. The effect depends on the density of matter in halo resonances and on how cold or hot the resonant material is. In order to absorb angular momentum a halo must be ``live\", as opposed to a rigid halo that cannot interact with other galaxy components. Cold, live halos can absorb so much angular momentum that a bar can grow very strong. Weaker bars form in warmer, smaller halos or rigid halos, while in hot disks, mainly ovals will form. The extreme effects of a live halo interaction has led to the possibility that bars might be found in systems lacking a background disk in the bar region. Gadotti and de Souza (2003) present two possible cases of this, although LSB05 present contrary evidence that does not support their claim.  Athanassoula computes relative Fourier mass profiles for her models that can be compared to the $I_m/I_0$ profiles presented in this paper. Provided the mass-to-light ratio in the $K_s$ band is relatively constant, such a comparison should be fair. The problems of comparing $n$-body models with real bars are summarized by Athanassoula \\& Misiriotis (2002), who discuss the types of bars that develop in massive halo (MH), massive disk (MD), and intermediate models. Considering that the models are not of any specific galaxy, and are only a few of a large number of models actually computed, the relative Fourier mass profiles well resemble the observed profiles for some of our galaxies. The MH models show strong $m$=2, 4, 6, and 8 components while these are much weaker for the MD models, with $m$=6 and 8 being in the noise. The MH profiles most resemble those for NGC 1452, 2983, 4608, and 4643 in our sample. The model bars have a fairly sharp outer edge, as is seen in these galaxies.  Athanassoula (2003) shows how the mass of the halo impacts the Fourier terms. More massive halos lead to stronger bars. Her most massive halo model (M$\\gamma$3) has Fourier profiles similar to those of NGC 1452 and 4643, two of the strongest bars in our sample. Model MH1, with a less massive halo, has profiles similar to NGC 936 and 2787, while model MH2, with the lowest halo mass in the three illustrated, shows mainly an oval bar and has profiles resembling those of NGC 1302 and 2781.  The good agreement between the models and the observations is further shown by the simulations of Athanassoula, Lambert, \\& Dehnen (2005), who evaluate the effects of a central mass concentration (CMC) on the evolution of the bar models in the previous papers. They show that a CMC has an effect on the higher $m$ terms, weakening their importance and leading to more $m$=2-dominated bars. Before the introduction of a CMC in their massive halo (MH) models, the relative Fourier mass profiles (their Figure 4) strongly resemble those we observe for NGC 4643, which has the sharpest outer edge of all the bars in our sample. In this model, the mass ratios $A_4/A_2$ = 0.69, $A_6/A_0$ = 0.47, and $A_8/A_0$ = 0.36, are very comparable with the intensity ratios found using the first gaussian component of NGC 4643: $I_4/I_2$=0.66, $I_6/I_2$=0.45, and $I_8/I_2$=0.29, respectively (from Table 3). Comparable values are found for NGC 1452 and NGC 4608. In their MH models with a CMC, the profiles are more symmetric and the ratios above are considerably reduced. This is as observed for most of the other galaxies in our sample whose bars are weaker. Massive disk (MD) models in this paper also show similar effects, but the bars are weaker than for the MH models and the profile asymmetries are less important.  Bureau \\& Athanassoula (2005) present three $n$-body models designed for deducing diagnostics of edge-on bars. These models are similar to those used by Athanassoula \\& Misiriotis (2002) and include a weak bar, an intermediate strength bar, and a strong bar, with each model bar surrounded by a strong (and largely circular) stellar inner ring. The Fourier decompositions of these barred galaxy models again resemble the galaxies described above. In the strong bar model, the outer end of the bar has a sharp edge, similar to what is seen in NGC 1452 and 4643 in our sample. The strongest circular inner rings in our sample are seen in NGC 1452 and 4608 (Figure~\\ref{images}). Weaker circular inner rings are seen in NGC 936, 1317, and 2787, but in several of our galaxies the rings are highly elongated. Early-type barred galaxies have a wide range in intrinsic inner ring shapes, as is typical of such features (Buta 1995).  The relative Fourier intensity profiles of the bars in our sample could also include some of the effects of bar destruction. The bars in our sample that include overlapping extended ovals could be cases where the bar destruction was in progress at the time the galaxies became relatively deficient in gas. Bournaud \\& Combes (2002) suggested that bar destruction in the absence of external gas accretion can lead to an oval lens feature (see also Kormendy 1979). An extreme example of this in our sample could be NGC 2859, where the impact of the oval lens extends considerably beyond the ends of the bar. Our Fourier analysis is efficient for separating the lens from the bar in this case, based on the disappearance of the higher order Fourier modes in the lens. Such an extended oval is only weakly seen in NGC 1452, 4608, or 4643, as if these bars were at their peaks and had not begun to disintegrate. These same three galaxies also have little or no central oval or secondary bar."
    },
    "0609/astro-ph0609630_arXiv.txt": {
        "abstract": "{The context of this paper is buoyant toroidal magnetic flux ropes, which is a part of flux tube dynamo theory and the framework of solar-like magnetic activity.} {The aim is to investigate how twisted magnetic flux ropes interact with a simple magnetized stellar model envelope---a magnetic ``convection zone''---especially to examine how the twisted magnetic field component of a flux rope interacts with a poloidal magnetic field in the convection zone.} {Both the flux ropes and the atmosphere are modelled as idealized 2.5-dimensional concepts using high resolution numerical magneto-hydrodynamic (MHD) simulations.} {It is illustrated that twisted toroidal magnetic flux ropes can interact with a poloidal magnetic field in the atmosphere to cause a change in both the buoyant rise dynamics and the flux rope's geometrical shape. The details of these changes depend primarily on the polarity and strength of the atmospheric field relative to the field strength of the flux rope. It is suggested that the effects could be verified observationally.} {} ",
        "introduction": "Buoyant magnetic flux tubes are an essential part of the framework of current theories of dynamo action in both the Sun and solar-like stars: it is widely believed that flux tubes are formed by some combination of rotational shear and turbulent convection near the bottom of the convection zones (CZ) of these stars. When sufficiently buoyant the magnetic flux rise in the form of tubular (toroidal) $\\Omega$-shaped loops of magnetic field lines. Rising under the influence of rotational forces, the loops finally emerge as slightly asymmetric and tilted bipolar magnetic regions at the surface, see e.g.\\ \\cite{Fan+ea94} and \\cite{Caligari+ea95}. During the last decade, it has been established that if they do exist, these flux tubes must in fact be flux {\\em ropes} of intertwined or twisted field lines; otherwise they would quickly be disrupted by a magnetic ``mushrooming'' instability, see e.g.\\ \\cite{Emonet+Moreno1998} and \\cite{Dorch+Nordlund1998}. It has been shown that the instability is inhibited if the degree of systematic twist is sufficiently high, corresponding to a critical value of the field line pitch angle $\\Psi_c$ approximately determined by equality between the energy density of the twisted field component and the ram pressure due to the rise of the flux rope e.g.\\ \\cite{Moreno+Emonet1996}.  A fundamental ingredient in all theories of solar-like dynamos is a process where a relatively weak poloidal magnetic field is turned into a toroidal field by differential rotation: it is this toroidal field that eventually gives rise to the flux ropes, which at the end of their existence return a poloidal field component to the CZ, thereby completing the magnetic activity cycle. This dynamo process is gradual and different cycles overlap.  Buoyant magnetic flux ropes have been extensively study through numerical simulations, cf.\\ the reviews in \\cite{Dorch2002} and \\cite{Fan2004}, and the most recent high resolution study in two dimensions by \\cite{Cheung+ea2006}. However, so far studies of the interaction of flux ropes with poloidal fields have solely dealt with flux ropes that emerge into the solar corona, e.g.\\ \\cite{Archontis+ea04}, i.e.\\ not with a magnetic CZ.  The question addressed in this paper is how the polarity and strength of the field in a magnetized convection zone may affect the rise and evolution of twisted flux ropes. When the predominantly toroidal buoyant flux ropes rise through the CZ, they encounter a poloidal magnetic field that has a component perpendicular to the ropes' axes. However, since the flux ropes are twisted, the transversal field components may be either parallel or anti-parallel to the average magnetic field in the CZ, i.e.\\ the field in the two-dimensional plane perpendicular their toroidal axis.  Several questions relevant to the theory of flux tube dynamos then emerge: e.g.\\ is the rise faster or slower and do more or less magnetic flux reach the surface in the presence of a magnetized convection zone; how does a magnetized convection zone affect the required amount of twist. In this paper I attempt to shed light on these questions by presenting results from idealized numerical 2.5-dimensional MHD simulations of the cross-sections of buoyant flux ropes that interact with a horizontal magnetic layer while they rise. ",
        "conclusions": "\\label{discussion.sec}  It can be argued that two-dimensional models have several restrictions compared to the full three-dimensional case: e.g.\\ there is a tendency for enhanced flows due to the lack of the extra degrees of freedom associated with three dimensions. Therefore, one can consider the questions posed in this paper to be addressed here only to a certain order.  On the one hand, the initial naive assumption that one could have---that the effect of the magnetized layer would be to destroy a flux rope with a twist anti-parallel to the polarity of the layer---is wrong. In fact, the opposite is true: even though reconnection within the magnetic layer reduces the twist of the flux ropes, it is easier for them to penetrate the poloidal layer, because the reconnection of field lines allows the ropes to ``carve'' their way through the layer. This is provided, however, that the flux ropes contain enough transversal flux.  On the other hand, the simple estimate of $\\chi_{\\rm c} \\approx 5$ in Eq.\\ \\ref{chi_c.eq} seems to hold in the simulations: in the models with $\\chi < \\chi_{\\rm c}$, the ropes' buoyancy are strongly damped. However, while ropes with $\\chi > \\chi_{\\rm c}$ are allowed to rise, they may still be halted in the CZ, but the more correct critical limit will have to be a function also of twist, i.e.\\ $\\chi_{\\rm c} = \\chi(\\epsilon)$.  It has been shown here that the flux ropes' dynamics depend on both the strength and sign of the poloidal field: for a relatively strong poloidal field the rise is slower and therefore the flux ropes reach a lower height in the same amount of time. In certain cases the rise can be completely halted. It turns out that anti-parallel twisted ropes reach higher and faster. When it comes to the flux ropes' topology, the geometrical shape of ropes also depends on the strength and sign of the poloidal field: the apex of anti-parallel twisted ropes are flatter and have steeper magnetic gradients in their axial field. Ideally this could be observed as a seemingly faster passage through horizontal layers (emergence) of the axial part of the flux rope, relative to cases with more moderate vertical gradients in the axial field: e.g.\\ \\cite{Archontis+ea04} have shown that in fact a steep axial gradient is required by flux emergence (their Eq. 10). As to the flux ropes' twisted field components, reconnection reduces the twist of anti-parallel ropes. If ropes on either sides of the equator are oppositely twisted, this could then in principle be falsified observationally by comparing observations of poloidal field strength in emerging northern and southern active regions.  The most fundamental problems remaining are those of the origin of the twist, and the question of how it arises. One may speculate that twisted field lines could be generated in large-scale flux bundles located near the bottom of the convection zone connecting across the solar equator: such flux bundles would experience a rotating motion since their lower parts are located in a region rotating slower than their uppermost parts. This rotation would twist the magnetic field lines and this twist could be transmitted to the parts of the flux bundle at slightly higher latitudes, thereby possibly giving rise to a twisted toroidal flux system. The sign of a twisted component generated in this geometric manner would always be pointing in the same direction (northwards in the Sun). Hence, in consecutive 11-year half-cycles of the solar dynamo, the twist would alternate between being anti-parallel and parallel to the poloidal field (from which the corresponding toroidal field was generated in the respective cycle). Such an alternation would result a 22-year cycle in the amount of twist in the emerging parts of the toroidal flux ropes, e.g.\\ in the helicity of bipolar magnetic regions.  Fully three-dimensional MHD simulations extending the present study are under way, using the numerical scheme of \\cite{Archontis+ea04}."
    },
    "0609/hep-ph0609305_arXiv.txt": {
        "abstract": "The LVD detector, located in the INFN Gran Sasso National Laboratory (Italy), studies supernova neutrinos through the interactions with protons and carbon nuclei in the liquid scintillator and interactions with the iron nuclei of the support structure. We investigate the effect of neutrino oscillations in the signal expected in the LVD detector. The MSW effect has been studied in detail for neutrinos travelling through the collapsing star and the Earth. We show that the expected number of events and their energy spectrum are sensitive to the oscillation parameters, in particular to the mass hierarchy and the value of $\\theta_{13}$, presently unknown. Finally we discuss the astrophysical uncertainties, showing their importance and comparing it with the effect of neutrino oscillations on the expected signal. ",
        "introduction": "There are many experimental works suggesting neutrino conversion among flavors in the recent few years, through the study of atmospheric \\cite{skatm}, solar \\cite{Cl} \\cite{Ga} \\cite{Sage} \\cite{Kam} \\cite{sksol} \\cite{sno}, reactor \\cite{KamLAND} and accelerator \\cite{K2K} neutrinos. The interpretation of all these phenomena in terms of neutrino oscillations is rather robust, because it is able to include all the experimental data (except the ``not yet confirmed'' LSND \\cite{lsnd} signal), even if the expected oscillatory behavior, in terms of the observable $L/E$, has not been yet experimentally observed (preliminary results that show a low significance hint for a oscillatory behavior have been found by a re--analysis of the SK data \\cite{skatm}). An interesting fact is that the inclusion of the MSW effect \\cite{msw} permits a consistent interpretation of KamLAND results and the `high energy' solar neutrino data \\cite{Kam,sksol,sno}.  In the standard three flavor scenario, six parameters must be determined by oscillation experiments: 3 mixing angles ($\\theta_{{\\rm sol}}$, $\\theta_{13}$, $\\theta_{{\\rm atm}}$ ), 2 squared mass differences ($\\Delta m^2_{{\\rm sol}}$ and $\\Delta m^2_{{\\rm atm}}$) and 1 CP-violation phase $\\delta$. A recent analysis of all the available experimental data \\cite{VissStrum05}  constrains the ``atmospheric'' and ``solar'' parameters to be in the following $99\\% ~C.L.$ ranges (compare also with the results in \\cite{FLglobal}):  \\begin{table}[h] $$\\begin{array}{lrlc} \\hbox{Oscillation parameter}&\\multicolumn{2}{c}{\\hbox{central value}} &\\hbox{$99\\%$ C.L. range}\\\\  \\hline \\hbox{solar mass splitting} & \\Delta m^2_{{\\rm sol}} ~= & (8.0\\pm 0.3)\\,10^{-5}~\\eV^2 &  (7.2\\div 8.9)\\,10^{-5}\\eV^2 \\\\ \\hbox{atm. mass splitting~~~~}  & |\\Delta m^2_{{\\rm atm}}| ~= & (2.5\\pm 0.3) \\,10^{-3}~\\eV^2~~~ & (1.7\\div 3.3) \\,10^{-3}\\eV^2\\\\ \\hbox{solar mixing angle} &  \\tan^2 \\theta_{{\\rm sol}} ~=&  0.45\\pm0.05 & 30^\\circ < \\theta_{{\\rm sol}}<38^\\circ \\\\ \\hbox{atm. mixing angle} & \\sin^2 2\\theta_{{\\rm atm}} ~= &   1.02\\pm 0.04 &36^\\circ <\\theta_{{\\rm atm}}< 54^\\circ\\\\ \\end{array}$$ \\label{tab1} \\end{table}  However the other parameters are not completely determined: the $\\theta_{13}$ mixing angle is only upper limited, mainly by the Chooz experiment data \\cite{Chooz} ($\\sin^2 \\theta_{13} < 3.~ 10^{-2}$  at the $99 \\%~ C.L.$), the sign of $\\Delta m^2_{{\\rm atm}}$ (that fixes the so--called mass hierarchy) is completely unknown, as well as the CP--violation phase $\\delta$.  Because of the wide range of matter density in the stellar envelope, a supernova explosion represents a unique scenario for further study of the neutrino oscillation mixing matrix. Indeed neutrinos can cross two resonance density layers and therefore the resulting possible mixing scenarios are different from the solar ones. The emerging neutrino spectra are sensitive to the sign of $\\Delta m^2_{{\\rm atm}}$ and to the value of $\\theta_{13}$.  Before proceeding, it is important to recall that, at present, there is not a unique theory of supernova explosions. Till now, numerical investigations of the ``standard model'' based on a delayed scenario of the explosion \\cite{Olga0} failed to reproduce the explosion. On top of that, other models are being studied where rotation \\cite{Olga1} or magnetic field \\cite{Olga2} play an essential role. In the following, we will use a simple description of the neutrino flux that does not contradict the SN1987A events seen by Kamiokande-II \\cite{Olga3}, IMB \\cite{Olga4} and Baksan \\cite{Olga5}, see e.g. \\cite{Olga6} for a discussion, although it is not able to take into account the events seen in Mont-Blanc observatory \\cite{Olga7,Oscar0}. This ``standard'' description, however, corresponds to the expected neutrino emission in the delayed scenario and in the last phase of the collapse with rotation \\cite{Olga8}. For this reason, we take it as a useful starting point for the investigation of the impact of oscillations in the neutrino signal from a supernova.  The main aim of this paper is to show how neutrino oscillations affect the signal detected by the LVD observatory in the INFN Gran Sasso National Laboratory, Italy. We also evaluate the impact on the signal of the astrophysical parameters of the supernova explosion mechanism, such as the total energy emitted in neutrinos, the star distance, the neutrino--sphere temperatures and the partition of the energy among the neutrino flavors.  In section \\ref{se:sn} we describe the characteristics of the neutrino fluxes emitted during a gravitational core collapse. In section \\ref{se:osc} the neutrino oscillation mechanism is shown, in particular the peculiarities of the MSW effect in the supernova matter and in the Earth. The LVD detector and the relevant neutrino interactions both in the liquid scintillator and in the iron support structure are described in section \\ref{se:lvd}. The impact of neutrino oscillations in the signal expected in the LVD detector is presented in section \\ref{se:res} while the uncertainties in the astrophysical parameters and their effect on the results are discussed in section \\ref{se:ap}. Finally, the conclusions are drawn in section \\ref{se:sum}. Two appendices complete this work, describing in more detail the MSW effect calculation in the Earth (A) and the neutrino interaction with the iron of the LVD support structure (B).  Preliminary results have been presented previously in \\cite{taup01}, \\cite{icrc03} and \\cite{Giacobbe}. ",
        "conclusions": "\\label{se:sum} The main aim of this paper was to show how neutrino oscillations affect the signal expected in the LVD detector at the occurence of the next galactic supernova.  The LVD detector has been described in its main components. It is able to detect neutrinos of all flavors, by studying them in the various CC and NC channels. All the neutrino interactions that occur in the liquid scintillator as well as in the iron support structure have been studied in detail taking into account the neutrino energy threshold, cross section and detection efficiency.  We assumed a galactic supernova explosion at a typical distance of $D = 10$~kpc, parametrized with a pure Fermi--Dirac energy spectrum ($\\eta = 0$) with a total energy $E_b = 3 \\cdot 10^{53}$ erg and perfect energy equipartition $f_{\\nu_e}=f_{\\bar\\nu_e}=f_{\\nu_x}=1/6$; we fixed $T_{\\nu_x} / T_{\\bar{\\nu}_e} = 1.5$, $T_{\\nu_e} / T_{\\bar{\\nu}_e}  = 0.8$ and $T_{\\bar{\\nu}_e} =5~{\\rm MeV}$.  We considered neutrino oscillations in the standard three-flavor scenario. The MSW effect has been studied in detail for neutrinos travelling through the supernova matter. We also considered the distortion in the expected neutrino spectra induced by a possible path inside the Earth before their detection.  For the chosen supernova parameters, it results that the expected number of events and their energy spectrum depend on the unknown oscillation parameters: the mass hierarchy and the value of $\\theta_{13}$.  In particular, the inverse beta decay interactions ($\\anue p, e^+ n$) are highly sensitive to the mass hierarchy: for adiabatic transition, the number of events increases of $\\sim 25 \\%$ in the IH case, with respect to the NH one, since the detected $\\anue$ completely come from the higher energy $\\nux$. The mean energy of the detected positrons is correspondingly increased.  The total number of $(\\nu_e + \\bar\\nu_e)$ CC interaction with $^{12}$C nuclei is highly increased taking into account neutrino oscillations, because of their high energy threshold. For adiabatic transition the expected number of events is higher than the non adiabatic one, because at least one specie (between $\\nu_e$ or $\\bar\\nu_e$) comes significantly from the original and higher--energy $\\nux$ in the star. However, if it is not possible to discriminate between $\\nu_e$ and $\\bar\\nu_e$, the normal and inverted hierarchy cases present similar results. Indeed, in the NH (IH) case, the increase in $\\nue$ ($\\anue$) is compensated by a decrease in $\\anue$ ($\\nue$).  The neutrino interactions with the iron of the support structure, which are studied in detail in this work, are also incread by the oscillations. The efficiency for the detection of the produced charged leptons and gammas in the active part of the detector has been calculated with a full simulation of the apparatus. The contribution of $(\\nu_e+\\bar \\nu_e)$ {\\rm Fe} interactions can be as high as $17\\%$ of the total number of events (in the adiabatic NH case) and they contribute mostly to the high energy part of the spectrum.  With respect to the previous detection channels, the number of NC interactions with $^{12}$C nuclei does not depend on oscillations. In principle they could be used as a reference to identify the $\\nux$--sphere temperature. However, this is partly limited by the uncertainties in the other astrophysical parameters.  We completed the calculations taking into account the effect of the passage of neutrinos through the Earth before their detection. This induces a characteristic modulation in the energy spectrum; however, given the expected number of events and the assumed oscillation parameters, the effect is quite weak.  In conclusion, for the choice of the astrophysical parameters adopted in this work, the expected signal of neutrinos in the LVD detector from a supernova core collapse greatly benefits of the neutrino oscillation mechanism, practically in all the possible detection channels, especially if the transition is adiabatic and the hierarchy inverted (since in LVD the most relevant signal is given by $\\anue$).  However, being aware of the fact that the astrophysical parameters of the supernova mechanism are up to now not well defined, we performed the same calculations using different values of them. The resulting differences are in fact important; they are mainly due to the poor theoretical knowledge of the physics of the gravitational collapse. This will be hopefully improved after the occurence and detection of the next galactic supernova, to which the LVD experiment can give a significant contribution, thanks to its cabability to observe and measure neutrino events of several types.  \\appendix  \\newpage {\\bf \\Large Appendices}"
    },
    "0609/nucl-th0609074_arXiv.txt": {
        "abstract": "A method for calculation of Gamow-Teller transition rates is developed by using the concept of the Projected Shell Model (PSM). The shell model basis is constructed by superimposing angular-momentum-projected multi-quasiparticle configurations, and nuclear wave functions are obtained by digonalizing the two-body interactions in these projected states. Calculation of transition matrix elements in the PSM framework is discussed in detail, and the effects caused by the Gamow-Teller residual forces and by configuration-mixing are studied. With this method, it may become possible to perform a state-by-state calculation for $\\beta$-decay and electron-capture rates in heavy, deformed nuclei at finite temperatures. Our first example indicates that, while experimentally known Gamow-Teller transition rates from the ground state of the parent nucleus are reproduced, stronger transitions from some low-lying excited states are predicted to occur, which may considerably enhance the total decay rates once these nuclei are exposed to hot stellar environments. ",
        "introduction": "The knowledge on weak interaction processes is one of the most important ingredients for resolving astrophysical problems. The first systematical work on stellar weak-interaction rates was performed by Fuller, Fowler, and Newman \\cite{FFN1,FFN2,FFN3,FFN4}, who recognized the decisive role played by the Gamow-Teller (GT) transitions. Due to their pioneering work, the study of stellar weak interaction rates requested by astrophysics becomes essentially a nuclear structure problem, in which the actual decay rates are determined by the microscopic inside of nuclear many-body systems. It has been suggested that the nuclear shell model, i.e. a full diagonalization of an effective Hamiltonian in a chosen model space, is the most preferable method for GT transition calculations. This was noticed early by Aufderheide {\\it et al.} \\cite{ABRM93}, and has recently been emphasized by Langanke and Mart\\'inez-Pinedo \\cite{LM03}.  For a theoretical model employed in GT transition calculations, it is generally required that the model can reproduce a wide range of structure properties of relevant nuclei. It has been shown that the state-of-the-art shell-model diagonalization method is indeed capable of performing such calculations. For example, Wildenthal and Brown \\cite{Wild84,BW85} obtained nuclear wave functions in the full $sd$-shell model space, which were successfully applied to calculation of GT rates in the $sd$ shell nuclei \\cite{Oda94}. By using the method developed by the Strasbourg-Madrid group \\cite{SM94}, Langanke and Mart\\'inez-Pinedo \\cite{LM01} made the shell-model GT rates available also for the $pf$ shell nuclei. Still, these sophisticated calculations are tractable only for nuclei up to the mass-60 region, and cannot be applied to heavier nuclei which play important roles in the nuclear processes in massive stars.  In the long history of the nuclear shell-model development, tremendous effort has been devoted to extending the shell-model capacity from its traditional territory to heavier shells. Despite the great progress made in recent years, it seems impossible to treat an arbitrarily large nuclear system in a spherical shell model framework due to the unavoidable problem of dimension explosion. One is thus compelled to seek judicious schemes to deal with large nuclear systems. The central issue has been the shell-model truncation. There are many different ways of truncating a shell-model space. While in principle, it does not matter how to prepare a model basis, it is crucial in practice to use the most efficient one. In this regard, we recognize the fact that except for a few lying in the vicinity of shell closures, most nuclei in the nuclear chart are deformed. This naturally suggests for a shell model calculation to use a {\\it deformed} basis to incorporate the physics in large systems. That is the philosophy that the Projected Shell Model (PSM) \\cite{PSM} and the several important generalizations \\cite{Sheikh99,Chen01,Sun02,Sun03,Gao06} are based on.  The present article reports on the new development for calculation of GT transition rates in the PSM framework. Before a detailed description of the work, we mention a few attractive features in our approach, which may be relevant for future astrophysical applications.  \\begin{itemize}  \\item The PSM utilizes single particle bases generated by deformed mean-field models yet carries out a shell-model diagonalization like the conventional shell model. Conceptually, the PSM bridges two important nuclear structure methods: the deformed mean-field approach and the conventional shell model, and takes the advantages of both. On the one hand, as a shell model, the PSM can be applied to any heavy, deformed nuclei without a size limitation. On the other hand, unlike the mean-field models or models with an average nature, the PSM wave functions contain correlations beyond mean-field and the states are written in the laboratory frame having definite quantum numbers such as angular-momentum and parity. These are needed properties when the wave functions are employed in transition calculations.  \\item Because of the way the PSM constructs its basis, the dimension of the model space is small (usually in the range of $10^2 - 10^4$). With this size of basis, a state-by-state evaluation of GT transition rates is computationally feasible. This feature is important because in stellar environments with finite temperatures, the usual situation is that the thermal population of excited states in a parent nucleus sets up connections to many states in a daughter by the GT operator. However, our current knowledge on GT transitions from excited nuclear states is very poor, and in many cases, it must rely on theoretical calculations.  \\item The calculation of forbidden transitions involves nuclear transitions between different harmonic oscillator shells and thus requires multi-shell model spaces. Such a calculation is not feasible for most of the conventional shell models working in one-major shell bases. The PSM is a multi-shell shell model. This feature is desired particularly when forbidden transitions are dominated.  \\item Isomeric states belong to a special group of nuclear states because of their long half-lives. The existence of isomeric states in nuclei could alter significantly the elemental abundances produced in nucleosynthesis. There are cases in which an isomer of sufficiently long lifetime can change the paths of reactions taking place and lead to a different set of elemental abundances \\cite{Sun05a,Sun05b}. It has been shown that the PSM is indeed capable of describing the detailed structure of isomeric states \\cite{Sun04}.  \\end{itemize}  The paper is organized as follows. In Sec. II, we briefly introduce the PSM concept and describe how shell model diagonalization is carried out in the angular-momentum-projected bases. Technical details of calculation of transition matrix elements in the projected bases are given in Sec. III. Our first example of a GT transition calculation is illustrated in Sec. IV, where we first validate the model by comparing the calculated structure properties (energy levels and electromagnetic transitions) with experiment. The obtained GT transition results are then discussed and the effects caused by the Gamow-Teller residual forces and by configuration mixing are studied. Finally, the work is summarized and an outlook on future applications is given in Sec. V. ",
        "conclusions": "In this article, we have presented the new development of a shell model method for calculation of Gamow-Teller transition rates. The method is based on the Projected Shell Model. Different from the conventional shell model, which builds its states in a spherical basis, the PSM constructs its states in a deformed basis in which important nuclear correlations are taken into account very efficiently. Therefore, it is possible for a shell model diagonalization in the PSM to be carried out in a manageable space for medium to heavy, and even for super-heavy nuclei.  We have shown how the GT transition matrix elements are calculated in the PSM framework. A computer code has been developed and been tested. One nontrivial test has been done through the Ikeda sum-rule. We have obtained a reasonable distribution of the B(GT) strength with a fulfillment of the sum-rule. We have presented the first example from the rare earth region. In the calculation of $^{164}$Ho $\\rightarrow$ $^{164}$Dy electron-capture process, we have predicted the GT transition rates for the excited states. Such rates should be included as part of the total rate when these states are thermally populated in hot stellar environments.  In the calculation discussed so far, we have considered only allowed $\\beta$-decay for the low-lying states. Calculation of allowed $\\beta$-decay for states with high excitation and the forbidden transitions is possible in the PSM framework. Study of GT giant resonance is also under consideration.  The method described in the present article can be applied to various fields such as nuclear astrophysics and fundamental physics, where weak interaction processes take place in nuclear systems \\cite{ALW05}. In particular, one may find interesting applications to cases where a laboratory measurement for certain weak interaction rates is difficult and where the conventional shell model calculations are not feasible. Potential applications in nuclear astrophysics are calculations of $\\beta$-decay rates for the r-process \\cite{r-process} and the rp-process \\cite{rp-process} nucleosynthesis, and electron-capture rates for the core collapse supernova modelling \\cite{LM03,SN}. In the double-$\\beta$ decay theory, theoretical calculations for the nuclear matrix elements are needed, for which one has relied on the Quasiparticle Random Phase Approximation \\cite{Civitarese98,double-beta}, particularly when heavy nuclei are involved. We expect that the method presented here can make important contributions to all these studies."
    },
    "0609/astro-ph0609156_arXiv.txt": {
        "abstract": "{} {We use 9.5-yr of BiSON Sun-as-a-star data to search for dependence of solar-cycle parameter changes on the angular degree, $\\ell$, of the data. The nature of the Sun-as-a-star observations is such that for changes measured at fixed frequency, or for changes averaged across the same range in frequency, any $\\ell$ dependence present carries information on the latitudinal distribution of the agent (i.e., the activity) responsible for those changes.} {We split the 9.5-yr timeseries into contiguous 108-d pieces, and determine mean changes in the damping of, power in, and energy supplied to the modes through the solar cycle. We also apply a careful correction to account for the deleterious effects of the ground-based BiSON window function on the results.} {From our full analysis we obtain a marginally significant result for the damping parameter, where the mean change is found to be weakest at $\\ell=0$. The other parameters show hints of some dependence in $\\ell$.} {Our main conclusion is that the mean fractional solar-cycle change in the $\\ell=0$ damping rates is approximately 50\\,\\% smaller than was previously assumed. It had been common practice to use an average over all low-$\\ell$ modes; our downward revision of the radial-mode value has implications for comparisons with models of the global solar cycle changes, which are usually based on a spherically symmetric geometry.} ",
        "introduction": "\\label{sec:intro}  The fact that damping rates and powers of the global p modes change through the solar cycle is now well established \\citep[e.g., ][]{chaplin00,komm00,salabert03,chano03,chano04,salabert06}. At least where the main part of the p-mode spectrum is concerned damping gets heavier, and observed power gets weaker, as the level of solar activity increases. Information on the damping and power parameters comes straightforwardly from the observations. The damping rates are assumed to be linearly related to the linewidths of the resonant peaks in the frequency power spectrum; while the powers are proportional to the product of the peak widths and heights (the latter more formally termed the maximum power spectral densities).  \\citet{komm02} took advantage of the large number of components available to analyze in the medium-degree range (from $\\ell=40$ to 80) the latitudinal dependence of the width and height changes. They found that the changes were concentrated in latitudes occupied by the active regions. The damping rates and powers, like the mode frequencies, therefore seemed to be responding in some fashion to changes wrought on the active regions by the changing magnetic fields.  Most of the low-$\\ell$ results have come from the Sun-as-a-star data (e.g., from BiSON, GOLF, IRIS and VIRGO/SPM; or from GONG and MDI data combined into a Sun-as-a-star-like proxy). These data show only the even $\\ell+m$ components; the odd components are so weak as to be unobservable. Since it is the outer, sectoral components (with $\\ell=|m|$) that appear most prominently, the Sun-as-a-star mode parameters estimated by the usual analysis methods are dominated by, and therefore \\emph{close to}, the sectoral values. This characteristic gives a noticeable change in the latitudinal sensitivity from $\\ell=0$ to $\\ell=1$, 2 or 3. The spatial sensitivity of Sun-as-a-star data at the latter three values of $\\ell$ is weighted toward the lower latitudes, where the active regions reside. The Sun-as-a-star parameters at these $\\ell$ are therefore more sensitive to the solar cycle than are the $\\ell=0$ data. Evidence for $\\ell$ dependence in the shifts of the Sun-as-a-star mode parameters therefore carries information on the latitudinal distribution of the agent responsible for those shifts. (Changes in inertia, at fixed frequency, between these $\\ell$ are very small indeed.)  It is much more difficult to uncover spatial dependence of the mode parameter changes in the low- than in the medium-$\\ell$ data because there are far fewer components to analyze, and uncertainties on the results are commensurately larger. When changes to damping and power were first uncovered in the Sun-as-a-star data, results were averaged over $\\ell$ to reduce errors. Given the modest precision in the results, it was assumed the $\\ell$-averaged values also provided a working proxy of the radial-mode shifts. And so these averages were used as meaningful comparisons for models of global changes in damping \\citep[see, e.g.,][]{houdek01}, models set up in spherical geometry, i.e., pertinent only to the $\\ell=0$ case. With better analysis, and more data, it has become possible to give results for each $\\ell$, and to therefore test whether the $\\ell=0$ shifts really are weaker than an average across $\\ell=0$ to 2 or 3.  We made a first cut at such an analysis in \\citet{chaplin03a}. Our results, from BiSON Sun-as-a-star data, suggested very strongly that the $\\ell=0$ linewidth changes were indeed significantly weaker than at $\\ell=1$ and 2. The clear implication was that results of the theoretical models would now need to be compared to this new, smaller shift. In the analysis of the BiSON data, we had to allow for the corrupting influence on the results of the ground-based BiSON window function. We did so by applying corrections that were designed originally for a different study, using data from a different period. In this paper we revisit our analysis. We design and implement a correction procedure for the 9.5-yr BiSON dataset in question. We find that implementation of this internally consistent, and more accurate, correction gives little change to the results, and reinforces our earlier conclusion. In summary, the mean fractional solar-cycle change in the $\\ell=0$ damping rates is approximately 50\\,\\% smaller than was previously assumed. ",
        "conclusions": "We have analyzed some 9.5-yr of BiSON Sun-as-a-star data -- collected in solar cycles 22 and 23 -- to search for dependence of the mode excitation and damping parameter changes on the angular degree, $\\ell$, of the data. The nature of the Sun-as-a-star observations is such that for changes measured at fixed frequency, or for changes averaged across the same range in frequency, any $\\ell$ dependence present carries information on the latitudinal distribution of the agent (i.e., the activity) responsible for those changes.  We split the 9.5-yr timeseries into contiguous 108-d pieces, and determined mean changes in the damping of, power in, and energy supplied to, the modes through the solar cycle. We also applied a careful correction to account for the deleterious effects of the ground-based BiSON window function on the results. This correction was calibrated by, and then fully tested on, artificial seismic data generated by the solarFLAG mode simulation code.  From our full analysis we obtained a marginally significant result for the damping parameter, where the mean change was found to be weakest at $\\ell=0$, and higher in data on the other $\\ell$. The result implies the damping is strongest in the active regions, confirming results on the more numerous higher-$\\ell$ data \\citep{komm02}. The other excitation and damping parameters we investigated showed hints of some dependence in $\\ell$, but nothing that could judged as statistically significant.  Our main conclusion is that the mean fractional solar-cycle change in the $\\ell=0$ damping rates is approximately 50\\,\\% smaller than was previously assumed. It had been common practice to use an average over all low-$\\ell$ modes (where mean solar-cycle shift values have averaged about 18\\,\\% \\citep{chaplin00,salabert03,chano03,chano04}. Our downward revision of the radial-mode value has implications for comparisons with models of the global solar cycle changes, which are usually based on a spherically symmetric geometry."
    },
    "0609/astro-ph0609683_arXiv.txt": {
        "abstract": "We explore the possibility of having a composite (self-conserved) dark energy (DE) whose dynamics is controlled by the quantum running of the cosmological parameters. We find that within this scenario it is feasible to find an explanation for the cosmological coincidence problem and at the same time a good qualitative description of the present data. ",
        "introduction": "Independent data from different observations\\,\\cite{Supernovae, WMAP3Y,LSS} provide strong support for the existence of DE and seem to agree that it presently constitutes $\\sim70\\%$ of the total energy density. Nevertheless, the nature of DE remains unclear. If we identify it with a cosmological constant (CC) arising from the quantum field theory (QFT) vacuum energy, as done in the $\\CC$CDM model\\,\\cite{Peebles84}, we are led to a value many orders of magnitude greater than the measured one, what has been called the CC problem \\cite{weinRMP,Copeland06}. This problem could be alleviated by means of a dynamical DE. This possibility is supported by some recent studies\\,\\cite{Alam,Jassal1} and has been exploited profusely in various forms \\cite{Copeland06}. Among them the scalar fields are the most paradigmatic scenario. It must be stressed though that the presence of a scalar field is not essential for a model to be described in terms of an effective EOS, $\\pD=\\we\\,\\rD$ (for instance, this has been proven for any model with variable cosmological parameters in \\cite{SS12}).  We present here a model in which the DE, in addition of being dynamical, is allowed to be composite. This model may offer an explanation to the ``cosmological coincidence problem''\\,\\cite{weinRMP} -i.e. to the fact that the DE and matter densities happen to be similar precisely at the present epoch- by keeping the ratio between these two densities bounded and of order 1 during the entire Universe existence. At the same time, the effective EOS of our model can match the available data. This feature is exemplified through specific numerical examples. ",
        "conclusions": "We have shown that the $\\CC$XCDM model can be in good agreement with present data and provide a solution to the cosmological coincidence problem as well as a clear signature. In our opinion the next generation of high precision cosmology experiments (DES, SNAP, PLANCK)\\,\\cite{SNAP} should consider the possibility of a composite DE with dynamics controlled by the running of the cosmological parameters. \\begin{theacknowledgments}  This work has been supported in part by Ministerio de Eduaci\\'on y Ciencia of Spain (MEC) and FEDER under project 2004-04582-C02-01, and also by DURSI under 2005SGR00564. JG was also supported by MEC under BES-2005-7803. The work of HS is financed by the MEC and he thanks the Dep. ECM of the UB for the hospitality.  \\end{theacknowledgments}  \\newcommand{\\JHEP}[3]{{\\sl J. of High Energy Physics } {JHEP} {#1} (#2)  {#3}} \\newcommand{\\NPB}[3]{{\\sl Nucl. Phys. } {\\bf B#1} (#2)  {#3}} \\newcommand{\\NPPS}[3]{{\\sl Nucl. Phys. Proc. Supp. } {\\bf #1} (#2)  {#3}} \\newcommand{\\PRD}[3]{{\\sl Phys. Rev. } {\\bf D#1} (#2)   {#3}} \\newcommand{\\PLB}[3]{{\\sl Phys. Lett. } {\\bf #1B} (#2)  {#3}} \\newcommand{\\EPJ}[3]{{\\sl Eur. Phys. J } {\\bf C#1} (#2)  {#3}} \\newcommand{\\PR}[3]{{\\sl Phys. Rep } {\\bf #1} (#2)  {#3}} \\newcommand{\\RMP}[3]{{\\sl Rev. Mod. Phys. } {\\bf #1} (#2)  {#3}} \\newcommand{\\IJMP}[3]{{\\sl Int. J. of Mod. Phys. } {\\bf #1} (#2)  {#3}} \\newcommand{\\PRL}[3]{{\\sl Phys. Rev. Lett. } {\\bf #1} (#2) {#3}} \\newcommand{\\ZFP}[3]{{\\sl Zeitsch. f. Physik } {\\bf C#1} (#2)  {#3}} \\newcommand{\\MPLA}[3]{{\\sl Mod. Phys. Lett. } {\\bf A#1} (#2) {#3}} \\newcommand{\\CQG}[3]{{\\sl Class. Quant. Grav. } {\\bf #1} (#2) {#3}} \\newcommand{\\JCAP}[3]{{\\sl JCAP} {\\bf#1} (#2)  {#3}} \\newcommand{\\APJ}[3]{{\\sl Astrophys. J. } {\\bf #1} (#2)  {#3}} \\newcommand{\\AMJ}[3]{{\\sl Astronom. J. } {\\bf #1} (#2)  {#3}} \\newcommand{\\APP}[3]{{\\sl Astropart. Phys. } {\\bf #1} (#2)  {#3}} \\newcommand{\\AAP}[3]{{\\sl Astron. Astrophys. } {\\bf #1} (#2)  {#3}} \\newcommand{\\MNRAS}[3]{{\\sl Mon. Not.Roy. Astron. Soc.} {\\bf #1} (#2)  {#3}}"
    },
    "0609/astro-ph0609360_arXiv.txt": {
        "abstract": "We present the age distributions for star clusters and individual stars in the Small Magellanic Cloud (SMC) based on data from the Magellanic Clouds Photometric Survey by Zaritsky and collaborators. The age distribution of the SMC clusters shows a steep decline, $dN_{cluster}/d\\tau \\propto \\tau^{-0.85\\pm0.15}$, over the period $10^7 \\lea \\tau \\lea 10^9$~yr.  This decline is essentially identical to that observed previously for more massive clusters in the merging Antennae galaxies, and also for lower-mass embedded clusters in the solar neighborhood.  The SMC cluster age distribution therefore provides additional evidence for the rapid disruption of star clusters (``infant mortality'').  These disrupted clusters deliver their stars to the general field population, implying that the field star age distribution, $dN_{fld star}/d\\tau$, should have an inverse relation to $dN_{cluster}/d\\tau$ if most stars form initially in clusters.  We make specific predictions for $dN_{fldstar}/d\\tau$ based on our cluster disruption models, and compare them with current data available for stars in the SMC.  While these data do not extend to sufficiently young ages for a definitive test, they are consistent with a scenario wherein most SMC stars formed in clusters. Future analyses of $dN_{fldstar}/d\\tau$ that extend down to ages of $\\sim$ few million years are needed to verify the age relationship between stars residing in clusters and in the field. ",
        "introduction": "The age distribution of a population of star clusters contains important information about the formation and disruption of the clusters\\footnote{We use the term ``cluster'' generically to mean any aggregate of stars -- regardless of mass, size, or age, and whether bound or unbound -- with a density significantly higher than the local stellar background, making it recognizable as a distinct entity.}. The number of embedded clusters in the solar neighborhood (with masses $\\sim10^2-10^3~M_{\\odot}$) declines rapidly with age, and it has been estimated that only $4-7$\\% of these clusters will survive for 100~Myr (see the review by Lada \\& Lada 2003).  We recently found that much more massive (compact) clusters ($\\geq3\\times10^4~M_{\\odot}$) in the merging ``Antennae'' galaxies also show a rapid decline in their age distribution, with $dN_{cluster}/d\\tau\\propto\\tau^{-1}$ (Fall, Chandar, \\& Whitmore 2005, hereafter FCW05).  We interpret this decline in numbers as due to a high rate of early disruption (``infant mortality'') for massive clusters in the Antennae.  Taken together, these age distributions suggest that compact, massive clusters in the chaotic environment of two merging disk galaxies disrupt at approximately the same rate as lower mass, embedded clusters in the more quiescent solar neighborhood.  This suggests that star clusters disrupt rapidly due to processes internal to the clusters themselves (e.g., the removal of interstellar material due to the mass and energy input from massive stars; Hills 1980; Fall 2004; FCW05), and implies that the age distribution should be roughly $dN_{cluster}/d\\tau \\propto \\tau^{-1}$ in all star-forming galaxies.  In contrast to the results for clusters in the Milky Way and in the Antennae, the cluster age distribution in the Small Magellanic Cloud (SMC) has been reported as $dN_{cluster}/d\\tau\\propto\\tau^{-2.1}$ (Rafelski \\& Zaritsky 2005; hereafter RZ05).  However, this age distribution is actually the number of clusters divided by the number of individual stars of the same age, i.e., it is a {\\it normalized} cluster age distribution (RZ05).  Does the proposed $\\tau^{-2}$ power law for the normalized cluster age distribution in the SMC differ from the $\\tau^{-1}$ form found in the Antennae and solar neighborhood because of the normalization procedure, because clusters in the SMC are forming and/or dissolving at a different rate, or for some other reason?  The motivation for normalizing the number of clusters by the number of individual stars was to assess whether the formation of stars in clusters parallels the formation of stars in the field.  An underlying assumption made by RZ05, then, is that star formation occurs in {\\it both} clusters and in the field simultaneously. However, our study of the Antennae galaxies suggests that at least 20\\% and possibly all of the $H\\alpha$ flux is emitted by compact clusters, which disrupt rapidly (FCW05).  This picture implies instead that at least 20\\% and possibly all stars form in clusters, with stars being subsequently delivered from clusters to the field through the process of cluster disruption.  This in turn implies that $dN_{cluster}/d\\tau$ and $dN_{fldstar}/d\\tau$ should have an inverse relationship to each other.  If the SMC cluster system is found to have an unnormalized age distribution $\\propto \\tau^{-2.1}$, then our picture of clusters disrupting at essentially the same rate in all star-forming environments may have to be significantly revised.  What is the unnormalized SMC cluster age distribution, does it differ from the normalized cluster age distribution, and what is the relationship between field and cluster stars?  Addressing these questions is the main focus of this Letter. ",
        "conclusions": "We have shown here that the (unnormalized) cluster age distribution in the SMC has the form $dN_{cluster}/d\\tau \\propto \\tau^{-0.85\\pm0.15}$ and that overall the field star age distribution is nearly flat (i.e., $dN_{fldstar}/d\\tau\\propto \\tau^{0}$, over the age range $10^7 \\lea \\tau \\lea 10^9$~yr).  The {\\it normalized} SMC cluster age distribution (the number of clusters divided by the number of individual stars of the same age) therefore must have the same form as the unnormalized distribution [i.e., $(dN_{cluster}/d\\tau)_{norm} \\propto \\tau^{-0.85\\pm0.15}$].  Our conclusion contradicts that of RZ05, who found that the normalized cluster age distribution is $(dN_{cluster}/d\\tau)_{norm} \\propto \\tau^{-2.1}$ from the same data. The difference in the power-law exponent between our work and that of RZ05 comes about because we divided the cluster age distribution ($dN_{cluster}/d\\tau$) directly by that for the field stars ($dN_{fldstar}/d\\tau$), whereas RZ05 divided $dN_{cluster}/d\\mbox{log}\\tau$ by $dN_{fldstar}/d\\mbox{log}\\tau$, and then appear to have introduced an extraneous factor of $\\tau^{-1}$.  Our result that the age distribution of SMC clusters has a power-law form $dN_{cluster}/d\\tau \\propto \\tau^{-0.85\\pm0.15}$ has important implications regarding the physical processes that disrupt star clusters.  The fact that clusters in vastly different environments (including the chaotic merging of two galaxies, a quiescent spiral disk, and a gas-rich dwarf galaxy) have age distributions with nearly identical forms provides strong evidence that star clusters disrupt rapidly due to processes {\\it internal} to the clusters themselves. Over the first few million years photoionization, stellar winds, jets and supernovae can remove much of the interstellar matter from clusters leaving the member stars freely expanding (e.g., Hills 1980; Boily \\& Kroupa 2003a,b).  These expanding clusters will be observable for $\\sim$~few~$\\times 10^{7}$~yrs after becoming unbound, until their surface brightness becomes so low that they become indistinguishable from statistical fluctuations of stars in the field (FCW05).  In addition to the removal of interstellar material, mass loss from stars themselves will likely contribute to the continued unbinding of clusters (e.g., Applegate 1986; Chernoff \\& Weinberg 1990; Fukushige \\& Heggie 1995).  Finally, two-body evaporation will destroy any remaining clusters on longer timescales (e.g., Vesperini 1997; Baumgardt 1998; Fall \\& Zhang 2001).  The early disruption (``infant mortality'') appears to operate relatively independent of cluster mass, while two-body evaporation is mass-dependent.  The stars from disrupted clusters add to the general field, requiring a close relationship between the age distributions of clusters and field stars if most stars form initially in clusters.  In this work we compared age distributions for clusters and individual stars in the SMC with model predictions to further test this general framework. Our models suggest that the rapid drop in the cluster age distribution should be reflected in a rapid rise in the field star age distribution.  However this rapid rise in the number of field stars is only observable for the first $\\sim10^7$~yrs {\\it regardless of the fraction of stars that initially formed in clusters}.  Because the present data on SMC stars begin just as the predicted field star distributions flatten, and because there is no discrimination between stars belonging to clusters versus to the field, they do not provide an unambiguous answer.  However the observed relations are consistent with a scenario where most stars form in clusters that are subsequently disrupted by the mechanisms discussed above.  The models also predict that even if all stars form initially in clusters, galaxies will contain a relatively small number of observable clusters at any given time.  Future studies with the {\\it Hubble Space Telescope} (which allow better discrimination between stars in clusters and in the general field), and in the infrared (which can probe very young, deeply embedded stars and clusters) could be used to search for the expected rise in the field star age distribution at $\\tau \\lea 10^7$~yrs.  A measurement of the rise in $dN_{fldstar}/d\\tau$ would provide a direct estimate of the fraction of stars which originally form in clusters versus in the field."
    },
    "0609/astro-ph0609010_arXiv.txt": {
        "abstract": "The Southern Galactic Plane Survey (SGPS) is a radio survey in the 21 cm \\ion{H}{1} line and in 1.4~GHz full-polarization continuum, observed with the Australia Telescope Compact Array and the Parkes 64m single dish telescope. The survey spans a Galactic longitude of $253\\degr < l < 358\\degr$ and a latitude of $|b|<1\\degr$ at a resolution of 100 arcsec and a sensitivity below 1~mJy/beam. This paper presents interferometer only polarized continuum survey data and describes the data taking, analysis processes and data products. The primary data products are the four Stokes parameters I, Q, U, and V in 25 overlapping fields of $5\\degr.5$ by $2\\degr$, from which polarized intensity, polarization angle and rotation measure are calculated.  We describe the effects of missing short spacings, and discuss the importance of the polarized continuum data in the SGPS for studies of fluctuations and turbulence in the ionized interstellar medium and for studying the strength and structure of the Galactic magnetic field. ",
        "introduction": "Galactic magnetism is one of the major components in the Milky Way, mostly in equipartition with gas and cosmic rays. Interstellar magnetic fields are believed to profoundly influence the ionized interstellar medium (ISM) through flux freezing and energy dissipation, affect star formation, determine the trajectories and acceleration of low and medium energy cosmic rays and play a major role in the turbulent gas dynamics (see e.g. reviews by Ferri\\`ere 2001, Scalo \\& Elmegreen 2004, Elmegreen \\& Scalo 2004).  Knowledge about the strength and structure of the Galactic magnetic field is still sketchy. Yet, the field has received increasing attention, not only with the objective of studying magnetic fields in galaxies but also because the Galactic magnetized ISM forms a polarized foreground which needs to be determined for Cosmic Microwave Background Polarization \\citep{dto03} and Epoch of Reionization studies \\citep{mh04}.  The only methods to probe Galactic magnetic fields in diffuse ionized gas over a large range of spatial scales are by way of radio polarization and Faraday rotation. Right and left circularly polarized components of radio emission experience birefringence while propagating through a magnetized and ionized medium. This causes the polarization angle of linearly polarized emission $\\phi$ to rotate as a function of wavelength $\\lambda$ as $\\Delta\\phi = \\mbox{RM}\\lambda^2$, where the rotation measure RM is RM~$=0.81\\int n_e \\mathbf{B}\\cdot\\mathbf{ds}$, $n_e$ is the thermal electron density in cm$^{-3}$, $\\mathbf{B}$ is the magnetic field vector in microGauss, $\\mathbf{ds}$ is the path length vector through the medium in parsecs, and the integral is along the line of sight from the observer to the source of polarized emission. Therefore, Faraday rotation measurements allow estimation of the magnetic field component along the line of sight, weighted by the electron density, and integrated over the pathlength. Depolarization characteristics can be used to determine the scale and amplitude of fluctuations in the medium.  The observed polarized radiation used to trace the Galactic magnetic field can come from pulsars, polarized extragalactic sources or diffuse Galactic synchrotron emission (including supernova remnants). All of these sources have their own advantages and disadvantages. Pulsars are unique because model-dependant distance estimates allow constraining of the path length, and because a dispersion measure can be calculated, which in combination with RM yields a direct measure of the magnetic field averaged over the path length. However, they are scarce and distributed mainly in the Galactic plane. Unresolved extragalactic sources, on the other hand, are distributed all over the sky. But they have an intrinsic RM contribution, and currently published datasets yield an average of one source~deg$^{-2}$ in the Galactic plane (Brown et al.\\ 2003, Brown et al., in prep), and only 0.02-0.03 source~deg$^{-2}$ in the rest of the sky (e.g.\\ Simard-Normandin et al.\\ 1981, Broten et al.\\ 1988). Only diffuse synchrotron emission provides a pervasive background of polarized radiation which can be used to form RM maps of large fields in the sky with high resolution \\citep{hkb03a, hkb03b, ulg03, r04}. Diffuse synchrotron emission does suffer from depolarization, which decreases the possible measurements of RM, but which in itself can yield information about the fluctuations in the magneto-ionized interstellar medium \\citep{gdm01}.  Recently the whole sky has been mapped in absolutely calibrated polarized continuum at 1.4~GHz (see Wolleben et al.\\ 2005 for the Northern sky, and Testori et al.\\ 2004 in the South) at a resolution of about half a degree. More than half of the Galactic plane is being surveyed at the much higher resolution of an arcmin, in two separate polarization surveys.  The Canadian Galactic Plane Survey (CGPS, Taylor et al.\\ 2003) covers the Northern Sky at Galactic longitudes $74.2\\degr <l< 147.3\\degr$ and latitudes $-3.6\\degr<b<+5.6\\degr$. The Southern Galactic Plane Survey (SGPS, McClure-Griffiths et al.\\ 2005; hereafter Paper I) extends from $253\\degr < l < 358\\degr$ (Phase~I) and $5\\degr <l<20\\degr$ (Phase~II) at longitudes $|b|<1\\degr$. In both these projects full-polarization continuum and \\ion{H}{1} were observed simultaneously. Although for the \\ion{H}{1} part of the SGPS data from the Australia Telescope Compact Array (ATCA) interferometer and Parkes 64-m single dish have been combined, the polarization measurements discussed here are derived from ATCA observations only.  The \\ion{H}{1} data in the SGPS have been discussed in detail in Paper~I, whereas a test region covering the region $325\\degr.5 < l < 333\\degr.5$ and $-0\\degr.5 < b < +3\\degr.5$ is described in \\citet{gdm01} and \\citet{mgd01}.  This paper describes the radio continuum data of the SGPS Phase~I, i.e.\\ at Galactic longitudes $253\\degr < l < 358\\degr$ (hereafter refered to as SGPS); the continuum data observed in the region $5\\degr<l<20\\degr$ will be discussed elsewhere. In Sect.~\\ref{s:data} we briefly summarize some details about the telescopes, data taking, and calibration. Furthermore, this Section describes the polarization calibration and the effect of missing short spacings.  In Section~\\ref{s:res} we present the data, and Sect.~\\ref{s:science} discusses some of the science done with the polarized SGPS data. ",
        "conclusions": "The polarized radio continuum part of the Southern Galactic Plane Survey (SGPS) provides a major new resource for studies of the ionized interstellar medium and Galactic magnetic field. The SGPS covers almost one third of the Galactic plane with a Galactic longitude coverage of $253\\degr < l < 358\\degr$, at Galactic latitudes $|b|<1\\degr$. The rms sensitivity of the SGPS continuum data is $\\sim$~0.3~mJy/beam increasing to double that number close to the Galactic Center, the instrumental polarization is below 0.1\\% and the resolution is 100\\arcsec. Due to beam depolarization at low resolutions, maps with this high resolution reveal a wealth of polarization structure which would be depolarized in low resolution maps.  The polarized radio continuum data in the SGPS are uniquely suited for studies of the ionized gas and magnetic field strength and structure in the inner Galactic plane, taking into account the constraints of missing large-scale structure. The polarized intensity can be used for depolarization studies to obtain the scale of fluctuations in the magneto-ionized medium. Rotation measure data give information about the electron-density-weighted magnetic field strength along the line of sight, and can be used for statistical analysis of these fluctuations, large-scale modeling of the Galactic magnetic field, or study of individual magnetic objects."
    },
    "0609/gr-qc0609057_arXiv.txt": {
        "abstract": "Cosmological perturbation equations are derived systematically in a canonical scheme based on Ashtekar variables. A comparison with the covariant derivation and various subtleties in the calculation and choice of gauges are pointed out. Nevertheless, the treatment is more systematic when correction terms of canonical quantum gravity are to be included. This is done throughout the paper for one example of characteristic modifications expected from loop quantum gravity. ",
        "introduction": "The backbone of most of current cosmology is the theory of perturbation equations for metric modes around an isotropic space-time \\cite{CosmoPert}. It is used in particular for cosmological structure formation and for testing alternative theories beyond general relativity such as quantum gravity candidates. The underlying equations of typical interest are the linearized Einstein's equations, and so it is straightforward to include corrections if they come from a Lagrangian modified by quantum or other effects. This is, in fact, the situation encountered in most studies of so-called trans-Planckian issues for the effect of quantum gravity on structure formation. Modifications derived from a Hamiltonian formulation as it is used in canonical quantizations can, however, not be implemented in this direct way. Since several effective modifications to Hamiltonians have been derived in recent years in particular within the framework of loop quantum gravity \\cite{Rov,ThomasRev,ALRev,LivRev}, it is of interest to re-derive cosmological perturbation equations in a purely Hamiltonian fashion starting from the gravitational Hamiltonian. This is done in detail in this paper in a derivation based, as loop quantum gravity itself, on real Ashtekar variables \\cite{AshVar,AshVarReell}. We will present a detailed derivation for scalar modes in longitudinal gauge around a spatially flat model, pointing out several subtleties compared to the Lagrangian derivation. Our analysis treats gravitational and matter terms on the same footing, showing how all of them can be obtained from the total Hamiltonian. This presents the classical basis for a systematic investigation of effective perturbation equations and cosmological implications resulting from canonical quantum gravity. In our calculations, only one type of corrections (from inverse powers of metric components in the Hamiltonian) is used, and their implications are discussed. ",
        "conclusions": "We have presented in detail the Hamiltonian derivation of cosmological perturbation equations for scalar modes in longitudinal gauge around a flat isotropic background. The same scheme, of course, applies to other gauges and also under inclusion of vector and tensor modes and for perturbations around different backgrounds. As in the case of scalar modes, due care has to be taken in deciding when a gauge or mode selection is to be specified.  Since our main interest is to compute corrections from a canonical quantization of gravity, typical such correction functions have been included. We have seen how simple modifications of the constraints can propagate to more involved corrections of equations of motion. We emphasize that we have not presented a complete set of effective equations including all possible correction terms. Alternative corrections can arise, and moreover gauge issues have to be studied.  For applications, it is important to note that not only coefficients in evolution equations are corrected, but also constraints are modified. Since constraints generate gauge transformations, the form of gauge invariant variables changes, too. For instance, it is not sufficient to take the classical expression of the gauge invariant curvature perturbation ${\\cal R}=\\psi + {\\cal H}\\delta\\varphi/\\dot{\\varphi}$ and use corrected equations of motion for all variables involved. A complete treatment requires correction terms in matter and metric equations as derived above, as well as in expressions for the relevant quantities to be related to observations. Ignoring any of the ingredients in general can lead to misleading conclusions. Nevertheless, some qualitative conclusions can be drawn. For instance, a modified evolution equation for ${\\cal R}$ can be derived which implies correction terms leading to a slight non-conservation of this curvature perturbation \\cite{InhomEvolve}. While the quantity ${\\cal R}$ itself will have to be corrected as the relevant gauge-invariant quantity, implying additional corrections to the evolution of curvature perturbations, this is unlikely to happen in such a way that all corrections from equations of motion and gauge invariance properties conspire to cancel each other. For the precise form of non-conservation, however, all these effects have to be taken into account. Still, interesting qualitative effects for cosmological phenomenology have already materialized.  We have started here a program to derive effects systematically and presented a first set of corrected constraints as well as evolution equations. A systematic study of different gauges and of observable implications is still to be done. The derivation in a Hamiltonian formulation as well as the use of Ashtekar variables are crucial for the inclusion of effective quantum gravity effects in modified perturbation equations if canonical quantum gravity, in particular loop quantum gravity, is employed.  Primary dynamical objects are then the constraints, rather than Lagrangians, which are modified by quantum effects. Without regarding gravitational parts and all matter energy-momentum terms, changes to the classical behavior in an inflationary context have been considered in \\cite{PowerLoop,ScaleInvLQC,CCPowerLoop} in a strongly modified regime of background correction functions and in \\cite{PowerPert} in a perturbative regime. Ignoring corrections in gravitational parts of the equations corresponds to choosing a flat gauge in which no metric perturbations are present. The availability of this gauge choice is based on classical reasoning, and has to be reconsidered with gauge transformations generated by the quantum corrected constraints. Our treatment is more general since we allowed metric perturbations $\\psi$ and matter perturbations $\\delta\\varphi$ to be independent before they are to be combined to a quantity gauge invariant under the quantum corrected transformations. This allowed us to discuss non-conservation of curvature perturbations as a new effect.  Following the lines of derivations in this paper, basic effects in the constraints then translate unambiguously into effects in perturbation equations. Since several of the variational equations have to be combined in different manners, even simple modifications in the constraints can have complex implications at the level of perturbation equations. Modifications one expects on general grounds are regular versions of any inverse power of metric variables such as those of $p$ in the Hamiltonian constraint, the spin connection and the matter Hamiltonian \\cite{QSDV,Ambig,AmbigConstr,Spin,SphSymmHam}, higher order corrections as powers of $k$ \\cite{SemiClass,DiscCorr,SemiClassEmerge} and higher derivative terms in space as well as in time \\cite{Karpacz}. All this gives rise to characteristic correction functions which can be computed at least qualitatively. In the perturbation equations, coefficients as well as the derivative order of the equations can then change and differ considerably from the classical ones in strong quantum regimes.  We have illustrated this throughout the paper with corrections which are expected from inverse power modifications. Those corrections are easiest to implement and to deal with because they change only coefficients but not the type of perturbation equations. They are also expected to be stronger in inhomogeneous situations \\cite{InhomLattice}. One expects four different correction functions, two for the gravitational Hamiltonian and two for the matter Hamiltonian. When they equal one, classical behavior is reproduced, while on small scales they can differ considerably from one and lead to modified and new coefficients. On very small scales, i.e.\\ in regimes where correction functions are not Taylor expandable around the value one, cosmological perturbation theory is more difficult to apply.  There are many effects from quantum gravity in combination, and even different implementations depending on the quantization scheme used for constraints. An effective analysis shows which of the terms are most crucial for physical consequences and should be fixed. Other corrections on which the behavior does not depend so sensitively can then first be ignored. We can clearly see this from our example, where the correction functions $\\beta$ from the spin connection and $\\sigma$ from the matter gradient term do not play as important roles as the functions $\\alpha$ and $D$. This is fortunate in particular for $\\beta$ because there is no tight prescription for its behavior in the full theory. Also the function $\\tilde{\\beta}$, which could equal $\\beta$ or simply one depending on how one deals with the diffeomorphism constraint, only appears once in the final equations (\\ref{PertMod3}) and in a way which does not significantly change the behavior given by $\\alpha$-corrections. (For large $\\bar{p}$, $1-\\tilde{\\beta}$ is negative while $-\\alpha'\\bar{p}/\\alpha$ is positive. Due to the perturbative form (\\ref{alphapert}) of the functions as a power series in $\\ell_{\\rm P}^2/\\bar{p}$, however, the correction from $\\alpha$ is dominant in this regime and determines the sign of the correction. For small $\\bar{p}$, on the other hand, effects from $\\tilde{\\beta}$ can be more pronounced but perturbation theory is more complicated.)  The most sensitive corrections at the level of linearized perturbations around flat space are thus those coming from $\\alpha$ and $D$. A phenomenological analysis then shows which behavior of these functions is preferred."
    },
    "0609/astro-ph0609821_arXiv.txt": {
        "abstract": "{We present an explanation of the Z-track phenomenon based on spectral fitting results of {\\it Rossi-XTE} observations of the source GX\\th 340+0 using the emission model previously shown to describe the dipping Low Mass X-ray Binaries. In our Z-track model, the soft apex is a quiescent state of the source with lowest luminosity. Moving away from this point by ascending the normal branch the strongly increasing luminosity of the Accretion Disc Corona (ADC) Comptonized emission $L_{\\rm ADC}$ provides substantial evidence for a large increase of mass accretion rate $\\dot M$. There are major changes in the neutron star blackbody emission, $kT$ increasing to high values, the blackbody radius $R_{\\rm BB}$ decreasing, these changes continuing monotonically on both normal and horizontal branches. The blackbody flux increases by a factor of ten to three times the Eddington flux so that the physics of the horizontal branch is dominated by the high radiation pressure of the neutron star, which we propose disrupts the inner disc, and an increase of column density is detected. We further propose that the very strong radiation pressure is responsible for the launching of the jets detected in radio on the horizontal branch. On the flaring branch, we find that $L_{\\rm ADC}$ is constant, suggesting no change in $\\dot M$ so that flaring must consist of unstable nuclear burning. At the soft apex, the mass accretion rate per unit area on the neutron star $\\dot m$ is minimum for the horizontal and normal branches and about equal to the theoretical upper limit for unstable burning. Thus it is possible that unstable burning begins as soon as the source arrives at this position, the onset of unstable burning being consistent with theory. The large increase in $R_{\\rm BB}$ in flaring is reminiscent of radius expansion in X-ray bursts. Finally, in our model, $\\dot M$ does not increase monotonically along the Z-track as often previously thought. ",
        "introduction": "The Z-track sources form the brightest group of Low Mass X-ray Binary (LMXB) sources containing a neutron star, consisting of 6 members: GX\\th 340+0, GX\\th 5-1, Cyg\\th X-2, \\hbox{Sco\\th X-1}, GX\\th 17+2 and GX\\th 349+2 (Hasinger \\& van der Klis 1989). All of these persistently radiate at about the Eddington limit for a neutron star and do not, in general, exhibit X-ray bursting characteristic of LMXBs with lower luminosities. The sources trace a Z-shaped pattern in X-ray hardness versus intensity or hardness versus softness demonstrating strong spectral evolution associated with major physical changes within the sources. However, the nature of these changes is not understood and this remains a significant astrophysical problem. Moreover, the Z-track sources are all detected as radio emitters showing relativistic jets to be present, but essentially on one branch only (e.g. Penninx 1989). The mechanism of jet formation also remains an important astrophysical problem. The occurrence of jets on one branch only may allow us to probe how jets are launched.  Hasinger et al. (1989) showed that the sources could be found on   % three different branches of a skewed Z-shape in a hardness-intensity diagram: the horizontal branch (HB), the normal branch (NB) and the flaring branch (FB). These diagrams demonstrated in a model-independent way that strong spectral evolution was taking place around the Z-track, suggestive of changes at the compact object and accretion disc, but did not reveal the nature of the changes. It was found that quasi-periodic oscillations (QPO) exhibit systematic changes around the Z-track (e.g. van der Klis et al. 1987; Hasinger \\& van der Klis 1989). Clearly, investigation of the spectral evolution along the Z-track is likely to reveal the nature of the changes; however, spectral studies of the Z-track sources have been hindered by a lack of agreement over which emission model for LMXBs to use.  \\subsection{Spectral fitting} The nature of X-ray emission in LMXBs has been controversial with two main types of model for the continuum: the first is the Eastern model comprising multi-colour blackbody emission from the inner accretion disc plus non-thermal emission from a central Comptonizing region (Mitsuda et al. 1989); the second model discussed below comprises neutron star blackbody emission plus Comptonization in an extended corona. Previous spectral fitting of the Z-track sources has been largely based on the Eastern model. Spectral evolution along the Z-track in several sources was investigated by Schulz and Wijers (1993) and Schulz et al. (1989). More recently, spectral fitting of the source Cyg\\th X-2 was carried out by Done et al. % (2002) using the Eastern model. Spectral fitting results were given in terms of the parameters of the disc blackbody and thermal Comptonization emission components, but it was not clear what caused the parameters to change in the way found. The same model was used by Agrawal \\& Sreekumar (2003) to fit {\\it Rossi-XTE} spectra of GX\\th 349+2. Spectral fitting of broadband {\\it BeppoSAX} data was carried out for GX\\th 17+2, GX\\th 349+2 and for Cyg\\th X-2 (di Salvo et al. 2000, 20001, 2002) who argued that all three sources could be fitted by the Eastern model. In GX\\th 17+2 and GX\\th 349+2, there was little change in blackbody parameters with position on the Z-track, however, in Cyg\\th X-2 there was an increase of blackbody temperature and decrease of blackbody radius moving from the horizontal to the normal branch, suggesting that the inner disc radius was shrinking although an explanation of this was not proposed.   Theoretical models for the Z-track sources have been proposed (Psaltis et al. 1995), % involving a magnetosphere at the inner accretion disc, and changes in the geometry and extent of the magnetosphere. However, it has not been possible to test this model in detail, as, for example, the geometry cannot be determined. This modelling also assumes the main element of the Eastern model, that the Comptonizing region is a small central region, and evidence (below) does not support this.  \\subsection{Approach of the present work}  Thus previous spectral fitting based on the Eastern model has not produced a clear explanation of the Z-track and there is now substantial evidence against this model provided by the dipping class of LMXBs. In these sources with high inclination (65 - 85$\\degmark$), absorption of the X-ray emission takes place on every orbital cycle in the bulge in the outer disc. Complex spectral evolution in dipping strongly constrains emission models providing substantial evidence for a model comprising simple, unmodified blackbody emission from the neutron star and Comptonized emission from an extended accretion disc corona of radial extent $R$ typically 50000 km, and height $H$ having \\hbox{$H/R$ $<$ 1} (the ``Birmingham model''; Church \\& Ba\\l uci\\'nska-Church 1995, 2004). The extended size of the ADC is demonstrated by dip ingress timing (Church \\& Ba\\l uci\\'nska-Church 2004) which allows measurement of the ADC radial extent. The spectra of the dipping LMXBs in the non-dip state and in every stage of dipping from many observations have been well-fitted using this model (Church et al. 1997, 1998a,b, 2005; Ba\\l uci\\'nska-Church et al. 1999, 2000; Smale et al. 2001; Barnard et al. 2001), and so the evidence for the model is substantial and hence argues against the Eastern model with a small central Comptonizing region.  It should not be thought that the spectral model used is {\\it ad hoc}, i.e. having a simple form easily implemented that happens to fit the spectra, but with possibly large inaccuracies. The form used is {\\sc bb + cpl}, where {\\sc bb} is simple blackbody emission from the neutron star and {\\sc cpl} is a cut-off power law representing Comptonization. The correct form for Comptonization is very dependent on the correct specification of the seed photon spectrum. With an extended ADC, the seed photons must originate in the disc below the corona with a large population of soft photons below 0.1 keV from large radii. We previously showed that a power law or cut-off power law represents Comptonization well allowing for this, and the Comptonized spectrum continues to rise below 1 keV (Church \\& Ba\\l uci\\'nska-Church 2004) whereas the {\\sc comptt} model often used incorporates a blackbody seed photon spectrum decreasing strongly below 1 keV. We have never been able to detect any departure of the thermal emission from simple blackbody form, and we previously reviewed the evidence for use of a simple blackbody and discussed in detail how possible modification in the neutron star atmosphere depends critically on the electron density which is very poorly known (Ba\\l uci\\'nska-Church et al. 2001). Kuulkers et al. (2002) also present evidence that X-ray burst emission is of simple blackbody shape.  Finally, when the spectra of LMXBs are fitted with the Eastern model in the form disc blackbody + cut-off power law (e.g. Church \\& Ba\\l uci\\'nska-Church 2001), it is often found that the inner radius of the disc blackbody is unphysically small with values less than 1 km, i.e. substantially less than the neutron star radius. Done et al. (2002) argued that this was purely an artefact of using an incorrect form (cut-off power law) for the Comptonized emission. The above argument shows that, on the contrary, the cut-off power law is a proper description of the Comptonized spectrum if the ADC is extended (at least below the high energy cut-off) as discussed fully in Church \\& Ba\\l uci\\'nska-Church (2004). Thus the conclusion that the Eastern model gives unphysically small inner disc radii is valid, and in addition, modelling based on assuming seed photons that are blackbody with $kT$ $\\sim$ 1 keV as in {\\sc comptt} is invalid.   Thus the approach used in the present work consists of applying for the first time the Birmingham emission model to the Z-track sources, as it provided a good explanation of the dipping LMXBs, to determine whether this model suggests an explanation of the Z-track phenomenon. We chose the source GX\\th 340+0 since high quality {\\it Rossi-XTE} data extending round the full Z-track were available. The source was classified as a Z-track source by Hasinger \\& van der Klis (1989). QPO were found in {\\it Exosat} data (van Paradijs et al. 1988) and in {\\it Ginga} data (Penninx et al. 1991) and kilohertz QPO were discovered by Jonker et al. (1998). The detection of radio by Penninx et al. (1993) is important indicating the presence of jets. Fender \\& Hendry (2000) summarized previous work on radio detection from X-ray binaries showing that radio emission is detected from black hole binaries and from the Z-track sources when on the Horizontal Branch, concluding that a requirement for jet formation was an X-ray luminosity approaching the Eddington limit, i.e. greater than 0.1 $L_{\\rm Edd}$. More recently, Migliari \\& Fender (2006) summarize the detection of relatively weak radio emission from a number of Atoll sources.  \\subsection{A monotonic increase of mass accretion rate ?} Finally, we address a widely-held view that the observed hardness changes in the Z-track sources are driven in some way by a mass accretion rate $\\dot M$ which increases monotonically around the Z-track in the direction HB-NB-FB. The movement round the track on timescales of hours, without jumping between tracks, suggested the variation of a single parameter, probably $\\dot M$ (Priedhorsky et al. 1986). In a multi-wavelength campaign (Hasinger et al. 1990), \\hbox{Cyg\\th X-2} was observed in X-ray, ultraviolet, optical and radio. The results of the {\\it IUE} observations (Vrtilek et al. 1990) showed an increase of intensity for motion HB-NB-FB. However, the identification of these data with the flaring branch depended on strong variability which was probably due to X-ray dipping, and omitting these data would mean that there is no very dramatic correlation of UV with track position. Van Paradijs et al. (1990) presented evidence for increase of optical emission line strength for motion HB-FB but noted that the number of data points was rather limited, while absorption line strength decreased, indicating that the processes involved are complex and can involve X-ray reprocessing. \\begin{figure*}                                                           % \\includegraphics[width=64mm,height=176mm,angle=270]{f1} \\caption{Top: Background-subtracted and deadtime-corrected PCA light curve of the 1997 September observation of GX\\th 340+0 with 64 s binning. Parts of the lightcurve are identified (see text) with branches of the Z-track (Fig. 2) and shown as horizontal bars. Most of the first half of the observation is on the Normal Branch except for two large flares as shown; the second half of the observation is mostly on the Horizontal Branch. The dotted lines emphasize the upper and lower intensity limits which are consistent across the observation: these limits correspond to the intensities of the Hard Apex and Soft Apex of the Z-track. Bottom: the corresponding variation of hardness ratio (7.3 -- 18.1 keV)/(4.1 -- 7.3 keV).} \\label{} \\end{figure*}  A monotonic increase of $\\dot M$ around the Z-track in the direction HB-NB-FB is, of course, inconsistent with the basic observational fact that the X-ray intensity {\\it decreases} between the hard apex and soft apex. A monotonic increase of $\\dot M$ is not an assumption of the Eastern model; however Agrawal et al. (2003) suggested that the results of fitting this model might be qualitatively consistent with a monotonic increase. It would be necessary for the Comptonized emission luminosity to decrease from hard to soft apex while $\\dot M$ increased. In the present work, the luminosity decrease by a factor of two seems unlikely if $\\dot M$ were increasing.  Kuulkers et al. (2002) provided another argument against the increase of $\\dot M$ around the Z-track. Unstable nuclear burning is expected at various mass accretion rates (Fujimoto et al. 1981; Fushiki \\& Lamb 1987; Bildsten 1998; Schatz et al. 1999) including accretion rates close to the Eddington $\\dot M$ due to He burning in a mixed H/He environment, i.e. $1.0\\times 10^{-9} <\\dot M < 2.6\\times 10^{-8}$ M$_{\\sun}$ yr$^{-1}$. (Bildsten 1998). At higher mass accretion rates corresponding to luminosities greater than $3\\times 10^{38}$ erg s$^{-1}$, burning is expected to be stable. Thus, Kuulkers et al. argued that unstable nuclear burning should be observed in the Z-track sources, and X-ray bursting is not, in general, seen. Bursts are seen in GX\\th 17+2 and Cyg\\th X-2 (Kahn \\& Grindlay 1984; Tawara et al. 1984; Sztajno et al. 1986; Kuulkers et al. 1995; Wijnands et al. 1997; Smale 1998). However, bursting in these sources showed no correlation of burst properties with Z-track position (Kuulkers et al. 1995, 1997) as expected if the mass accretion rate was changing substantially along the Z-track. It thus appears that the increase of $\\dot M$ around the Z-track cannot be regarded as proven, and the results presented here will be viewed in this context. ",
        "conclusions": "We have investigated spectral changes along the Z-track based on the use of a particular emission model. In the following we examine the implications of the results on the assumption that this model is correct. As discussed in Sect. 1, we do not know of any major inaccuracy in use of the spectral form {\\sc bb + cpl} within the model, and it is difficult to quantify any small level of inaccuracy. However, such inaccuracy will not change the {\\it pattern} of parameter variation around the Z-track and is not expected to lead to large errors in, for example, the blackbody radius at the soft apex. The results provide a plausible physical picture of the Z-track, of the observation of jet radio emission on the horizontal branch and of the nature of flaring, which is quite different from previous pictures. In the model presented here, the constant luminosity of the ADC emission on the FB implies $\\dot M$ is constant and so flaring must be thermonuclear. The soft apex is the quiescent state of the source, with $\\dot M$ increasing along the normal branch causing an increase of the ADC luminosity. The neutron star blackbody temperature then rises causing increased radiation pressure which has a strong effect on the inner disc, blowing away the outer layers (farthest from the orbital plane). It is not appropriate to discuss our model in relation to QPO in detail here, however, we note that the blowing away is not inconsistent with the formation of QPO in the inner disc. It has been suggested that the upper kHz QPO frequency may correspond to a Keplerian frequency in the inner disc (e.g. Jonker et al. 2000). The removal of the outer parts of the inner disc by radiation pressure but not parts close to the orbital plane is not inconsistent with this, and it may even be that the disturbance to the disc triggers the oscillation.    \\subsection{The reduction in blackbody radius}  Church \\& Ba\\l uci\\'nska-Church (2001) can suggest why the neutron star blackbody emitting area decreases on the normal and horizontal branches. It is unlikely that this is due to any process on the stellar surface, and is more likely to be due to processes at the inner disc. In our survey of 14 LMXBs using {\\it ASCA} and {\\it BeppoSAX}, we investigated the variation of blackbody luminosity with total luminosity. It is well-known that the blackbody fraction is small in faint sources rising to about 50\\% in sources close to the Eddington limit. The results were presented geometrically in terms of $h$, the half-height of the blackbody emitter, and $H$, the half-height of the inner disc. The blackbody luminosity depends on $h$, and is calculated from the blackbody radius (Sect. 3.2.1) while $H$ is proportional to the mass accretion rate in the standard theory of the inner, radiatively-supported, disc (Frank et al. 2002). % It was found that for all sources $h$ $\\simeq$ $H$, suggesting that the neutron star emitting area was determined by the height of the inner disc. The mechanism for this might be radial flow across the gap between disc and neutron star, or alternatively the mechanism of Inogamov \\& Sunyaev (1999) which predicts that the accretion flow spreads vertically on the star to a height depending on $\\dot M$. The survey results agreed approximately with this model (Church et al. 2002). %  In a radiatively-supported disc, the vertical height of the disc increases rapidly with radial distance until at $\\sim$10 km from the surface of the neutron star, it has its equilibrium value $H_{\\rm eq}$ (Frank et al. 2002). In the present work on GX\\th 340+0, $H_{\\rm eq}$ calculated from the source luminosity varies between 30 km at the soft apex and 50 km at the hard apex, so that the inner disc towers above the neutron star. If the radiation pressure of the neutron star is strong, we would expect some of the inner disc to be blown away above and below the disc, leaving a residual inner disc of reduced height. On the basis of the survey results, the height on the neutron star would respond to this by becoming much smaller which may explain $h$ decreasing from $\\sim$ 11 km at the soft apex (where the radiation pressure is small) to $\\sim$ 1 km on the horizontal branch.   \\subsection{The nature of the Horizontal Branch}  The spectral fitting results show that the luminosity of the Comptonized emission increases on the normal branch and then decreases by about the same amount on the horizontal branch, following the same path in Fig. 5, suggesting that the mass accretion rate first increases, then decreases towards the initial value and we might expect the changes on the neutron star to also reverse at the hard apex. But the neutron star blackbody does not reverse its behaviour at the hard apex, the temperature and radius changing monotonically on the NB and HB (Fig. 4). A possible reason that this does not happen could relate to time delays in the system. An increase in $\\dot M$ will affect the emission of the extended ADC first causing the source to move along the NB towards the hard apex, but changes in the blackbody emission would be delayed until the increase reaches the neutron star. Fig. 4 provides evidence for this: the Comptonized emission has increased before the blackbody changes; Fig. 1 shows that several hours are required for the source to move halfway along the NB from the soft apex. This timescale is approximately consistent with estimates of flow times through the disc based on the viscous timescale (Frank et al. 2002). When $\\dot M$ decreases in the disc then $L_{\\rm ADC}$ falls, but the increase of $\\dot M$ at the neutron star could continue for some time.  We next ask how the source leaves the HB end of the Z-track. Assuming as above that $\\dot M$ has already decreased, we would expect the neutron star to cool, radiation pressure to fall and $R_{\\rm BB}$ to increase. However, the effects in the ADC are not clear. If there was no change in the emission (as $\\dot M$ is constant), the total luminosity would be about constant and the source may be able to jump from the HB to the NB as the neutron star cools. However, jumps have not so far been detected, and Z-track sources are seen to move smoothly from HB to NB.  We can assume that the X-ray luminosity is a good indicator of $\\dot M$ when the radiation pressure is low (on the NB), but this is probably not true when the radiation pressure is high, so that the decreasing $L_{\\rm CPL}$ on the HB may not mean that $\\dot M$ falls. Firstly, the blowing away of outer layers of the inner disc by increased radiation pressure may reduce the ADC emission by a geometric factor (giving a reduction in $L_{\\rm CPL}$) while $\\dot M$ is constant or rising. In addition, the radiation pressure can divert a substantial fraction of $\\dot M$ into the jets causing a decrease of $L_{\\rm CPL}$ on the HB because of a decrease of accretion flow in the disc below the corona, reducing the electron density while the total $\\dot M$, including the part diverted to the jets, remains constant, or increases. Then the asymmetry between the blackbody and Comptonized emission disappears, since the reversal of behaviour of both components takes place at the end of the Z-track, both components reversing their evolution when the mass accretion rate falls, the source moving back along the Z.     \\subsection{The nature of the Flaring Branch}  We have suggested that on the flaring branch there is unstable nuclear burning. This is not unexpected since the theory of unstable nuclear burning (e.g. Bildsten 1998) shows that He burning in a hydrogen-rich environment should be unstable for $\\dot M$ $<$ $2.6\\times 10^{-8}$ M$_{\\sun}$ y$^{-1}$, while for higher $\\dot M$, burning is stable. The Z-track sources have luminosities at about the Eddington limit, but lack of spectral fitting has often not allowed a definitive statement to be made on whether unstable burning is expected. However, Bildsten (1995) argued that the timescale of flaring in Sco\\th X-1 is consistent with the timescale at which unstable burning starting in one part of the neutron star will spread across the surface. In sources of luminosity mid-way between burst sources and Z-track sources, we have previously detected bursting and flaring in the same observation, as in XB\\th 1254-690 (Smale et al. 2002).  On the flaring branch, the blackbody radius increases systematically from its value of 10.8$\\pm$1.0 km at the soft apex to $\\sim$16.1$\\pm$0.8 km. If we accept that the whole neutron star is emitting at the soft apex, then in flaring the emission must extend beyond the surface of the star. Similar values of $R_{\\rm BB}$ greater than the neutron star radius have been found previously for Z-track sources, such as a value of 27 km found by Christian \\& Swank (1997) for GX\\th 5-1. Lower luminosity X-ray burst sources can exhibit radius expansion (Lewin et al. 1985) in which % $R_{\\rm BB}$ increases indicating photospheric expansion due to radiation pressure, the luminosity remaining close to the Eddington limit. In the present work we also have an increase of $R_{\\rm BB}$ on the flaring branch which we suggest is unstable nuclear burning, and the emitted flux approaches the Eddington flux. Thus it is possible that radius expansion may take place during flaring in the Z-track sources. The implication is that the emission expands beyond the surface of the neutron star, and the deep hole where the height of the inner accretion disc falls to that of the neutron star is partly filled by plasma. The rise in blackbody temperature is however, only 10\\%, and so much less than the temperature rise in X-ray bursts.  We next ask whether unstable nuclear burning takes place on all branches of the Z-track or only on the flaring branch. Conditions on the other branches either prevent unstable burning, or mask the appearance of flaring. On the flaring branch flares develop from the lowest X-ray intensity of the source at the soft apex, and are very obviously seen in the lightcurve. We have examined the lightcurve for variations on timescales of a few thousand seconds in the other branches, but this was inconclusive. However, the present results suggest why unstable burning may occur only on the FB.  At the soft apex, $L_{\\rm Tot}$ has its smallest value of $3.7\\times 10^{38}$ erg s$^{-1}$, corresponding to a mass accretion rate per unit area on the neutron star $\\dot m$ = $1.5\\times 10^5$ g cm$^{-2}$ s$^{-1}$. The upper limit for unstable nuclear burning is $\\dot m$ = $1.3\\times 10^5$ g cm$^{-2}$ s$^{-1}$ (Bildsten 1998). Thus the actual $\\dot m$ is close to the theoretical upper limit which Bildsten estimates to be accurate to within 30\\%, and unstable burning can be possible. On the normal branch, the luminosity and mass accretion rate eventually approach a factor of 2 higher than at the soft apex suggesting that unstable burning is not possible and that flaring will not take place on the NB and HB. On the flaring branch, the emitting blackbody area increases, and so $\\dot m$ remains below the critical value, and nuclear burning remains unstable. Examination of the lightcurve (Fig. 1) shows that flaring begins {\\it exactly} at the lowest luminosity of the source. Thus if we move towards the soft apex along the NB, it may be that unstable burning begins immediately $\\dot m$ falls below the critical value. Thus the identification of the flaring branch with unstable nuclear burning appears to be consistent with theory, as is the lack of unstable nuclear burning on the NB and HB.   \\subsection{The formation of jets by radiation pressure}  Explaining the formation of relativistic jets from the neighbourhood of compact objects in both Galactic sources and in AGN has been a major astrophysical goal. Two main types of model have been proposed. First, it has long been suspected that radiation pressure may play an important role (Begelman \\& Rees 1984), and Lynden-Bell (1978) suggested that a radiatively-thickened inner accretion disc will define two funnels which allow jet formation perpendicular to the disc, driven by radiation pressure. Second, the other type of model depends on electromagnetic production via strong electric fields arising from magnetic fields in the disc around a black hole (Blandford \\& Znayek 1977). The present work reveals the link between jet formation and the position on the Z-track where the radiation pressure is high, strongly indicating the importance of radiation pressure.  It is well known that radio emission is detected from Z-track sources on the horizontal branch, and upper normal branch providing strong evidence that jets are present. On this part of the Z-track the high neutron star temperatures we determine mean that the radiation pressure of the neutron star is  high, the emitted flux approaching three times super-Eddington (Fig. 7). It thus appears significant that jets are formed on this part of the Z-track suggesting that the launching of jets is due to the strong radiation pressure. Our spectral fitting result that the emitting area on the neutron star decreases is consistent with matter being blown away from the inner disc, and the release of matter within the system is directly supported by the measured increase of column density to a maximum at the hard apex. We thus propose that jets are formed by the very strong radiation pressure of the neutron star at this part of the Z-track, where the radiation flux substantially exceeds the Eddington flux. The emission area is small implying it is from an equatorial strip around the neutron star, and this geometry favours upwards direction of the jet. Radiation pressure acting horizontally may disturb the inner disc. However, a line from the equatorial belt to the inner edge of the disc will rise steeply upwards, the radiation pressure acting in a direction approaching vertical.  Jets are not seen on the flaring branch, and this is consistent with the model proposed. The radiation pressure is not so strong and the blackbody emitter apparently expands beyond the neutron star surface. There is normally a conical funnel in the central radiatively-supported inner accretion disc (Frank et al. 2002) where the disc height falls rapidly towards the stellar surface providing an opening for jet formation on the NB and HB. However, in flaring this funnel may be partly or completely blocked by plasma this also opposing jet formation.     \\thanks{ This work was supported in part by the Polish KBN grant KBN-1528/P03/2003/25 and by PPARC grant PPA/G/S/2001/00052. We thank the referee for his very useful comments.}"
    },
    "0609/astro-ph0609494.txt": {
        "abstract": "The combined data of the University of Michigan Radio Astronomy Observatory and Mets\\\"ahovi Radio Observatory provide us with radio light curves for Active Galactic Nuclei monitored by both observatories from 4.8 to 37~GHz covering time intervals up to $\\sim$~25 years. We consider here such composite light curves for four $\\gamma$-ray blazars that have been nearly continuously monitored at both observatories: 0458--020, 0528+134, 1730--130 and 2230+114.  We have decomposed the most prominent outbursts in the light curves of these four blazars into individual components using Gaussian model fitting, and estimated the epochs, amplitudes, and half-widths of these components as functions of frequency. We attempt to distinguish ``core outbursts,'' which show frequency-dependent time delays and are associated with brightening of the core, from ``jet outbursts,'' which appear nearly synchronously at all frequencies and are accompanied by the emergence of new jet components and their subsequent evolution. The outbursts in 0528+134 and 2230+114 display fine structure and consist of individual sub-outbursts. Available 43~GHz VLBA images allow us to identify only one pure core outburst (in 2230+114) and one pure jet outburst (0458--020). Most of the outbursts analyzed are mixed, in the sense that they display frequency-dependent time delays (i.e., they are optically thick) and are associated with the eventual emergence of new jet components. The maxima of the jet and mixed outbursts probably correspond to epochs when newly ejected components become fully optically thin. These epochs are also marked by a significant increase in the angular velocities of the ejected components. There is evidence that the outbursts in 2230+114 repeat every $8.0\\pm0.3$~years, with the positions of individual sub-outbursts being preserved from one quasi-periodic eight-year cycle to another, even though their amplitudes vary by more than a factor of two.  Preliminary estimates of the total durations of possible activity cycles based on an analysis of total flux-density variations and all available VLBI data are given for the remaining sources. ",
        "introduction": "The term ``blazar'' is sometimes used to refer to the subset of core-dominated Active Galactic Nuclei (AGNs) that display high-amplitude variability over a wide range of wavebands, testifying that one of the relativistic jets in these objects is oriented close to the line of sight toward the Earth. The variability of blazars is a very complex phenomenon, and many different processes can take part, such as activity of the central engine, the propagation of primary perturbations down the jet, the evolution of the jet itself and various propagation effects (e.g. interaction with the surrounding medium) (Marscher 1996). During the last few years, important progress has been made in understanding many aspects of these phenomena through theoretical simulations (see Gomez et al. 2004 and references therein), investigations of total flux variability in the X-ray (Pian 2003), $\\gamma$ (Kranich 2003) and radio (Aller et al. 1999, Ter\\\"asranta, Wiren \\& Koivisto 2003) and various VLBI studies (e.g. Britzen et al. 1999a; Gomez et al. 2000; Lister \\& Homan 2005; Jorstad et al. 2005).  Nevertheless, our understanding of the global, long-term development of the activity remains very fragmentary.  The detection of periodicity in the activity of the blazars OJ287 ($P\\sim12$~yr, Pursimo et al. 2000), 0059+581 ($P\\sim4$~yr, Pyatunina et al. 2003), 0202+149 ($P\\sim4$~yr, Pyatunina et al. 2000) and 0235+164 ($P\\sim5.7$~yr, Raiteri et al. 2001) looks promising in this regard, but all attempts to estimate exact periods for the activity have been unconvincing, even for the best sampled case of OJ287 (Kidger 2000), indicating that the phenomenon is probably not strictly periodic, but instead quasi-periodic. Indeed, even if the phenomena triggering a blazar's activity were purely periodic in nature, the observed variability would likely be at best quasi-period, since it is the product of multiple complex processes associated with the generation and propagation of the radio emission. Numerical simulations of a binary black-hole model for OJ287 (Sundelius et al. 1997) show that the triggering of the activity itself can be a complex function of time (see below), making the observed variability even more complex. Nevertheless, analysis of such quasi-periodicity can be an important diagnostic tool for modelling the energy generation in AGNs and investigating binary black-hole systems.  Below we propose an alternative approach to investigating the global evolution of blazars.  Physically, the activity of a source can be described using two independent time scales: $T_{act}$, the characteristic time scale for the activity of the central engine, and $T_{ev}$, the time scale for evolution of the jet after it has been disturbed by a perturbation at it's base. The time scale $T_{ev}$ includes a whole sequence of events, from the initial appearance of the perturbation to the time when the propagating perturbation fades and merges into the quiescent jet emission. At present, it may be difficult to define this important time scale observationally. The time scale $T_{act}$ can be thought of as the duration of one complete ``activity cycle,'' which we propose to define as the characteristic time interval between two successive events marking onsets of phases of activity. With this definition, for example, a blazar that showed a pattern of being active for five years, then quiescent for ten years, then active five years, and so forth, would have $T_{act}\\simeq 15$~yrs. If the time scale for activity exceeds the time scale for the jet's evolution ($T_{act}\\ge T_{ev}$), two subsequent activity cycles should be easily distinguishable; otherwise, manifestations of core activity and jet evolution associated with different cycles can be superimposed and complicate the observed picture. It should also be borne in mind that the observed $T_{act}$ may be affected by the source's redshift, while the observed $T_{ev}$ may be affected by both the redshift and the Doppler factor of the jet (for more detail, see the discussion in Lister 2001).  According to theoretical arguments (Marscher 1996; Gomez et al. 1997), a shock induced by a primary excitation at the base of the jet is initially manifest in radio light curves as an outburst that is delayed at lower frequences due to the combined effects of the frequency stratification of the emitting electrons, non-zero opacity and light-travel delays. Such time-delayed outbursts are associated observationally on milliarcsecond (mas) scales with brightening of the VLBI core, and can be considered ``core'' outbursts. The propagation of the shock downstream in the jet is revealed via the emergence of a new optically thin jet component (or components) in the VLBI image, sometimes accompanied by ``jet'' outbursts, which evolve nearly synchronously at all frequencies in the radio light curve. The integrated monitoring data support the division of observed outbursts into such core and jet outbursts, at least in some sources (0202+149, Pyatunina et al. 2000; 0420--014 Zhou et al. 2000; 0059+581, Pyatunina et al. 2003). Thus, in practice, it is reasonable to approximate the duration of an activity cycle as the time interval between two successive ``core'' outbursts. % although this interval may not %coincide with the typical time of a jet's evolution defined above. The questions of how this interval varies from source to source and how stable it is for a particular source may be keys for our understanding of the activity's grand design.  The majority of blazars do not seem to display a well-defined activity cycle and associated time scale $T_{act}$, due in part to incompleteness of their monitoring data over sufficiently long time scales (decades) and in part due to the complexity or nature of the underlying physical processes giving rise to the variability (see, e.g., Lister 2001). In addition, the clarity with which an activity cycle in a given source is manifest may be very different at different frequencies, especially if they are fairly widely spaced. Combined analyses of multi-epoch VLBI observations and multi-frequency light curves for those cases when a reasonably well-defined cycle can be identified across a range of frequencies show that the duration of the activity cycle can vary from source to source over a wide range: $T_{act}\\sim 4$~yr for 0059+581 (z=0.643, Pyatunina et al. 2003), $\\sim 12$~yr for 0735+178 (z=0.424, Carswell et al. 1974; Agudo et al. 2002) and $\\sim 20$~yr for 1308+326 (z=0.996, Hewitt et al. 2003; Pyatunina et al. 2004). As has been noted above, the question of the stability of the activity time scale, or of the quasi-periodicity of the activity, in a particular source is more complex, not only because the relevant time scales can be rather long, but also due to several other important factors. It has been suggested that a supermassive binary black hole in the nucleus of the AGN may be responsible for both the existence of quasi-periodicity (Lehto \\& Valtonen 1996; Valtaoja et al. 2000) and the activity itself (Sillanp\\\"a\\\"a 1999), as well as for structural changes on parsec scales (Lobanov \\& Roland 2002).  In this picture, activity can be triggered by the periodic tidal perturbation of the primary black hole's accretion disk by the orbital motion of a secondary black hole or the secondary's accretion disk.  The eccentricity of the orbit (Lehto \\& Valtonen 1996; Sundelius et al. 1997) may divide the overall period in two unequal parts, while precession of the jet may change the viewing angle, resulting in modulation of the amplitudes of outbursts. Relativistic precession of the secondary's orbit will complicate the observed variability even more.  \\begin{figure} \\hspace*{-1.0cm} \\includegraphics[width=0.45\\textwidth,angle=-90]{pyatunina_fig1.ps} %\\vspace*{-0.5cm} \\caption{Example of our Gaussian decompositions of the light curves: light curve for 0528+134 at 22~GHz.} \\label{fig:gaussexam} \\end{figure}  In our approach, we carry out a joint analysis of the total-flux and spectral variability based on multi-frequency light curves, combined with the observed structural changes in VLBI images.  This paper presents our results for a first small sample of four selected sources. %In the future, we hope to enlargen our %sample and carry out more detailed analyses using all available VLBI data. The frequency-dependent time delays found here can also be used to test models of the nonthermal emission in blazars (Gomez et al. 1997; Lobanov 1998; Marscher 2001). ",
        "conclusions": "\\subsection{Types of outbursts}  Analysis of the outbursts combined with the detailed structural evolution at 43~GHz shows that only one of the outbursts can be classified as a pure core (2230+114, A) and one as a pure jet (0458-020, C) outburst in terms of both their time delays and their association with core or jet brightening. Most outbursts are mixed, in the sense that they display both frequency-dependent time lags and a connection with the appearance of new jet components. Note that previous classifications of outbursts (Pyatunina et al. 2000; Zhou et al. 2000) were based on low frequency (5 and 8~GHz) VLBI observations. %As has been already noted by Zhou et al. (2000), classification of outbursts %rely on their light curves as well as VLBI images and therefore %depends on observing frequency and resolution. However, a jet component resolved at 43~GHz would be merged with the core at lower frequencies, so that an outburst classified as a core type at low frequency might be identified as a mixed type at high frequencies. This illustrates the importance of high frequency VLBI data for probing the early stages of activity within the jet.  It seems likely that the pure core and mixed core--jet outbursts represent fundamentally the same phenomenon: activity in the core that gives rise to the ejection of material along the jet. In this case, the brightness of the ejected material associated with the pure core outburst A in 2230+114 faded before becoming distinct from the observed core at 43 GHz. It is easy to understand the larger number of pure core outbursts at lower frequencies as well: with the lower resolution available, the ejected components are more likely to have faded beyond detection by the time they become distinct from the core. The relatively small number of pure jet outbursts then represent brightening of an optically thin region of the jet well after its ejection from the core, e.g., due to the formation of shocks.  \\subsection{Fine structure of outbursts}  Some traces of fine structure can be found in all of the outbursts, at least at the highest frequencies. In some of the sources, the fine structure dominates, splitting outbursts into distinct sub-outbursts (0528+134, 2230+114), while in others (0458-020 and 1730-130), it degenerates into small fluctuations of the fluxes. Note that the appearance of fine structure in light curves is independent of resolution, but can be hidden due to high optical depth. %or viewing angle (0528+134, see above). Savolainen et al. (2002) suggested that outburst fine structure can be induced by shocks that grow and decay in the innermost few tenths of a milliarcsecond. According to theoretical simulations, shocks can be generated by several mechanisms: (a) multiple recollimation shocks, which are induced by the pressure mismatch between the jet and surrounding medium (Gomez et al. 1997, 2003); (b) interactions between helical instability modes that grow due to perturbations to the jet by jet--cloud collisions or precession, with elliptical (Lobanov, Hardee \\& Eilek 2003) or body waves (Hardee et al. 2001) growing into the jet; and (c) interactions of the jet with obstacles in its path. The preservation of the positions of individual sub-outbursts in outbursts belonging to different activity cycles in 2230+114 may provide important information for distinguishing between these mechanisms. However, it is necessary to carry out similar analyses for more sources, in order to investigate whether this property is common or peculiar to this source.  The presence of fine structure in the outbursts can significantly complicate the identification of a ``starting point'' of the activity. Strictly speaking, only information about flux variations at frequencies above $\\sim40$~GHz in combination with 43~GHz VLBI data can address this problem.  \\subsection{Time delay as a function of frequency}  The parameters of the time-delay function $\\Delta T=\\Delta T_0\\nu^{\\alpha}$ for all events defined for at least four frequencies are given in Table~\\ref{tab:summary}. According to theoretical estimates (Gomez et al. 1993, 1997), the position of the core %as the position of the maximum flux density is expected to depend on the frequency as $r_{core}\\propto\\nu^{-0.8}$. The frequency-dependent shift of the core position in 3C345 estimated by Lobanov (1998) turned out to be close to the theorical expectation $r\\propto\\nu^{(-1.04\\pm0.16)}$. If the flow velocity is independent of frequency, then the time delay is expected to behave similarly. The indices of the time-delay functions given in Table~\\ref{tab:summary} vary from --0.6 (0528+134, B3) to --2.2 (1730--130) with the mean value $\\alpha_{mean}=(-1.20\\pm0.49)$. The time-delay functions for individual events (Figs.~\\ref{fig:0458td}, \\ref{fig:0528td}, \\ref{fig:1730td} and \\ref{fig:2230td}) show significant dispersions for the time delays at separate frequencies. If we restrict our consideration to the approximating functions with the smallest residuals (0528+134 -- D, E; 2230+114 -- A3, B3), the corresponding mean value is $\\alpha_{mean}=(-0.85\\pm0.27)$. Both the ``pessimistic'' and ``optimistic'' estimates for the mean time-lag index are consistent with the expected value of --0.8, taking into account the fairly large errors. The only exception might be the index for 1730--130.  \\begin{table} \\begin{center} \\vspace{2mm} \\caption{\\bf{Parameters of the time-delay function for core outbursts}} \\begin{tabular}{lllll} \\hline %\\noalign{smallskip} Source &  Comp.  &  $T_{max}$(37GHz) & $\\alpha$ \\\\ 0458-020 &  A  &   1989.1  &  -1.46  \\\\ %0528+134 &  A  &   1992.2  &  -0.83  \\\\ %0528+134 &  B  &   1993.5  &  -0.36  \\\\ %0528+134 &  C  &   1994.4  &  -0.65  \\\\ 0528+134 &  D  &   1995.9  &  -0.87  \\\\ 0528+134 &  E  &   1998.4  &  -1.22  \\\\ 0528+134 &  F  &   1999.8  &  -1.31  \\\\ 1730-130 &  A  &   1996.3  &  -2.16  \\\\ 2230+114 &  A2 &   1989.6  &  -0.89  \\\\ 2230+114 &  B2 &   1990.5  &  -1.55  \\\\ 2230+114 &  A3 &   1997.5  &  -0.72  \\\\ 2230+114 &  B3 &   1997.9  &  -0.60  \\\\ \\hline \\end{tabular} \\end{center} \\label{tab:summary} \\end{table}  The shift of the core position between two frequencies is expected to depend on the luminosity of the source (Lobanov 1998): $$ \\Delta r\\propto L_{syn}^{2/3} %\\Delta r\\propto[\\frac{L_{syn} \\sin{\\theta}}{\\beta(1-\\beta \\cos{\\theta})}]^{2/3} $$ As above, we expect a similar dependence for the time lags. The relationship between the amplitudes of the outbursts in two subsequent activity cycles in 2230+114 and their time lags is in qualitative agreement with the expectation (Fig.~\\ref{fig:2230td}; the outbursts in cycle 3 display both higher amplitudes and larger time delays). However, the time lags for individual sub-outbursts within a single outburst do not display a similar relation with the sub-outburst amplitudes (0528+134). The time lags for the individual sub-outbursts %in an outburst may depend on the dynamical evolution of the underlying perturbation, rather than on their amplitude.  \\subsection{Outbursts and structural evolution}  As was already shown by Savolainen et al. (2002), the emergences of all bright components are correlated with local maxima of the light curves. A comparison of the 37~GHz light curves with the flux densities of the cores and new superluminal components as functions of time (Figs.~\\ref{fig:0458}, \\ref{fig:0528}, \\ref{fig:1730}, \\ref{fig:2230vlbi}) augments the correlation. The flux densities of both the cores and new-born jet components reach their highest values near the maxima of the 37~GHz light curves, after which they fade.  \\begin{figure*} \\centering \\begin{minipage}[c]{\\textwidth} \\centering \\includegraphics[width=0.45\\textwidth]{pyatunina_fig12left.ps} %   \\includegraphics[width=0.45\\textwidth]{0528_rf.ps} \\hspace{0.02\\textwidth} \\includegraphics[width=0.45\\textwidth]{pyatunina_fig12right.ps} %   \\includegraphics[width=0.45\\textwidth]{2230_mot.ps} \\end{minipage} \\vspace*{-0.5cm} \\caption{Distance from the core and flux as functions of time for component B4 in {\\bf 0528+134} ({\\it left}) and the components A1, B2 and B3 in {\\bf 2230+114} ({\\it right}) (Jorstad et al. 2005).} \\label{fig:0528+2230vlbi} \\end{figure*}  Some well studied events (Fig.\\ref{fig:0528+2230vlbi}) indicate that the times when the flux densities of new-born components reach their maxima nearly coincide with the times when they experience a sudden increase in their velocity of separation from the core. A similar increase in the velocity of the superluminal component C8$^*$ was observed near 1998 in the quasar 3C279 (Jorstad et al. 2005); note the close coincidence of this epoch with the maximum of the brightest outburst observed in 3C279 since 1965. This suggests that the flux density may reach its maximum when a shocked region that gives birth to a new jet component becomes optically thin and enters a free-expansion regime, as has been suggested by Gomez et al. (1997) (although changes in brightness occuring simultaneously with accelerations can also result from geometric effects, such as bending of the jet toward or away from the line of sight; see, e.g., Homan et al. 2003). Another intriguing phenomenon observed in 2230+114 is the interaction of the stationary component A1 with the moving component B3, when A1 is dragged downstream by B3, then returns upstream, increasing its flux density (Gomez et al. 2001).  \\subsection{Conclusions}  Our main results can be summarized as follows.  %Essentially all integrated optically thick outbursts associated %with birth of new components when verified with VLBI at high frequencies  (1) Frequency-dependent time delays for bright outbursts in four $\\gamma$-ray quasars have been estimated and approximated using an exponential function of frequency, $T_{max}(\\nu)\\propto\\nu^{\\alpha}$. The time delays in the sources range from $\\sim0.3$~yr to $\\sim1.0$~yr. The mean value of the index of the time-delay function is $\\alpha_{mean}=(-1.2\\pm0.5)$, and the maximum value in terms of its magnitude is $\\alpha=(-2.2\\pm0.3)$, for the source 1730--130.  (2) The observed outbursts in 0528+134 and 2230+114 reveal fine structure and consist of narrow ($\\sim1$ year) sub-outbursts. Some traces of unresolved fine structure are also seen in 0458--020 and 1730--130. The time delays and exponential-function indices vary from one sub-outburst to another.  (3) Due to the prominent fine structure observed in some outbursts, identification of the starting point of the activity requires information about the spectral evolution at frequencies higher than 40~GHz or the structural evolution at mm wavelengths.  (4) Lower limits for the durations of the activity cycles, or activity time scales, have been estimated for 0458--020 ($T_{act}\\ge16$~yr) and 1730--130 ($T_{act}\\ge9$~yr). The results of VLBI imaging at 8~GHz (Britzen et al. 1999b) and 43~GHz (Jorstad et al. 2001, 2005) suggest that the approximate duration of the activity cycle for 0528+134 is $T_{act}\\sim14$~yr. Bright outbursts in 2230+114 repeat with a quasi-period of $8.04\\pm0.30$~yr. Although the amplitudes of individual sub-outbursts vary from one quasi-period to another, their positions within the overall outburst are preserved. The time lags for two adjacent periods of activity in 2230+114 suggest a correlation with outburst amplitude: the brighter outburst shows a greater time delay. The first maximum of the following bright outburst in 2230+114 is expected to be near $2005.5\\pm0.3$.  (5) The maximum flux densities of new jet components often coincide with a sudden increase in their velocities from the core. This suggests that the component's flux reaches its maximum when the shocked region that corresponds to the jet component becomes optically thin and begins to freely expand, as was suggested by Gomez et al. (1997).  %We are planning to increase our sample of sources in the future, in order to %verify the suggestions made in this paper."
    },
    "0609/nucl-ex0609013_arXiv.txt": {
        "abstract": "The nuclear physics input from the $^3$He($\\alpha,\\gamma$)$^7$Be cross section is a major uncertainty in the fluxes of $^7$Be and $^8$B neutrinos from the Sun predicted by solar models and in the $^7$Li abundance obtained in big-bang nucleosynthesis calculations. The present work reports on a new precision experiment using the activation technique at energies directly relevant to big-bang nucleosynthesis. Previously such low energies had been reached experimentally only by the prompt-$\\gamma$ technique and with inferior precision. Using a windowless gas target, high beam intensity and low background $\\gamma$-counting facilities, the $^3$He($\\alpha,\\gamma$)$^7$Be cross section has been determined at 127, 148 and 169\\,keV center-of-mass energy with a total uncertainty of 4\\,\\%. The sources of systematic uncertainty are discussed in detail. The present data can be used in big-bang nucleosynthesis calculations and to constrain the extrapolation of the $^3$He($\\alpha,\\gamma$)$^7$Be astrophysical S-factor to solar energies. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609232_arXiv.txt": {
        "abstract": "{}{We present \\Nif~mass estimates for seventeen well-observed type~Ia supernovae determined by two independent methods.} {Estimates of the \\Nif~mass for each type~Ia supernova are determined from (1) modeling of the late-time nebular spectrum and (2) through the combination of the peak bolometric luminosity with Arnett's rule. The attractiveness of this approach is that the comparison of estimated \\Nif~masses circumvents errors associated with the uncertainty in the adopted values of reddening and distance. } {We demonstrate that these two methods provide consistent estimates of the amount of \\Nif~synthesized. We also find a strong correlation between the derived \\Nif~mass and the absolute $B$-band magnitude (M$_{B}$).} {Spectral synthesis can be used as a diagnostic to study the explosion mechanism. By obtaining more nebular spectra the \\Nif--M$_{B}$ correlation can be calibrated and can be used to investigate any potential systematic effects this relationship may have on the determination of cosmological parameters, and provide a new way to estimate extra-galactic distances of nearby type Ia supernovae.} ",
        "introduction": "Type~Ia supernovae (hereafter SNe~Ia) are a vital tool for our understanding of the universe. Due to their high luminosities and remarkable uniformity they are excellent cosmological distance indicators, and allow us to determine the evolution of cosmic expansion \\citep[see e.g.][$~$and references therein]{leibundgut01}.  SNe~Ia are used  as high precision cosmological probes \\citep[see][]{riess04,astier06}. To maximize the level of precision achieved from current and future experiments requires the ability to understand and reduce any source of systematic error. The main motivation of this work was thus to learn more about the relationship between the \\Nif~mass and peak absolute magnitude.  To understand the dependence of the absolute brightness on the \\Nif~mass requires knowledge of the progenitor system and explosion mechanism of SNe~Ia. Despite considerable efforts on the theoretical side to investigate the explosion physics, our current knowledge of the ignition mechanism and subsequent flame propagation through the progenitor star is still rudimentary, and in some instances debatable \\citep{hillebrandt00}. From an observational perspective, considerable effort has been put into supernova monitoring programs and many excellent data sets have been assembled. With these data it is now possible to conduct detailed investigations into the physical mechanism leading to the observed diversity in the SNe~Ia population.  Some first steps have been made in this direction, e.g. \\citet{cappellaro97,mazzali98,contardo00,mazzali01,suntzeff03,benetti05,stritzinger06a,branch06}.  The amount of \\Nif~produced in a SN~Ia explosion is a  key parameter that can be derived from observations and is important for our understanding of the luminosity decline-rate relation. In this article we demonstrate that the amount of \\Nif~computed from modeling the late-time nebular spectrum of a SN~Ia is consistent with estimates obtained from the combination of the UltraViolet Optical near-InfraRed (UVOIR) bolometric light curve with Arnett's rule \\citep{arnett82}. We also find a strong correlation between the derived \\Nif~mass and M$_{B}$. As we can now estimate the mass of synthesized \\Nif~% accurately from fitting the nebular spectrum this method may provide a means to understand possible sources of systematic error that could effect the analysis of data from future high-z SNe~Ia surveys.  The structure of this paper is as follows. In Sect.~2 we provide a brief description of the observational data considered in this study. Section~3 describes the two methods used to estimate the \\Nif~mass. Section~4 contains the results, and is then followed by a discussion in Sect.~5. ",
        "conclusions": "The agreement between the \\Nif~mass estimates shown in Fig.~2 is encouraging as these two methods are independent and use different observational data. The ability to accurately probe the explosion mechanism with one spectrum taken more than 100 days after the explosion reveals the power of the spectral synthesis method.  The range of a factor of ten in the amount of synthesized \\Nif, as indicated by the spectral synthesis method, is consistent with results obtained from previous studies of UVOIR light curves \\citep{suntzeff96,cappellaro97,contardo00,strolger02,suntzeff03,stritzinger06a}. The observed range in \\Nif~mass is difficult to reconcile with the current SN~Ia paradigm \\citep{nomoto84,woosley86,hillebrandt00}, i.e. the thermonuclear disruption of a Chandrasekhar-size white dwarf. This has lead to the introduction of the so-called delayed detonation model \\citep{khokhlov91,woosley90,woosley94a,hoflich96}, and most recently more exotic explosion mechanisms \\citep{bravo06}. An alternative to the Chandrasekhar mass model that addresses the observed diversity of SNe~Ia properties is one involving a variable progenitor mass. There is observational evidence that sub-Chandrasekhar mass white dwarfs are a plausible candidate for the less luminous events and at least for the sub-luminous SN~1991bg-like variety of SNe~Ia \\citep{mazzali97,cappellaro97,stritzinger06a}. Another possible progenitor maybe a super-Chandrasekhar mass progenitor \\citep{yoon04,yoon05}.     Efforts to model the light curves of SNe~Ia led to the suggestion that the absolute luminosity at maximum light [hence $\\Delta$m$_{15}$($B$)] is a direct consequence of the \\Nif~mass \\citep{arnett82,hoeflich96,cappellaro97,mazzali01,pinto01,kasen06}. The strong correlation shown in Fig.~3 between the \\Nif~mass and $\\Delta$m$_{15}$($B$) provides observational evidence that supports this hypothesis. This argument is strengthened by the excellent agreement between the Phillips M$_{B}$--$\\Delta$m$_{15}$($B$) relation and the results in Fig.~3. With the acquisition of more nebular spectra it should be possible to calibrate the \\Nif~mass--M$_{B}$ relation. With this calibration one could then use a nebular spectrum to (1) determine the distance to a nearby SN~Ia, and more importantly (2) address the relationship between the \\Nif~mass and the peak absolute magnitudes. This method is independent of the luminosity-decline rate relation, so distances determined through this method are not as susceptible to the secondary parameters that may affect the peak phase light curve \\citep{hamuy96,mazzali98,tripp99,hantano00,benetti04,benetti05}. The late-time bolometric luminosity, however, may be susceptible to the distribution of isotopes and/or the fraction of positron kinetic energy that is deposited within the ejecta. Nonetheless, by gaining a deeper understanding of how the peak absolute brightness relates to the \\Nif~mass we increase our understanding of the explosion mechanism and any possible systematic errors incurred via this relationship.  The scatter seen between $\\Delta$m$_{15}$($UVOIR$) and $\\Delta$m$_{15}$($B$) hints at the complexities associated with the transfer of radiation through expanding ejecta, and how time-dependent NLTE effects influence the flux evolution of the near-IR light curves. Indeed a comparison between the near-IR light curves for events with similar values of $\\Delta$m$_{15}$($B$) indicates a diversity in their morphology.  The knowledge afforded by the spectral synthesis method is significant and warrants further SNe~Ia observational programs to obtain well-calibrated nebular spectra."
    },
    "0609/astro-ph0609718_arXiv.txt": {
        "abstract": "{}{ We address the origin and evolutionary status of hot subdwarf stars by studying the optical spectral properties of 58 subdwarf O (sdO) stars. Combining them with the results of our previously studied subdwarf B (sdB) stars, we aim at investigating possible evolutionary links. } {We analyze high-resolution (R$>$18000), high-quality optical spectra of sdO stars obtained with the ESO VLT UVES echelle spectrograph in the course of the ESO Supernova Ia Progenitor Survey (SPY). Effective temperatures, surface gravities, and photospheric helium abundances are determined simultaneously by fitting the profiles of hydrogen and helium lines using dedicated synthetic spectra calculated from an extensive grid of NLTE model atmospheres. } { We find spectroscopic or photometric evidence for cool companions to eight sdO stars as well as a binary consisting of two sdO stars. A clear correlation between helium abundances and the presence of carbon and/or nitrogen lines is found: below solar helium abundance, no sdO star shows C or N lines. In contrast, C and/or N lines are present in the spectra of all sdO stars with supersolar helium abundance. We thus use the solar helium abundance to divide our sample into \\emph{helium-deficient} and \\emph{helium-enriched} sdO stars. While \\emph{helium-deficient} sdO stars are scattered in a wide range of the \\tefflogg-diagram, most of the \\emph{helium-enriched} sdO stars cluster in a narrow region  at temperatures between 40\\,000 and 50\\,000K and gravities between $\\log$ g=5.5 and 6.0. } { An evolutionary link between sdB stars and sdO stars appears plausible only for the \\emph{helium-deficient} sdO stars. They probably have evolved away from the extreme horizontal branch, i.e. they are the likely successors to sdB stars. In contrast, the atmospheric properties of \\emph{helium-enriched} sdO stars cannot be explained with canonical single star evolutionary models. Alternative scenarios for both single star (late hot flasher) as well as for binary evolution (white dwarf merger; post-RGB evolution) fail to reproduce the observed properties of \\emph{helium-enriched} sdO stars in detail. While we regard the post-RGB scenario as inappropriate, the white dwarf merger and the late hot flasher scenarios remain viable to explain the origin of \\emph{helium-enriched} sdO stars. } ",
        "introduction": "}  Hot subluminous stars are an important population of faint blue stars at high galactic latitudes closely related to the horizontal branch. A proper spectral classification of hot subluminous stars is rendered difficult by the diversity of the helium line spectra. They can roughly be grouped into the cooler sdB stars, whose spectra typically display no or only weak helium lines, and the hotter sdO stars, which have a higher helium abundance on average and can even be dominated by helium. The former have recently been studied extensively because they are sufficiently common to account for the UV excess observed in early-type galaxies. Pulsating sdB stars are important tools for asteroseismology \\citep{char04} and sdB stars in close binaries may qualify as Supernova Ia progenitors \\citep{maxt00,geie06}.  Subluminous B stars have been identified as Extreme Horizontal Branch (EHB) stars \\citep{heb86}, i.e. they are core helium burning stars with hydrogen envelopes too thin to sustain hydrogen burning (unlike normal HB stars). Therefore they evolve directly to the white dwarf cooling sequence avoiding the Asymptotic Giant Branch (AGB). While the sdB stars spectroscopically form a homogenous class, a large variety of spectra is observed among sdO stars \\citep{heb92,heb06}. Most subluminous B stars are helium poor, whereas only a relatively small fraction of sdO stars are.  Ever since the pioneering work by \\citet{green74}, the helium-rich sdO stars were believed to be linked to the evolution of the hydrogen-rich subluminous B stars. Any evolutionary link between subluminous B and O stars, however, is difficult to explain since the physical processes driving a transformation of a hydrogen-rich star into a helium-rich one remain obscure. The convective transformation has been explored by \\citet{wese81} as well as by \\citet{grot85}. While the former found helium convection to occur even at subsolar helium abundances which mixes helium from deeper layers into the photosphere, the latter concluded that a helium driven convection zone develops only in helium-rich atmospheres. If the latter is true, convective transformation would not work. Non-standard evolutionary models were introduced to explain the formation of sdO stars \\citep[e.g.][]{swei97,bro01,moe04}. In particular, the {\\it late hot flasher scenario} predicts that the core helium flash may occur when the star has already left the red giant branch (RGB) and is approaching the white dwarf cooling sequence (delayed He core flash). During the flash, He and C are dredged-up to the surface. Hydrogen is mixed into deeper layers and burnt. The remnant is found to lie close to the helium main sequence, i.e. at the very end of the theoretical extreme horizontal branch.  The fraction of sdB stars in short period binaries (periods less than ten days) is high. \\citet{max01} found 2/3 of their sdB sample (drawn from the Palomar Green survey) to be such binaries, whereas somewhat lower fractions of 40\\% and 48\\% were found recently for the SPY sample \\citep{napi04} and for a sample drawn from the Edinburgh-Cape catalog \\citep{mora06}, respectively. Obviously, binary evolution plays an important role for the formation of sdB stars and possibly also for that of the sdO stars. A recent population synthesis study \\citep{han03} identified three channels to form sdB stars: (i) one or two phases of common envelope evolution, (ii) stable Roche lobe overflow, and (iii) the merger of two helium-core white dwarfs. The latter could explain the population of single stars. The simulations by \\citet{han03}  cover the observed parameter range of sdB stars \\citep[see][henceforth paper~I]{lis05}.  The ESO Supernova Ia Progenitor SurveY \\citep[SPY,][]{spy0} has obtained VLT/UVES spectra for over 1000 white dwarf candidates to test possible scenarios for type Ia supernovae by searching for double degenerate white dwarf binary systems close to the Chandrasekhar mass limit. Many of the target stars of SPY came from the Hamburg ESO survey \\citep{wiso96}. SPY also observed 137 hot subluminous stars that entered the target sample because they were previously classified as white dwarfs. 76 of these stars were now classified as sdB/sdOB, and 58 as O-type subdwarfs \\citep{chri01,lis03}.  The data for our hot subdwarf sample are of unprecedented quality and homogeneity. Spectral analyses of all sdB stars from that sample have already been presented in paper~I in order to test evolutionary models, in particular the binary population models of \\citet{han03}. Two diagnostic tools -- the effective temperature ($\\teff$) vs.\\ surface gravity ($\\logg$) diagram and the cumulative luminosity function -- yielded conflicting results. Moreover, the models of \\citet{han03} predicted some EHB stars to be hotter than the sdB stars contained in the sample of paper\\,I. This led us in paper I to the conclusion that their sample of hot subdwarfs may not be sufficiently complete to describe the whole parameter range covered by the simulations, and that it needs to be complemented by a similar analysis of subdwarf O stars.  In this paper we focus on the subdwarf O spectra from the SPY sample and present the spectral analysis of the high resolution spectra using state-of-the-art NLTE model atmospheres. This enables us to address the still open question about the evolutionary status of hot subdwarfs. We combine our results for the SPY sdO stars with those for the SPY sdB stars from paper\\,I to cover the entire stellar atmospheric parameter range with the same high accuracy. This allows for the first time a detailed comparison with evolutionary model predictions, particularly in the regime of hotter temperatures and higher helium abundances than in previous studies of sdB stars alone.  The paper is organised as follows. In section~\\ref{sec:data} we outline the available data and explain our criteria for spectral classification. Section~\\ref{sec:nlte} briefly describes the construction of a new NLTE model atmosphere grid. In section~\\ref{sec:results} we detail the derivation of stellar parameters by fitting model atmospheres to observed spectral lines and present the results in section~\\ref{sec:result2}. Various evolutionary scenarios are tested in section~\\ref{sec:obstheo}, leading to a discussion and summary in section~\\ref{sec:discuss}. ",
        "conclusions": "}   We analyzed  high resolution optical spectra of 58 subluminous O stars. We found spectroscopic and/or photometric evidence for cool companions to eight subluminous O stars. While HE~1502$-$1019 and HE~1513$-$0432 display only weak helium lines, the others have strong helium lines. HE~0301$-$3039 is an unique binary consisting of two sdO stars with spectra dominated by helium lines \\citep{lis04}. Evidence for cool companions to the other helium-enriched stars rests on photometry only. Near-infrared spectroscopy is required for classifying them. As the photometric measurements for our sample are only available for 52 out of 58 stars, the fraction of sdO stars with cool companions may be slightly larger.  A grid of synthetic H-/He-line spectra calculated from NLTE model atmospheres was used to derive the stars' atmospheric parameters. Reliable atmospheric parameters were obtained for 13 \\emph{helium-deficient} sdO and 33 \\emph{helium-enriched} sdO stars. A clear-cut correlation between CN class and helium abundance was found. \\emph{Helium-deficient} sdO stars did not show any C and/or N lines, while all \\emph{helium-enriched} sdO stars do show either C or N lines or both, suggesting that they form two different populations of stars. The \\emph{helium-deficient} sdO stars are scattered in the \\tefflogg\\ diagram, whereas most \\emph{helium-enriched} sdO stars cluster in a narrow range.    Comparing the observed distribution to the predictions of evolutionary calculations for single as well as for close binary stars, we conclude that many  \\emph{helium-deficient} sdO stars can be explained as evolved sdB stars. Both classes of star are helium deficient and have weak metal lines which is caused by atmospheric diffusion processes.  Most of the \\emph{helium-enriched} sdO stars cluster in a narrow region of the \\tefflogg-diagram at temperatures between 40kK and 50kK. While diffusion is probably causing helium deficiency, it is unlikely to account for the helium enrichment. Non-standard evolutionary scenarios were therefore considered as well. The predictions from the late hot flasher scenario as well as the helium white dwarf merger scenario are roughly consistent with the observed distribution of \\emph{helium-enriched} sdO stars but do not match them in detail. The occurrence of a delayed helium core flash as well as the merger of two helium white dwarfs may explain the helium enrichment. In both cases carbon and/or nitrogen can be dredged up to the stellar surface, which would explain the strength of the C and/or N lines in \\emph{helium-enriched} sdO stars.    Some high gravity \\emph{helium-enriched} sdO stars may lie below the helium main sequence, which is at variance with any core helium burning model. Therefore, we considered models for post-RGB stars with inert helium cores which evolve through the sdB/sdO regime in the \\tefflogg\\ diagram into helium core white dwarfs. This is the only scenario that can explain stars to lie below the helium main sequence. However, it is not obvious how the helium enrichment is brought about as the hydrogen envelopes of post-RGB models has to be relatively thick or, otherwise, the stars would evolve too fast to be observable in large quantities. If the existence of a population of high gravity sdO stars below the helium main sequence could be confirmed by surveys with larger sample size, a possible explanation could be provided by the post-RBG scenario.  Our conclusions can be tested by measuring the binary frequency and the C and N abundances. If \\emph{helium-deficient} sdO stars are evolved sdB stars, their binary fractions should be the same. As about 40\\% of the sdB stars are in close binaries with periods below 10 days \\citep{napi04}, we expect 5 \\emph{helium-deficient} sdO stars in our sample to be radial velocity variable on time scales of 10 days or less. An investigation of the binary frequency is underway in our group. If \\emph{helium-enriched} sdO stars result from mergers, no radial velocity variations would be expected (except for objects arising from triple stars). If material processed by nuclear burning is dredged up, the helium enrichment is expected to be accompanied by enrichment of nitrogen and/or carbon. While our CN-classification scheme (section~\\ref{sec:class}) is based on the \\emph{presence} of C and/or N-lines in the spectrum, the latter issue requires a \\emph{quantitative} abundance analysis to test the late hot flasher and the merger scenario.  In addition, improvements to our analyses of the SPY-sample are required for a more detailed comparison with predictions of evolutionary calculations. The role of metal-line blanketing in the NLTE model atmospheres needs to be investigated. Due to the complexity of the problem, only few studies of NLTE metal line blanketing in sdO star atmospheres are available \\citep{haas96,lanz97,deet00} which indicate that the effective temperatures may have to be reduced. According to these investigations, the most important contribution to line blanketing stems from the iron group elements. However, their abundances can be derived from ultraviolet spectra only. The observed diversity of carbon and nitrogen line strengths indicates that the treatment of metals has to be done carefully on a star-by-star basis.  Apparently, the SPY sample suffers from observational selection biases. We found it likely that sdO stars were selected preferentially for the SPY target list. Hence sdO stars may be overrepresented with respect to sdB stars. To this end the Sloan Digital Sky Survey (SDSS) is a promising source of hot subdwarf stars, as its selection criteria are very different from that of SPY. We have already begun a spectroscopic analysis of SDSS sdO stars \\citep{heb06} which will contribute significantly towards discriminating between the various evolutionary scenarios and hypotheses outlined above."
    },
    "0609/astro-ph0609004_arXiv.txt": {
        "abstract": "We have developed a new three-dimensional general relativistic magnetohydrodynamic (GRMHD) code, RAISHIN, using a conservative, high resolution shock-capturing scheme. The numerical fluxes are calculated using the Harten, Lax, \\& van Leer (HLL) approximate Riemann solver scheme. The flux-interpolated, constrained transport scheme is used to maintain a divergence-free magnetic field. In order to examine the numerical accuracy and the numerical efficiency, the code uses four different reconstruction methods: piecewise linear methods with Minmod and MC slope-limiter function, convex essentially non-oscillatory (CENO) method, and piecewise parabolic method (PPM) using multistep TVD Runge-Kutta time advance methods with second and third-order time accuracy. We describe code performance on an extensive set of test problems in both special and general relativity. Our new GRMHD code has proven to be accurate in second order and has successfully passed with all tests performed, including highly relativistic and magnetized cases in both special and general relativity. ",
        "introduction": "Both magnetic and gravitational fields play an important role in determining the evolution of the matter in many astrophysical objects. In highly conducting plasma, the magnetic field can be amplified by gas contraction or shear motion. Even when the magnetic field is weak initially, the magnetic field grows short time scales and influences the gas dynamics of the system. This is particulary important for a compact object such as a black hole or a neutron star. Relativistic jets have been observed or postulated in various astrophysical objects, including active galactic nuclei (AGNs) (e.g., Urry \\& Pavovani 1995; Ferrari 1998; Blandford 2002), microquasars in our galaxy (e.g., Mirabel \\& Rodr\\'{i}guez 1999), and gamma-ray bursts (GRBs) (e.g., Zhang \\& M\\'{e}sz\\'{a}ros 2004; Piran 2005; M\\'{e}sz\\'{a}ros 2006). The most promising mechanisms for producing the relativistic jets involve the magnetohydrodynamic centrifugal acceleration  and/or magnetic pressure driven acceleration from the accretion disk with the compact objects (e.g., Blandford \\& Payne 1982; Fukue 1990), or the extraction of rotating energy from the rotating black hole (Penrose 1969; Blandford \\& Znajek 1977). In addition, the differential rotation of the plasma in the disk raises the magnetorotational instability (MRI), which plays an important role for the transportation of angular momentum from the disk due to the associated turbulence in accretion disks (Balbus \\& Hawley 1991, 1998). It is much more efficient than the internal disk transport. Magnetars are neutron stars with extremely large magnetic fields ($\\sim 10^{14}-10^{15} G$), as inferred from studies of anomalous X-ray pulsars and soft gamma-ray repeaters (e.g., Woods \\& Thompson 2006)). These intense magnetic fields are expected to affect the internal structure of the neutron star (e.g., Bocquet et al. 1995). The less intense fields of other neutron stars, $\\sim 10^{12} - 10^{13} G$, may also affect the internal structure.  There has been much progresses in the analytical studies of relativistic astrophysical phenomena (e.g., Camenzind 1990; Fendt 1997; Fendt \\& Ouyed 2004). However, the problem is complex, involving time-dependent, three-dimensional dynamics of magnetized plasmas in the relativistic potential. Therefore analytical solutions are rather limited to simplified cases that are time stationary and spatially symmetric. Because of these, numerical experiments play an important role for complimenting theoretical works.  A complete review of numerical approaches to relativistic hydrodynamics is given by Mart\\'{i} \\& M\\\"{u}ller (2003) and Font (2003). Numerical codes in special relativistic magnetohydrodynamics (SRMHD) have been developed by a growing number of authors (van Putten 1993; Koide et al. 1996; Komissarov 1999a; Balsara 2001; Koldoba et al. 2002; Del Zanna, Buccianti, \\& Londrillo 2003; Leismann et al. 2005; Mignone \\& Bodo 2006) and have been applied to the study of the relativistic jets (Koide et al. 1996, 1997; Koide 1997; Nishikawa et al. 1997, 1998; Kommisarov 1999b; Leismann et al. 2005; Mignone \\& Bodo 2006) and pulsar wind nebulae (Kommisarov \\& Lyubarsky 2003, 2004; Bucciantini et al. 2003, 2004, 2005, 2006; Del Zanna et al. 2004). In addition, Komissarov (1999a) and Balsara (2001) proposed a comprehensive set of tests to validate numerical codes in special relativity. The exact solutions of the Riemann problem in the SRMHD have been obtained by Giacomazzo \\& Rezzolla (2006).  In order to investigate relativistic magnetorotators (RMRs), general relativistic magnetohydrodynamics (GRMHD) codes with fixed spacetimes have been developed by some authors (Yokosawa 1993; Koide et al.1998; De Villiers \\& Hawley 2003; Gammie et al. 2003; Komissarov 2004; Ant\\'{o}n et al. 2005; Anninos et al. 2005).  These codes have been used to study the structure of accretion flows onto Kerr black holes and/or the formation of jets (Yokosawa 1995; Koide et al. 1998, 1999, 2000; Nishikawa et al. 2005; De Villiers et al. 2003, 2005a; Hirose et al. 2004; Krolik et al. 2005; Hawley \\& Krolik 2006; McKinney \\& Gammie 2004; McKinney 2005, 2006), the Blandford-Znajek effect near the rotating black hole (Koide et al. 2002; Koide 2003; Komissarov 2005), and the formation of GRB jets in collapsars (Mizuno et al. 2004a, 2004b; De Villiers et al. 2005b).  Few attempts have been made to simulate relativistic MHD flows in dynamical spacetimes (e.g., Wilson 1975, Baumgarte \\& Shapiro 2003). Recently new codes capable of evolving the Einstein-Maxwell-MHD equations have been developed by Duez et al. (2005), Shibata \\& Sekiguchi (2005) and Anderson et al. (2006). These codes have been applied to investigate rotating neutron stars (Duez et al. 2006a,b), and collapsing neutron stars for a central engine for short gamma-ray bursts (Shibata et al. 2006).  These new  GRMHD simulations with dynamical spacetimes will be applied to investigate various astrophysical systems such as merging neutron stars and black holes including gravitational waves and  relativistic jets (e.g., Duez et al. 2005).  Our previous GRMHD code developed by Koide has been applied to many high-energy astrophysical phenomena and showed pioneering results (Koide et al. 1998, 1999, 2000, 2002; Koide 2003, 2004, 2006; Mizuno et al. 2004a,b; Nishikawa et al. 2005). However, the code cannot perform calculations in highly relativistic or highly magnetized regimes. The critical problem of the previous GRMHD code is that it cannot guarantee a divergence-free magnetic field. Even though we introduce a divergence-cleaning step in the code, it cannot perform long-term evolution, except in special cases. In order to overcome these numerical difficulties, we have developed a new, three-dimensional, GRMHD code called RAISHIN, for RelAtIviStic magnetoHydrodynamic sImulatioN. (RAISHIN is the ancient Japanese god of lightning.) It uses a conservative, high resolution shock-capturing scheme.  The structure of the paper is as follows. In section 2 we describe the basic equations of GRMHD including the essentials of the 3$+$1 formalism, the description of the magnetic field, the induction equation and the conservation equations of particle number, and the stress-energy tensor in conservative form. In section 3 we describe the basic algorithm of our new GRMHD code based on the recent modern techniques such as HLL approximate Riemann solver, various reconstruction methods and flux-interpolated constrained transport. We show the performance of the code on a series of test problems in section 4. The summary and conclusions are given in section 5.  The new simulation results of jet formation using this code will be reported separately (Mizuno et al. 2006). ",
        "conclusions": "We have developed a new three-dimensional GRMHD code, RAISHIN, by using a conservative high resolution shock-capturing scheme. The numerical fluxes are calculated using the HLL approximate Riemann solver scheme. The flux-interpolated, constrained transport scheme is used to maintain a divergence-free magnetic field. Several reconstruction and time advance schemes can be chosen for the numerical accuracy and computational resources.  We have described several test problems in both special and general relativity. They have shown significant improvements over our previous GRMHD code (Koide 2003). Our new GRMHD code can perform in the regimes of high Lorentz factor ($ \\gamma >20 $) and high magnetic field ($\\sigma > 550$), and in the presence of large discontinuity of density, pressure and magnetic field. We have compared the results of several reconstruction schemes. The code is second-order accurate even when we use the higher order reconstruction schemes such as CENO and PPM. Nevertheless, higher-order reconstruction schemes can provide more accurate results for some applications. The PPM reconstruction scheme allows the well-resolved detection of sharp discontinuities. We found the limitation to handle the regimes of high Lorentz factor and of highly magnetized depends on the schemes to solve the GRMHD equations.  We have performed several simulations of non-rotating and rotating black hole systems with a thin accretion disk (Mizuno et al. 2006). The simulation results show the formation of jets driven by the Lorentz force and the gas pressure. It appears that the rotating black hole creates an additional, faster, and more collimated inner outflow beside an outflow generated by the rotating accretion disk in the non-rotating black hole. Thus, kinematic jet structure could be a sensitive function of the black hole rotation.  Our new GRMHD code has proven to be accurate in second order and has successfully passed all applied tests including highly relativistic cases , and highly magnetized cases in both special and general relativity. We plan to apply this code to a number of high-energy astrophysical phenomena involving highly relativistic flows or compact objects with strong gravitational fields and magnetic fields."
    },
    "0609/astro-ph0609374_arXiv.txt": {
        "abstract": "We use highly spectroscopically complete observations of the radio sources from the VLA 1.4~GHz survey of the Hubble Deep Field-North region to study the faint radio galaxy population and its evolution. The fraction of radio sources that can be optically spectroscopically identified is fairly independent of radio flux, with about $60-80$\\% identified at all fluxes. We spectrally classify the sources into four spectral types: absorbers, star formers, Seyfert galaxies, and broad-line AGNs, and we analyze their properties by type. We supplement the spectroscopic redshifts with photometric redshifts measured from the rest-frame ultraviolet to mid-infrared spectral energy distributions. Using deep X-ray observations of the field, we do not confirm the existence of an X-ray--radio correlation for star-forming galaxies. We also do not observe any correlations between 1.4~GHz flux and $R$ magnitude or redshift. We find that the radio powers of the host galaxies rise dramatically with increasing redshift, while the optical properties of the host galaxies show at most small changes. Assuming that the locally determined far-infrared (FIR)--radio correlation holds at high redshifts, we estimate total FIR luminosities for the radio sources. We note that the FIR luminosity estimates for any radio-loud AGNs, which we conservatively do not try to remove from the sample, will be overestimates. Considering only the radio sources with quasar-like bolometric luminosities, we find a maximum ratio of candidate highly-obscured AGNs to X-ray--luminous ($L_{0.5-2~{\\rm keV}}$ or $L_{2-8~{\\rm keV}}\\ge 10^{42}$~ergs~s$^{-1}$) sources of about 1.9. Finally, we use source-stacking analyses to measure the X-ray surface brightnesses of various X-ray and radio populations. We find the contributions to the $4-8$~keV light from our candidate highly-obscured AGNs to be very small, and hence these sources are unable to account for the light that has been suggested may be missing at these energies. ",
        "introduction": "\\label{secintro}  In the local universe, Seyfert~2 galaxies outnumber Seyfert~1 galaxies by 4 to 1 (Maiolino \\& Rieke 1995), and more than half of Seyfert 2 nuclei are Compton-thick with obscuring column densities $N_H>10^{24}$~cm$^{-2}$ (Maiolino et al.\\ 1998; Risalti et al.\\ 1999; Guainazzi et al.\\ 2005; Comastri 2004). In the high-redshift universe, X-ray background population synthesis models (refined after the discovery of a lower redshift distribution than expected for the hard X-ray sources; e.g., Barger et al.\\ 2002, 2003; Szokoly et al.\\ 2004) similarly predict a population of highly-obscured active galactic nuclei (AGNs) that are missed in the deep {\\em Chandra\\/} and {\\em XMM-Newton\\/} hard X-ray surveys (e.g., Franceschini et al.\\ 2002; Gandhi \\& Fabian 2003; Ueda et al.\\ 2003; Comastri 2004; Gilli 2004; Fabian \\& Worsley 2004; Treister \\& Urry 2005; Ballantyne et al.\\ 2006). Such theoretical predictions appear to be receiving some observational confirmation in the work of Worsley et al.\\ (2005). Using a source-stacking analysis on the {\\em Chandra} Deep Field-North (CDF-N) and {\\em Chandra} Deep Field-South (CDF-S) fields, Worsley et al.\\ (2005) found that the resolved fraction in the $6-8$~keV band was only about 60\\%. (Worsley et al.\\ 2004 did a similar analysis on the {\\em XMM-Newton\\/} Lockman Hole field and found a 50\\% resolved fraction above 8~keV.) According to these authors, the missing X-ray background (XRB) component has a spectral shape consistent with a population of highly-obscured AGNs at redshifts $z\\sim 0.5-1.5$ having obscuring column densities of $N_H\\sim 10^{23}-10^{24}$~cm$^{-2}$.  Of course, to derive these results, Worsley et al.\\ (2005) had to choose which estimate of the $2-8$~keV extragalactic XRB spectrum to use, and there is a great deal of uncertainty in these measurements (see Figure~15 of Hickox \\& Markevitch 2006 for a summary; Worsley et al.\\ 2005 used the XMM-Newton measurement made by Deluca \\& Molendi 2004, but Revnivtsev et al.\\ 2005 found that the intensity of the XRB measured by focusing telescopes is higher than that measured by collimated experiments by about $10-15$\\%, possibly due to the extreme complexity of measuring effective solid angles of focusing telescopes). Indeed, Barger et al.\\ (2002) found that the issue of whether there is a need for a substantial population of as-yet undetected highly-obscured AGNs depends critically on how the low-energy and high-energy XRB measurements tie together.  Regardless, the combination of the theoretical modelling predictions for a population of highly-obscured AGNs missing from the deep X-ray surveys and the possible observational evidence for a substantial unresolved fraction of the XRB has inspired observers to try to quantify the fraction of obscured AGNs that may be being missed. Deep mid-infrared (MIR) and radio images are an obvious avenue for searching for highly-obscured AGNs, since extinction in the MIR and radio is small. Various approaches using MIR or X-ray data (e.g., Alonso-Herrero et al.\\ 2006; Polletta et al.\\ 2006), the combination of MIR and radio data (e.g., Mart\\'{\\i}nez-Sansigre et al.\\ 2005), or the combination of MIR, radio, and X-ray data (e.g., Donley et al.\\ 2005) have been adopted, and a candidate population of highly-obscured AGNs has been identified. However, the various selection effects in these approaches mean that a reliable upper limit on the number of highly-obscured AGNs that are undetected even in the $2-8$~keV band has not yet been obtained.  In this paper, we take an alternative approach to look for highly-obscured AGNs at high redshifts using a pure microJansky radio survey selection. Ultradeep radio surveys are very useful for tracing the evolution of dust-obscured galaxies and AGNs. In contrast to optical surveys, which may omit dusty sources, and submillimeter surveys, which currently have low spatial resolution and are limited to small areas for any significant depth, the 1.4~GHz surveys do not suffer from extinction, provide subarcsecond positional accuracy, and cover large areas.  Although the Jansky and milliJansky radio source populations are dominated by powerful AGNs, at the microJansky level, the radio source population is increasingly dominated by star-forming galaxies, which produce non-thermal radio continuum at 1.4~GHz through synchrotron emission from supernova remnants. In local star-forming galaxies and radio-quiet AGNs, it is well-known that the radio power is tightly correlated with the far-infrared (FIR) luminosity (e.g., Helou et al.\\ 1985; Condon et al.\\ 1991; Condon 1992), probably as a result of both being linearly related to the massive star formation rate. If we assume that this correlation continues to hold at high redshifts (note that there is some observational support that the correlation does extend to $z\\ge 1$ from the work of Appleton et al.\\ 2004, though the number of galaxies at the higher redshifts is relatively low), then we can use this relation to estimate the total FIR luminosity of a galaxy from its radio power alone. Indeed, even with the advent of the {\\em Spitzer Space Telescope\\/}, this is still the most robust way to estimate a galaxy's total FIR luminosity, due to the limited MIPS sensitivities at 70 and 160~$\\mu$m and the uncertainties in the conversion from 24~$\\mu$m luminosity to total FIR luminosity. Although estimating total FIR luminosities from 24~$\\mu$m data has been advocated (e.g., Appleton et al.\\ 2004; Marcillac et al.\\ 2006), the conversion depends on the assumed shape of the spectral energy distribution (SED), which is expected to change as a function of dust temperature and metallicity.  Using their {\\em Spitzer\\/} Infrared Nearby Galaxies Survey (SINGS) data, Dale et al.\\ (2005) found a total range in rest-frame 8~$\\mu$m (which corresponds to observed-frame 24~$\\mu$m at $z=2$) to total FIR luminosity ratios of more than a factor of 20. When they excluded the metal-poor dwarfs, the total range dropped to a factor of 10, which they argued should be taken as a minimum systematic uncertainty for total FIR luminosities inferred from rest-frame 8~$\\mu$m fluxes alone. They cautioned that if the metal abundance of the population cannot be inferred a priori, then the uncertainties may be considerably higher. Likewise, they found that the rest-frame 24~$\\mu$m to total FIR ratio spans a factor of 5 in the SINGS sample.  In addition to our assumption that the FIR-radio correlation continues to hold at high redshifts, the other possible concern with our procedure of estimating the total FIR luminosities of the radio sources using the FIR-radio correlation is the presence of radio-loud AGNs. For these sources, the radio continuum will be in excess of the level expected from star formation, and we will overestimate their total FIR luminosities. Fortunately, however, since we are only interested in obtaining upper limits on the number of highly-obscured AGNs, this is not a major concern for our present analysis, and we have conservatively not made any effort to remove such sources from our sample.  Finally, we note that even in the case of the radio-quiet AGNs, we are not able to separate out the AGN contribution to the FIR luminosity from the star-forming contribution. However, since local radio-quiet AGNs are empirically observed to follow the FIR-radio correlation, we have an upper limit on the AGN FIR emission, which is sufficient for our analysis.  In this paper, we use the ultradeep 1.4~GHz observations of the Hubble Deep Field-North (HDF-N; Richards 2000), for which we have obtained very complete spectroscopic identifications, to study the properties of the radio sample. The HDF-N has a wealth of existing multiwavelength data, including the {\\em HST\\/} ACS Treasury Great Observatories Origins Deep Survey-North data (GOODS-N; Giavalisco et al.\\ 2004) and the {\\em Spitzer Space Telescope\\/} Legacy GOODS-N data (P.I. M. Dickinson). In addition, we have recently obtained deep wide-area $J-$ and $H-$band observations using the University of Hawaii 2.2~m telescope (Trouille et al.\\ 2006, in preparation), which enable us to assign robust photometric redshifts to the spectroscopically unidentified sources.  The HDF-N is also the site of the deepest X-ray image of the sky, the 2~Ms CDF-N. We use these data to examine the issue of whether there exists an X-ray--radio correlation for star-forming galaxies. In addition, we use these data to remove the X-ray--luminous ($L_{0.5-2~{\\rm keV}}$ or $L_{2-8~{\\rm keV}}\\ge 10^{42}$~ergs~s$^{-1}$) sources from our radio-identified candidate highly-obscured AGN population for tighter upper limits. Finally, we measure the X-ray surface brightnesses for our candidate highly-obscured AGN population and others, which we compare with the shape and intensity of the XRB.  The structure of the paper is as follows. In \\S\\ref{secdata}, we describe the radio sample and the available multiwavelength imaging and spectroscopic data. In \\S\\ref{secids}, we discuss the spectroscopic identifications of the radio sample. In \\S\\ref{secphotz}, we briefly describe our photometric redshift method, which uses the optical to MIR SEDs, and we show our redshift distribution for the radio sources. In \\S\\ref{sec24um}, we discuss the $24~\\mu$m properties of the radio sample and the difficulties with using these data to estimate FIR luminosities. In \\S\\ref{secspectral}, we describe our spectral classification scheme and examine the properties of the radio sample by spectral class. In \\S\\ref{secxray}, we discuss the X-ray properties of the radio sample and investigate whether there is an X-ray--radio correlation for star-forming galaxies at these higher redshifts, like there is locally. In \\S\\ref{secopt}, we describe the optical properties of the radio sample. In \\S\\ref{secuplim}, we determine the FIR luminosities of the radio sources using our redshifts and the FIR-radio correlation. We then put quantitative upper limits on the numbers of highly-obscured AGNs that could have gone undetected in this deepest of hard X-ray surveys, and we discuss the implications for the history of supermassive black hole growth. In \\S\\ref{secxrb}, we measure the X-ray surface brightnesses of our radio-identified candidate highly-obscured AGN population, as well as of other radio and X-ray populations, and we compare them with the shape and intensity of the XRB. We summarize our results in \\S\\ref{secsummary}.  We assume $\\Omega_M=0.3$, $\\Omega_\\Lambda=0.7$, and $H_0=70$~km~s$^{-1}$~Mpc$^{-1}$. All magnitudes are in the AB magnitude system. ",
        "conclusions": "\\label{secsummary}  In this paper, we analyzed the nature and evolution of microJansky radio sources using a highly spectroscopically complete, deep VLA survey of the HDF-N region. We supplemented the spectroscopic identifications with photometric redshifts measured from the rest-frame ultraviolet to MIR spectral energy distributions. Our results are as follows.  $\\bullet$ We found that the fraction of radio sources that can be optically spectroscopically identified is fairly independent of radio flux, with about $60-80$\\% identified at all fluxes. We spectrally classified the galaxies into four spectral types: absorbers, star formers, Seyfert galaxies, and broad-line AGNs.  $\\bullet$ We did not confirm the existence of an X-ray--radio correlation for star-forming galaxies in either the hard or soft X-ray bands. Previous claims of such a correlation appear to be the result of selection effects.  $\\bullet$ We did not observe any correlation between 1.4~GHz flux and $R$ magnitude, which helps to explain the uniformity of the spectroscopically identified fraction with radio flux. Previous claims of the existence of such a correlation were likely due to the much shallower optical observations that were used in those studies.  $\\bullet$ We also did not observe any correlation between 1.4~GHz flux and redshift for the spectroscopically and photometrically identified sources. The redshift distribution is consistent with being invariant with radio flux, though we may be missing the higher redshift tail of the redshift distribution function due to optical spectroscopic bias.  $\\bullet$ The observed invariance in the rest-frame colors and only small amount of change in the absolute magnitudes of the host galaxies suggests that radio observations sample the same type of host galaxies at all redshifts and that the host galaxy properties are not evolving much with redshift.  $\\bullet$ There is dramatic evolution in the radio powers of the host galaxies with increasing redshift, and similar optical galaxies are hosting more powerful radio galaxies at higher redshifts. A quantitative description the evolution of the 1.4~GHz luminosity function can be found in Cowie et al.\\ (2004a).  $\\bullet$ Assuming that the locally determined FIR-radio correlation holds at high redshifts, we estimated total FIR luminosities for the radio sources. We note that these will be overestimates in the case of radio-loud AGNs, but we conservatively did not try to remove the radio-loud AGNs from the sample, since we are only determining upper limits. For the radio sources with total FIR luminosities comparable to quasar-like bolometric luminosities, we obtained an upper limit of 1.9 for the ratio of candidate highly-obscured AGNs to X-ray--luminous sources.  $\\bullet$ We measured the X-ray surface brightnesses for various X-ray and radio populations within a $9.5'$ radius circle to determine their contributions to the X-ray light. For the known X-ray sources in the CDF-N, we found our measurements to be (coincidentally) consistent with the HEAO1/A2 XRB measurement made by Revnivtsev et al.\\ (2005). We found that the Richards (2000) radio sample contributes half of the $4-8$~keV light, which is also how much the X-ray sources without radio counterparts contribute. However, most of the light from the radio sample comes from the X-ray--luminous radio sources. The non--X-ray--luminous and X-ray--undetected subsamples contribute only $2.3\\pm 0.4$\\% and $1.2\\pm 0.3$\\%, respectively. This tells us that the current radio source population cannot account for the background light that one might think is missing at these energies. However, the shapes of the surface brightness measurements for these two samples both appear much harder than that of the XRB, which suggests that these two populations are consistent with containing highly-obscured AGNs. For the $z\\le 2$ radio-identified ULIRG sample, we found that the X-ray--luminous subsample contributes a very substantial one-quarter of the $4-8$~keV light, while the non--X-ray--luminous and X-ray--undetected subsamples each contribute less than a percent. Again, the shapes of the latter two subsamples are considerably harder than the XRB and hence are consistent with containing highly-obscured AGNs."
    },
    "0609/astro-ph0609697_arXiv.txt": {
        "abstract": "In this paper we present the results of the past two years' observations on the galactic microquasar LS I +61 303 with the Whipple 10m gamma-ray telescope. The recent MAGIC detection of the source between 200 GeV and 4 TeV suggests that the source is periodic with very high energy (VHE) gamma-ray emission linked to its orbital cycle. The entire 50-hour data set obtained with Whipple from 2004 to 2006 was analyzed with no reliable detection resulting. The upper limits obtained in the 2005-2006 season covered several of the same epochs as the MAGIC Telescope detections, albeit with lower sensitivity. Upper limits are placed on emission during the orbital phases of 0$\\rightarrow$0.1 and 0.8$\\rightarrow$1, phases which are not included in the MAGIC data set. ",
        "introduction": "\\label{sec:1}  LS I +61 303 is a high-mass X-ray binary system located at a distance of roughly 2kpc \\cite{Ref6}. The system is comprised of a compact object, most likely a neutron star (although the presence of a black hole has not been ruled out), in a highly eccentric orbit around a rapidly rotating B0-B0.5 main sequence star \\cite{Ref7}. The most accurate determination of the orbital parameters of the system comes from optical spectroscopy measurements \\cite{Ref6}. In that paper, the authors use the period of 26.496 $\\pm$ 0.0028  days \\cite{Ref8} along with their optical measurements to derive an orbital phase of $\\phi$= 0.22 $\\pm$ 0.02 for periastron passage (see figure 1). From the high energy perspective, interest in this system began in 1978 \\cite{Ref9},  when the authors proposed the association of LS I +61 303 with the COS B gamma-ray source 2CG 135+01.  Later, it would also be associated with the EGRET unidentified source 3EG J0241+6103 \\cite{Ref4}. The identification of LS I +61 303 as a microquasar occurred in 2001 \\cite{Ref3} when relativistic, precessing radio jets were discovered extending roughly 200 AU from the center of the source.   \\begin{figure} \\begin{center} \\includegraphics[width=75mm,height=70mm]{smith_fig1.eps} \\end{center} \\caption{The elliptical orbit of the compact object in LS I +61 303 and associated radio, X-Ray, and TeV observations in terms of its orbital phase. The hashed circles represent X-ray flares, whereas the solid circles represent radio flares. Modified from \\cite{Ref11}.} \\label{fig:1}       % \\end{figure} An intriguing feature of this system is the presence of periodic modulation of outbursts in both radio and X-ray frequencies on a time scale of roughly 26 days, coincident with the orbital period of the system \\cite{Ref6}. In the case of the radio emission, the amplitude and the phase of these outbursts is also modulated, with a period of roughly 4.6 years \\cite{Ref8}. The exact mechanism for this long term modulation is unclear, but it is believed to be linked to periodic changes in the stellar wind structure of the Be star. Multi-wavelength campaigns undertaken by ASCA, ROSAT, VLA, GBI and RYLE  indicate that most hard X-ray and radio emission is localized between $\\phi$ = 0.35 and 0.65 (\\cite{Ref1}\\cite{Ref2}\\cite{Ref10}); however, radio outbursts have been detected all the way out to $\\phi \\approx$ 0.9 \\cite{Ref8}. The persistent feature of these observations is that hard X-ray outbursts tend to precede the radio outbursts by a few days. It has been suggested \\cite{Ref11} that the most likely physical mechanism for this correlation is the increased disk accretion rate (resulting in a spectral shift in X-rays to a hard state) which then fuels radio-loud jet ejection events. It is asserted that if these events take place in a region of the orbit where a large fraction of stellar UV photons is present, gamma-ray emission can take place via inverse Compton scattering of energetic jet electrons off this photon population. Recently, the MAGIC collaboration has detected VHE gamma-ray emission between $\\phi$= 0.4 and 0.7, which is consistent with the previous multi-wavelength results \\cite{Ref5}.  If we assume that the VHE emission in this system is definitely linked to accretion increases in the disk (fueling jets that scatter UV photons), we should look for regions of the orbit where jet production can take place and stellar UV photon exposure is readily available. There are two regions in the orbit in which this is believed to be possible \\cite{Ref11}, the first at periastron passage where the UV photon and stellar wind densities are the highest. The second region occurs at a point in the orbit where the decrease in orbital velocity with respect to the stellar wind compensates for the decreased stellar wind density ($\\frac{dM}{dt} \\propto \\frac{\\rho}{v^{3}}$), resulting in a second accretion rate peak ($\\approx \\phi$= 0.5 $\\rightarrow$ 0.8) \\cite{Ref14}. If, during this second accretion peak, the jet is exposed to a significant population of stellar UV photons, then it is expected that VHE gamma-ray emission should take place.  It is during and around these expected accretion peaks that we would expect to see radio, X-ray, and gamma-ray emission from LS I +61 303. This assertion is given credence by \\cite{Ref4}, in which the authors detail EGRET observations showing significant gamma-ray flux above 100 MeV both at periastron and around $\\phi$ = 0.5. However, this conflicts with the fact that the MAGIC data shows no VHE emission around periastron. Morever, no radio outbursts have ever been detected around periastron. This may indicate that there some secondary effect inhibits or obscures emission near periastron passage, but the EGRET detection contradicts this. If it is still to be maintained that the broadband emission from this system can be explained as resulting from accretion powered jets, these observational discrepancies must be addressed.  There has also been an alternate model for VHE emission from LS I +61 303 which postulates that the system is not a microquasar at all, but rather a binary pulsar system \\cite{Ref15}. In this model, the gamma-ray emission results from stellar photons upscattered by electrons originating from the rotationally powered pulsar wind. Although binary pulsar systems (such as PSR 1259-63) have been detected in TeV gamma-rays, it remains unclear whether the current observational data on LS I +61 303 can support such a hypothesis. After decades of observations on LS I +61 303, the exact nature of emission from this system remains unknown. ",
        "conclusions": ""
    },
    "0609/hep-ph0609257_arXiv.txt": {
        "abstract": "We explore the prospects for indirect detection of neutralino dark matter in supersymmetric models with an extended Higgs sector (NMSSM). We compute, for the first time, one-loop amplitudes for NMSSM neutralino pair annihilation into two photons and two gluons, and point out that extra diagrams (with respect to the MSSM), featuring a potentially light CP-odd Higgs boson exchange, can strongly enhance these radiative modes. Expected signals in neutrino telescopes due to the annihilation of relic neutralinos in the Sun and in the Earth are evaluated, as well as the prospects of detection of a neutralino annihilation signal in space-based gamma-ray, antiproton and positron search experiments, and at low-energy antideuteron searches. We find that in the low mass regime the signals from capture in the Earth are enhanced compared to the MSSM, and that NMSSM neutralinos have a remote possibility of affecting solar dynamics.  Also, antimatter experiments are an excellent probe of galactic NMSSM dark matter. We also find enhanced two photon decay modes that make the possibility of the detection of a monochromatic gamma-ray line within the NMSSM more promising than in the MSSM. ",
        "introduction": "Numerous theoretical and phenomenological motivations exist for a Minimal Supersymmetric extension of the Standard Model  (MSSM).  At the same time one of the attractive by-products of low energy supersymmetry is the natural occurrence in the particle content of the theory of a stable weakly interacting massive particle, the lightest neutralino, which could be the microscopic constituent of the as yet unobserved galactic halo dark matter. Another strong motivation comes from the SM hierarchy problem, originating from the large fine-tuning required by the stability of the electroweak scale to radiative corrections, originating from the large number of orders of magnitude occurring between the GUT, or the Planck, scale, and the electroweak scale itself.  Although very appealing, the MSSM has been challenged by various pieces of experimental information, and by some arguments of more theoretical nature. Among these, the LEP-II limit on the mass of the lightest CP-even Higgs~\\cite{Bastero-Gil:2000bw}, the constraints on the masses of supersymmetric (\\susy) charged or colored particles from direct searches at LEP and at the Tevatron \\cite{pdg}, and the so-called $\\mu$ problem, {\\em i.e.} the fundamental reason why the \\susy\\ Higgsino mass term $\\mu$ appearing in the MSSM superpotential lies at some scale near the electroweak scale rather than at some much higher scale.   The addition of a new gauge singlet chiral multiplet, $\\s$, to the particle content of the MSSM can provide an elegant solution to the mentioned $\\mu$ problem of the MSSM~\\cite{Kim:1983dt}. The so-called Next to Minimal Supersymmetric Standard Model (NMSSM)~\\cite{Ellis:1988er} is an example of one such minimal extension that also alleviates the {\\it little fine tuning problem} of the \\mssm, arising from the non-detection of a neutral CP-even Higgs at LEP-II~\\cite{Bastero-Gil:2000bw}.  A further motivation to go beyond the MSSM comes from Electro-Weak Baryogenesis (EWB), {\\em i.e.} the possibility that the baryon asymmetry of the Universe originated through electro-weak physics at the electro-weak phase-transition in the Early Universe. Although still a viable scenario within the MSSM \\cite{ewb}, EWB generically requires the Higgs mass to be in the narrow mass range above the current LEP-II limits and below $\\simeq 120$ GeV, a rather unnatural mass splitting between the right-handed and the left-handed stops (the first one required to lie below the top quark mass, and the other in the multi-TeV range), CP violation at levels sometimes at odds with electric dipole moment experimental results, and, generically, a very heavy sfermion sector \\cite{ewbmssm}. In contrast, the NMSSM provides extra triscalar Higgs couplings which hugely facilitate the occurrence of a more strongly first-order EW phase transition, and extra CP violating sources, relaxing most of the above mentioned requirements in the context of the MSSM~\\cite{ewbnmssm,Funakubo:2005pu}.   One of the chief remaining cosmological issues associated with the NMSSM,  the cosmological domain wall problem~\\cite{Abel:1995wk}, caused by the discrete $\\mathbb{Z}_3$ symmetry of the \\nmssm, can be circumvented by introducing non-renormalizable Planck-suppressed operators~\\cite{Panagiotakopoulos:1998yw}.  The Higgs sector of the \\nmssm\\ contains three CP-even and two CP-odd scalars, which are mixtures of \\mssm-like Higgses and singlets. Also, the neutralino sector contains five mass-eigenstates, instead of the four in the \\mssm, each of which has, in addition to the four \\mssm\\ components, a singlino component, the latter being the fermionic partner of the extra singlet scalars. The extended Higgs and neutralino sectors weaken the mass bounds for both the Higgs bosons and the neutralinos. Very light neutralinos and Higgs bosons, even in the few GeV range, are in fact not excluded in the \\nmssm~\\cite{Franke:1994hj} (see also~\\cite{Ellis:1988er,Franke:1995tc}; the particle spectrum with the dominant 1-loop and 2-loop corrections to the Higgs sector is available via the numerical code \\nmhdecay~\\cite{Ellwanger:2005dv}).  The cosmology of Dark Matter singlinos has been addressed long ago in Ref.~\\cite{Greene:1986th}, while the LSP relic abundance in the \\nmssm\\ was first calculated in~\\cite{Abel:1992ts}. Constraints from electroweak symmetry breaking and GUT scale universality were added in~\\cite{Stephan:1997ds}. More recently, the computer code \\micro~\\cite{Belanger:2006is}, has been extended to allow for the relic density calculation in the \\nmssm~\\cite{Belanger:2005kh}\\footnote{ The phenomenology of the lightest neutralino in a different extension of the \\mssm, the Left-Right \\susy\\ model, has been recently surveyed in \\cite{Demir:2006ef}.}. %  The implications for the direct detection of \\nmssm\\ neutralinos were first studied in~\\cite{Bednyakov:1998is}, where light neutralinos ($\\sim$ 3 GeV) with acceptable relic abundance and sufficiently large expected event rates for direct detection with a ${}^{73}$Ge-detector were found in different domains of the parameter space, when the gaugino unification relation~\\cite{Stephan:1997ds} was relaxed. In~\\cite{Cerdeno:2004xw}, the theoretical predictions for the spin-independent neutralino-proton cross section, $\\sigma_{\\lsp-p}$, were reevaluated, and all available experimental constraints from LEP on the parameter space were taken into account. Values within reach of present dark matter detectors were obtained in regions with very light Higgses, $m_{h_1^0} \\lesssim 70 \\gev$, with a significant singlet contribution. The lightest neutralino, in those regions, features a large singlino-Higgsino composition, and a a mass in the range $50 \\gev \\lesssim m_\\lsp \\lesssim 100 \\gev$. More recently, \\nmssm\\ neutralinos as light as $100 {\\rm MeV} \\lesssim m_\\lsp \\lesssim 20 \\gev$, satisfying accelerator constraints and with the right relic density, have been shown to occur in~\\cite{Gunion:2005rw}, where it was argued that the \\nmssm\\ can, moreover, provide neutralinos in the mass range that would be required to reconcile the DAMA claim of discovery with the limits placed by CDMS.  So far, theoretical studies have not addressed the possibility that \\nmssm\\ neutralinos making up the galactic dark matter can manifest themselves {\\em indirectly}. For instance, pair annihilations of neutralinos in the galactic halo can produce sizable amounts of antimatter, which current and forthcoming space-based antimatter search experiments can possibly detect; neutralinos trapped in the Sun or in the Earth~\\cite{Krauss:1985ks, Krauss:1985aa} can give rise to a coherent flux of energetic neutrinos from the center of the Sun or of the Earth; pair annihilation of neutralinos, either in nearby large-dark-matter-density sites, or from the cumulative effect of annihilations outside the Galaxy, can produce gamma-ray fluxes at a level detectable by GLAST or by ground-based air Cherenkov Telescopes; and last, but not least, the exciting possibility of peculiar gamma-ray spectral features, like a sharp monochromatic peak at $E_\\gamma\\simeq m_\\chi$ (where $\\chi$ indicates the lightest neutralino, assumed to be the lightest supersymmetric particle), from loop-induced $\\chi\\chi\\rightarrow\\gamma\\gamma$ processes, can also be, in principle, very promising.  A first motivation for looking into indirect dark matter detection within the \\nmssm\\ comes from the possibility that thermally produced neutralinos, in this context, can be {\\em very light}. Since the pair annihilation rate of thermal relics is roughly fixed by requiring that the thermal neutralino abundance coincides with the cold dark matter abundance inferred by astrophysical observations \\cite{Spergel:2006hy}, the indirect detection rates generically scale as $1/m_\\chi^2$: light neutralinos are therefore expected to give significantly enhanced rates with respect to the standard case.  A more technical point has provided us with a second motivation to look into indirect detection prospects for \\nmssm\\ neutralinos: loop induced pair annihilations of neutralinos into two photons or two gluons (respectively contributing to the mentioned monochromatic gamma-ray line and to, {\\em e.g.}, antimatter fluxes) are predicted to be increased, within the \\nmssm, by the presence of extra diagrams mediated by the (potentially light) extra CP-odd Higgs boson. We therefore extend here, for the first time, the MSSM results for these loop-induced neutralino pair-annihilation amplitudes \\cite{Bergstrom:1997fh} to the NMSSM.  Our results suggest several signatures, including muons resulting from neutralino annihilation in the Earth, and antiparticle and gamma ray production from neutralino annihilation in the galaxy, where the NMSSM produces signals that are enhanced compared to those predicted in the MSSM.  The outline of this article is as follows: we  first introduce the theoretical framework and set our notation in Sec.~\\ref{sec:nmssm}; we devote Sec.~\\ref{sec:paramspace} to a discussion of the viable \\nmssm\\ parameter space. Sec.~\\ref{sec:indirectdet} contains our central results on indirect \\nmssm\\ neutralino dark matter detection, while the appendices provide the reader with details on the relevant \\nmssm\\ neutralino pair-annihilation amplitudes and on neutralino-nucleon scattering cross sections. ",
        "conclusions": "Gauge singlet extensions of the Higgs sector of the minimal supersymmetric Standard Model provide well motivated theoretical and phenomenological laboratories. Besides offering an elegant solution to the supersymmetric $\\mu$ problem, they provide a viable way out of the difficulties connected to encompassing a mechanism of electroweak baryogenesis in the MSSM. In the present analysis, we focused on one specific such extension, the so-called NMSSM, and investigated, for the first time, the prospects for neutralino dark matter indirect detection.  The present study is motivated by two basic observations: first, in the NMSSM, unlike the MSSM, the lightest neutralino can be naturally very light, as a result of the possibility of it annihilating through a potentially very light extra CP-odd, mostly singlet-like Higgs boson; second, the extended Higgs sector leads to extra diagrams in the loop-amplitude relevant for the pair annihilation of neutralinos in photon or gluon pairs. An enhancement of the monochromatic $\\chi\\chi\\rightarrow\\gamma\\gamma$ gamma-ray line is therefore generically expected within the NMSSM, as opposed to the minimal supersymmetric setup.  We found that the rate of neutrinos produced by the annihilation of neutralino dark matter particles captured inside the Earth and the Sun is in general large in the NMSSM; unlike the MSSM, we found that most models give a larger signal from annihilations in the core of the Earth rather than in the Sun, at a level which can, in certain cases, be constrained by current available data from SuperKamiokande and MACRO.  This is presumably due in part to additional low velocity contributions to the local neutralino density in the region of the Earth which can result for light neutralinos. Future neutrino telescopes with increased sensitivity for low neutrino energies will be able to probe a sizable part of the parameter space by looking at signals from both the Earth and the Sun.   The dynamics of the Sun could also be modified due to energy transport by neutralinos. Our estimates show that, especially bino-like, neutralinos below $m_\\chi \\lesssim 30\\ \\gev$ might contribute a significant fraction of the total solar luminosity. More detailed studies, using self-consistent solar models, could unveil large enough modifications on the sound speed or on the boron neutrino flux to significantly disfavor light neutralino scenarios.    We showed that within the NMSSM the expected antimatter yield from neutralino pair annihilations in the galactic halo can be sizable, although the absolute normalization of the flux depends on specific assumptions about the dark matter halo profile. In particular, we found that signals at low-energy antideuteron search experiments such as GAPS, and at space-based antimatter search experiments such as PAMELA, are expected, though not guaranteed, for very light neutralinos ($m_\\chi\\lesssim 20$ GeV) or for intermediate mass mixed singlino-bino neutralinos ($60\\lesssim m_\\chi\\lesssim 90$ GeV).  We worked out for the first time the loop-induced pair annihilation cross section for NMSSM neutralinos into two photons and two gluons, pointing out that the expected branching ratio, with respect to tree-level neutralino pair annihilation into other Standard Model particles, is typically large, especially when compared to the MSSM case. The reason for this enhancement is traced back to diagrams which are resonant when $2 m_\\chi\\simeq m_{a_1}$, the latter quantity indicating the mass of the lightest, extra CP-odd Higgs boson.  We analyzed in detail the prospects for the detection of the monochromatic gamma-ray line resulting from $\\chi\\chi\\rightarrow\\gamma\\gamma$ annihilation processes in the Galactic Center, in a nearby dwarf spheroidal Galaxy (Draco) and the coherent effect of annihilations in any dark matter halo contributing to the extra-galactic gamma-ray radiation. We pointed out that most models are not expected to give any detectable signal at GLAST, although this detection channel looks significantly more promising than in the usual MSSM setup.  Finally, with the purpose of making the present study a useful and complete starting point for future research in the field, and in order to sort out and clarify some notational ambiguities and inconsistencies, we collect in the Appendix the details of the one-loop computation of the $\\chi\\chi\\rightarrow\\gamma\\gamma,\\ gg$ amplitudes and other quantities relevant for the estimate of indirect detection rates."
    },
    "0609/astro-ph0609047.txt": {
        "abstract": "We consider a newly-born millisecond magnetar, focusing on its interaction with the dense stellar plasma in which it is initially embedded. We argue that the confining pressure and inertia of the surrounding plasma acts to collimate the magnetar's Poynting-flux-dominated outflow into tightly beamed jets and increases its magnetic luminosity. We propose this process as an essential ingredient in the magnetar model for gamma-ray burst and asymmetric supernova central engines. We introduce the ``pulsar-in-a-cavity'' as an important model problem representing a magnetized rotating neutron star inside a collapsing star. We describe its essential properties and derive simple estimates for the evolution of the magnetic field and the resulting spin-down power. We find that the infalling stellar mantle confines the magnetosphere, enabling a gradual build-up of the toroidal magnetic field due to continuous twisting. The growing magnetic pressure eventually becomes dominant, resulting in a magnetically-driven explosion. The initial phase of the explosion is quasi-isotropic, potentially exposing a sufficient amount of material to $^{56}$Ni-producing temperatures to result in a bright supernova. However, if significant expansion of the star occurs prior to the explosion, then very little~$^{56}$Ni is produced and no supernova is expected. In either case, hoop stress subsequently collimates the magnetically-dominated outflow, leading to the formation of a magnetic tower. After the star explodes, the decrease in bounding pressure causes the magnetic outflow to become less beamed. However, episodes of late fallback can reform the beamed outflow, which may be responsible for late X-ray flares. ",
        "introduction": "\\label{sec-intro}  %**********************************************************************    In the classical collapsar scenario for long-duration gamma-ray bursts (GRBs), the core of a rotating massive star collapses to form a black hole, whereas the overlying stellar material possesses enough angular momentum to form an accretion disk that persists for at least several seconds, long enough for its jet to breakout from the star (Woosley~1993; Paczynski~1998; MacFadyen \\& Woosley~1999). This accretion disk---black hole system then acts as a central engine for the GRB. The power for the explosion comes both from accretion energy, released via neutrinos and perhaps via a magnetic mechanism (e.g., the magnetic tower mechanism as proposed by Uzdensky \\& MacFadyen 2006), and from the black-hole rotational energy, released via the Blandford--Znajek (1977) mechanism.  In the present paper we investigate an alternative scenario, in which the central object formed as a result the of core collapse is not a black hole, but rather a rapidly-rotating (millisecond) magnetar with a large-scale poloidal magnetic field of the order of $10^{15}$~G. Such a strong magnetic field can be produced, for example, by a turbulent $\\alpha-\\Omega$ dynamo driven by convection in a proto-neutron star (PNS) subject to neutrino cooling (Duncan \\& Thompson~1992; Thompson \\& Duncan~1993). An alternative possibility is that the progenitor core of about $10^4$~km has a magnetic field of order $10^9$~G, similar to the field levels actually observed in some white dwarf of similar size. When such a highly-magnetized core collapses into a neutron star of 10~km radius, flux freezing leads to amplification of the magnetic field to $10^{15}$~G, as discussed in Uzdensky \\& MacFadyen (2006). In addition, calculations by Akiyama~et~al. (2003) have shown that turbulent dynamo driven by the magneto-rotational instability (MRI) in the collapsing differentially-rotating core is capable of producing $10^{15}-10^{16}$~G fields within about 100~km from the center, on the timescale of just a few tens of milliseconds after the bounce (see also Ardeljan~et~al.\\ 2005). Whatever the origin of the very strong magnetic field in the PNS is, in this paper we shall take it for granted. Our main goal will be to investigate the role such a strong field plays in the explosion dynamics.  The idea of using a millisecond magnetar as a central engine for gamma-ray bursts has been first proposed by Usov (1992) in the context of accretion-induced collapse of a highly-magnetic ($10^9$~G) white dwarf and, independently, by Duncan \\& Thompson (1992). It has been further developed and applied to different explosion scenarios by several authors (e.g., Thompson~1994; Yi \\& Blackman 1998; Nakamura~1998; Spruit~1999; Wheeler~et~al.\\ 2000, 2002; Ruderman~et~al.~2000; Lyutikov \\& Blandford 2002; Thompson~et~al.~2004; Metzger~et~al.\\ 2007). The present paper is also devoted to investigating the millisecond-magnetar scenario, but viewed within the overall context of a collapsing star.  At the most basic level, the main idea is that a GRB (or a supernova) explosion is powered by the magnetic extraction of rotational energy of the newly-born rapidly-rotating magnetar. This magnetic luminosity operates alongside the much stronger neutrino cooling, which is the main avenue for releasing the gravitational binding energy of the young, still-contracting neutron star. However, most of the neutrinos escape to infinity without sharing their energy with the stellar-envelope gas (unless their spectrum is strongly modified by coronal processes, see Ramirez-Ruiz \\& Socrates~2005). Magnetic fields, on the other hand, couple to the gas tightly and this makes them a very efficient explosion agent. Energetically, magnetic GRB models are usually quite plausible. For example, assuming a typical surface magnetic field $B_*=10^{15}$~G, a rotation rate $\\Omega_*=10^4\\ {\\rm sec}^{-1}$, and a radius $R_*=10$~km, it is easy to see that the resulting basic energetics and timescales fall just in the right ballpark to make the millisecond magnetar a plausible candidate for a GRB central engine (e.g., Thompson~1994). Indeed, the total rotational energy of a millisecond-period neutron star is $E_{\\rm rot}\\simeq 5 \\cdot 10^{52}\\ {\\rm erg}$, which is more than enough to drive a long-duration GRB. The time scale for the energy extraction can be estimated by dividing this available energy by the total magnetic luminosity (the spin-down power), $L_{\\rm magn}$. The latter can be roughly estimated by the usual pulsar luminosity formula $L_{\\rm magn}\\sim B_*^2\\, R_*^6\\, \\Omega_*^4\\, c^{-3}$, which for the above parameters yields $\\sim 3\\cdot 10^{50}\\ {\\rm erg\\ sec}^{-1}$, corresponding to the characteristic timescale of order $100$~sec. Thus, from the point of view of the overall energetics and timescales, the millisecond-magnetar central engine is just a scaled-up version of the Ostriker \\& Gunn (1971) model for regular pulsar-powered supernovae (with the magnetic field scaled up by three orders of magnitude and the time scaled down by six orders of magnitude).  It has to be noted, however, that the plausible overall energetics and timescales are, by themselves, not sufficient for making a good GRB central engine model. This is because there are some extra physical requirements mandated by observations. In particular, to make a successful GRB, the central engine has to be capable of producing an energetic outflow that is (1) ultra-relativistic; (2) highly-collimated; and (3) baryon-free. The pioneering works cited above have focused  mostly on the energetics and timescales, but not on the mechanisms for producing an outflow that satisfies these requirements (see, however, Wheeler~et~al.\\ 2000 and Bucciantini~et~al.~2006 for a discussion of collimation). Also, most of these previous models, with the notable exceptions of Wheeler~et~al.\\ (2000) and Arons (2003), see~\\S~\\ref{subsec-const-pressure}, have considered a magnetar in isolation; that is, they have completely ignored the effect of any surrounding stellar gas on shaping the outflow. This may be a good approximation for the accretion-induced collapse of a white dwarf, but it is not appropriate in the collapsing-star scenario.  In contrast, in this paper we stress that the infalling stellar gas is still present during the explosion and needs to be taken into account. Thus, an important new element that distinguishes our model from those previous works is the consideration of the interaction between a newly-born magnetar and the stellar plasma in which it is initially embedded. Specifically, we argue that the pressure and inertia (i.e., the ram pressure) of the surrounding stellar gas acts as a natural collimator forcing the magnetized outflow into two tightly beamed jets. It also plays a crucial role in magnetic extraction of rotational energy from the magnetar.  In order to illustrate these ideas we introduce the ``Pulsar-in-a-Cavity'' problem as a basic-physics paradigm for this scenario. We describe this problem in detail in~\\S~\\ref{sec-pulsar-cavity}. We first give a general description of the problem and its various versions. Then, in \\S~\\ref {subsec-fixed-cavity}, we consider the simplest special case of a rotating force-free magnetosphere inside a fixed rigid cavity. In that section, we first demonstrate that differential rotation of the magnetic field lines is inevitably established inside the cavity, even if the pulsar itself is rotating uniformly; as a result, a strong toroidal magnetic field gradually builds up. We then study the long-term evolution of the magnetic field inside the cavity and show that the magnetic luminosity increases with time. We also show that a massive, non-force-free plasma strip unavoidably arises in the equatorial plane beyond the light cylinder. In \\S~\\ref{subsec-collimation} we discuss the subtle issue of hoop-stress collimation and argue that external confinement and differential rotation are two important ingredients for collimating relativistic Poynting-flux dominated outflows. We then consider, in \\S~\\ref{subsec-const-pressure}, the case of a magnetosphere surrounded by a cavity with a fixed external pressure (instead of a fixed radius).  In~\\S~\\ref{sec-stalled-shock} we discuss a specific example relevant to the core collapse of a massive star: a cavity formed behind the stalled bounce-shock at the center of the collapsing star. The radius of the shock stays roughly stationary on the timescale for magnetic fields in the cavity to grow. At the same time, both the ram pressure of the gas falling onto the cavity and the neutrino energy deposition inside it decrease with time. We therefore argue that at some point, a fraction of a second after bounce, the magnetic field will inevitably start to dominate the force balance, leading to a magnetically-driven explosion.   In~\\S~\\ref{sec-discussion}, we further explore some of the astrophysically-interesting aspects of our model. Thus, in \\S~\\ref{subsec-stability}, we discuss the possibility of MHD instabilities (e.g., kink) developing in the twisted magnetosphere and their implications for our model. In \\S~\\ref{subsec-nickel} we address an important issue of $^{56}$Ni production and argue that the two-phase nature of the explosion in our model is well-suited to explain a large amount of $^{56}$Ni inferred from observations. In~\\S~\\ref{subsec-restarting} we briefly discuss the possibility of restarting the GRB engine by the fall-back of the post-explosion material. In~\\S~\\ref{subsec-preshaping} we describe an extension to our model: a ``magnetar-in-a-tube'', motivated by the fact that the material along the rotation axis does not experience a centrifugal barrier and hence falls onto the PNS faster. In section~\\S~\\ref {subsec-pulsar-kicks} we discuss the implications of our model for pulsar kicks. Finally, in~\\S~\\ref{subsec-numerical}, we suggest some directions for future numerical simulations of this problem. We draw our conclusions in~\\S~\\ref{sec-conclusions}.    %******************************************************************** ",
        "conclusions": "\\label{sec-conclusions}  In this paper we have investigated the millisecond-magnetar scenario for the central engine of Gamma-Ray Bursts and core-collapse Supernovae. We have focused on the interaction between the rapidly-rotating magnetar magnetosphere and the surrounding infalling stellar envelope. We have argued that the stellar material provides a confining (ram) pressure that has a strong effect on both the size and the shape of the magnetosphere. In particular, it can channel the highly-magnetized outflow originating from the proto-neutron star into two collimated magnetic towers.  More specifically, we suggest that the stalled bounce shock --- a common feature in models of core-collapse supernovae --- effectively plays a role of a cavity that confines the magnetosphere. The cavity's radius, determined by the balance between the pressure of the hot neutrino-heated gas and the ram pressure of the infalling material, stays quasi-stationary at $R_0\\simeq 200$~km during the first few hundreds of milliseconds after the bounce. To get a feeling for what happens to the magnetar magnetosphere during this stage, we introduce a simplified fundamental-physics problem that we call the {\\it Pulsar-in-a-Cavity} problem. A large part of our paper (\\S~\\ref{sec-pulsar-cavity}) is devoted to investigating this problem. We show that since the radius of the cavity is larger than the pulsar light-cylinder radius, the magnetic field inside the cavity continuously winds up. Correspondingly, both the toroidal field strength and the magnetic spin-down luminosity of the pulsar increase roughly linearly with time. The magnetic energy in the cavity then grows quadratically with time. We then demonstrate that in the context of a millisecond magnetar inside a collapsing star the magnetic field becomes dynamically important a fraction of a second after the bounce. This leads to a subsequent revival of the stalled shock and may result in a successful magnetically-driven explosion. As long as the expansion of the cavity is non-relativistic, the toroidal magnetic field inside it remains larger than the poloidal magnetic and electric fields. As a result, the hoop-stress collimates the Poynting-flux-dominated outflow into two vertical channels that are similar to Lynden-Bell's (1996) magnetic towers (see Uzdensky \\& MacFadyen~2006).  Finally, we discuss the implications of model for several observationally-motivated questions relevant to GRBs and core-collapse supernovae, such as $^{56}$Ni production, late-time X-ray flares, and pulsar kicks. We also outline a set of numerical studies that we feel need to be done.   %-------------------------------------------------------------------"
    },
    "0609/astro-ph0609138_arXiv.txt": {
        "abstract": "We have investigated the abundances of heavy neutron-capture elements, including osmium (Os) and iridium (Ir), in the two Carbon-Enhanced Metal-Poor (CEMP) subgiants {\\CS} and {\\LP}. {\\CS} is known to be a so-called CEMP-r/s star, which exhibits large excesses of s-process elements such as barium (Ba) and lead (Pb), as well as a significant enhancement of europium (Eu) that cannot be explained by conventional s-process production in Asymptotic Giant Branch star models. Our analysis of the high-resolution spectrum for this object has determined, for the first time, the abundances of Ir and Os, elements in the third peak of the r-process nucleosynthesis. They also exhibit significant excesses relative to the predictions of standard s-process calculations. These two elements are not detected in a similar-quality spectrum of {\\LP}; the derived upper limits on their abundances are lower than the abundances in {\\CS}. We compare the observed abundance patterns of neutron-capture elements, including Os and Ir, in these two stars with recent model calculations of the s-process, and discuss possible interpretations. ",
        "introduction": "\\label{sec:intro}  Recent abundance analyses of Carbon-Enhanced Metal-Poor (CEMP) stars have provided a unique opportunity to investigate the abundance patterns produced by s-process nucleosynthesis at low metallicity \\citep[e.g., ][]{norris97,hill00}.  Of particular importance, in the past several years abundance measurements of neutron-capture elements have been extended to cover a wide range of atomic numbers, including heavy species such as lead \\citep[Pb; e.g. ][]{aoki00, vaneck01} and bismuth \\citep[Bi; e.g. ][]{ivans05}. One interesting class of objects that have recently been identified are the so-called r/s stars \\citep[e.g. ][]{hill00, cohen03, beers05, jonsell06}. These stars exhibit large enhancements of the neutron-capture elements, whose overall abundance patterns can be explained by the s-process, but whose abundance of the r-process element Eu is much higher than the value predicted by conventional s-process nucleosynthesis models \\footnote{The solar-system isotopic abundances of neutron-capture elements are decomposed using s-process models. The fraction of r-process contribution is estimated by subtracting the s-process contribution from the total abundances. The terms s-process elements and r-process elements mean the elements whose solar-system abundances are mostly explained by corresponding processes. For instance, more than 94\\% of Eu in the Solar System is yielded by the r-process, according to \\citet{arlandini99} and \\citet{simmerer04}.}. All of the presently recognized r/s stars also exhibit large carbon enrichment ([C/Fe] $\\gtrsim + 1.0$); \\citet{beers05} refer to these as CEMP-r/s stars.   The observed excesses of s-process elements in CEMP stars are usually explained by the yields of Asymptotic Giant Branch (AGB) stars (see Herwig 2005 for a recent review). Since most CEMP stars observed today are not in their AGB phase, it is usually assumed that the AGB star responsible for the s-process (and carbon enhancement) was once the primary of a binary system, and the elements formed by it were transferred by a stellar wind to a lower-mass companion (the star now observed) prior to the primary evolving to become a faint white dwarf. One possible scenario to explain CEMP-r/s stars is to assume that the parent cloud of the binary system was already enriched in the r-process elements. The r-process nucleosynthesis source(s), presumably core-collapse supernovae, would leave an elemental imprint of its products on such very low-metallicity objects, according to the above interpretation. One important observational constraint is that a significant fraction of very metal-poor s-process-enhanced stars are in fact r/s stars. For instance, \\citet{jonsell06} list 17 r/s stars among 24 s-process-enhanced objects. This stands in contrast to the fact that r-process-enhanced ([Eu/Fe]$\\gtrsim +1.0$) stars with no excess of s-process elements are quite rare among very metal-poor stars ($\\sim 5$\\%; Barklem et al. 2005). The large fraction of r/s stars among s-process-enhanced stars is not explained by assuming that these stars accidentally had large overabundances of r-process elements, but suggests that a close relationship may exist between the s- and r-processes that form the abundance patterns of r/s stars.    The interest in r/s stars has been rapidly growing in recent years. However, the classification of such objects mostly relies on the relatively easily measured Eu abundance, although the abundances of elements between Gd and Lu also provide constraints on the origin of neutron-capture elements for a small number of stars (e.g., CS~29497-030; Ivans et al. 2005). In order to better estimate the relative contributions of the s- and r-processes in r/s stars, it is important to increase the number of such stars with available abundance measurements for elements in the third peak of the r-process nucleosynthesis (Os and Ir, whose r-process fractions in solar-system material are 91.6 and 98.8\\%, respectively, according to \\citet{simmerer04}).  We have investigated high-resolution ultraviolet-blue spectra of two CEMP subgiants, {\\CS} and {\\LP}, whose abundance patterns were studied relatively well by previous work \\citep{norris97,aoki00,aoki02a,aoki02b,johnson04}. These stars have quite similar atmospheric parameters and carbon enhancements. Moreover, variations of radial velocities have been detected in both of these stars, indicating that they both belong to binary systems. In this Letter, the abundances of Os and Ir in {\\CS}, as well as upper limits on these elements in {\\LP}, are reported. The abundance patterns of these stars are then compared to the predictions of modern s-process nucleosynthesis models. ",
        "conclusions": "There is no astrophysical model that can well explain the abundance pattern of CEMP r/s stars. The model is required to explain the following two observational facts: (1) More than ten CEMP-r/s stars that exhibit very large excesses of Eu ([Eu/Fe]$>+1$) are already known, while only a few r-process-enhanced stars without excesses of s-process elements are known; (2) The CEMP-r/s stars exhibit larger enhancements of heavy neutron-capture elements, on average, than the stars having only s-process enhancements \\citep[][]{jonsell06}. To account for these constraints, a modified s-process model which yields abundance patterns that are quite different from those of standard models might be preferred, as has been discussed by previous authors \\citep[see ][ and references therein]{jonsell06}. However, the high Os and Ir abundances, as well as the high Eu abundance, found in {\\CS} appear to exclude this possibility, because no s-process model known to date predicts such high yields of these elements. We confirmed that Os and Ir are not efficiently produced by our AGB models, even if model parameters (initial mass, efficiency of mixing to produce the neutron source $^{13}$C) are significantly changed. If a single astrophysical process capable of producing the observed abundance patterns of r/s stars exists, it must be significantly different than s-process production by AGB stars known so far.    As mentioned in the previous section, the abundance pattern of neutron-capture elements in {\\CS} is better accounted for by a model that assumes an initial composition for the binary system with large excesses of heavy elements having a scaled solar r-process abundance pattern.  Thus, we argue that this system was formed from a parent cloud polluted by a supernova that yielded r-process elements, which might have also triggered the formation of next-generation low-mass binary stars \\citep{vanhala98}, while the s-process elements were provided by the former primary star during its AGB phase. A possible alternative scenario is to assume that pollution of {\\CS} with r-process-enhanced material took place after s-process nucleosynthesis. This might be the case if the progenitor was an 8--10~M$_{\\odot}$ star that underwent s-process nucleosynthesis during its AGB phase, followed by the production of r-process elements during its subsequent supernova explosion \\citep{wanajo06}. However, it is not yet clear how much of the s-process material may be mixed within the envelope of these stars (Doherty 2006; Siess 2006). These two scenarios are, unfortunately, not easily distinguished by the elemental abundance pattern of the low-mass stars that are currently observed, because the majority of the r-process nuclei for which excesses are assumed for the initial composition are not subject to the s-process. Hence, the r- and s-processes behave as almost independent contributors to the final yields.   Finally, we point out that, although {\\CS} and {\\LP} have similarly large excesses of carbon and neutron-capture elements, their elemental abundance patterns exhibit non-negligible differences.  {\\CS} exhibits a large excess of Pb, as expected from s-process models \\citep{aoki02b,johnson04}, while {\\LP} has a smaller enhancement of Pb \\citep{aoki00,aoki02a}. The [La/Eu] ratios show some difference, as mentioned in \\S 3. In addition, the sodium (Na) abundance is significantly different between the two stars \\citep{aoki02a,aoki06}. These differences suggest that the process(es) responsible for the production of these elements may not be the same in these two stars. More detailed comparisons of the elemental abundances for these two stars are strongly desired."
    },
    "0609/astro-ph0609412_arXiv.txt": {
        "abstract": "A comparison of X-ray observations of the Cassiopeia A supernova remnant taken in 2000, 2002, and 2004 with the {\\it Chandra} ACIS-S3 reveals the presence of several small scale features ($\\leq 10\\arcsec$) which exhibit significant intensity changes over a 4 year time frame. Here we report on the variability of six features, four of which show count rate increases from $\\sim$ 10\\% to over 90\\%, and two which show decreases of $\\sim$ 30\\% -- 40\\%. While extracted 1-4.5 keV X-ray spectra do not reveal gross changes in emission line strengths, spectral fits using non-equilibrium ionization, metal-rich plasma models indicate increased or decreased electron temperatures for features showing increasing or decreasing count rates, respectively. Based on the observed count rate changes and the assumption that the freely expanding ejecta has a velocity of $\\sim$ 5000 km s$^{-1}$ at the reverse shock front, we estimate the unshocked ejecta to have spatial scale variations of $0.02 - 0.03$ pc, which is consistent with the X-ray emitting ejecta belonging to a more diffuse component of the supernova ejecta than that seen in the optically emitting ejecta, which have spatial scales $\\sim$ 10$^{-3}$ pc. ",
        "introduction": "Cassiopeia A (Cas A) is currently the youngest known Galactic supernova remnant (SNR) with an estimated explosion date no earlier than around 1670 AD \\citep{thorstensen01,Fesen06a}.  The remnant consists of optical, infrared, and X-ray emission arranged in a shell roughly 4$\\arcmin$ diameter ($\\simeq$ 4 pc at 3.4$_{-0.1}^{+0.3}$ kpc; \\citealt{reed95}) consisting of undiluted supernova (SN) ejecta rich in O, Si, Ar, Ca, and Fe heated by the remnant's  $\\simeq$ 3000 km s$^{-1}$ reverse shock \\citep{chevalier78,chevalier79,douvion99, hughes00,willingale03,hwang03,laming03,morse04}.  Viewed in X-rays, the remnant's brightest emission is concentrated in a 30$ \\arcsec$ thick, $\\simeq$ 110$\\arcsec$ radius emission ring dominated by thermal plasma arising from shock heated SN debris. Farther out beyond this emission ring, lies fainter filamentary X-ray emission (radius $\\simeq$ 150$\\arcsec$) associated with the current position of the remnant's $\\approx $ 6000 km s$^{-1}$ forward blast wave as it moves through the surrounding circumstellar material \\citep{gotthelf01,delaney03}.  If the remnant's expanding SN ejecta were uniform in density and chemical composition, X-ray emission from the shocked ejecta would appear as a uniform ring of emission with little spectral variation. What is actually observed, however, is a patchy, irregular ring of clumpy emission reflecting the inhomogeneous nature of the expanding SN debris as it interacts with the reverse shock front. Due both to strong ejecta clumping caused by radiative cooling instabilities and the turbulent mixing of different chemical layers of the progenitor star during the SN explosion, the resulting X-ray emission morphology is structurally and chemically complex \\citep {hughes00, hwang00}.  Although Cas A's reverse shock front has been located to arcsecond accuracy in a few regions through high-resolution optical images \\citep{Fesenetal01, morse04}, it has only been broadly localized through a sharp increase in X-ray emission in azimuthal averages of selected portions of the remnant \\citep{gotthelf01}. On small spatial scales, evidence for the reverse shock interacting with the ejecta would be manifested in X-rays as a rise in emission as the SN ejecta is shocked and decelerated. Post-shock cooling of the metal-rich ejecta could lead to equally dramatic localized emission declines.  Here we present an analysis of archival {\\sl Chandra} images of Cas~A which show a number of substantial, small-scale brightness changes between 2000 and 2004 which we interpret as marking regions of reverse shock passage. We report on spectral analyses for six specific regions and find some plasma temperature changes in the X-ray emitting ejecta. In $\\S$2 we present the relevant {\\it Chandra} observations and subsequent data reduction. In $\\S$3 we discuss the results from our analysis, and summarize our conclusions in $\\S$4. ",
        "conclusions": "With the exception of the rapidly evolving remnant associated with  SN~1987A \\citep{Burrows00,Park04}, no remnant has been previously reported to  show significant flux changes across portions of its X-ray emitting  structure. This is due in part to the limited number of high-resolution images available on young SNRs covering a significant time span, and the relatively small number of very young Galactic or LMC/SMC SNRs known.  Inspection of X-ray data taken of the Cassiopeia A supernova remnant  with {\\it Chandra} ACIS-S3 in 2000, 2002, and 2004 reveals the presence of  several small scale features ($\\leq 10\\arcsec$) which exhibit significant intensity  changes over this 4 year time frame. Here we described six such features four  of which had count rate increases from $\\sim$ 10\\% to over 90\\%, while two  others showed decreases of $\\sim$ 30\\% -- 40\\%. Whereas the majority of the six variable flux features exhibited no gross spectral changes in the 1--4 keV energy band during the four year period, one region which had the greatest increase in brightness also showed hints of slight emission line changes as well as an increase in its 6 cm radio emission. Spectral fits using non-equilibrium ionization, metal-rich plasma models indicate increased or decreased electron temperatures for features showing increasing or decreasing count rates, respectively.  A fading emission filament  projected near the remnant's interior (Region 5) showed a completely nonthermal spectrum  and thus may be associated with the forward shock front. It decreasing  intensity may be the result of changing path lengths along the edges of a wavy forward  shock as it interacts with the local, inhomogeneous circumstellar medium.  Based  on the observed count rate changes and the assumption of freely expanding 5000 km s$^ {-1}$ ejecta at the reverse shock front, we estimate the  unshocked ejecta to have spatial scale variations of $0.02 - 0.03$ pc,  consistent with the X-ray emitting ejecta belonging to a more diffuse component of the supernova ejecta.  In general, bright X-ray and optical knots are not coincident. This is not unexpected given the differences in their temperatures and densities. Moreover, the sizes of the X-ray knots ($\\sim1''-5''$) appear to be an order of magnitude larger than the optical knots ($\\sim0 \\farcs2-0 \\farcs3$; \\citealt{Fesenetal01}).  However, the remnant's optical knots can also appear and brighten on timescales as short as a few years, not unlike the timescales we observed for these X-ray knots.  Given the newly observed short term variability seen here in the X-ray, future X-ray and optical observations of selected knots and filaments might provide valuable information regarding the dynamical and radiative evolution of supernova ejecta knots over a range of densities, sizes, and postshock temperatures.  Cassiopeia A is currently the youngest and one of the most studied Galactic SNRs. The small yet noticeable changes we found combined with future X-ray observations of this bright X-ray remnant may help shed light not only on the remnant's near-term X-ray evolution, but also help us better understand the fine-scale structure, dynamics and chemistry of debris from core-collapse supernovae in general."
    },
    "0609/astro-ph0609248_arXiv.txt": {
        "abstract": "Narrow dust rings observed around some young stars (e.g., HR 4796A) need to be confined. We present a possible explanation for the formation and confinement of such rings in optically thin circumstellar disks, without invoking shepherding planets. If an enhancement of dust grains (e.g., due to a catastrophic collision) occurs somewhere in the disk, photoelectric emission from the grains can heat the gas to temperatures well above that of the dust. The gas orbits with super(sub)-Keplerian speeds inward (outward) of the associated pressure maximum. This tends to concentrate the grains into a narrow region. The rise in dust density leads to further heating and a stronger concentration of grains. A narrow dust ring forms as a result of this instability. We show that this mechanism not only operates around early-type stars that have high UV fluxes, but also around stars with spectral types as late as K. This implies that this process is generic and may have occurred during the lifetime of each circumstellar disk. We examine the stringent upper-limit on the H$_2$ column density in the HR 4796A disk and find it to be compatible with the presence of a significant amount of hydrogen gas in the disk. We also compute the OI and CII infrared line fluxes expected from various debris disks and show that these will be easily detectable by the upcoming Herschel mission. Herschel will be instrumental in detecting and characterizing gas in these disks. ",
        "introduction": "\\label{sec:intro}  Observations using infrared space telescopes, such as the {\\it Infrared Space Observatory} ({\\it ISO}), the {\\it Infrared Astronomical Satellite} ({\\it IRAS}) and recently the {\\it Spitzer Space Telescope}, have demonstrated that many stars exhibit the {\\it``Vega-phenomenon,''} --- excess infrared radiation by orbiting dust grains heated by the central star \\citep{Zuckerman}. These grains are believed to lie in optically-thin, gas-poor disks that are likely the remains of more massive protoplanetary disks \\citep{Backman,SE}. Such disks are termed `debris disks' because the grains have short life-times and must be continuously regenerated by larger bodies.  Famous examples of this phenomenon include Vega \\citep{Aumann84} and $\\beta$ Pic \\citep[][and references therein]{Aumann85}. \\citet{MeyerPPV} presents the most recent review of debris disk observations.  When spatially resolved, many of these disks display peculiar morphologies, including rings, holes, clumps and warps. Properties of some well-known debris disks and their host stars are summarized in Table~\\ref{tbl-1}. Ring-like features appear to be prevalent. In this article, we focus on the phenomenology of narrow rings, with the two most prominent examples being disks around the two A-type stars HR 4796A \\citep{RayJay,Koerner98,Schneider} and Fomalhaut \\citep{Holland,kalas2,Marsh}. In the former case, a dust ring is seen in scattered light by HST at a distance of 70 AU away from the star with an un-resolved width $< 17$ AU \\citep{Schneider}. The dust disk in Fomalhaut was resolved by HST to be distributed in a narrow belt (width $\\sim 25$ AU) at a distance of 135 AU from the star \\citep{kalas2}.  Narrow dust rings also appear around stars of later spectral types: the dust disks of two F-type stars, HD 139664 (F5) and HD 181327 (F5/F6) were recently resolved by HST as narrow belts, the former appears as a ring of width $\\sim 26$ AU centered $\\sim 90$ AU away from the star \\citep{Kalas2006}, while the latter a width $36$ AU at $\\sim 86$ AU away \\citep{Schneider2006}.  In addition, model fitting using Spitzer spectrophotometry data indicate that dust around the solar-type star HD 105 may lie in a narrow belt (width $< 4$ AU) located $\\sim42$ AU away from the star \\citep{Meyer}.  While HR 4796A, HD 181327 and HD 105 are relatively young (a few Myrs), the ages for Fomalhaut and HD 139664 are estimated to be $\\sim 10^8$ years.  Barring issues related to observational sensitivity limits, the presence of a clear outer edge for these rings is puzzling.  The surface brightness of debris disks is likely dominated by grains just above the radiative blow-out size ($\\beta=1/2$, see eq. [\\ref{eq6}]). Even if these grains are produced by parent bodies on circular orbits that lie within a ring, radiation pressure launches them to eccentric orbits upon creation (with $e = \\beta/(1-\\beta)$). The narrow ring geometry is destroyed within a dynamical time and one observes a greatly broadened disk with no sharp outer edge \\citep[][also see \\S \\ref{subsec:eccentric}]{augereau,Thebault,strubbe,Krivov2006,Wyatt2006}. This likely explains the distribution of grain populations in AU Mic, $\\beta$ Pic and Vega: small grains are observed in a broad disk, whereas larger grains (detected in thermal infrared in the cases of $\\beta$ Pic and Vega) are observed in ring-like geometry. So why do some systems appear ring-like in scattered-light images?    \\begin{deluxetable*}{ccccccccccc} \\tabletypesize{\\scriptsize} \\tablecaption{Properties of some spatially resolved debris disks} \\tablewidth{0pt} \\tablehead{ \\colhead{Star} & \\colhead{Spectral} & \\colhead{$L_\\ast$} & \\colhead{Age} & \\colhead{$\\chi$\\tablenotemark{a}} & \\colhead{$M_{\\rm dust}$\\tablenotemark{b}} & \\colhead{$L_{\\rm IR}$} & \\colhead{Peak} & \\colhead{Ring Width} & \\colhead{Refs.} \\\\ \\colhead{} &\\colhead{Type} & \\colhead{($L_\\odot$)} & \\colhead{(Myr)} & \\colhead{} & \\colhead{($M_\\oplus$)} & \\colhead{$(L_\\ast)$} & \\colhead{emission (AU)} & \\colhead{(AU)} & \\colhead{} } \\startdata HR 4796A & A0 & 21 & 8 $\\pm$2 & 538 & $\\ge$0.25 & $5\\times 10^{-3}$ & $70$ & $< 17$ & 1,2,3,4 \\\\ Vega & A0 & 60 & $350$ & \\nodata & 0.003 & $2\\times 10^{-5}$ & $> 86$ & \\nodata & 5,6,7,8 \\\\ Fomalhaut & A3 & 13 & 100-300 & \\nodata & 0.017 & $4.6\\times 10^{-5}$ & $135$ & $25$ & 6,9,12,13 \\\\ $\\beta$ Pic & A5 & 8.7 & 12$^{+8}_{-4}$ & 0.17 & $0.04$ & $3\\times 10^{-3}$ & $120$ & {$90$\\tablenotemark{c}} & 10,12,15 \\\\ HD 139664 & F5 & 3.3 & $300^{+700}_{-200}$ & \\nodata & $> 8.7\\times10^{-4}$ & $10^{-5}$ & $\\sim 90$ & 26 & 16 \\\\ HD 181327 & F5/F6 & 3.1  & $\\sim10$ & \\nodata & 0.05 & $2\\times 10^{-3}$ & 86 & 36 & 17,18 \\\\ HD 105 & G0 & 1.3 & $30\\pm 10$ & \\nodata  & $\\sim 0.03$ &  $3.9\\times 10^{-4}$ & 42 & 4 & 19\\\\ HD 107146 & G2 & 1.1 & 80-200 & $3.2\\times10^{-6}$& 0.1-0.4 & $1.2\\times 10^{-3}$ & $130$ & $85$ &5,11 \\\\ HD 53143 & K1 & 0.7 & $1000\\pm200$ & \\nodata & $1.2\\times 10^{-4}$ & $> 1.6\\times10^{-5}$ & $\\sim55-110$ & $>55$ & 7,16\\\\ $\\epsilon$ Eridani & K2 & 0.33 & $850$ & \\nodata & 0.002 &$8\\times 10^{-5}$ & $65$ & $50$ & 12,14 \\\\ \\enddata \\tablenotetext{a}{The stellar FUV ($11.2$ - $12.42$ eV) photon flux, normalized by the average interstellar value of $1.2 \\times 10^7/\\cm^2/\\s$ \\citep{habing}, and evaluated at 70 AU from the star.  The high resolution spectra are provided by Kamp \\& Hauschildt \\citep{Hauschildt}, and the flux calibration is performed by \\citet{Fernandez}.} \\tablenotetext{b}{Dust masses are estimated from sub-millimeter fluxes assuming optically thin radiation, except for HD 139664 and HD 53143 for which the dust masses were calculated from the optical depth as described in [16].} \\tablenotetext{c}{We adopt here the radial profile inversion by Artymowicz (private communication) based on observations from \\citet{heap}.} \\tablerefs{ (1) \\citet{Jura98}; (2) \\citet{Stauffer}; (3) \\citet{Greaves2}; (4) \\citet{Schneider}; (5) \\citet{Williams}; (6) \\citet{Dent}; (7) \\citet{Song}; (8) \\citet{Su}; (9) \\citet{Barrado}; (10) \\citet{Zuck2001}; (11) \\citet{Ardila}; (12) \\citet{DiFolco}; (13) \\citet{kalas2}; (14) \\citet{greaves05}; (15) \\citet{heap}; (16) \\citet{Kalas2006}; (17) \\citet{Schneider2006}; (18) \\citet{Mannings}; (19) \\citet{Meyer}.} \\label{tbl-1} \\end{deluxetable*}  The most popular explanation for ring confinement invokes one or more massive shepherding planets orbiting at distances comparable to that of the ring.  As grains spiral inward due to Poynting-Robertson (P-R) drag, the gravitational tug from the planets may temporarily stall the grains at special radii (mean-motion resonances, MMR), giving rise to an apparent narrow ring \\citep[see, e.g.][]{Ozernoy, Wilner, Quillen, kuchner,Moran}. The outward migration of a large planet may also trap grains in MMRs, without the need for P-R drag \\citep[see, e.g.][]{Wyatt2003,Wyatt2006}. This process is similar to the resonance capture of Kuiper Belt objects by Neptune. Alternatively, gravitational scattering by massive planets may prevent grains from migrating across the planet's orbit, leading to a dust distribution with a sharp inner edge \\citep{moromartin}.  However, there are problems with the planet hypothesis, even if we accept the existence of giant planets at such large radii.  Firstly, grain-grain collisions typicaly induce radial diffusion at a rate orders of magnitude above that from P-R drag, thus weakening the capture in MMRs. The sharp edges formed by planet scattering may also be smeared out by collisions.  Secondly, the ability of planets to confine grains that are launched to elliptic orbits may be limited. \\citet{Moran} noted that at eccentricities larger than that of the planet's orbit, grains spiralling inwards due to P-R drag can no longer be trapped in MMRs. \\citet{Wyatt2006} further illustrated that even grains generated by a resonant parent population will become radially spread out and azimuthally axisymmetric, so long as they are small enough to experience significant radiation pressure ($\\beta \\sim 1/2$). Further studies taking these effects into account are clearly warranted.  Interestingly, \\citet{wyatt} argued that the process of grain-grain collisions can itself act as a potential confining mechanism for the narrow ring. For most observed debris disks, small dust grains collide destructively on timescales much shorter than that of P-R drag. Consequently, dust generated at large distances from the central star cannot penetrate the collisional barrier to reach the inner region of the system. This produces a dust ring with a sharply defined inner-edge. If this dust ring is also optically thick in the radial direction, so that radiation pressure cannot  disperse the disk (see above), it is possible for the ring to remain narrow.  \\citet{KB} introduced an alternative theory for ring formation. In a planetesimal disk, runaway growth occurs at successively later times further away from the star \\citep[for an overview see][]{NagasawaPPV}. \\citet{KB} suggested that these newly formed protoplanets gravitationally stir the left-over planetesimals, inducing violent collisions that erode the planetesimal disk.  As such, at any given time the disk will appear as a broad ring with a dark inner hole (or gap). The observed rings herald the locations of recent runaway growth and the dark holes where growth has completed, or shadows where planets have yet to form \\citep{KB127}. The dust masses and luminosities they obtained are comparable to those in the observed debris disks \\citep[see also][]{KB602}. However, if these disks are optically thin in the radial direction, the ring should not have a well-defined outer edge since the freshly produced grains will spread out in a dynamical timescale.   \\citet{krauss} recently considered a novel effect called photophoresis. They showed that in optically thin circumstellar disks with gas densities comparable to that in a minimum mass solar nebula (MMSN), grains can be pushed outward to a critical radius where the gas `photophoretic' force equals the gas drag. The width of the resulting dust ring is, as of yet, unexplored in this model. Even so, this model faces the following two difficulties: (1) For the critical radius to be at 70 AU, the gas density needs to be twice that in the MMSN. This creates a problem -- there can be no photophoresis effect once the gas has blocked out the direct starlight. At this density, the optical depth caused by hydrogen Rayleigh scattering alone is $\\tau_{\\rm Ray} = 0.4$ if integrated from 10 AU to 70 AU, and $24$ if integrated from 1 AU to 70 AU \\citep[Rayleigh opacity $\\kappa \\sim 10^{-3} \\cm^2/\\g$ at 4000 \\AA,][]{mayer}. This model thus requires the gas disk to be truncated inward of $\\sim 10$ AU. (2) The assumption of twice the MMSN density for the observed debris disks has been ruled out by the observed upper limits on the column densities of multiple gas species (not just hydrogen) in these disks \\citep{Greaves2,Chen}. If we assume a gas density of $2\\times 10^{-5}$ MMSN instead (like that in our model), and that grains and gas are tightly coupled, the critical radius for photophoresis should be at $\\sim 0.2 \\AU$ (eq. [13] of \\citet{krauss}). This radius is even smaller if grains and gas are not tightly coupled. For these two reasons, we believe that photophoresis cannot explain rings in debris disks.  Along a different line, \\citet[][hereafter referred to as {\\bf TA01}]{TA} suggested that rings arise naturally in disks passing through the transitional phase from gas-dominated (proto-planetary) disks to dust-dominated (debris) disks. Residual gas in these disks can be rotating with either super- or sub-Keplerian speeds depending on the local radial pressure gradient, whereas dust grains have sub-Keplerian speeds that depend on grain size. The grains therefore feel either a head-wind or a tail-wind, forcing them to migrate radially until they are in co-rotation with the gas. TA01 showed that this effect causes a local concentration of grains; in particular, when the gas pressure experiences a sharp drop-off, a narrow dust ring forms locally. In their work this pressure drop-off is provided by an artificial cut-off in gas density over a radial distance of order the disk vertical scale height.  In this study, we propose an extension to TA01's work. We lift the undesirable assumption of an artificially-truncated gas disk and show that narrow rings form under more general initial conditions. We assume that the residual gas is smoothly distributed within the disk, but that the dust is locally enhanced. The enhancement can occur over a distance of tens of AU, e.g., as the result of a major collisional event in the underlying planetesimal disk (\\S \\ref{sec:observemass}). We relax the conventional assumption of temperature equality between the gas and dust components in the disk and show that the gas heats up due to the dust enhancement. We demonstrate that this gas response leads to the confinement of the dust grains. The mechanism discussed in this paper is an instability: a slight confinement of the dust heats the gas further which leads to a stronger confinement.    \\citet{Klahr} (hereafter KL) have proposed a similar mechanism for the creation of narrow rings. Their analytical work assumes a simple relationship between the gas temperature and dust density, $T_{\\rm gas} \\propto \\rho_{\\rm dust}^\\beta$ with $\\beta > 0$. Our thermal analysis validates this assumption for a large range of gas/dust masses and stellar spectral types. While KL study grains of a single size and argue that narrow rings will arise from infinitesimally small perturbations of a smooth dust distribution, we study a continuous spectrum of grain sizes and a large perturbation (a dust belt). KL also speculate that the dust ring ``freezes'' after it has formed, despite the disappearance of the gas as the disk evolves. This speculation needs to be assessed by taking into account radiation pressure and frequent grain collisions (see above).        In the following, we discuss the observations related to dust and gas masses in known debris disk systems (\\S \\ref{sec:observemass}); these motivate our model for the transitional disk (\\S \\ref{sec:fiducial}). In \\S \\ref{sec:real}, we present calculations of gas heating/cooling rates and temperatures in such a disk. Narrow dust rings are shown to arise in our fiducial model, as well as in a wide range of parameter space (\\S \\ref{sec:instability}). We discuss the best tracers for gas in these disks (\\S \\ref{subsec:howto}), as well as major complications to our model (\\S \\ref{sec:caveats}). We conclude in \\S \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We have studied the thermal and dynamical consequences of an optically thin dust disk embedded in a gas disk of comparable mass. We showed that photoelectric emission from the dust grains heats up the gas and this modifies the grains' orbital motion: if the dust is produced in a belt-like region (like that of our Kuiper belt), the heated gas will collect dust particles into a narrower region with sharp edges. This resulting dust distribution may be associated with the dust rings observed in various debris disks, such as that around HR 4796A, and possibly HD 105.   This mechanism operates around stars with a spectral type at late as K2 and is valid for gas and dust masses over $\\sim 2$ orders of magnitude in range. So, if every circumstellar disk goes through a transitional stage as we have described here, a narrow dust ring is expected to form. However, it is likely that the ring will disperse as the gas dissipates.   We have restricted ourselves to consider only axi-symmetric rings; an interesting extension of this work will be to explore whether non-axisymmetric features may also derive from gas-dust interaction. We have also considered the impact of eccentricities induced by radiation pressure on small grains that are released from larger bodies on circular orbits. This process complicates our theory, but does not diminish the instability.  We conclude that the presence of narrow dust rings in optically thin, transitional-stage disks, such as that observed around HR 4796A, does not necessarily imply the presence of ``shepherding'' planets. Instead, we argue that the detection of narrow dust rings may indicate the presence of gas. We have shown that this conclusion is compatible with the stringent upper limit on the H$_2$ column density in the HR 4796A disk.  Lastly, most of the gas cooling occurs via the OI or CII infrared lines. We have computed the expected fluxes for various systems and point out that the upcoming Herschel mission will be instrumental in observing and characterizing the gas."
    },
    "0609/astro-ph0609762_arXiv.txt": {
        "abstract": "We combine the results from several {\\it HST} investigations of the central structure of early-type galaxies to generate a large sample of parameterized surface photometry. The studies selected for inclusion were those that used the ``Nuker law\" to characterize the inner light distributions of the galaxies. The sample comprises WFPC1 and WFPC2 $V$ band observations published earlier by our group, $R$ band WFPC2 photometry of Rest et al., NICMOS $H$ band photometry by Ravindranath et al.\\ and Quillen et al., and the brightest cluster galaxy WFPC2 $I$ band photometry of Laine et al.  All parameters are transformed to the $V$ band and a common distance scale. The distribution of the logarithmic slopes of the central brightness profiles strongly affirms that the central structure of elliptical galaxies with $M_V<-19$ is bimodal, based on both parametric and non-parametric analysis. At the {\\it HST} resolution limit, most galaxies are either power-law systems, which have steep cusps in surface brightness, or core systems, which have shallow cusps interior to a steeper envelope brightness distribution. A rapid transition between the two forms occurs over the luminosity range $-22<M_V<-20,$ with cores dominating at the highest luminosities, and power-laws at the lowest.  There are a few ``intermediate'' systems that have both cusp slopes and total luminosities that fall within the core/power-law transition, but they are rare and do not fill in the overall bimodal distribution of cusp slopes.  These results are inconsistent with the Ferrarese et al. Virgo Cluster Survey (VCS) analysis, which did not find a bimodal distribution of cusp slopes. While the VCS ACS/WFC images have lower angular resolution than the WFPC2/PC F555W images used for much of the sample, the basic galaxy profiles measured with either camera agree well after deconvolution. However, using galaxies common to the present and VCS samples, we demonstrate that the VCS models of the cusps are either a poor match to the observations or consist of forms fitted to the galaxy envelopes and extrapolated inward to the {\\it HST} resolution limit. ",
        "introduction": "Long before massive black holes were accepted to be present in the centers of galaxies, \\citet{bbr} argued that a binary black hole created in the merger of two galaxies would eject stars from the center of the newly created system as the binary slowly hardened. N-body simulations support this hypothesis \\citep{ebi, mak97, mnm}. We now believe that nearly every elliptical galaxy or spiral bulge has a black hole at its center \\citep{mag}. We also believe that the most massive elliptical galaxies were formed by merging pre-existing galaxies.  It is possible that their central structure still bears witness to such events. The ``cores'' seen at the centers of the most luminous elliptical galaxies may indeed be the signatures of gravitational stirring and heating (``core scouring'') by binary black holes \\citep{f97}.  Cores were initially seen in ground-based observations of nearby luminous elliptical galaxies as a central region of nearly constant surface brightness \\citep{l85, k85}. The cores were seen to be non-isothermal, but their true form as $r\\rightarrow0$ was unknown. {\\it HST} images, however, later showed that nearly all galaxies have singular starlight distributions in the sense that surface brightness diverges as $\\Sigma(r)\\sim r^{-\\gamma}$ \\citep{l91, l92a, l92b, crane, k94, f94, l95}. Typically, in low luminosity early-type galaxies, $\\gamma$ decreases only slowly as the center is approached and a steep $\\gamma>0.5$ cusp continues into the {\\it HST} resolution limit; \\citet{l95} classified these systems as ``power-law'' galaxies.  In more luminous galaxies, however, the steep envelope profile transitions to a shallow inner cusp with $\\gamma<0.3$ at a ``break radius,'' $r_b.$ This behavior was seen in galaxies that had cores from ground-based observations; the break radius roughly corresponded to what had been measured as the core radius. We thus chose to continue to call these ``core galaxies,'' even though the shallow cusps in projected brightness in most\\footnote{\\citet{l02} and \\citet{l05} identified a number of galaxies in which the brightness profiles actually decreased as $r\\rightarrow0.$  These systems thus have $\\gamma<0$ at small radii and do not have cusps in space density.} of these systems imply steep and singular cusps in luminosity density \\citep{l95}.  A strong justification of this classification schema was that the distribution of $\\gamma'$ over a sample of early-type galaxies was seen to be strongly bimodal \\citep{f97}, where $\\gamma'$ is the local cusp slope at the {\\it HST} resolution limit. Initially no galaxies were seen to have $0.3<\\gamma'<0.5.$ It was thus sensible to apply the separate core and power-law classifications to the two distinct and cleanly separated peaks in the $\\gamma'$ distribution. The bimodal distribution was initially inferred from fits to the surface photometry using the ``Nuker law'' (\\citealt{l95}; see equation \\ref{eqn:nuker}), which has the asymptotic form $\\Sigma(r)\\sim r^{-\\gamma},$ as $r\\rightarrow0,$ however, \\citet{g96} demonstrated that non-parametric treatment of the surface photometry recovered an identical bimodal distribution, and further that bimodality was a feature of the luminosity {\\it density} distributions of early type galaxies, not just the surface brightness distributions.  The core and power-law classifications are fundamental because they correlate with global physical properties that are diagnostic of galaxy formation. Core galaxies are on average more luminous than power-law galaxies, with power-laws dominating below $M_V\\sim-21$ and cores greatly dominating at higher luminosities \\citep{f97}. Core galaxies generally have boxy isophotes, while power-law galaxies are disky \\citep{n91, k96, f97}. Core galaxies rotate slowly, while power-law galaxies rotate more rapidly \\citep{f97}; similarly, boxy galaxies include slow rotators while disky galaxies rotate rapidly \\citep{b88, b89, b94, k96}. Core-boxy galaxies show minor-axis rotation indicative of triaxiality, while coreless-disky galaxies do not \\citep{k96}. And core galaxies are systematically rounder than power-law galaxies \\citep{jaffe, f94, ryden2, l05}, just as bright (boxy) galaxies are systematically rounder than faint (disky) galaxies \\citep{tm}. Significantly, core and power-law galaxies of identical luminosity can still be separated by all of these secondary characteristics \\citep{f97}. Thus central and global properties together point to differences in the formation of core and power-law ellipticals. \\citet{k96} propose a revision of the Hubble sequence for elliptical galaxies that recognizes this physical dichotomy, using isophote shape as a convenient surrogate for the above physical properties.  More recent work continues to show that power-law and core galaxies have physically distinct properties.  Cores \\citep{cap, bal}, like boxy isophotes \\citep{b89}, correlate with radio-loud active nuclei. And cores \\citep{pell}, like boxy isophotes \\citep{b89}, correlate with strong x-ray emission.  The properties of the power-law versus core galaxies suggest that core systems resulted from the mergers of less luminous power-law systems. Simulations by \\citet{mnm}, for example, showed that the merger of two power-law galaxies, each with a central massive black hole, would produce a core galaxy. The bimodal cusp distribution would thus be one key to understanding the central structure of elliptical galaxies over a history of mergers.  As larger samples of elliptical galaxies were observed with {\\it HST}, the bimodal slope distribution remained robust \\citep{quil, rest, rav, l05}. While all of these studies identified a small number of ``intermediate'' galaxies, which have limiting cusp slopes with $0.3<\\gamma'<0.5,$ they are not common enough to fill the ``valley'' between the power-law and core cusp-slope distributions.  This is a critical point that has often been misunderstood in the literature.  The sample of galaxies available to \\citet{f97} had no systems with $0.3<\\gamma'<0.5,$ thus the distributions of slopes was not only bimodal, but disjoint. The subsequent discovery of a small number of intermediate galaxies shows that the distribution of cusp slopes is truly continuous, but it has remained bimodal in its overall character in all the studies cited above.  This picture has been challenged by the recent distribution of cusp slopes presented by \\citet{lf}, who do not find bimodality in their sample of 100 galaxies observed by their {\\it HST} Virgo Cluster Survey (VCS), although they still identify a separate class of core galaxies, which they define as having brightness profiles that fall below \\citet{sersic} forms in the center. While the use of \\Ser profiles is a new approach to identifying core galaxies, the class of VCS core galaxies identified in this way is essentially consistent with the definition of core galaxies in earlier studies. The distribution of slopes for the non-core galaxies in \\citet{lf}, however, does not match that of the power-law galaxies presented in all the studies cited above, even though this set includes several galaxies previously classified as power-laws.  The difference between the VCS work and all earlier studies lies not in whether or not there are two classes of galaxies, core and \"non-core\", but in the form of the central structure of the non-core class. Resolving the different pictures presented by \\citet{f97} and the VCS results is crucial to understanding how the central structure of elliptical galaxies varies with luminosity, and how it may change over mergers.  This paper has three goals.  We first construct a large sample of galaxies that have had their surface brightness distributions characterized by Nuker law fits.  These are reduced to a common photometric and distance scale.  This sample provides the basic material for an updated investigation into the relationships between the central structure of galaxies and their more global properties, as was presented in \\citet{f97}; this new analysis is presented in \\citet{l07}. The second goal is to reinvestigate the distribution of cusp slopes in early type galaxies.  We find that the distribution of slopes remains strongly bimodal, a conclusion that we ratify with a repeat of the non-parametric analysis presented by \\citet{g96}. The final task is to bolster this conclusion with a comparison of the analysis of {\\it HST} imagery by \\citet{l95, l05} to that of \\citet{lf}.  We show that the VCS surface photometry models do not accurately recover the distribution of $\\gamma'$ for galaxies with $M_V<-19.$ ",
        "conclusions": "The distribution of surface brightness cusp slopes as the {\\it HST} resolution limit is approached is strongly bimodal.  This confirms once again the usefulness of classification of galaxies into core or power-law systems based on their central brightness profiles. Bimodality is still evident for galaxies within 50 Mpc, or for those restricted in luminosity to $-22.5<M_V\\leq-20.5.$  This last result is of particular interest as it is the interval over which the typical central structure of early-type galaxies changes in form.  Very few core galaxies have $M_V\\geq-20.5,$ while only a small number of power-law galaxies have $M_V\\leq-22.5.$  There are a few galaxies that have profiles with $0.3<\\gamma'<0.5,$ thus placing them in an intermediate position between core and power-law systems.  Intriguingly, these galaxies are most common in the luminosity range over which the core/power-law transition occurs.  Intermediate galaxies are rare, however, and they do not fill in the ``valley'' in the distribution of cusp slopes.  As a check on our conclusions that the distribution of cusp slopes is bimodal, we repeated the non-parametric analysis of the {\\it HST} brightness profiles presented by \\citet{g96}. This provides a strong demonstration of bimodality in the space {\\it density} profiles without reference to any fitting function or any physical assumption beyond adopting the radius of maximum curvature of the profiles as a common physical scale.  The concordance of this very simple approach with the results from the Nuker law fits, shows that inferences of bimodality from Nuker law fits alone are not due to artifacts in or poor behavior of its parametric form,  These conclusions are inconsistent with the results of \\citet{lf}, who presented a strongly different distribution of $\\gamma'$ values based on their VCS galaxies. By examining galaxies in common with the VCS sample, we find that profiles derived from ACS/WFC, WFPC2, and WFPC1 agree extremely well to their stated resolution limits. In contrast, we find that the VCS profile models, which the VCS uses to characterize cusp slopes, are either poor representations of the data as the {\\it HST} resolution limit is approached, or consist of inward extrapolations of forms fitted to the envelope. While the latter case, which requires the assumption of large nuclear clusters, may be a plausible model, it is neither unique nor compelling enough to demonstrate that the previous modeling of profiles conducted by our group and others is ``untenable.''  We do note our large composite sample is largely limited to galaxies with $M_V<-19.$ The core/power-law transition occurs at still larger luminosities, and our conclusions on bimodality are limited to this context. This luminosity limit corresponds to roughly the median VCS galaxy luminosity, however, and many of the faintest VCS galaxies have profiles that are nearly flat at the center. \\citet{lf} are concerned that under our criteria such systems would be classified as cores, something that we have not actually done. The envelope profiles of these galaxies are nearly exponential and have no well-defined breaks as the center is approached\\footnote{Indeed, the maximum logarithmic second derivative, and hence the location of $r_b,$ is reached only in the limit $r\\rightarrow\\infty$ for a pure exponential profile.}. Their overall forms are completely different from the luminous core galaxies and thus they are trivially distinguished from them. \\citet{k85}, \\citet{k87}, and \\citet{kb94} showed that ellipticals and spheroidals form well defined, disjoint sequences in parameter space. A recent analysis of the profiles of low-luminosity ellipticals and a comparison to a large sample of Sph galaxies strongly confirms the E/Sph dichotomy \\citep{k07}. The resolution of this question, however, is not relevant to the strong physical dichotomy between power-law and core galaxies advanced in \\citet{f97}. Cores are a phenomenon of the most luminous galaxies.  Their likely immediate progenitors are the power-law galaxies of slightly lower luminosity. The distribution of $\\gamma'$ for $M_V<-19$ is the relevant diagnostic for how this transition occurred."
    },
    "0609/astro-ph0609281_arXiv.txt": {
        "abstract": "High-resolution X-ray spectroscopy has addressed not only various topics in coronal physics of stars, but has also uncovered important features relevant for our understanding of stellar evolution and the stellar environment. I summarize recent progress in coronal X-ray spectroscopy and in particular also discuss new results from studies of X-rays from pre-main sequence stars. ",
        "introduction": "X-rays from cool stars are a manifestation of magnetic fields generated by the internal dynamos. This by now classical view has not changed despite the discovery of new X-ray phenomenology and the development of new theoretical concepts with which to explain the production of X-rays around stars (e.g., accretion, outflows, fluorescence, etc).   However, with  increasing detector sensitivities and the advent of high-resolution spectroscopy, ``stellar environmental issues'' have become of central interest in the area of pre-main sequence stars. The interplay between accretion disks and stellar magnetic fields leads to disk instabilities, ionization and heating of circumstellar material, chemical reactions in cool molecular gas, accretion flows from the disk to the star, outflows, and jets,  in all of which X-rays either play a key role or at least provide important diagnostic access.  Limitations of present-day instrumentation are, however, evident. The following sections summarize selected topics of X-rays from cool stars that have been addressed in recent years using high-resolution spectroscopy  from {\\it XMM-Newton} and {\\it Chandra}. ",
        "conclusions": "High-resolution X-ray spectroscopy has provided novel access to stellar coronae by resolving emission lines from which the composition and structure of magnetically confined coronal plasmas has been derived. There have also been little anticipated developments such as X-rays related to the {\\it stellar environment}, e.g. accretion disks, accretion flows, and jets. An encouraging array of new emission models and of implications for stellar evolution have been described in the literature. Nevertheless, limitations have rapidly been reached with the presently available grating spectrometers. Only few T Tau stars are accessible to spectroscopic investigations, and some crucial classes of stars (protostars, brown dwarfs) are out of reach of these gratings. The next generation of X-ray satellites should improve on two parameters: wavelength resolution and sensitivity. A resolving power of a few 1000 will systematically access orbit and rotational velocities of many stars and binaries. Flare plasma flow velocities of a few 100~km~s$^{-1}$ can be systematically studied. Further, the problem with unrecognized blends would be dramatically reduced. Sensitivity will be required to study objects in star forming regions, in nearby open clusters, but also to access close but faint objects such as evolved main-sequence stars or brown dwarfs.    \\small"
    },
    "0609/gr-qc0609039_arXiv.txt": {
        "abstract": "In this paper we study the duality in two-field Quintom models of Dark Energy. We find that an expanding universe dominated by Quintom-A field is dual to a contracting universe with Quintom-B field. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609358_arXiv.txt": {
        "abstract": "We model the stochastic attenuation by \\hi{} absorbers in the intergalactic medium (IGM), such as \\lya{} Forest clouds, and absorbers associated with galaxies, such as Lyman Limit systems (LLS) and Damped Lyman Alpha absorbers (DLAs), and compute an ensemble of $4 \\cdot 10^3$ attenuated Spectral Energy Distributions (SEDs) in the Johnson system for the spectrum of a galaxy with a constant star formation rate (CSFR). Using these, we asses the impact of the stochastic attenuation on the estimates of photometric redshifts for this type of galaxy by comparison with model SEDs that include only a mean attenuation. ",
        "introduction": "The stochastic distribution of the (high column density) \\hi absorbers causes a significant scatter in the broadband colors of galaxies, especially at redshifts close to the drop-out redshift in the corresponding passband where the photoelectric absorption is more severe, and this is expected to drastically affect the estimates of photometric redshifts. ",
        "conclusions": "As can be seen in figure 1, we find that, using the mean attenuation only, the median photometric redshifts are severe underestimated by typically $\\Delta z_{phot} = 0.3$, especially in the range \\mbox{$ z > 3.2$}. In addition, we observe a substantial scatter of \\mbox{$\\Delta z_{phot} \\approx \\pm 0.5$}.  We predict that any method for determining photometric redshifts that only takes into account the mean attenuation at high redshifts should show a bias such as the one seen in Figure \\ref{fig:zp}. We hence emphasise the need for an accurate modeling of the attenuation including stochastic effects, in order to correctly interpret in combination with evolutionary synthesis models such as {\\sc galev}, the observations of high-redshift galaxies in deep multi-band imaging surveys.\\\\  \\begin{figure}\\label{fig:zp} \\begin{center} \\includegraphics[scale=0.4, angle=0]{tepper_median_zphot_3} \\caption{Estimate of the photometric redshift $z_{phot}$ for an Sd-type spectrum with \\emph{stochastic} attenuation, using as input the same spectrum with \\emph{mean} attenuation. The solid line shows the median, and the shaded area gives the $3\\sigma$-range around the mean for the entire ensemble of 4000 SEDs.} \\end{center} \\end{figure}"
    },
    "0609/astro-ph0609391_arXiv.txt": {
        "abstract": "{Clusters of galaxies are believed to be at the intersections of cosmological filaments and to grow by accreting matter from these filaments.  Such continuous infall has major consequences not only on clusters but also on the physics of cluster galaxies. Faint galaxies are particularly interesting as they are very sensitive to environmental effects, and may have a different behaviour from that of bright galaxies.} {The aim of this paper is to sample the Coma cluster building history, based on the analysis of galaxy luminosity functions (LFs) and red sequences (RSs) in the Color Magnitude Relation (CMR) down to faint magnitudes, which are privileged tools for this purpose.} {The present analysis is based on deep (R$\\sim$24), wide ($\\sim$0.5 deg$^2$) multiband (BVRI Vega system) images of the Coma cluster obtained with the CFH12K camera at the CFHT. We have derived LFs and CMRs in twenty 10$\\times$10~arcmin$^2$ regions and in larger regions.} {In all photometric bands, we found steeply rising LFs in the north-northeast half of the cluster (due to early type galaxies at bright magnitudes and due to late type galaxies at the faint end), and much flatter LFs in the south-southwest region. Although the fine behaviour of the CMR RS is different in these two regions, a good agreement is found in general between the RS computed for faint and for bright galaxies.} {All these results can be interpreted consistently in the framework of the building up process proposed by Adami et al. (2005b). The Northern Coma area is a relatively quiescent region while the southern area experiences several infalls.} ",
        "introduction": "\\label{sec:intro}  Hierarchical building of clusters is a key ingredient of cosmological models, since clusters are believed to be at the nodes of cosmological filaments and to grow by accreting matter (gas, galaxies, groups etc.) from filaments (e.g.  Lanzoni et al. 2005).  Such continuous infall has major consequences on the shape and dynamics of clusters, as well as on the physics of cluster galaxies (e.g. Sarazin 1986) as revealed by the cluster galaxy mass function (usually traced by the luminosity function, LF hereafter) and by the cluster early type galaxy Red Sequence (RS hereafter).  The influence of infall on the LF is relatively well understood for bright galaxies in nearby to intermediate redshift clusters (e.g. Hansen et al.  2005), but only little is known for low mass galaxies because they are difficult to observe. It is also important to study galaxies at the faint end of the RS because the Color Magnitude Relation (CMR hereafter), which traces their star formation history, is directly related to their metal abundances and is poorly known. These faint galaxies are of major interest as their evolutionary paths are sometimes different from those of bright galaxies: they are very sensitive to environmental effects (e.g. fusion with bright galaxies or tidal disruptions); they also keep their dynamical memory longer (e.g.  Sarazin 1986) and their spatial distribution is different from that of bright galaxies (e.g. Biviano et al. 1996, but see also Edwards et al. 2002 for a different result).  To study the properties of galaxy clusters through their LF and CMR requires to sample the whole cluster galaxy populations from the bright cD to the faintest dwarfs (M$_R \\sim -10$). Very demanding observations are needed both in terms of photometric depth and size of field, and were not possible until the arrival of wide-field facilities such as the CFH12K on the CFHT. With such a 4m-class telescope, access to the faintest objects is still limited to nearby clusters. In this respect, the Coma cluster is possibly the best target: bright objects are well studied in all wavelengths (e.g. Biviano 1998) and its low redshift allows to go deep enough towards the faint end of the galaxy population.  LFs and CMRs have been recently computed for bright to intermediate magnitudes in Coma (e.g. Lobo et al. 1997, Terlevich et al. 2001, Andreon $\\&$ Cuillandre 2002, Beijersbergen et al. 2002, Iglesias-P\\'aramo et al. 2003) in large fields. Deeper studies in rather limited areas also exist (Bernstein et al. 1995 or Trentham 1998). All these surveys generally have a wavelength coverage too limited to provide strong constraints on the Coma cluster building history. For the first time, by using the CFH12K camera at the CFHT, we obtained a set of deep (R$\\sim$24), wide ($\\sim$0.5 deg$^2$) multiband (BVRI) images of the Coma cluster.  The key point to compute a LF or a CMR is to select objects in the cluster and then galaxies, as compared to intergalactic Globular Clusters and galactic stars. For the bright part of our sample, we had enough spectroscopic redshifts to discriminate Coma members.  As it is prohibitive to obtain spectra for the faint part, we used statistical comparisons with empty fields (containing no rich nearby clusters). This is a classical method which requires very homogeneous observations of the cluster and of the empty fields.  The empty fields we used were parts of two of the VVDS CFH12K imaging survey fields (McCracken et al. 2003: MC03 hereafter) completely re-analyzed in order to have catalogs homogeneous with the Coma data. This allows to treat properly possible pitfalls such as diffuse light in the cluster, galactic extinction, differences in the star counts, large scale structure of the background and crowding effects.  Other problems which could affect the background counts along the Coma line of sight are the presence of dust in the cluster or the gravitational effects on background galaxies but they have been shown to be negligible in Coma (Bernstein et al. 1995).  In Section 2, we review the data and catalogs and we describe the homogenization process of these catalogs in Section~3. In Sections 4 and 5, we compute the LFs for bright and faint galaxies. Section 6 is dedicated to the nature of the faintest Coma cluster objects. In Section 7, we describe the LF characteristics. Section 8 describes the RS in the Coma CMR. We discuss in Sections 9 and 10 the inferences of the Coma LF and CMR RS on the cluster building up scenario. Finally, Section 11 is the conclusion.  We assume a distance to Coma of 95 Mpc, H$_0$ = 75 km s$^{-1}$ Mpc$^{-1}$. The distance modulus is 34.89, and therefore the scale is 0.46 kpc arcsec$^{-1}$. We give magnitudes in the Vega system except if specified otherwise. ",
        "conclusions": "In this paper, we have analyzed the deep galaxy LF in the B, V, R and I bands and the faint B-R/R CMR for the Coma cluster. We applied a single Schechter law to model our LFs. In order to sample possible environmental effects on the LF, we derived twenty such functions in 10$\\times$10~arcmin$^2$ areas. We found a dichotomy between steeply rising LFs in the north-northeast Coma region, which appear to be due to red/early type galaxies at bright magnitudes, and to blue/late type objects at very faint magnitudes, and much flatter LFs in the south-southwest region, both for all objects and for blue/late type or red/early type galaxies. This result is valid for all photometric bands and all observation epochs.  We also found that the LFs differ around the main Coma cluster galaxies and groups. LFs in the NGC~4889 field are steeply rising in all photometric bands while around NGC~4874, they are rising moderately in B and V and steeply in R and I, with a maximum reached around R$\\sim$23 and I$\\sim$22.5. Around NGC~4911, the LF is moderately rising in B and V up to B$\\sim$23.75 and V$\\sim$23.5 and oscillating and rising up to R$\\sim$23 and over the whole I magnitude range.  We have discussed the observed shapes of the LFs in various regions in the framework of the building up process that we have proposed in a previous paper (Adami et al.  2005b). In this scenario, several infalls onto the Coma cluster are taking place, coming from three directions: northeast-east, northwest-west and south-west. NGC~4874 is the original cluster dominant galaxy, while NGC~4889 experienced a later capture. How does this match our observations?  The GC expelling scenario explains quite well the high fraction of GCs found close to NGC~4874.  The south-southwest regions showing quite flat LFs are also those which are directly located on the main field infall directions (in particular the infalling NGC~4839 group). The observed paucity of faint galaxies in the southern region is only partially explained by the fact that faint galaxies may have been swallowed by relatively bright galaxies on their way to the cluster center. Rather, north counts could be enhanced by these south infalls to explain their steeply rising LFs (the north part has a LF slope significantly different from the field, e.g. Ilbert et al. 2005).  However, we can note that faint galaxies are blue/late type while bright galaxies are red/early types in the north. This suggests that these blue and faint galaxies could also partially be created in the cluster from material expelled from the envelopes of spiral-like parent galaxies.  We have detected several dips in the LFs in different areas of the cluster and shown that some of them are probably due to mergings/fusions, while others are probably due to the inhibition of star-formation activity by cluster environmental effects.  Finally, we have also clearly detected a RS down to R=23.5, close to the bright galaxy RS computed in A06a. This shows that most of the faint Coma cluster galaxies have probably followed the same evolutionary path as the bright ellipticals in terms of mass-metallicity relation. However, part of the faint cluster population followed a different evolution on the blue and red sides. We suggest that blue objects present over the whole Coma cluster field are partially due to recent galaxy formation from material expelled from late type galaxies. It is particularly interesting to note that the CMR extends over such a large range of magnitudes, thus making it a very useful tool for several purposes, such as selecting galaxies with a high probability of belonging to a cluster from imaging data only.  In a near future, we intend to compute and discuss the Coma cluster luminosity functions based on photometric redshifts as soon as we will have acquired the  U band data lacking for our survey (see A06a)."
    },
    "0609/astro-ph0609672_arXiv.txt": {
        "abstract": "We present new atmosphere models for Wolf-Rayet stars that include a self-consistent solution of the wind hydrodynamics. We demonstrate that the formation of optically thick WR winds can be explained by radiative driving on Fe line opacities, implying a strong dependence on metallicity ($Z$). $Z$-dependent model calculations for late-type WN stars show that these objects are very massive stars close to the Eddington limit, and that their formation is strongly favored for high metallicity environments. ",
        "introduction": "The Potsdam Wolf-Rayet (PoWR) hydrodynamic model atmospheres combine fully line-blanketed non-LTE models with the equations of hydrodynamics \\citep[for details see][]{gra1:05,ham1:03,koe1:02,gra1:02}. The wind structure ($\\rho(r)$ and $v(r)$) and the temperature structure $T(r)$ are computed in line with the full set of non-LTE populations, and the radiation field in the co-moving frame (CMF). In contrast to all previous approaches, the radiative wind acceleration $a_{\\rm rad}$ is obtained by direct integration \\begin{equation} \\label{eq:arad} a_{\\rm rad} = \\frac{1}{c} \\int \\chi_\\nu F_\\nu {\\rm d}\\nu, \\end{equation} instead of making use of the Sobolev approximation.  In this way, complex processes like strong line overlap, or the redistribution of radiation, are automatically taken into account. Moreover, the models include small-scale wind clumping \\citep[throughout this work we assume a clumping factor of $D=10$, for details see also][]{ham1:98}. The models describe the conditions in WR\\,atmospheres in a realistic manner, and provide synthetic spectra, i.e.\\ they allow for a direct comparison with observations.  Utilizing these models, we recently obtained the first fully self-consistent Wolf-Rayet wind model, for the case of an early-type WC star with strong lines \\citep{gra1:05}.  Moreover, we have examined the mass loss from late-type WN stars and its dependence on metallicity \\citep{gra3:06,gra1:06}. ",
        "conclusions": ""
    },
    "0609/astro-ph0609444_arXiv.txt": {
        "abstract": "{}{We aim at detecting and determining the properties of the host galaxy of the dark GRB 050223.}{We use VLT optical/NIR images coupled to {{\\em Swift}} X-ray positioning, and optical spectra of the host galaxy to measure its properties.}{We find a single galaxy within the {{\\em Swift}} error box of GRB 050223. It is located at $z = 0.584$ and its luminosity is $L \\sim 0.4$~$L^*$. Emission lines in the galaxy spectrum imply an intrinsic ${\\rm SFR} > 7$~$M_{\\sun}$~yr$^{-1}$, and a large extinction $A_{\\rm V} > 2$~mag within it. We also detect absorption lines, which reveal an underlying stellar population with an age between 40~Myr and 1.5~Gyr.}{The identification of a host galaxy with atypical properties using only the X-ray transient suggests that a bias may be present~in the former sample of host galaxies. Dust obscuration together with intrinsic faintness are the most probable causes for the darkness of this burst.} ",
        "introduction": "\\label{intro}  Long gamma-ray bursts (LGRBs) are brief pulses of $\\gamma$-ray radiation occuring about once per day at random positions in the sky, and lasting more than 2~s. LGRBs are accompanied by fading afterglows detected at longer wavelengths, which can be observed for months or even years. Optical and NIR afterglows are usually observed, except for a subpopulation of dark LGRBs (DGRBs). This can be attributed, in some cases, to the lack of early observations reaching faint magnitudes. In other cases deep and prompt follow-ups indicate that true DGRBs indeed exist, and that they comprise a fraction of $\\sim$10\\% of all bursts (Jakobsson et~al. \\cite{Jak05}).  The lack of optical afterglows in DGRBs could be explained in different ways. DGRB afterglows could be intrinsically underluminous, or extincted by intergalactic Ly-$\\alpha$ absorbers (if located at high redshift) or by the dusty interstellar medium of their host galaxies (HGs). While some authors find that the intrinsic faintness hypothesis could explain some cases (e.g., Pedersen et~al. \\cite{Ped06}), for other DGRBs the existence of at least some obscuration inside the HG has been claimed (e.g., Piro et~al. \\cite{Pir02}). A particularly interesting case is GRB 970828 (Groot et~al. \\cite{Gro98}; Djorgovski et~al. \\cite{Djo01}), which has been proposed as the prototype of dust-enshrouded bursts, and is the only DGRB whose HG has been detected in the FIR by {\\em Spitzer} (Le Floc'h et~al. \\cite{LeF06}).  The hypothesis of dust extinction of DGRB afterglows is particularly interesting in the context of the relationship between LGRBs and star formation (SF). The connection of LGRBs with core-collapse supernovae (Hjorth et~al. \\cite{Hjo03}; Stanek et~al. \\cite{Sta03}; Malesani et~al. \\cite{Mal04}; Campana et~al. \\cite{Cam06}), the SF activity observed in their HGs (Le Floc'h et~al. \\cite{LeF02}; Prochaska et~al. \\cite{Pro04}), and their massive stellar progenitors proposed by theoretical models (Fryer et~al. \\cite{Fry99}), suggest that LGRBs might be good SF tracers. Their detectability at high redshift (Tagliaferri et~al. \\cite{Tag05}; Haislip et~al. \\cite{Hai06}; Kawai et~al. \\cite{Kaw06}) would then transform them in a promising tool to probe SF in the early Universe. However, the sample of HGs obtained so far shows a bias to underluminous and blue objects (Le Floc'h et~al. \\cite{LeF03}), while there is evidence that a significant fraction of the SF at high redshift took place in luminous, reddened and dust-enshrouded IR starbursts (Elbaz et~al. \\cite{Elb02}). If the obscuration hypothesis is correct, a natural explanation for this bias arises from the fact that, in most previous investigations, LGRB HGs have been identified through the accurate positioning given by optical afterglow measurements. Hence, mostly HGs of unobscured bursts have been identified, leading to a natural exclusion of dusty starburst galaxies. The detection of such galaxies associated to LGRBs would make a strong case for the obscuration hypothesis, and would also support the LGRB-SF connection.  GRB 050223 was discovered by the {\\em Swift} satellite at 03:09:06~UT on February 23, 2005 (Mitani et~al. \\cite{Mit05}), and also detected by {\\em INTEGRAL} (Mereghetti et~al. \\cite{Mer05}). It was a LGRB ($T_{90} = 23$~s), with a fluence of $7.4 \\times 10^{-7}$~erg~cm$^{-2}$ (15--350~keV). The {\\em Swift} X-ray telescope (XRT) detected its afterglow (Giommi et~al. \\cite{Gio05}), which was also observed by {\\em XMM-Newton} (De Luca \\& Campana \\cite{DeL05}). The afterglow presented a power-law decay with a slope $\\alpha = 0.99^{+0.15}_{-0.12}$. Its spectrum was well fitted by an absorbed power law with a photon index $\\Gamma = 1.75^{+0.19}_{-0.18}$, and a hydrogen density column $N_{\\rm H} \\sim 10^{21}$~cm$^{-2}$, consistent with the Galactic value, $N_{\\rm H} = 7 \\times 10^{20}$~cm$^{-2}$ (Page et~al. \\cite{Pag05}). XRT gave also a precise positioning of the afterglow, $\\alpha_{\\rm J2000} = 18^{\\rm h}05^{\\rm m}33\\fs08$, $\\delta_{\\rm J2000} = -62\\degr28\\arcmin20\\farcs5$ (5\\farcs4 error radius; Moretti et~al. \\cite{Mor06}).  Searches for the ultraviolet and optical afterglow were made by {\\em Swift} Ultraviolet and Optical Telescope (Gronwall et~al. \\cite{Gro05}), {\\em XMM} Optical Monitor (Blustin et~al. \\cite{Blu05}), ROTSE~III (Smith \\cite{Smi05}), PROMPT (Nysewander et~al. \\cite{Nys05}) and the Mount John University Observatory (Gorosabel et~al. \\cite{Gor05}). No observations were reported in the IR or radio domains. The deepest observations put a stringent limiting magnitude $R > 21.2$ only 4.1 hours after the burst. Assuming that the decay of the optical afterglow follows the same law as the X-ray afterglow, and taking into account the Galactic absorption of 0.2~mag at the position of the burst (Schlegel et~al. \\cite{Sch98}), this limit implies a magnitude $R > 23.8$ at 2 days after the burst. Hence, following the criterion of Djorgovski et~al. (\\cite{Djo01}), GRB 050223 can be classified as a DGRB. The aforementioned observations put also an upper limit $\\beta_{\\rm OX} < 0.8$ for the mean optical to X-ray spectral index of GRB 050223, hence it can not be classified as dark according to the criterion of Jakobsson et~al. (\\cite{Jak04}), which requires $\\beta_{\\rm OX} < 0.5$ for a burst to be dark. This is a direct consequence of the faintness of the X-ray afterglow, which was used by Page et~al. (\\cite{Pag05}), together with the value of $N_{\\rm H}$, to argue that GRB 050223 has an intrinsically underluminous afterglow.  In this Letter we report the detection of the host galaxy of GRB 050223 and the determination of its properties. Sects.~\\ref{obs} and \\ref{res} present our observations and results respectively. Sect.~\\ref{disc} discusses the properties of the galaxy and their implications on the nature of DGRBs. ",
        "conclusions": "\\label{disc}  Combining the accurate X-ray positioning of LGRB afterglows given by {\\em Swift} XRT, and very deep VLT optical and NIR images and optical spectra, we identified the likely HG of the dark GRB 050223. It is a 0.4~$L^*$, blue galaxy at $z = 0.584$, with strong SF (${\\rm SFR} > 7$~$M_{\\sun}$~yr$^{-1}$, SSFR~$> 17.5$~$M_{\\sun}$~yr$^{-1}$~$L^{*-1}$), a stellar population dominated by A0V stars with an age of 40~Myr--1.5~Gyr, and a large extinction ($A_{\\rm V} > 2$~mag).  This galaxy has some typical properties of LGRB HGs. It is underluminous and blue, with strong emission lines indicating SF, although it is brighter and redder than other LGRB HGs at its redshift (Le Floc'h et~al. \\cite{LeF03}). It shows also some atypical features. The [\\ion{Ne}{iii}] 3869~{\\AA} to [\\ion{O}{ii}] 3727~{\\AA} ratio is low, and the inferred SFR and extinction are among the highest ones derived from optical observations (much higher SFRs were determined from radio observations, e.g. by Berger et~al. (\\cite{Ber03}), although radio SFRs are in some cases higher than optical ones for the same galaxy). Another atypical property is the presence of a strong stellar population dominated by A0V stars, implied by the observed absorption lines, and the e(a) classification that follows from them. We point out that our finding of an HG spectrum that differs from those previously reported is the result of a new HG selection based only on X-ray transients. This may imply the presence of a bias in the HG sample that has been considered so far.  The large observed extinction suggests that this DGRB took place in a dust-enshrouded starburst. The e(a) type spectrum is usually interpreted as produced in galaxies with strong ongoing star formation, together with older stellar populations and a high selective obscuration affecting the youngest stars (Poggianti et~al. \\cite{Pog01}). Star forming regions, inside which LGRBs occur, are expected to be highly obscured in these galaxies. Hence, our results support the obscuration hypothesis for DGRBs, which was also strengthened by the recent classification of the HG of GRB 030115 as an extremely red object (Levan et~al. \\cite{Lew06}). Our picture is consistent with that of Nakagawa et~al. (\\cite{Nak06}), who proposed that the dark GRB 051022 took place inside a dusty molecular cloud, given the large extinction ($A_{\\rm V} = 49$~mag) derived from X-ray data. We point out that a very high dust extinction has been also proposed as a possible reason for the discrepancy between the rate of core-collapse SNe (which are related to LGRBs) determined from NIR surveys, and that estimated from the FIR luminosity of their parent galaxies (Mannucci et~al. \\cite{Man03}). However, we note that the only LGRB with a HG possibly showing a similar extinction ($A_{\\rm V} = 3.4$~mag according to Vreeswijk et~al. \\cite{Vre01}, but only 0.15~mag according to Christensen et~al. \\cite{Chr04}) to that of our candidate, is not a DGRB (GRB 990712 at $z = 0.43$). If its high extinction is confirmed, this could be the result of a different dust distribution in both galaxies, obscuration not affecting the star-forming regions in the HG of GRB 990712 (not an e(a) type). Or, this might imply that dust extinction alone is not a sufficient condition to produce a DGRB. On the other hand, the confirmation of a low extinction for the GRB 990712 HG would strengthen the hypothesis that dusty strabursts produce DGRBs. It is important to note, however, that there are DGRB HGs which are not representative of dusty starbursts (e.g., GRB 000210; Gorosabel et~al. \\cite{Gor03}).  Our measurement of a redshift $z = 0.584$ for the likely host galaxy of GRB 050223 casts a shadow over the high redshift hypothesis proposed by Page et~al. (\\cite{Pag05}) to explain the faintness of its X-ray afterglow, leaving the possibility of a large jet opening or viewing angle. Our detection of obscured star formation in its HG does not contradict their conclusion of the GRB 050223 afterglow being underluminous. Intrinsic faintness might well be complementary to dust absorption in producing a DGRB, either by requiring a modest dust mass to suppress the afterglow, or by being ineffective in destroying the circumburst dust. We note that GRB 050223 and GRB 990712 lie at opposite ends in the X-ray afterglow brightness \\-distribution.  Although it shows a high optical extinction, the likely HG of GRB 050223 does not seem to be one of the galaxies with luminous, reddened IR starbursts and very high SF activity. Particularly, it was not detected by {\\em Spitzer} (Le Floc'h 2006, priv. comm.) in the FIR. The constraint set by these observations, ${\\rm SFR} < 10$~$M_{\\sun}$~yr$^{-1}$, is consistent with our results. However, it is noteworthy that the position of this HG in Figs.~4 and 5 of Le Floc'h et~al. (\\cite{LeF03}) is consistent with the loci of the NIR counterparts of high redshift ISO galaxies. Further investigation of this HG would be important to get a deeper insight into the nature of DGRBs and the LGRB-SF connection."
    },
    "0609/astro-ph0609734_arXiv.txt": {
        "abstract": "Deep surveys have recently discovered galaxies at the tail end of the epoch of reionization. In the near future, these discoveries will be complemented by a new generation of low-frequency radio observatories that will map the distribution of neutral hydrogen in the intergalactic medium through its redshifted 21cm emission.  In this paper we calculate the expected cross-correlation between the distribution of galaxies and intergalactic 21cm emission at high redshifts.  We demonstrate using a simple model that overdense regions are expected to be ionized early as a result of their biased galaxy formation. This early phase leads to an anti-correlation between the 21cm emission and the overdensities in galaxies, matter, and neutral hydrogen. Existing Ly$\\alpha$ surveys probe galaxies that are highly clustered in overdense regions. By comparing 21cm emission from regions near observed galaxies to those away from observed galaxies, future observations will be able to test this generic prediction and calibrate the ionizing luminosity of high-redshift galaxies. ",
        "introduction": "An important question that should be addressed by successful models of reionization concerns whether overdense or underdense regions become ionized first.  In regions that are overdense, galaxies will be over-abundant for two reasons; first because there is more material per unit volume to make galaxies, and second because small-scale fluctuations need to be of lower amplitude to form a galaxy when embedded in a larger-scale overdensity (the so-called {\\it galaxy bias}; see Mo \\& White~1996). Regarding reionization of the intergalactic medium (IGM), the first effect will be compensated by the increased density of gas to be ionized. Furthermore, the increase in the recombination rate in overdense regions will be counteracted by the galaxy bias in overdense regions. However the latter effects need not cancel, and could either lead to enhanced or delayed reionization in overdense regions.  The process of reionization also contains several layers of feedback. Radiative feedback heats the IGM and results in the suppression of low-mass galaxy formation (Efstathiou, 1992; Thoul \\& Weinberg~1996; Quinn et al.~1996; Dijkstra et al.~2004). This delays the completion of reionization by lowering the local star formation rate, but the effect is counteracted in overdense regions by the biased formation of massive galaxies. The radiation feedback may therefore be more important in low-density regions where small galaxies contribute more significantly to the ionizing flux.   \\begin{figure*} \\includegraphics[width=15cm]{fig1.eps} \\caption{ \\textit{Left}: Fluctuations in the hydrogen neutral fraction [$\\delta_{\\rm HI}\\equiv x_{\\rm HI}(\\delta)/\\langle x_{\\rm HI}\\rangle-1$] versus overdensity $\\delta$ (for $R\\gg R_{\\rm min}$). The solid, dot-dashed and dashed lines represent values of $C=10$, $C=2$ and $C=20$. For comparison, we show three lines of slope $d\\delta_{\\rm HI}/d\\delta=-0.5$, $-1.5$ and $-2.5$. \\textit{Right}: The dependence of the predicted brightness temperature ($T$) on overdensity. The solid, dot-dashed and dashed lines represent values of $C=10$, $C=2$ and $C=20$. For comparison, the three lines show Equation~(\\ref{neut}) with values of $\\beta=-0.5$, -1.5 and -2.5. The upper and lower rows correspond to observations at $z=6.57$ and $z=8$ respectively. } \\label{fig1} \\end{figure*}   In this paper we use a simple model to evaluate the relative significance of the above effects and compute the correlation between the local overdensity and quantities such as the hydrogen neutral-fraction, the brightness temperature of redshifted 21cm emission, and the overdensity of massive galaxies.  In particular, we evaluate the prospects for finding whether reionization is enhanced or suppressed in overdense regions based on the combination of high redshift galaxy surveys and redshifted 21cm surveys.  For illustration, the problem may be parameterised by relating the neutral fraction [$x_{\\rm HI}^{R}(\\delta)$] in regions of linear overdensity $\\delta$ (smoothed over spheres of large radius $R$), to the average neutral fraction ($x_{\\rm HI}$) using a power-law with index $\\beta(R)$ \\begin{equation} \\label{neut} x_{\\rm HI}^{R}(\\delta)=x_{HI}\\left(1+\\delta\\right)^{\\beta(R)}\\approx x_{HI}\\left[1+\\beta(R)\\delta\\right], \\end{equation} where we assume $\\delta\\ll 1$. The resulting deviation in the brightness temperature of a region of overdensity $\\delta$ at $z\\ga 6$ is \\begin{eqnarray} \\label{temp} \\nonumber T(\\delta)&\\approx&22\\mbox{mK}x_{\\rm HI}^{R}\\left(1+\\frac{4}{3}\\delta\\right)\\\\ &\\approx&22\\mbox{mK}x_{HI}\\left[1+\\left(\\frac{4}{3}+\\beta(R)\\right)\\delta\\right], \\end{eqnarray} where we have assumed hydrogen in the IGM to have a spin temperature well in excess of the CMB temperature. In Equation~(\\ref{temp}) the pre-factor of 4/3 on the overdensity refers to the spherically averaged enhancement of the brightness temperature due to peculiar velocities in overdense regions (Bharadwaj \\& Ali 2005; Barkana \\& Loeb 2005). In a state where the ionization fraction in the IGM is independent of density, the value of the index is $\\beta=0$ for all scales. If $\\beta<0$, then ionization is enhanced in overdense regions. Conversely, reionization is suppressed in overdense regions if $\\beta>0$.  High redshift galaxies are preferentially located in large-scale regions with $\\delta>0$. The departure of the mean 21cm brightness temperature in these regions from the average IGM provides the opportunity to measure the value of $\\beta$, and hence to determine whether overdense or underdense regions were reionized first.  In \\S 2 we describe a simple model that predicts $\\beta<0$. Regions of higher density should therefore be more ionized and possess reduced levels of fluctuations in redshifted 21cm emission.  We later describe how the dependence of 21cm emission on overdensity could be extracted based on the fact that more massive galaxies populate higher density regions. Throughout the paper we adopt the set of cosmological parameters determined by {\\it WMAP} (Spergel et al. 2006) for a flat $\\Lambda$CDM universe. ",
        "conclusions": "In this paper we have calculated the expected cross-correlation between the distribution of galaxies and the intergalactic 21cm emission at high redshifts.  We constructed a simple model for reionization that accounts for both galaxy bias and an enhanced recombination rate in overdense regions, and used this model to compute the ionization fraction as a function of large scale overdensity in the IGM. Our model predicts that overdense regions will be ionized early as a result of their biased galaxy formation.  This early phase of reionization in overdense regions leads to anti-correlations between the 21cm emission and the overdensity of baryons, and between the 21cm emission and the overdensity of neutral hydrogen. In addition, because galaxies are biased towards overdense regions, our model also predicts an anti-correlation between 21cm emission and the galaxy population.  To explore the detectability of any correlation between 21cm emission and galaxy properties, we also constructed a simple model for the Ly$\\alpha$ emission corresponding to a dark-matter halo of a given mass, and hence for the Ly$\\alpha$ luminosity function. We compared this model to an existing Ly$\\alpha$ survey in the Subaru Deep Field. Through this comparison we demonstrated that current surveys probe galaxy masses that are larger than $10^{10}M_\\odot$. Due to their biased formation, galaxies of this mass are highly clustered in overdense regions of the IGM. We have shown that by comparing 21cm emission from regions near observed galaxies to those away from observed galaxies, future redshifted 21cm observations will be able to test the generic prediction that overdense regions are reionized first.    {\\bf Acknowledgments} The research was supported by the Australian Research Council (JSBW) and Harvard University grants (AL).  \\newcommand{\\noopsort}[1]{}"
    },
    "0609/astro-ph0609819_arXiv.txt": {
        "abstract": "We present results from the first 2D and 3D simulations of the collapse of rotating stellar iron cores in general relativity employing a finite-temperature equation of state and an approximate treatment of deleptonization during collapse.  We compare fully nonlinear and conformally flat spacetime evolution methods and find that the conformally flat treatment is sufficiently accurate for the core-collapse supernova problem. We focus on the gravitational wave (GW) emission from rotating collapse, core bounce, and early postbounce phases. Our results indicate that the GW signature of these phases is much more generic than previously estimated. In addition, we track the growth of a nonaxisymmetric instability of dominant $ m = 1 $ character in one of our models that leads to prolonged narrow-band GW emission at $ \\sim 930 \\mathrm{\\ Hz} $ over several tens of milliseconds. ",
        "introduction": " ",
        "conclusions": "\\end{table}     \\emph{Results.}---% In Fig.~\\ref{fig:cfc_cactus} we compare GW signals computed with \\textsc{CoCoNuT} in 2D-CFC and those computed with our 3D-full-GR approach. Model s20A2B2 (red lines) is a moderate rotator with a $ \\beta_\\mathrm{i} = 0.50\\% $, rotating rigidly in its central region. It stays axisymmetric throughout its numerical evolution. The agreement of 2D-CFC with 3D-full-GR is excellent for this model: Both waveforms match almost perfectly at bounce and during the very early postbounce phase. A few ms after bounce, when convection in the region behind the stalling shock sets in due to a negative entropy gradient, the signals begin to differ quantitatively while remaining in phase. We attribute this small mismatch to the choice of coordinate grids and to differences in the growth and scale of vortical postbounce motions between 2D and 3D. Model s20A1B5 rotates with constant $ \\Omega $ in the entire core. Despite its very large $ \\beta_\\mathrm{i} = 4\\% $ it remains essentially axisymmetric during the time covered by our simulation, since most of its angular momentum is attached to material at large radii that falls inward and spins up only slowly. The waveforms in CFC and full GR agree very well. Again, both waveforms match best for the strong burst related to core bounce during which more than $ \\sim 90\\% $ of the total GW energy are emitted in an axisymmetric model. The overall excellent agreement of CFC with full GR confirms results of~\\cite{shibata_04_a,cerda_05_a} and proves that CFC is a very good approximation to full GR in the core-collapse scenario.  \\begin{figure} \\includegraphics[width=8.7cm]{f1.eps} \\caption{GW strain $ h_+ $ along the equator for models s20A2B2 and s20A1B5. We compare 2D-CFC and 3D-full-GR results.} \\label{fig:cfc_cactus} \\end{figure}  \\begin{figure} \\includegraphics[width=8.7cm]{f2.eps} \\caption{Equatorial GW strain $ h_+ $ for representative models.} \\label{fig:waveforms} \\end{figure}  In Fig.~\\ref{fig:waveforms} we present waveforms of models with varying initial degree of differential rotation $ A $ and rotation rate $ \\beta_\\mathrm{i} $. We find that the inclusion of a microphysical finite-temperature EOS and of electron capture during collapse yields results that differ considerably from those obtained in previous, less sophisticated studies. Fig.~\\ref{fig:waveforms} exemplifies that largely independent of the initial rotational configuration, the GW signal of rotating collapse has a \\emph{generic shape}: a slow rise in the pre-bounce phase, a large negative amplitude at bounce when the motion of the inner core is reversed, followed by a ring-down. This so-called ``Type~I'' signature corresponds to a baryonic pressure dominated bounce~\\cite{moenchmeyer_91_a,zwerger_97_a,dimmelmeier_02_a_b,ott_04_a}. Thus all our models undergo core bounce dominated by the stiffening of the EOS at nuclear density.  This is in stark contrast to the studies using the hybrid EOS~\\cite{zwerger_97_a,dimmelmeier_02_a_b,shibata_04_a,shibata_05_a}, where initial models with rotation rates in the range investigated here develop sufficient centrifugal support during contraction to stop collapse at subnuclear densities, resulting in several consecutive centrifugal bounces separated by phases of coherent re-expansion of the inner core. While in GR models exhibiting such multiple centrifugal bounce and the corresponding ``Type~II'' GW signals are only rarer compared to Newtonian gravity~\\cite{dimmelmeier_02_a_b}, we do not observe \\emph{any} such model in this study.  An evident example is model s20A2B4: In previous studies without detailed microphysics, the corresponding model with identical initial rotation parameters (A2B4G1) showed clear ``Type~II'' behavior in both Newtonian and GR calculations~\\cite{zwerger_97_a,dimmelmeier_02_a_b}.  The suppression of the multiple centrifugal bounce scenario is due to two physical effects: (i) In contrast to the simple hybrid EOS, in our case the mass and dynamics of the inner core (which is most important for the GW emission) is not merely determined by the adiabatic index $ \\gamma = d \\ln P / d \\ln \\rho $ (at constant entropy) of the EOS, but also by deleptonization during collapse.  This leads to \\emph{considerably smaller inner cores with less angular momentum and weaker pressure support}. (ii) Since multiple centrifugal bounce was observed for a model with initially moderately fast rotation in a previous Newtonian study~\\cite{moenchmeyer_91_a} where both a microphysical finite-temperature EOS and a deleptonization scheme were employed, the absence of this collapse type in our study is not only caused by microphysical effects, but also by the effectively \\emph{stronger} gravity in GR. This is in accordance with simulations using the simple hybrid EOS~\\cite{dimmelmeier_02_a_b}. A detailed analysis of the interplay and quantitative influence of the above two effects responsible for the elimination of multiple centrifugal bounce in the rotational stellar core-collapse scenario will be presented in a future publication.  \\begin{figure} \\includegraphics[width=8.1cm]{f3.eps} \\caption{Normalized mode amplitudes $ A_m $ at postbounce times (upper panel) and GW strains $ h_+ $ and $ h_\\times $ along the poles (lower panel) for model E20A.} \\label{fig:bar_mode_signal} \\end{figure}  Model E20A possesses the largest GW amplitude of all our models. In addition, it reaches a high $ \\beta_\\mathrm{b} $ at core bounce (see Table~\\ref{tab:model_summary}) and settles at a postbounce $ \\beta_\\mathrm{f} $ of $ \\sim 9\\% $. In a previous Newtonian study, \\cite{ott_05_a} have found a low-$ T / |W| $ nonaxisymmetric instability for a PNS with similar $ \\beta_\\mathrm{f} $. In order to verify their findings, we trace the evolution of model E20A to $ 70 \\mathrm{\\ ms} $ after bounce and perform an analysis of azimuthal density modes $ \\propto e^{im\\varphi} $ in the equatorial plane by computing complex Fourier amplitudes $ C_m = \\frac{1}{2\\pi} \\int^{2\\pi}_0 \\! \\rho(\\varpi, \\varphi, z = 0) \\, e^{im\\varphi} \\, d\\varphi $ on rings of constant coordinate radius with respect to the coordinate center of mass. The latter stays within the innermost zone at all times. In the top panel of Fig.~\\ref{fig:bar_mode_signal} we display the normalized mode amplitudes $ A_m = \\frac{|C_m|}{C_0} $ extracted at $ 15 \\mathrm{\\ km} $ radius. Without adding artificial seed perturbations to model E20A, discretization errors trigger $ m = \\{1,2,3\\} $ modes, which rise to a level of $ \\sim 10^{-5} $ during the $ \\sim 220 \\mathrm{\\ ms} $ collapse phase. After bounce, the m=1 mode exhibits the fastest growth. This growth on a dynamical time scale, lasting over tens of ms until saturation, is closely followed by a growth of $ m = \\{2,3\\} $ daughter modes~\\cite{ott_05_a,ou_06_a}. Note that the $ m = 4 $ Cartesian grid mode starts out on the $ \\sim 10^{-4} $ level and remains constant until the plunge phase of collapse during which all modes are amplified. After core bounce, model E20A remains dynamically stable to the $ m = 4 $ grid mode. In the lower panel of Fig.~\\ref{fig:bar_mode_signal} we plot the  GW strains $ h_+ $ and $ h_\\times $ as seen along the polar  axis. The rotational symmetry of E20A before and early after bounce times is reflected in the fact that $ h_+ $ and $ h_\\times $ are practically zero until E20A develops considerable nonaxisymmetry at $ \\sim 25 \\mathrm{\\ ms} $ after bounce when the $ m = 1 $ mode becomes dominant and its $ m = 2 $, GW-emitting harmonic reaches a sizable amplitude. In remarkable agreement with expectations for a spinning bar, $ h_+ $ and $ h_\\times $ oscillate at the same frequency ($ \\sim 930 $ Hz) and are phase-shifted by a quarter cycle.   \\emph{Discussion.}---% Our results indicate that the GW signature of the collapse, bounce, and early postbounce phases of the core-collapse supernova evolution is much more generic than previously thought. We find that the dynamics of core bounce are dominated by mainly gravity and microphysics, reducing the importance of centrifugal support for the rotation rates considered here. Importantly, for our model set we do not observe rotationally induced multiple core bounce at subnuclear density as proposed by previous studies that did not include a microphysical finite-temperature EOS and electron capture treatment in combination with GR. Thus we predict that the core-bounce waveform of models in a large parameter space of initial rotation rate and degree of differential rotation will likely both qualitatively and quantitatively resemble those presented in Fig.~\\ref{fig:waveforms}.  Model E20A, which we evolve to later postbounce times, exhibits the dynamical growth of a nonaxisymmetric low-$ T / |W| $ corotation-type $ m = 1 $ instability \\cite{saijo_06_a,ou_06_a,ott_05_a}. We also find $ m = \\{2,3\\} $ contributions and thus significant GW emission from the quadrupole components of these modes. We emphasize that we observe this instability not only in E20A, but also in other models with comparable values of $ \\beta_\\mathrm{f} $. Our results, which remove the limitations of previous studies~\\cite{centrella_01_a,shibata_02_a,shibata_05_a,ott_05_a}, demonstrate that the development of nonaxisymmetric structures is neither limited to Newtonian gravity, simple matter models, equilibrium configurations, nor high values of $ \\beta $, but may rather be a phenomenon occurring generically in differentially rotating compact stars.  \\begin{figure} \\includegraphics[width=7.9cm]{f4.eps} \\caption{Spectra of the characteristic GW strain $ h_\\mathrm{char} $ of all models and the LIGO (optimal) rms noise curves~\\cite{shoemaker_06_a}.} \\label{fig:spectrum} \\end{figure}  For the GW signals from the axisymmetric collapse and core-bounce phase we obtain peak amplitudes of up to $ h \\sim 10^{-20} $ at $ 10 \\mathrm{\\ kpc} $, while the nonaxisymmetric structures in model E20A developing later emit GWs with $ h $ only a factor $ \\sim 5 $ smaller. However, since the latter emission process operates over several tens of ms, the total energy $ E_\\mathrm{gw} $ emitted in GWs is \\emph{larger} than that from the core-bounce signal. This is evident in Fig.~\\ref{fig:spectrum}, where we display the characteristic GW strain spectra $ h_\\mathrm{char} = R^{-1} \\sqrt{2 \\pi^{-2} dE_\\mathrm{gw} / df} $~\\cite{ott_04_a} for all models, evolving E20A for $ 70 \\mathrm{\\ ms} $ after bounce (see also Table~\\ref{tab:model_summary}). Considering only the core-bounce waveforms, $ h_\\mathrm{char} $ has its maximum between $ 300 $ and $ 800 \\mathrm{\\ Hz} $, while for model E20A it peaks at $ \\sim 930 \\mathrm{\\ Hz} $, the frequency of the GW-emitting component of its nonaxisymmetric structures. We conclude that the core-bounce GW signals of all models investigated here may be detectable by current and future LIGO-class detectors from anywhere in the Milky Way. Models that develop nonaxisymmetric instabilities may be detectable out to much larger distances if the instability persists for a sufficiently long time.  We point out that due to the nature of the approximation used for the neutrino effects in this study, we can only accurately model the GW emission in the collapse, bounce, and early postbounce epoch of the core-collapse supernova scenario, but not much later than the neutrino burst at shock breakout a few ms after bounce. In future work we plan to improve upon this and carry out longer-term postbounce evolutions, where additional GW emission mechanisms may play an important role~\\cite{mueller_04_a,ott_06_a}.   We thank A.~Burrows, E.~M\\\"uller, S.~Ou, L.~Rezzolla, E.~Seidel, D.~Shoemaker, N.~Stergioulas, and J.~Tohline for help and stimulating discussions.  This research was partially supported by the DFG (SFB/Transregio~7) and by the NCSA under grant No.~AST050022N.   \\vspace*{-.6cm}"
    },
    "0609/astro-ph0609028_arXiv.txt": {
        "abstract": "Possible explanations of the observed accelerated expansion of the Universe are the introduction of a dark energy component or the modifications of gravity at large distances. A particular difference between these scenarios is the dynamics of the growth of structures. The redshift distribution of galaxy clusters will probe this growth of structures with large precision. Here we will investigate how proposed galaxy cluster surveys will allow one to distinguish the modified gravity scenarios from dark energy models. We find that cluster counts can distinguish the Dvali-Gabadadze-Porrati model from a dark energy model, which has the same background evolution, as long as the amplitude of the primordial power spectrum is constrained by a CMB experiment like Planck. In order to achieve this, only a couple of hundred clusters in bins of width $\\Delta z = 0.1$ are required. This should be easily achievable with forthcoming Sunyaev-Zel'dovich cluster counts, such as the South Pole Telescope in conjunction with the Dark Energy Survey. ",
        "introduction": "During the course of the last ten years evidence for an accelerated expansion of the Universe has been mounting. Initially driven by Type Ia Supernovae (SNe) observations by the Supernovae Cosmology Project (SCP) and the High-z Supernovae Search Team \\cite{Perlmutter:97,Perlmutter98,Riess:98}, the combination of large scale structure surveys like the 2 degree field (2dF) and the Sloan Digital Sky Survey (SDSS) \\cite{Percival01,Abazajian:03}, cosmic microwave background (CMB) \\cite{Spergel03,Spergel:06} and x-ray observations of clusters of galaxies \\cite{Allen:02,Rapetti:04} confirmed these findings. In recent years even larger samples of Type Ia SNe \\cite{Knop:03,Riess:04,Astier:06} indicate strongly that the expansion of the Universe is speeding up. While the observations are getting better and better, theoretical models of accelerated expansions are still in its infancy. The most natural candidate seems to be a cosmological constant, however this comes with a drawback of fine tuning of the initial conditions to unnatural values.  In general there are three possible approaches to obtain accelerated expansion of the Universe. Firstly there could be an additional component in the energy-momentum tensor, which dominates the late time dynamics of the Universe and leads to accelerated expansion. This component requires negative pressure and an example of such a fluid could be a nearly homogeneous scalar field prevailing in the Universe \\cite{Wetterich:88, Peebles:88, Ratra:88, Ferreira:98, Zlatev:98}. The second possibility, is that Einstein's equations of gravity require modification on large scales \\cite{Dvali:00,Carroll:04,Capozziello:05,Carroll:05,Lue:06}. The third possible explanation is that the Universe is inhomogeneous on large scale, although it is hard to explain the boundary conditions for such a scenario in a natural way \\cite{Kolb:05}. In this article we concentrate on the first two possibilities.\\\\  Numerous dark energy models have been suggested in the past, where canonical scalar fields with various potentials are only a sub-class \\cite{Wetterich:88, Peebles:88, Ratra:88,Ferreira:98,Zlatev:98}. In all these models the background evolution can be described by the equation of state, i.e. the relation between the pressure and density of the dark energy fluid. The standard dark energy models also have a standard evolution of the growth of structures. Theories of modified gravity are only emerging as possible explanations for the observed acceleration. One of the most prominent models is one, which is motivated by brane cosmology involving extra dimensions and hence lead to an effectively modified Friedman equation on the 3+1 dimensional brane \\cite{Dvali:00,Lue:06}. Another possibility is to have inverse curvature terms in the Lagrangian of gravity. The additional terms become relevant at the late time evolution of the Universe and can lead to accelerated expansion and hence explain the Supernovae data \\cite{Carroll:04, Capozziello:05, Mena:06}. When modifying gravity, one has to be very careful so as not to violate high precision tests of gravity in the local Universe, like the Solar System, and also not to introduce unphysical features in the theory, like negative energy eigenstates, namely ghosts, or super-luminal modes \\cite{Luty:04,Nicolis:05,Koyama:05,Lue:06,Chiba:03,Soussa:04,Navarro:05,Navarro:06,DeFelice:06}. For the rest of this article we will concentrate on the Dvali-Gabadadze-Porrati (DGP) model as an example for a modified gravity model. The main reason for this is that in this model the dynamics of linear perturbations has been worked out \\cite{Song:05,Koyama:06a,Koyama:06}. Although, some care is required on large scales \\cite{Koyama:06,Sawicki:06}, this is not important for the cluster count test, which probes power spectra on sub-horizon scales. ",
        "conclusions": "We have shown that the background evolution of the DGP model can be entirely mimicked by a dark energy model with a suitably chosen equation of state, given in Eqn.~\\ref{eqn:wa}. Due to this fact Type Ia SNe alone can not distinguish modified gravity models from dark energy. However, due to the different time evolution of the linear perturbations, galaxy cluster counts can distinguish the DGP model from a dark energy model with the same background evolution. The difference between these two models is more then at the $2\\sigma$ level, even for relatively low number of clusters $\\approx 7000$. This result is only slightly weakened if we include parameter degeneracies in the mimic model. We have presented a very simplified analysis in the article, but the striking difference between the mimic dark energy and the DGP model, raises hopes that this result also stands up in a more careful analysis, which for example includes uncertainties in the gas physics of clusters in order to establish the mass limit \\cite{Younger:06}.  A further interesting question would be to see whether our findings also hold up for other modified gravity scenarios, like the inverse curvature model \\cite{Carroll:04}. In order to answer this question a perturbative analysis of these models is required, which is not an easy task. It would be very surprising if the standard perturbation result holds for these models. In general we would expect that modified gravity models lead to a modification of the growth of structure. Clusters of galaxies in combination with a geometrical probe, like SNe distances, could provide a uniquely sensitive window for the distinction between general modifications of gravity and dark energy.  {\\bf Acknowledgement} We thank Filipe Abdalla and Richard Battye for useful discussions."
    },
    "0609/astro-ph0609502_arXiv.txt": {
        "abstract": "We present Keck/LRIS spectra of a candidate damped Ly$\\alpha$ (DLA) galaxy toward the QSO 3C196 (z$_{em}$ = 0.871). The DLA absorption system has a redshift of z$_{DLA}$ = 0.437, and a galaxy at 1.5$''$ from the QSO has been identified in high resolution imaging with WFPC2/HST. We have detected emission lines of [O~II] $\\lambda$3727, H$\\beta$, [O III] $\\lambda$5007, H$\\alpha$ and [N II] $\\lambda$6584 at the absorption redshift. Based on the emission lines, we have found the redshift of the galaxy to be z$_{em}$ = 0.4376 $\\pm$ 0.0006. The emission lines also enabled us to calculate the extinction-corrected luminosities and metallicity indicators using established indices based on line strengths of different emission lines. These indicators suggest that the ISM of the DLA galaxy has a high metallicity comparable to or perhaps twice as much as solar (e.g. 12+ $\\log$(O/H) = 8.98 $\\pm$ 0.07). Based on the strengths of H$\\alpha$ and on the reddening derived from the relative strengths of H$\\alpha$ and H$\\beta$, the star formation rate is 4.7 $\\pm$ 0.8 M$_{\\sun}$ yr$^{-1}$. This places the galaxy in the range of gas-rich spiral galaxies. ",
        "introduction": "{\\label{intro}}  Damped Lyman-alpha (DLA) absorption systems, observed in spectra of quasars (QSOs), provide a powerful tracer of the chemical history of high redshift galaxies (Pettini {\\it et al.} 1997; Prochaska \\& Wolfe 2002) and represent a unique sample for studying the interstellar medium (ISM) of distant galaxies. These systems contain mostly neutral gas, with neutral hydrogen column densities $\\log$ N$_{\\rm HI}$ $\\geq$ 20.3, and thus provide an important link in our understanding of the star formation history of the Universe. DLAs remain the only class of galaxies at cosmologically significant redshifts with detailed measurements of element abundances. Unfortunately, the DLA host galaxies are faint and located at small projected distances from the QSOs. The background  QSOs are much brighter than the DLA and thus can contaminate the background at the position of the DLA, making spectroscopic studies of the DLA galaxies a challenging task.  Recent studies combining high angular resolution images and high SNR spectra are providing key evidence revealing the nature of DLAs. So far, most imaging studies of high-z and low-z DLAs have failed to detect large luminous disks (e.g. Jimenez {\\it et al.} 1999; Kulkarni {\\it et al.} 2000, 2001; Chun {\\it et al.} 2006 and references therein). The lack of substantial chemical evolution found in studies of element abundances in DLAs (e.g. Kulkarni {\\it et al.} 2005; P\\'eroux {\\it et al.} 2006 and references therein) suggests that the currently known population of DLAs seems to be dominated by metal-poor objects. Thus DLAs may consist of dwarf or low-surface brightness galaxies with modest star formation rate (SFR). Indeed, most emission-line imaging searches for DLAs also indicate low SFR (e.g. Kulkarni {\\it et al.} 2006 and references therein). It is possible however that the metal-rich DLAs are systematically underrepresented because of dust obscuration of background QSOs (e.g. Fall \\& Pei 1993; Boisse {\\it et al.} 1998).  SFRs can be estimated with several different methods. Emission line imaging provides a direct method to determine the SFR in low-z DLAs, although it is subject to the possibility of dust extinction. An alternative method based on CII$^*$ absorption lines has been suggested by \\cite{wolf03a}. However this method is less direct, since it depends on several assumptions. Furthermore, the CII$^*$ method is not accessibble for low-z DLAs, since the CII$^*$ absorption lines fall in the UV. SFRs determined from optical emission lines have been extensively studied in nearby galaxies (e.g. Kennicutt 1998; Kormendy \\& Kennicutt 2004) and DLAs (Chen {\\it et al.} 2005; Kulkarni {\\it et al.} 2006 and references therein). The hydrogen recombination lines are directly related to the UV emission of the shortest-lived stars and thus measuring H$\\alpha$ in the optical wavelength range provides a good tool to estimate the SFR. Assuming that the shape of the initial mass function is known, the main uncertainty in the measurement of optical emission line luminosities is extinction. However we can estimate the amount of extinction by measuring the Balmer decrement (e.g. Pei 1992; Cardelli {\\it et al.} 1989).  So far, SFR estimates for most of the absorbers appear to be small: 63\\% of the detections and about 73\\% of the limits in the sample studied by \\cite{kulk06} show SFR $<$ 5 M$_{\\sun}$ yr$^{-1}$. This can be explained by one of the following possible cases: 1) The SFRs in DLAs are in fact low; 2) dust extinction is high preventing the line  emission to escape from the HII region; 3) the star-forming regions are too close to the QSO line-of-sight masking the true SFRs. In order to distinguish the dominant effect, more emission line imaging studies should be undertaken.  We have carried out a spectroscopic study of a spiral galaxy (galaxy \\# 4 from Le Brun {\\it et al.} 1997) 1.5$''$ away from 3C196 (Q0809+483). This QSO has a 21 cm absorption line (Brown \\& Mitchell 1983) as well as a DLA system at z$_{abs}$ = 0.4368 with $\\log$ N$_{\\rm HI}$ = 20.8 $\\pm$ 0.2 (Boiss\\'e {\\it et al.} 1998). The spiral galaxy with AB(F702)\\footnote[1]{AB magnitude will be used throughout the paper. AB(F702) refers to the AB magnitude in the HST/WFPC2 filter F702W ($\\bar{\\lambda }$ = 6940.1 \\AA).}= 19.9 has been reported as a candidate host galaxy for the DLA (Le Brun {\\it et al.} 1997). The galaxy was suggested to be responsible for the DLA in the spectroscopic program of \\cite{chen05}, based on the detection of H$\\beta$ and [O~II] emission lines in the spectra of the galaxy at z = 0.43745 $\\approx$ 138 km s$^{-1}$ away from the redshift of the DLA. They also noted a narrow Ca II H and K absorption doublet in the spectrum of the background QSO at 21 km s$^{-1}$ blueshifted from the systematic redshift of the galaxy. No additional emission features had been reported prior to the present work.  HST/WFPC2 images showed that galaxy \\# 4 has the morphology of a barred spiral galaxy with B$-$R = 2.5 (Le Brun {\\it et al.} 1997; Boiss\\'e \\& Boulade 1990). Typical B$-$R colors of normal galaxies at z = 0.5 range from 3.2 for ellipticals to 1.2 for irregular galaxies with B$-$R = 2.1 for Sbc galaxies (Fukugita {\\it et al.} 1995). Galaxy \\# 4 is thus comparable in B$-$R (though possibly little redder than) to a typical Sbc galaxy at z = 0.5. \\cite{lebr97} also detected a compact galaxy (galaxy \\# 3) 1.1$''$ away from 3C196. This galaxy has B$-$R = 2.3 (Boiss\\'e \\& Boulade 1990), which is comparable to the typical B$-$R of a Sbc galaxy at z = 0.4 or z = 0.9. We note however that \\cite{lebr97} have suggested that galaxy \\# 3 may be associated with the QSO. In this paper, we study the spiral galaxy \\# 4, and put limits on the line emissions for galaxy \\# 3.  The paper is organized as follows. The observations and data reduction are described in \\S \\ref{obs+red}. The results are presented in \\S \\ref{res} and discussed in \\S \\ref{dis}. Throughout this paper we assume $\\Omega_{\\Lambda}$ = 0.7, $\\Omega_{m}$ = 0.3, and h = 0.7.\\\\ ",
        "conclusions": "{\\label{dis}} Table \\ref{tab:sum} summarizes the estimated properties of the DLA galaxy. It is instructive to compare these properties with those of other galaxies and DLAs.  \\subsection{Comparison with other galaxies} We obtained an oxygen abundance of 12 + $\\log$(O/H) = 8.98 $\\pm$ 0.07\\footnote[3]{All the metallicity values we use in this section for the DLA galaxy and other galaxies, are based on the R$_{23}$ index.} for this galaxy, for which the absolute magnitude has been estimated to be M$_B$(AB) = -21.1 (Chen {\\it et al.} 2005). Figure \\ref{fig:f4} shows the luminosity-metallicity relation for a large sample of various galaxies at z $<$ 0.25 (Tremonti {\\it et al.} 2004). It is clear that the DLA galaxy falls close to the luminosity-metallicity relation of local spiral galaxies (Kobulnicky \\& Zaritsky 1999; Zaritsky {\\it et al.} 1994).  Recently \\cite{sava05} investigated the metallicity of 69 galaxies at 0.4 $<$ z $< 1$ from the Gemini Deep Deep Survey and the Canada-France Redshift Survey. Using the R$_{23}$ metallicity indicator, they found the mean metallicity to be 12 + $\\log$(O/H) = 8.78 $\\pm$ 0.17, 0.05 dex higher than the value reported by \\cite{kk04}. These surveys included all types of galaxies and did not show any redshift evolution. Our result from the R$_{23}$ indicator that the DLA galaxy has 12 + $\\log$(O/H) = 8.98 $\\pm$ 0.07 suggests that the DLA is a metal-rich galaxy and its metallicity lies at the high end of the metallicity distribution in \\cite{sava05}.  The SFRs in various types of galaxies have been derived and discussed by several authors. These data show an enormous range, from virtually zero in gas-poor elliptical, S0, and dwarf galaxies to 10-20 M$_{\\sun}$ yr$^{-1}$ in giant Sc galaxies. Much larger SFRs, up to  1000 M$_{\\sun}$ yr$^{-1}$, can be found in starburst galaxies. The highest SFRs are associated with strong tidal interactions and mergers. SFRs of local blue compact galaxies, estimated by \\cite{petr97}, are about 0.3-0.5 M$_{\\sun}$ yr$^{-1}$. For LSB galaxies, \\cite{vand00} estimated SFRs of about 0.03-0.2 M$_{\\sun}$ yr$^{-1}$. For dwarf and LSB galaxies, \\cite{vanz01} estimated SFRs from 10$^{-5}$ to 10$^{-1}$ M$_{\\sun}$ yr$^{-1}$.  In a study of star formation vs. galaxy morphology for 334 spiral and irregular nearby galaxies, \\cite{jame04} found that the most strongly star-forming normal galaxies are those of Hubble types Sbc and Sc, with a wide range in SFR, up to 15 M$_{\\sun}$ yr$^{-1}$. For local Sbc galaxies, the derived mean SFR is 2.4 $\\pm$ 0.4 M$_{\\sun}$ yr$^{-1}$ (see Figure \\ref{fig:f5}) which is close to the mean calculated by \\cite{kenn83} 3.1 $\\pm$ 0.4 M$_{\\sun}$ yr$^{-1}$ for this type\\footnote[4]{Converted to the adopted Hubble constant of H$_o$ = 75 km s$^{-1}$ Mpc$^{-1}$ by \\cite{jame04}.}. Although many of the galaxies in their survey appeared to have high SFRs, the derived mean is rather low, since very low SFRs are common for all Hubble types.  The SFR distribution of a sample of 168 low-mass IR-selected star-forming galaxies out to z=1.5 was explored by \\cite{baue05}. In the redshift bin 0.4 $<$ z $<$ 1.0, they estimated SFR from 0.04 to 10 M$_{\\sun}$ yr$^{-1}$, based on the [O II] $\\lambda$3727 line. Their result suggests that the mean and the maximum SFR in each redshift bin increases with redshift.  Our result for the SFR of the DLA galaxy ($\\approx$ 4.7 $\\pm$ 0.8 M$_{\\sun}$ yr$^{-1}$ for the SMC type extinction correction and 5.4 $\\pm$ 1.0 M$_{\\sun}$ yr$^{-1}$ for the MW type) agrees well with the expected SFR values for local Sbc galaxies and IR-selected star-forming galaxies in the redshift bin 0.4 $<$ z $<$ 1.0. This DLA appears to be a large galaxy with high SFR that lies at the high end of SFR distributions of \\cite{jame04} and \\cite{baue05}.  \\subsection{Comparison with other DLAs and DLA models} Metallicity studies indicate that most DLAs have lower mean metallicity than samples of other types of galaxies (Savaglio {\\it et al.} 2005), and that they show little if any chemical evolution (e.g., Kulkarni {\\it et al.} 2005; P\\'eroux {\\it et al.} 2006). The N$_{\\rm HI}$-weighted mean Zn metallicity in the redshift bin of 0.1 $\\lesssim$ z $\\lesssim$ 1.4 reported by Kulkarni {\\it et al.} 2005 is -0.94 $\\pm$ 0.16. This would imply 12 + $\\log$(O/H) = 7.75 $\\pm$ 0.16, if we assume the Zn/O in DLAs is solar. Our DLA galaxy with 12 + $\\log$(O/H) = 8.98$\\pm$ 0.07 is much more metal-rich than the N$_{\\rm HI}$-weighted mean.  The star formation rates of DLAs are not well-understood. Emission line imaging searches for Ly$\\alpha$, H$\\alpha$, H$\\beta$, [O II], and [O III] suggest low SFRs in a large fraction of absorption-selected galaxies (e.g., Kulkarni {\\it et al.} 2006; see Figure \\ref{fig:f6}). \\cite{weat05} estimated SFRs from 9.5 to 28 M$_{\\sun}$ yr$^{-1}$ from [O III] $\\lambda$5007 emission in three DLAs at redshifts 1.9 $<$ z $<$ 3.2. Also, Wolfe {\\it et al.} (2003a, 2003b) using the C II* absorption line method, estimated the SFRs per unit area of about 30 DLAs at z $\\gtrsim$ 2 to be in the range 10$^{-3}$ to 10$^{-1}$ M$_{\\sun}$ yr$^{-1}$ kpc$^{-2}$. These authors suggested that at high-redshift, DLAs are similar to the population of Lyman break galaxies. On the other hand, they noted that such large SFRs are inconsistent with the low metallicities seen in high-redshift DLAs. Indeed, the results of the C II* method are somewhat model dependent.  In fact, \\cite{hopk05}, using N$_{\\rm HI}$ to estimate the SFR suggested that DLAs have a dominant role in the SFR density only in the late universe (z $<$ 0.6), while at higher redshift their contribution is much less important. They suggested that the DLA population may be dominated by dwarf-like systems with low average SFRs or a late onset of star formation; in other words the DLAs identified so far may not account for the majority of the neutral gas at high redshift. However the assumption of \\cite{hopk05} that the local SFR vs. surface density relation holds for DLAs has not been verified.  Recently, \\cite{okos05} have investigated models for DLA absorbing galaxies at z $<$ 1, using a semi-analytic galaxy formation model. They predicted a broad SFR range spanning from 10$^{-6}$ to 10$^2$ M$_{\\sun}$ yr$^{-1}$ and a mean value of 10$^{-2}$ M$_{\\sun}$ yr$^{-1}$ for DLAs with z $<$ 1. Their study predicted that DLA galaxies brighter than L/L* $\\sim$ 0.5 may have SFRs larger than 10 M$_{\\sun}$ yr$^{-1}$. They concluded that most of DLAs at z $<$ 1 are LSB dwarf galaxies with low SFRs, although some are massive spiral galaxies.  Our SFR estimation of 4.7 M$_{\\sun}$ yr$^{-1}$ for the DLA galaxy toward 3C196 lies within the broad SFR range predicted by \\cite{okos05} model for DLA absorbing galaxies at z $<$ 1. However, the value is higher than that for a large fraction of DLAs: 63\\% of the detections and about 73\\% of the limits in the SFR sample studied by \\cite{kulk06} have SFRs below 5 M$_{\\sun}$ yr$^{-1}$. It is interesting to note that our value for the disk galaxy toward 3C196 agrees well with the ``large disk'' calculations of \\cite{bunk99} (see Figure \\ref{fig:f6})."
    },
    "0609/astro-ph0609787_arXiv.txt": {
        "abstract": "{The origin of the spiral pattern of magnetic fields in disc galaxies is an open question. Comparison of the regular magnetic field orientation with the gaseous spiral arm pitch angles can tell us whether spiral shock compression is responsible for the magnetic spirals. We also wish to see whether the ridges of different components of the ISM show the large-scale, systematic shifts expected from density wave theory. We have developed a technique of isolating elongated structures in galactic images, such as spiral arms, using anisotropic wavelets and apply this to maps of the CO, infrared and radio continuum emission of the grand-design spiral galaxy M51. Systematic shifts between the ridges of CO, infrared and radio continuum emission that are several $\\kpc$ long are identified, as well as large variations in pitch angle along spiral arms, of a few tens of degrees. We find two types of arms of polarized radio emission: one has a ridge close to the ridge of CO, with similar pitch angles for the CO and polarization spirals and the regular magnetic field; the other does not always coincide with the CO arm and its pitch angle differs from the orientation of its regular magnetic field. The offsets between ridges of regular magnetic field, dense gas and warm dust are compatible with the sequence expected from spiral density wave triggered star formation, with a delay of a few tens of millions of years between gas entering the shock and the formation of giant molecular clouds and a similar interval between the formation of the clouds and the emergence of young star clusters. At the position of the CO arms the orientation of the regular magnetic field is the same as the pitch angle of the spiral arm, but away from the gaseous arms the orientation of the regular field varies significantly. Spiral shock compression can explain the generation of one type of arm of strong polarized radio emission but a different mechanism is probably responsible for a second type of polarization arm. ",
        "introduction": "\\label{sec:intro}  Disc galaxies often display spiral patterns in their distributions of stars, gas, dust and magnetic fields. Spiral structure is present in both the \\emph{distribution} of the radio synchrotron intensity, and the \\emph{orientation} of the $\\kpc$-scale regular field, as measured by the $B$-vectors of radio polarization (e.g. Beck\\ \\cite{Beck05}). Comparing the location and orientation of the spiral patterns in different components of the interstellar medium (ISM) can provide important information on the astrophysical connections between interstellar gas, dust and magnetic fields in galaxies.  Usually the prominent arms in disc galaxies are treated as having a logarithmic spiral pattern, where the pitch angle of the arm is constant along its length. However, Kennicutt (\\cite{Kennicutt81}) found that the logarithmic spiral form is no better a mathematical description of real spirals than hyperbolic forms, and that ``\\dots small scale distortions preclude the possibility of \\emph{any} universal shape for galactic spirals''. Deviations from the logarithmic spiral may be caused by local disturbances or arise from systematic global effects, such as the gravitational force of a nearby galaxy or the presence of more than one spiral density wave (Elmegreen et al.~\\cite{Elmegreen89}). If the regular magnetic field spirals are being aligned with the gaseous spiral arms by shock compression, the magnetic field orientation should be closely related to the local spiral arm pitch angle. In order to quantify how well the different spirals are aligned a robust method of measuring local pitch angles is required.  M51 is probably perturbed by a recent encounter with its companion galaxy NGC~5195. Such interactions often result in increased star formation rates, either localised or global, as tidal forces and spiral density waves compress the interstellar medium. The two spiral arms of M51 can be traced through more than 360\\degr\\ in azimuthal angle in numerous wavebands and several authors have investigated their structure. Elmegreen et al. (\\cite{Elmegreen89}) de-projected images taken in the optical B and I bands into the log(radius)--azimuth plane and used the observed amplitude variations along the arms to locate resonances of spiral modes. They concluded that M51 contains an inner and an outer system of spiral arms with a conjunction of resonances at 2\\arcmin\\,-- 3\\arcmin\\ galacto-centric radius, triggered by the companion. Measurements of the pitch angles of the main dust lanes by Howard \\& Byrd (\\cite{Howard90}) and of the variation in star formation efficiency along the arms by Knapen et al. (\\cite{Knapen92}) confirmed this double spiral structure. It is clear from a cursory visual inspection that even the well defined --- grand design --- spiral arms of M51 have many variations in pitch angle along their length.  Radial shifts between the locations of the crest of the spiral arms in different constituents are expected from density wave theory. Petit et al. (\\cite{Petit96}) observed a 300-400~pc shift between H$\\alpha$ and UV arms. Rand \\& Kulkarni (\\cite{Rand90}), Scoville et al. (\\cite{Scoville01}), and Tosaki et al. (\\cite{Tosaki02}) noted that the H$\\alpha$ arms are located outside the CO arms. Henry et al. (\\cite{Henry03}) found that the Paschen-$\\alpha$ line emission is offset from the CO emission along part of the arms. All these shifts refer to the inner spiral system and are expected for arms triggered by density waves. Tilanus et al. (\\cite{Tilanus88}) found that thermal radio continuum and H$\\alpha$ arms are significantly shifted outwards from the dust lanes, whereas the non-thermal radio continuum arms occur just inside the dust lanes.  Most of the above analyses were made by overlaying the contours of one image on a grey-scale map of the other or by making azimuthal cuts through the data at different radii. Both methods, while showing the most significant local displacements in the position of the arms, make it difficult to see the extent along the arm of any systematic shift. In the present paper we describe an objective method, using wavelet analysis, to determine the location of the spiral arm ridges, and measure their local pitch angles, along the entire length of an arm. We apply this method to images of the total and polarized radio continuum emission at \\wav{6} (Fletcher et al. in preparation), the dust emission at 15$\\mum$ (Sauvage et al.~\\cite{Sauvage96}) and the CO(1--0) line emission (Helfer et al.~\\cite{Helfer03}).  In Section~\\ref{sec:method} we describe the anisotropic wavelet transform and how we use this method to identify the location of a spiral arm and to measure its localised pitch angles. The observational data we use are briefly discussed in Section~\\ref{sec:data} and the results of applying the anisotropic wavelet method to the data are given in Section~\\ref{sec:results}. The astrophysical implications of the results are discussed in Section~\\ref{sec:discuss}. ",
        "conclusions": "\\label{sec:discuss}  New observations of the spiral galaxy M51, obtained with the VLA and Effelsberg radio telescopes, and published maps of the molecular (CO) and mid-infrared (15$\\mu$m) emissions, were used to compare the spiral patterns present in the magnetic field, dense gas and dust distributions.  1. We have shown that anisotropic wavelet functions are appropriate tools to locate the positions of spiral arms in images of galaxies and to measure their local pitch angles.  2. A systematic offset between the spiral ridges of molecular gas (CO) emission and of mid-infrared (ISO 15$\\mum$) emission can be followed over many kpc along the arms of M51 (Figs.~\\ref{fig:ridges} \\& \\ref{fig:log_ridges}). The typical offset is $100\\pc$ corresponding to a typical time delay of $\\sim 10^{7}\\yr$ for gas well inside the co-rotation radius, assuming a circular rotation velocity of $200\\kms$ (Garcia-Burillo et al. \\cite{GarciaBurillo93}), pattern speed of $38\\kms\\kpc^{-1}$ (Zimmer et al. \\cite{Zimmer04}) and arm inclination of $20\\degr$. The delay can be interpreted as the interval between the formation of Giant Molecular Clouds (GMCs) and the appearance of newly formed stars. A delay of $\\sim 10^{7}\\yr$ between dense gas accumulation and star formation was also estimated by Tosaki et al. (\\cite{Tosaki02}) from an analysis of the location of $^{12}$CO and $^{13}$CO emission lines.  We do not see a shift in the relative positions of the ISO and CO arms, i.e. the CO arm does not move from the concave to the convex side of the ISO arm, near to the co-rotation radius of $r\\simeq 6.2\\kpc$ determined by Elmegreen et al.~(\\cite{Elmegreen92}). This may be due to the non-trivial dynamical effects of two, overlapping, spiral arm patterns in M51 (Elmegreen et al.~\\cite{Elmegreen89}).  3. We also found a systematic offset between the spiral ridges of total radio continuum emission and those of mid-infrared and CO emission, indicating a shift between regions of strong thermal and non-thermal radio continuum emission. While the thermal component is expected to closely follow the distribution of star-forming regions as traced by mid-infrared emission, the non-thermal component is probably enhanced due to compression of magnetic fields in a spiral shock front, located upstream of the mid-infrared spiral arm. Cosmic ray electron density is expected to be higher on the downstream side of the arms, where supernova remnants should be more common, but on scales of a few hundred parsecs this may not have a significant effect on the synchrotron intensity; we are assuming that the distribution of cosmic ray electrons is mainly determined by the magnetic field, rather than the distribution of their sources. The purely non-thermal, polarized radio ridges lie the furthest upstream of all the data as expected in this scenario. This is in agreement with the results of Tilanus et al. (\\cite{Tilanus88}) who found a general arm sequence of: non-thermal emission -- (cold) dust lane -- thermal emission -- H$\\alpha$.  Inside co-rotation our observations reveal the following sequence from upstream to downstream: polarized radio emission, total radio emission, CO, infrared emission, interpreted as a result of the sequence of components: shock-compressed magnetic field, molecular gas clouds, and UV heated (warm) dust. This can arise if the spiral shock is located slightly ahead of the CO ridge, a plausible consideration if molecule formation and/or cloud condensation are \\emph{triggered} by the spiral shock. The compression of magnetic field, resulting in stronger polarized radio emission, occurs at the shock, but the formation of dense CO clouds is only completed a few hundred parsec downstream. Following similar arguments to those used in point 1 above, we can estimate the time delay between gas entering the shock front and the accumulation of dense GMCs: a typical polarized ridge offset of $300\\pc$ corresponds to an interval of $\\sim 10^{7}\\yr$. The total time between gas entering the spiral shock and the emergence of newly formed stars will then be a few tens of millions of years.  The above discussion only applies to a diffuse component of the CO in the inter-arm region, which can be compressed by a large scale shock at the spiral arm. Molecular gas that is already gravitationally bound in clouds will not respond to the shock. Extended, faint CO(1-0) emission in the inner inter-arm region of M51 is clearly seen in the combined single-dish and interferometer data of Helfer et al.~(\\cite{Helfer03}), suggesting that at least a part of the molecular gas is diffuse. However the higher resolution, interferometer only, CO(1-0) maps of Aalto et al.~(\\cite{Aalto99}) reveal small clumps of inter-arm molecular gas, so that some --- difficult to estimate --- fraction of the upstream gas is probably already confined in clouds. This means that the shock strength will be difficult to infer from gas density ratios in the inter-arm and arm regions. On the other hand, Aalto et al. could identify steep velocity gradients both in and, particularly clearly, between giant associations of molecular gas which are in broad agreement with theoretical expectations of spiral density wave shocks. The detailed physics of the gas response to possible shocks in M51 and the effect this has on the spiral arm morphology requires more sophisticated modelling than we can attempt here, especially the apparent upstream shift of the shock in the magnetic field; there are, however, strong indications that large scale spiral shocks occur in the vicinity of the gas arms and that these compress diffuse gas.  4. The maintenance of both radial and azimuthal components of the magnetic field, and hence a non-zero pitch angle, is an important prediction of dynamo theory (Beck et al. \\cite{Beck96}, Shukurov\\ \\cite{Shukurov00}) but is very difficult to explain if the magnetic field is purely passive and stretched by shear in the galactic differential rotation, since after a small number of galactic rotations the field would be completely circular, i.e. it would have a pitch angle $\\pB\\simeq 0$.  5. Although the B-vectors of polarized radio emission from M51 seem to have the same orientations as the CO arms (Fig.~\\ref{fig:orig_maps}), they smoothly change by about $\\pm15\\degr$ around any fixed radius, without a systematic trend between the arm and inter-arm regions (Fig.~\\ref{fig:co+psi}). Smaller pitch angles in inter-arm regions are predicted by some dynamo models and have been observed e.g. in NGC~6946 (Rohde et al.~\\cite{Rohde99}). In M51, however, the regular magnetic field structure is probably dominated by streaming motions and compression in the spiral shock (see below). The field is brought into good alignment at the CO arms \\emph{or vice versa} (Fig.~\\ref{fig:co+psi}) and then relaxes into a different configuration in the inter-arm regions that depends on e.g. the amount of compression in the arms, the energy of turbulent flows, the orientation of dynamo-generated field modes etc.  6. We found a close alignment of the CO pitch angles and the regular magnetic field orientation \\emph{at the position of the arms} (Fig.~\\ref{fig:co+psi}). This could be due to shock amplification of the regular field component parallel to the shock.  If the field is frozen into the gas then $\\Brpa^{(s)}=\\epsilon\\Brpa$ and $\\Brpe^{(s)}=\\Brpe$, where $\\epsilon=\\rho^{(s)}/\\rho$ is the compression ratio in the gas density $\\rho$ and the superscript $(s)$ labels quantities at the shock front. The deflection of the magnetic field $\\Delta\\theta$, where $\\theta$ is the angle between the field and the shock front upstream of the shock front, is given by \\begin{eqnarray} \\label{eq:align} \\Delta\\theta&=&\\theta-\\theta^{(s)}=\\arctan\\left(\\frac{\\Brpe}{\\Brpa}\\right)- \\arctan\\left(\\frac{\\Brpe}{\\epsilon\\Brpa}\\right) \\nonumber \\\\ & \\simeq & \\frac{\\Brpe}{\\Brpa}\\left(1-\\frac{1}{\\epsilon}\\right) \\end{eqnarray} where the last line holds for $\\Brpe/\\Brpa\\ll 1$. For $\\epsilon=4$ (maximum degree of gas compression by an adiabatic shock) and $\\Brpe/\\Brpa=0.5$ (i.e. the magnetic field is initially inclined at $30\\degr$ to the shock front, a conservative estimate c.f. Fig.~\\ref{fig:co+psi}) we obtain $\\Delta\\theta\\simeq 20\\degr$. Thus a strong spiral shock can be expected to align the regular magnetic field rather well with the spiral arm.  If the shock is weak ($\\epsilon<4$) or the upstream field makes a larger angle with the shock front then an additional source of alignment is required. The conversion of an isotropic random magnetic field component upstream of the shock into an anisotropic random field in the shocked region (where the anisotropic random field will be perfectly aligned with the shock) provides such a mechanism. Strong polarized emission arising from anisotropic random magnetic fields formed in shearing shocks has recently been identified in the barred galaxies NGC~1097 and NGC~1365 (Beck et al.~\\cite{Beck05a}).  7. Some of the polarized radio emission (and hence some of the regular magnetic fields) forms ridges which are interlaced with the gaseous arms: the polarized ridge of arm~2 clearly lies in between gas arms~1 and 2 south-east of the central region. The magnetic field is \\emph{not} oriented parallel to this polarization arm but is inclined to the axis of the arm by up to $20\\degr$ (Fig.~\\ref{fig:pi+psi}). In contrast to polarization arm~1, spiral shock compression does not seem to be the origin of these enhanced regular magnetic fields. The field orientation seems to be between that of the polarization arm and that of the nearby CO arm~2, like finding a compromise between two different forces. One force causes the polarization arm to cross CO arm~2 to join the outer arm~1, the second tries to align the field with the gas arm."
    },
    "0609/astro-ph0609578_arXiv.txt": {
        "abstract": "We discuss the prospects of using Gamma Ray Bursts (GRBs) as high-redshift distance estimators, and consider their use in the study of two dark energy models, the Generalized Chaplygin Gas (GCG), a model for the unification of dark energy and dark matter, and the XCDM model, a model where a generic dark energy fluid like component is described by the equation of state, $p= \\omega \\rho$. We find that this test yields rather disappointing results for the GCG model, being mainly sensitive to the total amount of matter present in the Universe in the case of the XCDM model. We also find that, within the framework of the XCDM model, a large sample of GRBs ($\\geq 200$) may turn out to be quite useful to improve the forthcoming type Ia supernovae data. ",
        "introduction": "Recently, there has been a great deal  of activity on attempts of using GRBs as cosmological probes~\\cite{bertolami06}. In the original proposal~\\cite{schaefer03}, it has been suggested that the magnitude versus redshift plot could be extended to redshifts up to $z\\simeq 4.5$, via correlations found between the isotropic equivalent luminosity, $L_{iso}$, and two GRB observables, namely the time lag ($\\tau_{lag}$)~\\cite{norris00} and variability ($V$)~\\cite{reichart01}. The isotropic equivalent luminosity is the inferred luminosity (energy emitted per unit of time) of a GRB if all its energy is radiated  isotropically, the time lag measures the time offset between high- and low-energy arriving GRB photons, while the variability is a measure of the complexity of the GRB light curve. Using these correlations, one can infer two estimates of the absolute isotropic equivalent luminosity, which are combined through a weighted average. Knowing the absolute isotropic equivalent luminosity together with the observed fluence yields an estimate of the luminosity distance to the GRB.  Unfortunately, these correlations are affected by a large statistical (or intrinsic) scatter. This statistical spread affects not only the cosmological precision via its direct statistical contribution to the distance modulus uncertainty, $\\sigma_\\mu$, but also through the calibration uncertainty given that the suitable GRB sample with known redshift is rather small. In what follows we show that a relatively small sample of GRBs with low redshifts is sufficient to greatly reduce the systematic uncertainty thanks to a more robust and precise calibration.  More recently, a new correlation has been suggested~\\cite{ghirlanda04a}, which is subjected to a much smaller statistical scatter. The so-called Ghirlanda relation is a correlation between the peak energy of the gamma-ray spectrum, $E_{peak}$ (in the $\\nu - \\nu F_{\\nu}$ plot), and the collimation-corrected energy emitted in gamma-rays, $E_\\gamma$. This collimation-corrected energy is a measure of the energy released by a GRB taking into account that it is beamed into a narrow jet. Unlike the $L_{iso}-\\tau$ and $L_{iso}-V$ relations, the Ghirlanda relation is not affected by large statistical uncertainties, however, it depends on poorly constrained quantities related to the properties of the medium around the burst. Indeed, to infer $E_\\gamma$ one must estimate the  angular opening of the jet, which can be performed assuming a density profile for the medium around the burst (or circum-burst medium for short), where a fraction $\\eta_\\gamma$ of the fireball kinetic energy is emitted in the prompt gamma-ray phase, and where one has measured the jet break time, $t_{jet}$ \\cite{sari99}. Assuming that the circum-burst medium has a constant density, the simplest possible assumption, requires one additional parameter. This constant density has been measured for a few bursts~\\cite{ghirlanda04a}, and it exhibits a wide variation from burst to burst.  Another difficulty involving GRBs is that they tend to occur at rather large distances, which makes it impossible to calibrate any relationship between the relevant variables in a way that is independent from the cosmological model. The method that is usually employed consist in fitting both, the cosmological \\emph{and} the calibration parameters, and then use statistical techniques to remove the undesired parameters. In here, we follow a different procedure \\cite{bertolami06,takahashi}. We consider that the luminosity distance for $z<1.5$ was previously measured using type Ia supernovae, and divide the GRBs sample in two sets; the low redshift sample, with  $z<1.5$, and the high redshift one, with $z>1.5$. Since the luminosity distance of GRBs in the range $z<1.5$ is already known, one can calibrate the luminosity estimators independently of the cosmological parameters and use the high redshift sample as a probe to dark energy and dark matter models. This method also allows us to verify whether the larger redshift range probed by GRBs can compensate for the larger uncertainty associated with the distance estimates thus obtained.  We have analyzed the use of these correlations in order to study of GCG, a model that unifies the dark energy and dark matter in a single fluid \\cite{bento02} through the equation of state $ p_{ch} = - A / \\rho_{ch}^\\alpha$, where $A$ and $\\alpha$ are positive constants. The case $\\alpha=1$ describes the the Chaplygin gas, that arises in different theoretical scenarios. If the curvature is fixed, there are only two free variables, $A$ and $\\alpha$, although it is more convenient to use the quantity $A_s \\equiv A/ \\rho_{ch,0}^{1+\\alpha}$ instead of $A$. Thus, we consider two free parameters, $\\alpha$ and $A_s$. A great deal of effort has been recently devoted to  constrain the GCG model parameters~\\cite{bertolami05}, which include, for instance, supernovae~\\cite{bertolami04,bento05b}, cosmic microwave background radiation~\\cite{bento03a,bento03b,bento03c}, gravitational lensing~\\cite{silva03} and cosmic topology~\\cite{bento05a}.  In addition to the GCG model, we also study the more conventional flat XCDM model. Likewise the GCG model, the XCDM model is also described by two free parameters, the  dark energy equation of state, $\\omega \\equiv p / \\rho$, and the fraction of non-relativistic matter, $\\Omega_m$. Testing these models is particularly relevant since they are degenerate for redshifts $z<1$~\\cite{bertolami04,bento05b}. ",
        "conclusions": ""
    },
    "0609/astro-ph0609322_arXiv.txt": {
        "abstract": "We present two-dimensional (2D) gas kinematics and excitation of the  inner 300\\,pc of the Seyfert galaxy ESO\\,428-G14 at a sampling of  14\\,pc$^2$, from near-infrared spectroscopic observations at  R$\\approx$6000 obtained with the Integral Field Unit of the Gemini Near-Infrared Spectrograph. From measurements of fluxes and  profiles of the emission lines [Fe\\,{\\sc ii}]$\\lambda 1.257 \\mu$m,  Pa$\\beta$, H$_2\\lambda 2.121 \\mu$m and Br$\\gamma$, we construct 2D maps of line intensities and ratios, radial velocities and velocity dispersions.  Emission lines ``tomography\" is provided by velocity slices obtained across the line profiles, a unique capability of IFUs, which allows the mapping of not only of peak velocities but  including also the wings. We compare these maps with a previously published high spatial resolution radio map and find a tight relation between the  radio structure and the emission-line flux distributions and  kinematics, revealing that the radio-jet plays a fundamental role not  only in shaping the NLR but also in the imprint of its kinematics. Blueshifts of up to 400\\,km\\,s$^{-1}$ and velocity dispersions of up  to 150\\,km\\,s$^{-1}$ are observed in association with the radio jet at  position angle PA\\,=\\,129$^\\circ$, which is also the PA of the  photometric major axis of the galaxy.  We  conclude that the radio jet is launched at a small angle relative to  the galactic plane, with the NW side slightly oriented toward us. This  angle is small enough for the radio jet to shock and compress the gas  in the plane of the galaxy, and for the nuclear continuum to ionize and heat it. The distinct kinematics and flux distributions observed for the different emission lines suggest different origins for their emission. The [Fe\\,{\\sc ii}] shows the largest  blueshifts and velocity dispersions and its flux distribution is concentrated along the jet, while the H$_2$ shows the lowest velocity dispersions and has additional flux contribution from regions beyond the jet.  Both X-rays emitted by the active galactic nucleus and shocks produced by the radio jet can excite the H$_2$ and [Fe\\,{\\sc ii}] emission lines. We use the 2D  velocity dispersion maps to estimate upper limits to the contribution of the radio jet to the excitation of [Fe\\,{\\sc ii}] and H$_2$ which may reach 90\\% for [Fe\\,{\\sc ii}] and 80\\% for H$_2$ in the jet region. The [Fe\\,{\\sc ii}]/Pa$\\beta$ emission-line ratios and the association of the [Fe\\,{\\sc ii}] flux distribution and kinematics with the radio structure supports a stronger contribution of the radio  jet to the [Fe\\,{\\sc ii}] excitation than to that of H$_2$.  In the regions beyond the jet the observations favor X-ray excitation. ",
        "introduction": "The Narrow-Line Region (NLR) of Seyfert galaxies is one of the best probes of the mechanisms in operation in the surrounding of accreting supermassive black holes in galaxies. The excitation and dynamics of the inner NLR gas can reveal how radiation and mass outflows from the nucleus interact with circumnuclear gas. Until recently, the measurement of these properties had to rely on either narrow-band images \\citep[e.g.][]{wilson93, capetti96, schmitt96} or long-slit spectroscopy. Long-slit spectroscopic studies provide information on the origin and excitation of the emission lines \\citep[e.g.][]{veilleux97, rodriguez-ardila05} as well as on the gas kinematics \\citep[e.g.][]{ winge97,winge99,crenshaw00,hutchings98, kaiser00,nelson00}, but are restricted to the locations covered by the long-slit.  \\citet{mulchaey96} and \\citet{ferruit00} have shown that obscuration can affect the optical morphology of the emitting gas region, a problem that can be alleviated by using infrared lines to map the NLR. Relevant emission-lines in the near-IR include [Fe\\,{\\sc ii}]$\\lambda$1.257$\\mu$m and 1.644$\\mu$m, HI lines such as Pa\\,$\\beta$ and Br\\,$\\gamma$ and molecular hydrogen lines such as H$_2\\lambda$1.957$\\mu$m and H$_2\\lambda$2.121$\\mu$m, which can be used to map the gas kinematics and excitation \\citep{storchi-bergmann99,winge00}.  Studies based on the above lines \\citep{forbes93,blietz94} have revealed a correlation between the [Fe\\,{\\sc ii}] and the radio emission, indicating that shock excitation by the radio jets is a likely mechanism for production of the [Fe\\,{\\sc ii}] emission in these objects, although other works have favoured photoionization by the nuclear source \\citep{simpson96,das05,das06}. H$_2$ emission is also strong in AGNs \\citep[e.g.][]{veilleux97,rodriguez-ardila04,rodriguez-ardila05} and can be used to probe the distribution of warm molecular gas.  To further progress in this area and investigate the origin of the near-IR emission lines it is necessary to spatially resolve the kinematics and excitation of the emitting gas.  The multiplexing capability necessary to map, simultaneously, the gaseous distribution, excitation and kinematics at such small scales with sufficient spectral resolution is now provided by the integral field units (IFU) in operation in large telescopes. In this work we use the IFU of the Gemini Near-Infrared Spectrograph which provides two dimensional (hereafter 2D) mapping in the near-IR (thus minimizing the effect of reddening) to study the circumnuclear emitting gas of the Seyfert galaxy ESO\\,428-G14. We present 2D flux, radial velocity and velocity dispersion maps in four strong emission lines: [Fe\\,{\\sc ii}]$\\lambda 1.257 \\mu$m, Pa$\\beta\\lambda 1.282\\mu$m, H$_2 \\lambda 2.121 \\mu$m and Br$\\gamma\\lambda 2.166\\mu$m. The R$\\approx$6000 spectral resolution has allowed us also to obtain velocity slices across the emission-line profiles in order to better probe the velocity fields.  We present in addition  broad-band {\\it J} and {\\it K} images and a narrow-band image in the line [Fe\\,{\\sc ii}]$\\lambda 1.257 \\mu$m obtained with the CTIO Blanco Telescope. We explore the relation between reconstructed images in the near-IR emission lines, their kinematic maps, the narrow-band [Fe\\,{\\sc ii}] image and previously published optical [O\\,{\\sc iii}]$\\lambda 5007$ narrow-band and radio images, in order to investigate the excitation mechanism of the near-IR lines, in particular [Fe\\,{\\sc ii}] and H$_2$.  ESO\\,428-G14 is an S0 galaxy which was classified as Seyfert 2 by \\citet{bergvall}. High resolution VLA maps show predominantly a two-sided and asymmetric radio jet oriented approximately along the galaxy line of nodes at position angle PA$\\approx\\,129^{\\circ}$, in which the NW side is less extended and terminates in a bright hot spot at 0$\\farcs$75 from the nucleus, while to the SE the radio emission is fainter, the jet bends to N and its extent is at least three times longer   \\citep{ulvestad,falcke96,falcke98}. [O\\,{\\sc iii}] and H$\\alpha$+[N\\,{\\sc ii}] images obtained with the Hubble Space Telescope show extended emission well aligned with the radio jet with stronger emission to the NW as observed in the radio. The ratio  [O\\,{\\sc iii}]/(H$\\alpha$+[N\\,{\\sc ii}] ) shows a bipolar structure with larger values to the SE \\citep{falcke96,falcke98}. We adopt  a distance for this galaxy of 19\\,Mpc \\citep{falcke96} for which 1$^{\\prime\\prime}$ corresponds to 92\\,pc at the galaxy.  This paper is organized as follows: in Section 2 we describe the observations and data reduction. In Section 3 we present the near-IR spectra and images, a comparison with previous optical narrow-band and radio images and the kinematics of the gas. In Section 4, we discuss the results and their implications for the origin of the [Fe\\,{\\sc ii}] and  H$_2$ emission and in Section 5 we present the conclusions of this work.  \\section[]{Observations and Reductions}  \\subsection{Integral Field Spectroscopy}  The spectroscopic data were obtained with the Gemini Near Infra-Red Spectrograph (GNIRS) \\citep{elias98} Integral Field Unit (IFU) in December 2004 under the instrument science verification program GS-2004B-SV-26,  and comprise two sets of observations centred at $\\lambda$=2.21$\\mu$m and 1.27$\\mu$m. We have used the 111 l/mm grating with the Short Blue Camera (0$\\farcs$15/pixel) which gives a resolving power of R=5900. The GNIRS IFU has a rectangular field of view, of approximately 3$\\farcs$2\\,$\\times$\\,4$\\farcs$8, divided into 21 slices. At the detector, the slices are divided along their length into 0$\\farcs$15 square IFU elements.  For our observations, the major axis of the IFU was oriented along the PA $\\approx129^{\\circ}$, which is the orientation of the radio jet. The observing procedure followed the standard Object--Sky--Sky--Object dither sequence, with off-source sky positions since the target is extended, and individual exposure times of 600s. One of the IFU slicing mirrors was damaged during assembly and presents a lower (20\\% of the nominal value) transmission region, about 1$\\farcs$8 in length, offset 0$\\farcs$5 from the center, therefore small spatial offsets in both directions were added between exposures to ``fill in'' the signal in this region. Telluric standard stars were observed immediately after the target. The basic observing information is shown in Table \\ref{obs}. Conditions during the observations were good for the night of Dec. 27 (clear skies, image quality (IQ) $\\sim$0$\\farcs$60), and patchy clouds (but stable guide counts) with similar IQ for the Dec. 28 dataset.  The data reduction was accomplished using tasks contained in GNIRS package which is part of the GEMINI IRAF package as well as generic IRAF tasks. Through the reduction tasks we have performed trimming, flat-fielding, sky subtraction, wavelength and s-distortion calibrations. We have also removed the telluric bands and flux calibrated the frames using the star HR4023 (A2V spectral type) as a ``relative\" flux standard. The cosmic ray cleaning was done before the sky subtraction using  the algorithm  described by \\citet{vandokkum}.  In Fig.\\,\\ref{espectro} we present the nuclear spectra   and the spectra at 0$\\farcs$8\\,NW from the nucleus, which is the position of the peak of the radio continuum emission within an aperture of 0$\\farcs$45 diameter.  \\begin{table} \\caption{Observations} \\centering \\begin{tabular}{c c c c c} \\hline Instrument/& Date &         & Total Exp.&  \\\\ Telescope  & (UT) &  Filter & time (s)   \\\\ \\hline GNIRS/Gemini  & 2004/12/27 & K$_-$G0503 & 2400  \\\\ GNIRS/Gemini  & 2004/12/28 & J$_-$G0505 & 1200  \\\\ OSIRIS/Blanco & 2000/02/15 &  K         &   80  \\\\ &            &  J         &   80   \\\\ &            & 1.06       &  270 \\\\ &            &[Fe\\,{\\sc ii}] 1.257      & 1800  \\\\ \\hline \\end{tabular} \\label{obs} \\end{table}  \\subsection {Imaging}  Broad-band {\\it J} and {\\it K} images and narrow band images centered on the emission line [Fe\\,{\\sc ii}]$\\lambda1.257$ and adjacent continuum were obtained in February 2000 at the Cerro Tololo Interamerican Observatory (CTIO) using the 4-m Blanco telescope.  The observations were performed with the infrared imager/spectrometer OSIRIS using a 1024\\,$\\times$\\,1024 HgCdTe array at an angular scale of 0\\farcs16\\,pix$^{-1}$. Details of the observations are presented in Table \\ref{obs}. The angular resolution was 0\\farcs8-0\\farcs9 (FWHM of standard star images).  The observing procedure consisted of taking a sequence of 4--9 dithered images on source with sky images taken before or after the object sequence.  Individual sky and object images were first divided by the flat-field. We then median combined the  sky images and subtracted them from the corresponding object images. The individual images in the same filter were then aligned using as many common point sources (stars and the nucleus) as  possible and finally averaged. We have used the task {\\tt imcoadd} from the GEMINI package to  perform this alignment/averaging procedure.  Photometric calibration was performed for both broad and narrow band filters using standard stars. Correction for extinction was performed using the mean extinction coefficients for February \\citep{frogel98}.  The narrow band [Fe\\,{\\sc ii}] image was calibrated by multiplying each filter transmission curve by the black body spectrum corresponding to the associated standard star spectral type and then using this integrated flux as the flux of the standard star in the filter. After calibration, the emission line image was created by subtracting the aligned continuum image (Filter 1.06$\\mu$m) from the continuum plus emission line image.  \\begin{figure} \\centering \\includegraphics[angle=-90, scale=0.55]{espectro.eps} \\caption{Top: Nuclear spectra for ESO\\,428-G14 in the {\\it J}-band (left) and  {\\it K}-band (right). Bottom: Spectra at 0$\\farcs$8\\,NW from the nucleus, which is the position of the radio continuum emission peak.}  \\label{espectro} \\end{figure} ",
        "conclusions": "We have analyzed near-IR {\\it J} and {\\it K} band  2D spectra of the inner $\\approx$300\\,pc of the Seyfert galaxy ESO\\,428-G14 obtained with the Gemini GNIRS IFU. The fine angular sampling of this instrument and the high spectral resolution of the observations provided a spatial sampling of 14\\,pc$^2$ at the galaxy and a velocity resolution of $\\approx$20\\,km\\,s$^{-1}$. These characteristics and 2D coverage has allowed us to obtain gas flux distributions and ratios as well as to map the gas kinematics in the four strong emission lines [Fe\\,{\\sc ii}]$\\lambda 1.257 \\mu$m, Pa$\\beta$, H$_2 \\lambda 2.121 \\mu$m and Br$\\gamma$. Such 2D mapping at the spatial and spectral resolution of our observations and in the near-IR is unique in the literature so far.  We have used the spectroscopic data in conjunction with a [Fe\\,{\\sc ii}]$\\lambda 1.257 \\mu$m narrow-band image and $J-K$ color map obtained with OSIRIS at the CTIO Blanco Telescope. The [Fe\\,{\\sc ii}] narrow-band image shows a bi-polar structure which correlates well with the structure observed in a previous published radio map, as well as with that observed in a narrow-band [O\\,{\\sc iii}] HST image. This bi-polar structure is oriented approximately along the photometric major axis of the galaxy at PA$\\approx$129$^\\circ$. Under the assumption that redder color is due to dust screening, we conclude from the $J-K$ color map that NE is the near side of the galaxy.  The IFU observations show that not only the [Fe\\,{\\sc ii}] flux distribution, but also Pa$\\beta$ and Br$\\gamma$ correlate with the radio emission, all showing elongated morphologies along the radio axis. The H$_2$ emission has an additional component more spread over the disk of the host galaxy.  The gas kinematics shows a pattern which, at first sight, seems to be due to rotation in a disk, but which is clearly disturbed by non-circular motions. The analysis of the gas kinematics is complicated by the fact that the galaxy major axis is approximately aligned with the radio axis. The radio emission morphology suggests that a radio jet is the cause of outflows observed in velocity ``slices\" obtained along the gas emission line profiles: blueshifts of $\\sim$ 300--400\\,km\\,s$^{-1}$ are observed in emission associated with a radio hot spot at $\\sim$0\\farcs8 NW of the nucleus, while mostly redshifts and some blueshifts observed to the opposite side of the nucleus are associated to fainter radio emission observed up to 2$^{\\prime\\prime}$ SE. From this velocity pattern, as well as from the emission-line ratios we conclude that the radio jet leaves the nucleus at a small angle relative to the plane of the galaxy, with the NW side slightly directed toward us and the counter-jet directed away from us, being partially hidden by the disk of the galaxy. Some redshift observed to the NW and some blueshift observed to the SE shows that the line emission is observed to both sides of the plane of the galaxy, confirming that the angle between the radio jet and the plane of the galaxy is small and emission associated with the jets occurs on both sides of the galaxy plane. The reddening distribution obtained from the Br$\\gamma$/Pa$\\beta$ emission-line ratio coincides with that obtained from the continuum, confirming that the line emission comes from circumnuclear gas located essentially in the galactic plane.  The velocity dispersion maps show the highest values between the nucleus and the regions of strongest radio emission (the radio hot spots). We interpret this result as due to kinetic energy deposited by ratio jet in the circumnuclear ISM, producing a compression in the gas and emission enhancement just beyond this compressed region.  While the [Fe\\,{\\sc ii}] emission is dominated by the highest velocities at the two hot spots in opposite sides of the nucleus, the H\\,{\\sc i} emission has an important contribution also from lower velocities, observed all along the radio emission structure. The H$_2$ emission has significant emission contribution at low velocities spread throughout the observed region, apparently coming from the disk of the host galaxy.  From the observed kinematics, we conclude that the radio jet has a fundamental role in shaping the emission line region as it interacts with the galaxy ISM surrounding the galaxy nucleus. % We have used the 2D velocity dispersion maps  to estimate the kinetic energy deposited in the circumnuclear ISM by the radio jet relative to regions away from the jet. Assuming that the H$_2$ and [Fe\\,{\\sc ii}] excitation  in the latter regions is dominated by X-rays, and that the X-rays excitation is the same in the jet region, we obtain contributions by shocks of up to 80\\,-\\,90\\% for the [Fe\\,{\\sc ii}] and up to 70\\,-\\,80\\% for H$_2$ in the jet region. These are however, upper limits due to the fact that the X-rays contribution along the jet axis may be larger than away from it. The stronger association of the [Fe\\,{\\sc ii}] emission and kinematics with the  radio structure supports a larger contribution of the radio jet to the excitation of [Fe\\,{\\sc ii}] than to  that of H$_2$."
    },
    "0609/astro-ph0609290.txt": {
        "abstract": "{} {\tWe present an analysis of \\object{OGLE~2004--BLG--254}, a high-magnification ($A_{\\circ} \\simeq 60$) and relatively short duration ($\\tE \\simeq 13.2$ days) microlensing event in which the source star, a Bulge K-giant, has been spatially resolved by a point-like lens. We seek to determine the lens and source distance, and provide a measurement of the linear limb-darkening coefficients of the source star in the $I$ and $R$ bands. We discuss the derived values of the latter and compare them to the classical theoretical laws, and furthermore examine the cases of already published microlensed GK-giants limb-darkening measurements. } { We have obtained dense photometric coverage of the event light curve with OGLE and PLANET telescopes, as well as a high signal-to-noise ratio spectrum taken while the source was still magnified by $A \\sim 20$, using the UVES/VLT spectrograph. We have performed a modelling of the light curve, including finite source and parallax effects, and have combined spectroscopic and photometric analysis to infer the source distance. A Galactic model for the mass and velocity distribution of the stars has been used to estimate the lens distance. } { From the spectrum analysis and calibrated color-magnitude of the event target, we found that the source was a K3~III Bulge giant, situated at the far end of the Bulge. From modelling the light curve, we have derived an angular size of the Einstein ring $\\ThE \\simeq 114\\,\\mu$as, and a relative lens-source proper motion $\\mu = \\ThE/\\tE \\simeq 3.1$ mas/yr. We could also measure the angular size of the source, $\\ThS \\simeq 4.5\\,\\mu$as, whereas given the short duration of the event, no significant constraint could be obtained from parallax effects. A Galactic model based on the modelling of the light curve then provides us with an estimate of the lens distance, mass and velocity as $\\DL\\simeq 9.6$~kpc, $M\\simeq 0.11\\,\\msun$ and $v \\simeq 145~\\mbox{km}\\,\\mbox{s}^{-1}$ (at the lens distance) respectively. Our dense coverage of this event allows us to measure limb darkening of the source star in the $I$ and $R$ bands.  We also compare previous measurements of linear limb-darkening coefficients involving GK-giant stars with predictions from ATLAS atmosphere models. We discuss the case of K-giants and find a disagreement between limb-darkening measurements and model predictions, which may be caused by the inadequacy of the linear limb-darkening law. } {} ",
        "introduction": "\\label{sec:intro} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  The microlensing technique is one of a few methods (together with interferometry, eclipsing binaries and transiting extra-solar planets) that can be used to measure brightness profiles and limb-darkening coefficients of stars at distances exceeding a few kpc. The apparent stellar disk is a projection of the near stellar hemisphere and therefore shows an axisymmetric variation of the observed brightness as the projection maps the variation of the emergent radiation with depth. The continuum is on average formed at larger depth at the disk center and at smaller depth at the limb \\citep[Eddington-Barbier effect, see e.g.:][]{WM1994,Dim95,Sasselov1997, HSL2000,Heyrovsky2003}. The microlensing method requires the source to transit a region with a large magnification gradient over its face, as present in the vicinity of caustics.  While a single lens creates a point-like caustic at its angular position, binary lenses create extended caustic patterns formed of lines (fold caustics) which merge at cusps.  PLANET observations of the microlensing event \\object{MACHO~1997--BLG--28} \\citep{Albrow1999b} containing a cusp passage, constituted the first limb-darkening measurement of a Bulge giant. The majority of subsequent limb-darkening measurements such as MACHO~1998--SMC--1 \\citep{Albrow1999a, Afonso2000}, \\object{MACHO~1997--BLG--41} \\citep{Albrow2000b}, \\object{OGLE~1999--BLG--23} \\citep{Albrow2001a}, \\object{EROS--2000--BLG--5} \\citep{Fields2003,An2002}, \\object{OGLE~2002--BLG--069} \\citep{Cassan2004,Kubas2005,Thurl2006} resulted from fold-caustic passages. The limb-darkening measurement in the solar-like star \\object{MOA~2002--BLG--33} \\citep{Abe2003} constituted a very special case with the source enclosing several cusps of the caustic at the same time. In the single lens case,  the extended size of the source star has a significant effect on the light curve if the angular source size is of the order or larger than the angular separation between source center and lens point-like caustic. However, only a small fraction of microlensing events provide angular separations that are small enough, and so far very few cases have been observed \\citep{Alcock1997,Yoo2004,Jiang2004}.  For events showing evidence of this effect, the limb darkening of the source star can be determined, and constraints on the physical properties of the lens can be derived using estimates of its distance, spectral type and event model parameters \\citep{Gould1994, Witt1995, Dim95, HSL2000, Heyrovsky2003}.  This work treats the case of OGLE~2004--BLG--254, a high magnification event from a point-like lens showing extended source effects, for which a dense photometric follow-up was performed at four PLANET observing sites, making it one of the best observed events of this kind to date.  Using high-resolution spectra collected with UVES (VLT), we derive the characteristics of the source star, a K3 giant. We search for a suitable single-lens limb-darkened source microlensing model, which we use together with the properties derived from the high resolution spectra to discuss the source center-to-limb variations and constraints on the lens properties.  We compare our limb-darkening measurements to previous observations of such stars, especially to EROS--2000--BLG--5 for which a large discrepancy was found between observations and atmosphere models.  We also note that the derived characteristics provide the basic information necessary for a forthcoming study of element abundance in the Bulge giant source star of OGLE~2004--BLG--254.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{sec:discussion} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  We have performed dense photometric monitoring of the microlensing event OGLE~2004--BLG--254, a relatively short duration and small impact parameter microlensing event generated by a point-like lens transiting a giant star. The peak magnification was about $A_{\\circ}\\sim 60$, effectively multiplying the diameters of our network telescopes by a factor $\\sim 8$.  High-resolution spectra were taken using the UVES spectrograph at La Silla while the source was magnified by a factor $A \\sim 20$, just after the end of the transit of the source over the caustic. This yielded a precise measurement of the characteristics of the star, a K3~III giant, although a discrepancy is found between photometric and spectroscopic temperature. Using a calibrated colour-magnitude diagram analysis and isochrones, we find a source angular radius $\\ThS \\simeq 4.5 \\pm 0.7 \\,\\mu\\mbox{as}$, and a physical radius $\\RS \\simeq 10.2 \\rsun$ with the source distance $\\DS = 10.5$~kpc. A Galaxy model together with the event time-scale $\\tE = 13.23~\\mbox{days}$ and the source-size parameter $\\rhoS = 0.04$ yielded a lens mass $M\\simeq 0.11\\,\\msun$, a lens distance $\\DL = 9.6~\\mbox{kpc}$, and a velocity $v \\simeq 145~\\mbox{km}\\,\\mbox{s}^{-1}$ at the lens distance. From our photometric data, we have derived measurements of the source limb-darkening coefficients for the $I$ and $R$ broadband filters, and provided arguments for a discussion about a lack of relevant physics in K-giants atmosphere models, or an inadequacy of the linear limb-darkening law to model.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0609/astro-ph0609114_arXiv.txt": {
        "abstract": "We discuss astrophysical scenarios relevant to the generation of gravitational waves (GW) and effects expected to arise from the interaction of GW and electromagnetic (EM) radiation. A strong programme of coordinated GW and EM astrophysical studies must be established in order to ensure the exploitation of the full scientific potential of the LISA mission. We describe on-going astrophysical work, and suggest alternative approaches to current studies, which are relevant to these considerations. ",
        "introduction": "{\\sl LISA} will return unprecedented data on GW sources; however, its full scientific potential will be realised only by matching the sources with astrophysical counterparts and correlating their properties over the EM spectrum. Different types of sources require different approaches: direct identification with known EM sources in some cases (like for ultra-compact binaries), and statistical estimation from the systems EM characteristics in others (such as for rates of Super Massive Black Hole, or SMBH, mergers, and Extreme Mass Ratio Inspirals, or EMRIs). In turn, knowledge of the EM properties (e.g. for the ultra-compact binaries) will be crucial in constructing accurate waveforms to aid {\\sl LISA's} signal processing. The consequences of the interaction between GW and EM radiation may also have important implications, for instance on events such as gamma-ray bursts. It is clear that a large degree of synergy is needed between the GW and EM astrophysical communities, in order to build a strong programme of coordinated studies targeted to the needs of {\\sl LISA}. Below we discuss on-going astrophysical work relevant to such ideas. ",
        "conclusions": ""
    },
    "0609/hep-ph0609059_arXiv.txt": {
        "abstract": "{We consider the split extended ($N=2$) supersymmetry scenario recently proposed by Antoniadis {\\it et al.}~[hep-ph/0507192] as a realistic low energy framework arising from intersecting brane models. While all scalar superpartners and charged gauginos are naturally at a heavy scale, the model low energy spectrum contains a Higgsino-like chargino and a neutralino sector made out of two Higgsino and two Bino states. We show that the lightest neutralino is a viable dark matter candidate, finding regions in the parameter space where its thermal relic abundance matches the latest determination of the density of matter in the Universe by WMAP. We also discuss dark matter detection strategies within this model: we point out that current data on cosmic-ray antimatter already place significant constraints on the model, while direct detection is the most promising technique for the future. Analogies and differences with respect to the standard split SUSY scenario based on the Minimal Supersymmetryc Standard Model are illustrated.} ",
        "introduction": "In recent years cosmological observations have provided increasingly convincing evidence that non-baryonic dark matter (DM) is the building block of all structures in the Universe~\\cite{Spergel:2006hy}. Consequently when defining standard model (SM) extension candidates of elementary particles, the possibility to embed a DM candidate has become a compelling guideline. In such respect the formulation of split supersymmetry~\\cite{Arkani-Hamed:2004fb,Giudice:2004tc} is no exception.  Split supersymmetry (SUSY) labels a generic realization of a SUSY SM extension with a \"split\" superpartner spectrum. On one side all sfermions are assumed to be very heavy (at some intermediate scale between, say, 100 TeV and the GUT scale): this feature explains why SUSY contributions to flavor and CP violation are small at the cost of invoking some mechanism, not related to SUSY, to stabilize the weak scale~\\cite{Arkani-Hamed:2004fb}, but still allowing for a successful unification of gauge couplings. On the other hand (at least) some of the fermionic superpartners need to be light: the cosmological measurements of the matter density set an upper bound to the thermal relic abundance of the lightest neutralino, the lightest SUSY  particle (LPS)  in this framework, which in turn can be translated into an upper bound on the LSP mass (about 340~TeV for a generic thermal relic~\\cite{Griest:1989wd},  about 2.1~TeV for a pure Wino-like neutralino, see e.g.~\\cite{mpu}). The split SUSY framework is thus the minimal SUSY setup which can accommodate a DM candidate: taking as strong prior the condition that the LSP accounts for all DM in the Universe sets tight constraints on the model.  A large class of scenarios predict a split SUSY spectrum~\\cite{Arkani-Hamed:2004yi, Antoniadis:2004dt,Kors:2004hz,Babu:2005ui,Dutta:2005zz}. We will focus here on the model arising in a string-inspired framework with intersecting branes, recently introduced in Refs.~\\cite{splitn2}. This model has gauge and Higgs sectors defined in multiplets of extended supersymmetry. Its low energy spectrum cannot be described as a subset of the spectrum in the minimal SUSY SM extension (MSSM), which is the standard lore for most phenomenological studies of low energy SUSY models.  We will consider the LSP relic density calculation as a guideline to examine the structure of the model, and discuss the perspectives of testing neutralino DM in this scenario. Previous studies of DM in split SUSY have assumed the MSSM as working framework, see e.g.~\\cite{Pierce:2004mk,Arvanitaki:2004df,mpu,Senatore:2004zf, Arkani-Hamed:2006mb} and so we will point out differences and analogies with respect to the present case.  The structure of the paper is as follows. In Section~\\ref{Sec:model} we introduce the framework and discuss its low energy limit. In Section~\\ref{Sec:omega} we compute the LSP relic density.  We then discuss in Section~\\ref{sec:phenoguide} current constraints on the model from direct and indirect DM searches, and prospects to test it at future facilities. Section~\\ref{Sec:concl} summarizes our results. ",
        "conclusions": "\\label{Sec:concl}  We have discussed the dark matter phenomenology for a split SUSY model arising from a high energy intersecting brane model with $N=2$ supersymmetry. Its active states at low energy differ from those in the standard split SUSY scenario based on the MSSM.  Analogies and differences compared to the standard case have emerged. In both frameworks the lightest neutralino is a dark matter candidate for: i) Intermediate mass neutralinos with a sizable Bino-Higgsino mixing, or ii) A pure Higgsino state at the TeV scale. The fine-tuning parameter for such configurations is in general rather small. In the split $N=2$ SUSY model direct detection is very promising, covering a significantly larger portion of the parameter space than in the standard case; in particular all results in the present analysis are essentially independent of $\\tan\\beta$, and the spin independent scattering cross section does not get suppressed in the large $\\tan\\beta$ limit as in the MSSM. Concerning indirect detection searches with neutrino telescopes are not relevant in this specific extended SUSY scenario, while antimatter measurements have already excluded a relevant portion of the parameter space. In the next few years the Pamela detector will measure the antiproton and positron fluxes with higher accuracy, allowing for further tests of the model in a regime which is complementary to direct detection searches."
    },
    "0609/astro-ph0609336_arXiv.txt": {
        "abstract": "We determine the relative rates of short GRBs in cluster and field early-type galaxies as a function of the age probability distribution of their progenitors, $P(\\tau)\\propto \\tau^n$.  This analysis takes advantage of the difference in the growth of stellar mass in clusters and in the field, which arises from the combined effects of the galaxy stellar mass function, the early-type fraction, and the dependence of star formation history on mass and environment.  This approach complements the use of the early- to late-type host galaxy ratio, with the added benefit that the star formation histories of early-type galaxies are simpler than those of late-type galaxies, and any systematic differences between progenitors in early- and late-type galaxies are removed.  We find that the ratio varies from $R_{\\rm cluster}/R_{\\rm field}\\sim 0.5$ for $n=-2$ to $\\sim 3$ for $n=2$. Current observations indicate a ratio of about $2$, corresponding to $n\\sim 0-1$.  This is similar to the value inferred from the ratio of short GRBs in early- and late-type hosts, but it differs from the value of $n\\approx -1$ for NS binaries in the Milky Way.  We stress that this general approach can be easily modified with improved knowledge of the effects of environment and mass on the build-up of stellar mass, as well as the effect of globular clusters on the short GRB rate.  It can also be used to assess the age distribution of Type Ia supernova progenitors. ",
        "introduction": "\\label{sec:intro}  Gamma-ray bursts (GRBs) are divided into two broad classes of short/hard and long/soft bursts \\citep{kmf+93}, which appear to have different progenitor populations.  Observations of Type Ic supernovae in association with long GRBs provide a direct confirmation that they arise from the death of massive stars \\citep{hsm+03,smg+03}.  Short GRBs, on the other hand, have long been suspected to arise from the merger of neutron star and/or black hole binaries (NS-NS, NS-BH; e.g., \\citealt{elp+89,npp92,rr02,ajm05}).  During the last year, follow-up observations of \\swift\\ and \\hete\\ short GRBs have provided initial confirmation to this idea, based in particular on the localization of some short GRBs to elliptical galaxies \\citep{bpc+05,gso+05,bpp+06}, their lack of association with bright supernovae \\citep{ffp+05,hwf+05,sbk+06}, and their lower energy release and wider beaming angles compared to long GRBs \\citep{bgc+06,sbk+06}.  Despite this observational progress we are still missing a clear understanding of the progenitor population, due to the lack of direct observations (circumburst chemical abundances, gravitational waves, or a sub-relativistic, radioactive component: \\citealt{lp98,kul05}). Thus, statistical studies of the burst properties can be highly effective in understanding their progenitor population(s) (e.g., \\citealt{gno+05,ngf05,gp06}).  Recently, \\citet{zr06} investigated the ratio of short GRBs in early- and late-type galaxies as a constraint on the age distribution of the progenitors.  Their analysis combines the global star formation rate in each galaxy type with a local galaxy stellar mass function, assuming the formation process of short GRB progenitors does not depend on any other physical parameters.  Since each galaxy type has experienced a globally different star formation history, the ratio of bursts in each type is predicted to vary as a function of the progenitor age distribution.  A complementary way to constrain the progenitor age distribution is to use the rates of short GRBs in clusters and the field (\\citealt{bsm+06}; hereafter Paper I).  This approach takes advantage of the following differences between cluster and field environment. First, the galaxy stellar mass function of clusters is more heavily dominated by massive galaxies than in the field \\citep{cfn+05,bbb+06}.  Since the star formation history is mainly determined by galaxy mass, the overall growth of stellar mass is in turn affected by the large-scale environment. Second, the fraction of early-type galaxies is larger in clusters than in the field \\citep{dre80,wgj93,bbb+06}.  Finally, there appears to be a systematic offset in the star formation histories of cluster and field early-type galaxies of $\\sim 1-3$ Gyr (e.g., \\citealt{brc+98,ksc+02,tmb+05}).  These effects lead to an overall difference in cluster and field star formation histories, which combined with the progenitor age distribution, is expected to affect the relative fraction of short GRBs in each environment.  In this paper we quantify these effects and show how the ratio of short GRBs in cluster and field early-type galaxies can be used to understand the age distribution of the progenitors.  This is part of our on-going systematic study of galaxy clusters hosting short GRBs, using multi-slit optical spectroscopy and X-ray observations (Paper I).  We find that the current observations, albeit with a small number of events, favor $P(\\tau)\\propto \\tau^n$ with $n\\sim 0-1$.  Finally, we provide a comparison of the systematic effects in this approach and the approach of \\citet{zr06}. ",
        "conclusions": "\\label{sec:disc}  We now turn to a comparison of our model with observation of short GRBs.  To date two short bursts have been localized to clusters at $z\\sim 0.2$: GRB\\,050509b \\citep{gso+05,bpp+06} and GRB\\,050911 (Paper1).  We note that in the latter case the large error circle prevents an association with a specific cluster galaxy, but the large early-type fraction of $80\\%$ \\citep{bsm+06} suggests that the burst was likely hosted by an early-type galaxy.  On the other hand, only one short GRB has been localized to a field early-type galaxy, GRB\\,050724 \\citep{bpc+05}.  We do not consider GRB\\,050813, which was hosted by a cluster at a much higher redshift ($z\\sim 1.8$; \\citealt{ber06}), and GRB\\,060502b, which may be hosted by an early-type galaxy \\citep{bpc+06}, but whose large-scale environment has not been fully explored yet\\footnotemark\\footnotetext{The limit on the X-ray luminosity at the redshift of the putative host galaxy, $L_X\\lesssim 6\\times 10^{42}$ erg s$^{-1}$, is lower by a factor of eight than that of the cluster hosting GRB\\,050509b, but it is somewhat higher than that of the cluster hosting GRB\\,050911 (Paper I).}.  Thus, the current ratio of short GRBs in cluster versus field early-type galaxies is about $2\\!:\\!1$, with a large uncertainty due to the small number of events.  From Figure~\\ref{fig:full}, we find that this ratio corresponds to $n\\sim 0-1$. This value is lower than $n\\gtrsim 3/2$ claimed by \\citet{zr06}, but is in rough agreement with $-1 \\lesssim n \\lesssim 0$ found by \\citet{bfp+06} based on their revised redshift distribution with $1/4-2/3$ of all short GRBs at $z\\gtrsim 1$.   Clearly, in both methods of estimating the age of short GRB progenitors the uncertainty in the inferred value of $n$ is currently dominated by the small number of bursts with a known redshift, host galaxy, and large-scale environment type.  Since this uncertainty will eventually diminish with a larger sample of events, it is interesting to consider systematic uncertainties in both theoretical approaches. Our analysis suffers from the somewhat poor definition of cluster and field environments.  We have used an overall cluster mass fraction of $20\\%$, as indicated by several researchers, but this number may range from $10$ to $30\\%$.  Second, we have used the simplified bimodal star formation history model of \\citet{tmb+05}, but these authors do not use the same quantitative definition of galaxy environment that was used for the mass functions by \\citet{bbb+06}.  Since we have used representative mass functions, and then scaled the results by the overall mass fraction in clusters and the field, this effect should not be significant. Third, the uncertainty in the definition of field galaxy environment leads to an overall uncertainty of about $20\\%$ in our calculated ratio. Finally, since for $n\\lesssim -1$ the ratio depends on $M_l$, it is essential to understand the appropriate low mass limit for early-type galaxies.  With the current inferred value of $n\\sim 0-1$, however, this may not be a relevant issue for short GRBs.  Similarly, the analysis of the short GRB rate in early- and late-type galaxies performed by \\citet{zr06} also suffers from systematic effects.  First, the ratio is affected by the uncertain star formation history at high redshift when early-type galaxies formed most of their stars.  Second, their estimation of star formation rate does not consider environmental effects that appear to be important in current observations from SDSS and 2dF, and which we have accounted for here. Finally, as noted by \\citet{zr06}, it is possible that late-type galaxies have an altogether different age distribution of short GRB progenitors than early-type galaxies.  This problem is overcome by our method since it considers only early-type host galaxies.  We note that if both derivations are in fact correct, then the estimated values of $n$ can be used to assess any systematic differences of short GRB progenitors in early- and late-type galaxies, as suspected to exist for type Ia supernovae \\citep{slp+06}.  Future applications of our approach will include the effect of globular clusters, which are thought to provide an efficient environment for the production of NS-NS binaries, and may account for a substantial fraction of all short GRB progenitors \\citep{gpm06,hopman+06}.  We expect that since the specific frequency of globular clusters increases significantly with galaxy mass \\citep{har91}, an association with globular clusters will increase the fraction of short GRBs in galaxy clusters compared to the trend shown in Figure~\\ref{fig:full}. Similarly, our approach can be extended to higher redshift to investigate the evolution in the fraction of short GRBs in clusters. We expect that the lack of strong evolution in the last several Gyr likely makes our analysis applicable out to $z\\sim 1$.  However, if some short GRBs are in fact associated with clusters at $z\\sim 2$, this presents an opportunity to assess any systematic changes in the value of $n$ with redshift.  We end with the following conclusion.  If our current estimate of $n\\gtrsim 0$ continues to be supported by future observations, then this implies that the majority of short GRBs in early-type galaxies will occur in clusters.  This therefore suggests that short GRBs can provide an efficient tool for finding forming galaxy clusters at high redshift, as already appears to be the case for GRB\\,050813 \\citep{ber06}.  Continued near-IR imaging and optical spectroscopic observations of short GRB fields may therefore provide an efficient method for finding the highest redshift clusters."
    },
    "0609/astro-ph0609046_arXiv.txt": {
        "abstract": "Stars within 0.1\\,pc of the supermassive black hole \\sag\\ at the Galactic centre are expected to encounter a cluster of stellar-mass black holes (BHs) that have segregated to that region.  Some of these stars will scatter off an orbiting BH and be kicked out of the Galactic centre with velocities up to $\\sim\\,2000\\,$\\kms.  We calculate the resulting ejection rate of hypervelocity stars (HVSs) by this process under a variety of assumptions, and find it to be comparable to the tidal disruption rate of binary stars by \\sag, first discussed by \\hills.  Under some conditions, this novel process is sufficient to account for all of the hypervelocity B-stars observed in the halo, and may dominate the production rate of all HVSs with lifetimes much less than the relaxation time-scale at a distance $\\sim 2\\,$pc from \\sag\\ ($\\gtrsim 2\\,$Gyr).  Since HVSs are produced by at least two unavoidable processes, the statistics of HVSs could reveal bimodal velocity and mass distributions, and can constrain the distribution of BHs and stars in the innermost $0.1\\,$pc around \\sag. ",
        "introduction": "Hypervelocity stars (HVSs) have velocities so great ($\\sim1000\\,$\\kms) that they are gravitationally unbound to the Milky Way galaxy.  Since the discovery of the first HVS by \\browna, six additional HVSs have been found \\citep{2005Apj...634L.181E, 2005A&A...444L..61H, 2006ApJ...640L..35B, 2006ApJ...647..303B}.  The age and radial velocities of all but one HVS are consistent with them originating from the Galactic centre\\footnote{In the case of \\citet{2005Apj...634L.181E}, the star appears to originate from the Large Magellanic Cloud.}, the most natural site for producing them (Hills 1988).  As more HVSs are found, they can be used to constrain many properties of both the Galactic centre as well as the Milky-Way galaxy as a whole, providing information about the Galactic potential \\citep{2005ApJ...634..344G,2006astro.ph..8159B}, the merger history of the \\sag\\ \\citep{2005astro.ph..8193L,2006astro.ph..7455B}, as well as the mass distribution of stars near \\sag.  Nearly 17 years before the initial discovery by \\browna, \\hills\\ first proposed that HVSs should populate the galaxy and provide indirect evidence for the existence of a supermassive black hole (SMBH) in the Galactic centre. \\hills\\ showed that when a tight stellar binary (with a separation $<0.1\\,$AU) gets sufficiently close to a SMBH, it will be disrupted by its strong tidal field. Consequently, one member of the binary could be ejected from the Galactic centre with sufficient energy to escape the gravitational potential of the entire galaxy \\citep{1991AJ....102..704H,2005MNRAS.363..223G}.  The other binary member is expected to remain in a highly eccentric orbit around the SMBH \\citep[see, e.g.,][]{2003ApJ...592..935G,2006MNRAS.368..221G}. \\yutremaine\\ analysed the production rate of HVSs in more detail, specifically for stars originating near the SMBH in our galaxy, \\sag.  The authors also corrected the initial calculation of \\hills\\ and accounted for the diffusion of hard binaries into the ``loss--cone'', finding the production rate of HVSs to be $\\sim\\,10^{-5}\\,$yr$^{-1}$, nearly three orders of magnitude below the previous estimate.  \\yutremaine\\ also looked at two additional mechanisms for producing HVSs.  They found that the rate could easily be higher ($\\sim\\,10^{-4}\\,$yr$^{-1}$) if \\sag\\ had a massive binary companion \\citep[see also][]{2005astro.ph..8193L,2006astro.ph..7455B,2006astro.ph..4299S}, and also determined that star-star scattering near \\sag\\ resulted in a nearly undetectable rate, due to physical collisions among the two stars.  In examining the rate estimate from \\yutremaine\\, one may adopt a simple mass function (MF) for stars, $\\rmd n/\\rmd m \\propto m^{-\\beta}$ to estimate the total number of B--type HVSs in the Galactic halo.  Using the Salpeter MF slope $\\beta = 2.35$ and a minimum mass limit of $0.5\\,\\msun$ one expects there to be about $\\sim\\,80$ HVSs with masses between $3$ and $5\\,\\msun$, consistent with the observed rate for such stars \\brownsp. However this is an overestimate.  The Salpeter MF overestimates the total number of stars with mass $>\\,3\\,\\msun$, since star formation occurs continuously over time near the Galactic centre and massive stars have short lifetimes \\citep{astern,\\perets}. Accounting for continuous star formation as well as a more realistic distribution of binary parameters, \\cite{\\perets} found a lower total HVS rate of $5\\times 10^{-7}\\,$yr$^{-1}$ yielding $\\sim 1$ HVS in the entire galaxy with mass between $3$ and $5\\,\\msun$. The discrepancy between this calculation and the number of observed HVSs may be overcome by massive perturbers in the Galactic centre, which reduce the relaxation time and increase the rate to that observed \\citep{\\perets}. Nevertheless, their results are highly sensitive to the mass distribution of the perturbers.  In their calculations the authors use a very flat slope for the distribution of masses which has not been well constrained by observations.  In their estimate, they assume that all stars are initially relaxed and therefore fill the entire energy-momentum space.  The lifetimes of the stars observed by \\browns\\ are very similar to the enhanced relaxation time, except in the most optimistic case of \\citet{\\perets}.  We argue then that the total number of B-type HVSs given by \\cite{\\perets} may still be an overestimate, and an additional source of HVSs may be required.  In this \\letter, we propose a novel source of HVSs in the Milky Way: in the dense stellar cusp near \\sag\\, {\\it stars scatter off stellar-mass black holes} (BHs) that are segregated there \\citep{1993ApJ...408..496M} and recoil out of the Galactic centre. The existence of stars on radial orbits originating from the Galactic centre was first proposed by \\cite{2000ApJ...545..847M} as evidence for a segregated BH cluster. Here, we show that there should exist a high-velocity tail of stars, significant enough in number to account for some if not all of the HVSs observed by \\browns.  This \\letter\\ is organised as follows.  In \\S \\ref{stars} we describe the distribution of stars and BHs around \\sag. We describe in detail our assumptions and calculations in \\S \\ref{theory}, and present our results in \\S \\ref{results}.  Finally, in \\S \\ref{discussion}, we discuss the observational implications and prospects for future work. ",
        "conclusions": "\\label{discussion} We constructed a variety of models for the distribution of stars and BHs in the innermost $0.1\\,$pc of the Galactic centre.  Assuming that the cusp and disk of stars we observe today represent a steady state over the past $\\sim\\,100\\,$Myr, we calculated the rate at which stars will scatter off BHs and populate the Milky Way halo. We showed that the total ejection rate of HVSs can be comparable to that produced by the tidal disruption of binaries, and may exceed it for intermediate mass stars. Our results are consistent with the total number of ejected stars from Fokker-Plank simulations by \\citet{2006astro.ph..3280F}. Assuming that the B-stars are fully relaxed as is observed with the younger S-Stars \\citep{\\ghez}, we can account for most, if not all, of the stars observed by \\browns.  We demonstrated that the velocity distribution as well as the mass distribution of HVSs should be truncated at high values for the BH scattering process compared to binary disruption events (see Figs. 1 and 2).  Better statistics of HVS detections could therefore determine the relative significance of these two plausible channels.   In our study, the ejected stars originate from the inner $\\sim 0.1\\,$pc near \\sag, whereas the tidally disrupted binaries that produce HVSs originally come from $\\gtrsim\\,2\\,$pc \\citep{2003ApJ...599.1129Y,2006astro.ph..6443P}. The observations of the young disk of stars \\citep{2006ApJ...643.1011P}, as well as the cluster of S-stars, suggest that the population of stars at 0.1\\,pc from \\sag\\ should be younger and more massive than at $2\\,$pc.  This is further supported by the steep drop in B-stars outside of 0.5\\,pc \\citep{2006ApJ...643.1011P}. In addition, binaries with B-stars ($t_{\\rm ms} \\lesssim 300\\,$Myr) may also not be fully relaxed as has been assumed in many previous calculations, and therefore the diffusion rate of the most massive binaries into the loss-cone may be lower by orders of magnitudes.  In the survey of \\citet{2006ApJ...647..303B}, the observed HVSs have the same colour as blue horizontal branch stars (BHBs) and have yet to be distinguished from B-stars with high resolution spectroscopy.  The calculations shown here suggest that low mass stars, such as BHBs and their progenitors, can become HVSs with high efficiency if they are in the inner 0.1\\,pc of \\sag. Although tidal disruption rates may be enhanced by massive perturbers, BHBs must undergo significant mass loss, causing them to widen or even disrupt, and therefore are unlikely to be found in the tight binaries which are disrupted \\citep{2006astro.ph..6443P}. The spectroscopic identification of one hypervelocity BHB star would be strong evidence in support of the mechanism presented here, or an inspiralling massive BH \\citep{2003ApJ...599.1129Y,2005astro.ph..8193L,2006astro.ph..7455B}. However, the expected rate of hypervelocity BHBs depends on uncertain details of stellar evolution and goes beyond the scope of this work.  A population of BHBs ejected through this mechanism may be spun up through repeated encounters in the BH cusp \\citep{2001ApJ...549..948A}, as well as during the final encounter which ejects them (H.\\ Perets, private communication).  Therefore, a rotating BHB star may appear to be more like a B-star without detailed spectroscopic evidence.  In fact, the HVS HE 0437-5439 \\citep{2005Apj...634L.181E}, should not be ruled out as being a BHB star based on its required spin value alone.  There is a curious link between the number of HVSs observed by \\citet{2006ApJ...647..303B} and the S-stars observed orbiting \\sag. Some S-stars may be the former companions to the HVSs from tidally disrupted binaries \\citep{2003ApJ...592..935G,2006MNRAS.368..221G,\\perets}. Interestingly, there exists a similar connection between HVSs scattered off of the BHs and the S-stars as well, although perhaps not one--to--one. For every strong encounter that produces a HVS, there is likely another encounter which can bring the star closer to \\sag, perhaps kicking out the BH instead, similar to the scenario proposed by \\citet{2004ApJ...606L..21A}.  In this scenario, the scattering of a small fraction of young stars from many previous disks into the inner $0.04\\,$pc may result in a population similar to the S-stars (O'Leary, R.\\ M.\\ \\& Loeb, A.\\ 2007, in preparation).  However, we should note that if the S-star are in fact the remnants of tidally disrupted binaries, then this mechanism will only be secondary, as it is only ejects a small fraction of such stars.  There is obviously considerable uncertainty in our models. The rate of diffusion of stars both close to and far from \\sag\\ into eccentric orbits is important to understanding both the mechanisms as well as the source of the observed B-type HVSs in the Galactic halo. In both cases, it is most likely that the stars formed on relatively circular orbits, whether in a disk around \\sag\\ or in an inspiralling cluster. Large scale simulations similar to those already done by \\citet{2006astro.ph..3280F} can help resolve the uncertainties of relaxation, and with modifications, may be able to account for the conditions of continuous star formation over long periods of time. For low mass stars, physical collisions may deplete the total number of stars, and reduce the rate.  In our discussion we have not considered the effects of resonant relaxation \\citep{1996NewA....1..149R,2006ApJ...645.1152H} near \\sag, which may have two counteracting effects on our rate calculation.  Resonant relaxation may flatten the cusp and deplete the number density of stars and BHs in the innermost 0.01\\,pc; at the same time, it may also drive more massive stars into the same region producing more B-type HVSs \\citep{2006ApJ...645.1152H}.  In our analysis, we also neglected the migration of massive objects near \\sag.  If, as suggested by \\citet{2006ApJ...641..319P}, many intermediate mass BHs (IMBHs with masses $ \\gtrsim 10^3\\,\\msun$) populate the inner pc of the Galactic centre and merge with \\sag\\ every $\\sim 10^7$--$10^8\\,$yr, then the cusp of stellar-mass BHs and stars would not regenerate fast enough to produce HVSs through BH-star encounters \\citep{2006astro.ph..7455B} but instead could produce them through IMBH-star encounters \\citep{2005astro.ph..8193L}.  BHs with masses between most stellar mass BHs and IMBHs (e.g. $\\sim 50\\,\\msun$), may also form through standard binary evolution \\citep{2004ApJ...611.1068B}.  Such BHs could considerably increase the rate of HVSs, even if they are relatively rare."
    },
    "0609/astro-ph0609441.txt": {
        "abstract": "{This is the third of a series of papers on investigating the non-gravitational effects in the motion of long-period comets. The influence of the non-gravitational effects on the original and future orbital elements, and in particular on reciprocals of semimajor axes is analyzed.The new orbit determinations were performed in a homogenous way using the  basic collection of 50 nearly-parabolic comets of first class orbits. The data consist of 23 hypothetically hyperbolic comets (sample~I), 8~Oort spike comets (sample~II), and 19~'old' comets (sample~III). The original and future reciprocals of semimajor axes for each comet were calculated for strictly gravitational case and the non-gravitational case in a fully consistent way. Clear shift of the original reciprocals of semimajor axes toward the more elliptical orbits is demonstrated. The average values of the original and future energy changes due to the non-gravitational acceleration are derived for each of three samples independently. Omission of the NG~effects significantly affects the position and width of the Oort spike. The shift of the Oort spike position from the standard $1/a_{\\rm ori}$ value of 2--4~$\\cdot 10^{-5}$\\,AU$^{-1}$ to the value greater than $5.5\\cdot 10^{-5}$\\,AU$^{-1}$ is discussed.}  % ",
        "introduction": "More than half century ago Oort \\cite*{Oort} postulated that at distances of tens of thousands AU existed roughly spherical cloud of comets. He has built his famous idea on just 19 well determined cometary orbits. This classical vision of Oort cloud remains still valid on the basis of almost four hundred %high-quality orbits of nearly-parabolic comets discovered up to now (Marsden and Williams Catalogue of Cometary Orbits, \\cite*{MWC}, hereafter MW~Catalogue)\\footnote{Almost all orbits of nearly parabolic comets are derived as purely gravitational}. Almost 30\\% of long-period (LP) comets have $E_{\\rm ori}\\equiv 1/a_{\\rm ori} < 10^{-4}$\\,AU$^{-1}$ where about one fourth of them %(38~objects, see Tab.~\\ref{tab:new}) have negative $E_{\\rm ori}$. However, recent determinations of the typical $E_{\\rm ori}$ gave significantly smaller values than original indicated by Oort. Oort estimated that most of nearly parabolic comets have semimajor axes of about 25--75 thousands AU.  %\\footnote{The sign convention for binding energy, E, is the opposite %of that used for orbital energy per unit mass} %$-{\\rm G}\\cdot {\\rm M}_{\\odot }\\cdot (2{\\rm a})^{-1}$  The median semimajor axis of the 109~comets with the first quality orbits and $0 < E_{\\rm ori} < 10^{-4}$\\,AU$^{-1}$ selected from the MW~Catalogue is 27\\,800\\,AU where all orbits are strictly gravitational. If all available non-gravitational (NG) orbits of $0 < E_{\\rm ori} < 10^{-4}$\\,AU$^{-1}$ are used the sample increases to 120~comets and the median semimajor axis declines to 26\\,000\\,AU. This result strongly suggests that taking the NG~effects into account may reduce the Oort Cloud distance from the Sun significantly below 50\\,000\\,AU.  In the above estimates the comets with a negative value of the binding energy were excluded. These objects could be interstellar intruders passing through the inner solar system. However, the modern estimates of the expected rate for extrasolar comet detection give at most a few on average every 450 yrs \\cite{SeR93,Hug91}. Thus the majority of 'hiperbolic' comets have to have local origin. This discrepancy between observations and theoretical expectations disappears when the NG~effects are included for orbit determination (Marsden et al. \\cite*{MSY,MSE}, Kr{\\'o}likowska \\cite*{Kro1}(Paper~I) \\cite*{Kro2}(Paper~II)). It is due to a fact, that NG~forces make 'hiperbolic' and elliptical osculating orbits on the average less eccentric than in the case of purely gravitational motion. Consequently, omission of the NG~effects may significantly affect the position and width of the Oort spike. It means also that the apparent inner edge of the Oort cloud should be shifted towards the more positive $1/a_{\\rm ori}$. In the present paper I try to give quantitative estimates for the magnitude of these shifts.  In Paper~I the sample of 33 comets with negative value of reciprocals of original semimajor axis were analyzed and for 16 of them the NG-effects could be determined. It was shown that for 14 of these 16 comets the incoming barycentric orbits changed from hyperbolic to ecliptic when the NG~effects were included. In sections~2~and~3 of the present paper the analysis of the NG~effects for the almost complete sample of comets that have the catalogue value\\footnote{By catalogue value I mean the value given in MW~Catalogue of Cometary Orbits, 2005} of $1/a_{\\rm ori} < 10^{-5}$AU$^{-1}$ is performed. For those comets the term 'dynamically new' (in short 'new') is applied, while comets with $10^{-5}$AU$^{-1}< 1/a_{\\rm ori} < 10^{-4}$AU$^{-1}$ are defined as the 'Oort spike' comets. This differentiation between 'new' and 'Oort spike' comets, however, does not imply that the Oort spike comets cannot be dynamically new objects. In fact, simulations by Dybczynski \\cite*{Dyb05} have shown that the value of $4\\cdot 10^{-5}$AU$^{-1}$ suggested previously by Marsden and Sekanina \\cite*{MS} is the most appropriate for definition of dynamically new comets.  In Section~4 two other samples of LP comets with NG~effects have been constructed. The first sample contains comets with the catalogue orbit inside the Oort spike (8 objects), the second -- comets with catalogue $1/a_{\\rm ori} > 10^{-4}$AU$^{-1}$ (19 objects). It is important to stress that the NG~effects were derived simultaneously with osculating orbital elements from the positional data taken near perihelion passage and were then consistently applied to determine the original/future orbits and original/future reciprocals of semimajor axis. The variations of the original and future energy due to the NG~effects are widely discussed in Section~6, where the average values of differences between the original orbital energies derived from NG~orbit and strictly gravitational orbit, $\\Delta E_{\\rm ori} = 1/a_{\\rm ori, NG} - 1/a_{\\rm ori, GR}$, were determined separately for all three basic samples. These average values of $\\Delta E_{\\rm ori}$ allow to correct the observed Oort spike for the NG-effects (section~7).  Modern simulations of the Oort Cloud evolution predict that the observed sample of new comets that reach the inner solar system is biased to objects with the semimajor axes greater than 20\\,000--30\\,000\\,AU because of the 'Jupiter barrier'. New comets must have decreased their perihelion distances from more than 10\\,AU to less than 3\\,AU in one orbital period. Otherwise, they had a chance to encounter Jupiter and/or Saturn during their earlier evolution, which would remove them from the Oort spike. Heisler \\cite*{He90} has derived that the binding energies of Oort spike comets are peaked at $E_{\\rm ori} = 3.5\\cdot 10^{-5}$\\,AU$^{-1}$ ($a \\sim 29\\,000$\\,AU), assuming a local density of 0.185\\,M$_{\\odot}$/pc$^3$. Taking currently accepted value of 0.1\\,M$_{\\odot}$/pc$^3$ for local density, Levison et al. \\cite*{Lev01} argued that the maximum should be rather at $\\sim $~34\\,000\\,AU, whereas the inner edge of the outer Oort cloud -- at $\\sim $28\\,000\\,AU. Thus, one should notice some discrepancy between the theoretical predictions and apparent inner edge at 20\\,000\\,AU. However, both the theoretical calculations and orbit determinations do not include the NG-effects. Hence it is open question if this discrepancy is real \\cite{Ric04}.  \\begin{table*} \\begin{center} \\caption{{\\small The samples of the investigated LP comets with first quality orbits}} % The asterisk$^*$ denotes that the most 'hyperbolic' comet C/1853~R1~Bruhns with %poorly determined catalogue orbit was removed from the sample~A and AA \\label{tab:3sa} \\vspace{0.10cm} %{\\setlength{\\tabcolsep}{1.0mm} {\\small \\begin{tabular}{cccc} \\hline \\hline %\\begin{tabular}{ccc} %~\\hspace*{5.5cm}& {\\hspace*{4.0cm}}           & {\\hspace*{4.0cm}}      \\\\ Comets             & designation & criterion                                  & Number     \\\\ &             &                                            & of comets \\\\ \\hline % &&&  \\\\ \\multicolumn{4}{c}{\\bf Basic samples (contain comets with the detectable NG effects)} \\\\ 'new'              & I           & $1/a_{\\rm ori} < 10^{-5}$\\,AU$^{-1}$                      & 23         \\\\ Oort spike         & II          & $10^{-5}$\\,AU$^{-1} < 1/a_{\\rm ori} < 10^{-4}$\\,AU$^{-1}$ &  8         \\\\ 'old'              & III         & $1/a_{\\rm ori} > 10^{-4}$\\,AU$^{-1}$                      & 19         \\\\ \\hline % &&&  \\\\ \\multicolumn{4}{c}{\\bf Global samples constructed directly from the MW Catalogue;} \\\\ \\multicolumn{4}{c}{ contain LP comets with aphelions $Q > 250$\\,AU }              \\\\ 'new'+Oort spike   & A           & $1/a_{\\rm ori} < 10^{-4}$\\,AU$^{-1}$           & 131        \\\\ 'old'              & B           & $1/a_{\\rm ori} > 10^{-4}$\\,AU$^{-1}$           & 131        \\\\ %all                & C           &                                                & 376             \\\\ \\hline % &&&  \\\\ \\multicolumn{4}{c}{\\bf Corrected  global samples} \\\\ \\multicolumn{4}{c}{ contain LP comets with aphelions $Q > 250$\\,AU }              \\\\ 'new'+Oort spike   & AA          & $1/a_{\\rm ori} < 10^{-4}$\\,AU$^{-1}$           & 144$-$7  \\\\ 'old'              & BB          & $1/a_{\\rm ori} > 10^{-4}$\\,AU$^{-1}$           & 142$+$7  \\\\ %all                & CC          &                                                & 398      \\\\ \\hline \\end{tabular}}  %}} \\end{center} \\end{table*}  \\begin{figure} %\\vspace{0.3cm} \\begin{center} \\includegraphics[width=9.0cm]{fig1nb.eps} \\end{center} %\\includegraphics[width=9.0cm]{all_376a.eps} %\\epsfig{file=fig_hip.eps,height=12.0cm,width=12.0cm} \\caption{{\\small Distributions of cometary energies, $E_{\\rm ori} \\equiv 1/a_{\\rm ori}$ (top panel), and $E_{\\rm fut} \\equiv 1/a_{\\rm fut}$ (middle panel), $\\delta E\\equiv 1/a_{fut} - 1/a_{\\rm ori}$ (bottom panel). The filled histograms show the respective energy distributions derived for strictly gravitational orbits for all of the investigated comets (50 objects from all three samples: I+II+III; Table~\\ref{tab:3sa}). The shaded histogram drawn in the top panel represents the 11 comets with $E_{\\rm ori} < 10^{-5}$\\,AU$^{-1}$, where the NG-effects were indeterminable (part (i) of Table~\\ref{tab:new}; sample AA). The solid line histograms represent global distributions of all comets with first quality orbits and aphelions $Q>250$\\,AU taken from MW~Catalogue (sample~A+B, Table~\\ref{tab:3sa})}}\\label{fig:376} \\end{figure}  %%%\\end{document} ",
        "conclusions": "In the present study the NG~effects are derived for 50~LP~comets, where 31 of them have $E_{\\rm ori, GR} < 10^{-4}$\\,AU$^{-1}$. It was confirmed that NG~forces make 'hiperbolic' and elliptical osculating orbits less eccentric than by assuming purely gravitational motion. Such number of comets with the detectable NG~effects allow us to give some estimates of the average value of $\\Delta E_{\\rm ori} = 1/a_{\\rm ori, NG} - 1/a_{\\rm ori, GR}$ for three subsamples of the LP~comets. Due to large dispersion of the $\\Delta E_{\\rm ori}$ distributions (within each subsample) the median values of $\\Delta E_{\\rm ori}$ seem to represent more adequately typical corrections for $E_{\\rm ori, GR}$ of the LP~comets with the non-detectable NG~effects.  Applying these corrections for the all LP~comets, the estimates of the Oort spike shift from its classical position (i.e. position derived on the basis of pure gravitational model of comet's motion) have been obtained.  Unfortunately, the available data warrant only simplified type of models of the $1/a$~corrections. In the first type it was assumed that the NG~effect are present only for $q<3.1$\\,AU (Fig.~\\ref{fig:spike}), while in the second type no upper limit for $q$ was imposed (Fig.~\\ref{fig:spike2}). Although one could expect smooth decline of the NG~effects with the increase of the perihelion distance, no such relationship is implied by the observational data. In particular it is not present in the our basic sample of 50~comets (the bottom panel in Fig.~\\ref{fig:mean1} and section 6.1). On the other hand, the active comets at large perihelion distances have been observed (e.g. C/2004~T3~Siding Spring with $q=8.9$\\,AU, C/2000~A1~Montani with $q=9.7$\\,AU, and C/2003~A2~Gleason with $q=11.4$\\,AU). In such cases another formula for the NG~acceleration due to sublimation of more volatile ices is necessary. Such modeling of the NG~effects in the LP~comets with large perihelion distance based on Yabushita CO-function \\cite{Yab96} was performed in Paper~II. However, no reliable NG~effects have been found for comets with the perihelion distance greater than 3.5\\,AU.  Using the model of the NG~effects limited to comets with $q\\leq 3.1$\\,AU, the Oort spike is transformed into significantly more spread and less peaked distribution with the maximum at $E_{\\rm ori, NG} = 5.5 \\cdot 10^{-5}$AU$^{-1}$, whereas if no upper limit on $q$ is imposed the Oort spike peaks near $E_{\\rm ori, NG} \\simeq 8.5\\cdot 10^{-5}$\\,AU$^{-1}$ (middle panels of Figs.~\\ref{fig:spike} and \\ref{fig:spike2}). Assuming conservatively $E_{\\rm ori, NG} = 5.5 \\cdot 10^{-5}$AU$^{-1}$ the outer Oort Cloud distance from the Sun is approximately 37\\,000\\,AU. Consequently the inner edge of the Oort Cloud is situated at a distance significantly smaller than the 'classical' 20\\,000\\,AU assumed in the literature. The middle panel of Fig.~\\ref{fig:spike} suggests that this inner edge could be less spectacular than in the classical point of view, i.e. in the strictly gravitational case (top panel of Fig.~\\ref{fig:spike}). %suggested by simulations of dynamical evolution.  On the other hand, the theoretical simulations predict that the energies of new comets should peak near $E_{\\rm ori}\\approx 2.9 \\cdot 10^{-5}$AU$^{-1}$, i.e. at a semimajor axis of about 34\\,000\\,AU for the currently accepted value of a local mass density of 0.1~M$_{\\odot}$/pc$^3$ \\cite{Lev01}. Assuming that Sun formed in a denser galactic environment than it occupies now, it is possible to obtained the Oort Cloud bound more tightly \\cite{Egg99,FeB00}. For example, taking an initial local density of 100 times greater than the current local density, and assuming that the early Sun spent in such environment first 20~m.y. yrs, Eggers \\cite*{Egg99} modeled the Oort cloud formation. She has found that the tightly bound cloud of comets might have been formed at a heliocentric distance of about 6--7~thousand\\,AU.  It is obvious that such significant reduction of the Oort cloud distance as postulated in the present paper has a deep consequences in many aspects of theoretical modeling of the Oort Cloud formation as well as simulations of a long-term evolution of the Oort Cloud comets, and cometary fading and disruption. In the last decades, all these three blocks of issues have been intensively explored (see reviews by Dones et al. \\cite*{DWLD04}, and Rickman~\\cite*{Ric04}). However, evolutionary correlations between the observed cometary populations and their reservoirs are still incomplete and are inconsistent in many aspects \\cite{EBa98,NVZR,Ric05} %(Emel'yanenko \\& Bailey~\\cite{EBa98},Nurmi et al. \\cite{NVZR}, Rickman~\\cite{Ric05}). One of the most important problems raised by the recent simulations %(Levison et al. \\cite{Lev01}, Duncan et al. \\cite {DQT87}, Dones et al. \\cite{DWLD04}) \\cite{Lev01,DQT87,DWLD04} are correlations between the outer and inner Oort Cloud populations and their characteristics. Still more realistic simulations are needed.   %\\end{document}"
    },
    "0609/astro-ph0609270_arXiv.txt": {
        "abstract": "We investigate the relation between the mass of super-massive black holes ($\\mbh$) in QSOs and the mass of the dark matter halos hosting them ($\\mdh$).  We measure the widths of broad emission lines (\\mgii\\ \\lam2798, \\civ\\ \\lam1549) from QSO composite spectra as a function of redshift.  These widths are then used to determine virial black hole mass estimates.  We compare our virial black hole mass estimates to dark matter halo masses for QSO hosts derived by Croom et al. (2005) based on measurements of QSO clustering.  This enables us to trace the $\\mbh-\\mdh$ relation over the redshift range $z=0.5$ to 2.5. We calculate the mean zero-point of the $\\mbh-\\mdh$ relation to be $\\mbh=10^{8.4\\pm0.2}\\,{\\rm \\msun}$ for an $\\mdh=10^{12.5}\\,{\\rm \\msun}$.  These data are then compared with several models connecting $\\mbh$ and $\\mdh$ as well as recent hydrodynamical simulations of galaxy evolution. We note that the flux limited nature of QSO samples can cause a Malmquist-type bias in the measured zero-point of the $\\mbh-\\mdh$ relation. The magnitude of this bias depends on the scatter in the $\\mbh-\\mdh$ relation, and we reevaluate the zero-point assuming three published values for this scatter.   We create a subsample of our data defined by a constant magnitude interval around $L^{*}$ and find $(1+z)^{3.3\\pm1.3}$ evolution in $\\mbh$ between $z\\sim0.5$ and $2.5$ for typical, $L^{*}$ QSOs. We also determine the Eddington ratios ($L/\\ledd$) for the same subsample and find no significant evolution: $(1+z)^{-0.4\\pm1.1}$. Taken at face value, our data suggest that a decrease in active black hole mass since $z\\sim2.5$ is the driving force behind luminosity evolution of typical, $L^{*}$, optically selected AGN. However, we note that our data are also consistent with a picture in which reductions in both black hole mass and accretion rate contribute equally to luminosity evolution. In addition we find these evolution results are strongly affected by the virial black hole mass estimators used. Changes to the calibration of these has a significant effect on the evolution results. ",
        "introduction": "\\label{sec:intro}  Over recent years it has become apparent that massive black holes lie at the centre of the majority of local galaxies. In addition, correlations between the mass of these central black holes ($\\mbh$) and host galaxy properties suggests that the evolution of galaxies must be intimately related to the growth of their central black holes.  Galaxy properties which display a correlation with $\\mbh$ include spheroid luminosity \\cite{mag98}, spheroid mass \\cite{k+r95,fer02} and stellar velocity dispersion \\cite{f+m00,geb00}.   In this paper we investigate the relationship between the mass of QSO black holes and the mass of the dark matter halos that host them. Dark matter halo masses derived from QSO clustering were calculated for QSOs in the 2dF QSO Redshift Survey (2QZ; Croom et al. 2004) by Croom et al. (2005). These were evaluated for QSOs in 10 redshift bins from $z\\sim0.5$ to $2.5$. We take the identical sample to that used by Croom et al. (2005) and calculate average virial black hole masses for the QSOs in each redshift bin. This is achieved by constructing a composite spectrum for each redshift bin from all of the QSOs in that bin. The average virial black hole masses are then estimated from the widths of broad emission lines in the composite spectra.   In Section~\\ref{sec:bhmass} we describe the virial black hole mass estimators used in this paper, section~\\ref{sec:d+a} describes our data and analysis including composite making and line width measurement. In Section \\ref{sec:dhmass} we briefly review the clustering measurements which are used to derive dark halo host mass.  Our results are presented in Section \\ref{sec:res} and we present our conclusions in Section \\ref{sec:conc}.  Throughout this paper we assume a flat $(\\Omega_{\\rm m},\\Omega_{\\Lambda})=(0.3,0.7)$, $H_{0}=70\\,{\\rm km\\,s}^{-1}\\,{\\rm Mpc}^{-1}$ cosmology. ",
        "conclusions": "\\label{sec:conc}  We make composite spectra of QSOs from the 2QZ to find average virial black hole mass estimates for the QSOs in 10 redshift bins for which Croom  et al.~(2005) had already calculated $\\mdh$ via clustering analysis. Comparing the black hole and dark halo masses we find evidence for $\\sim(1+z)^{2.5\\pm1.8}$ evolution in the $\\mbh-\\mdh$ relation, although large errors are such that we can not exclude the possibility of there being no evolution. We derive the zero-point of the $\\mbh-\\mdh$ relation (averaged over redshift) and find it to be $\\mbh=10^{8.4\\pm0.2}\\,{\\rm \\msun}$ for a dark matter halo of mass $\\mdh=10^{12.5}\\,{\\rm \\msun}$.  This is most consistent with a model using a Seljak (2004) dark matter profile (under the assumption of no evolution in $\\mbh-\\mdh$), however, uncertainties are such that we are unable to definitively distinguish which model is preferred.  We compare our measured $\\mbh-\\mdh$ relation to that derived from hydrodynamical simulations of galaxy evolution \\cite{dsh05,rob05} and find good agreement.  We note that because QSOs are selected above a given luminosity this will tend to select objects above a given $\\mbh$.  This results in a Malmquist-type bias such that the observed mean $\\mbh$ will lie above the true $\\mbh-\\mdh$ relation.  The level of bias is crucially dependent on the amount of scatter in the $\\mbh-\\mdh$ relation.  We take a subsample of QSOs in a constant magnitude interval around $\\mb^{*}$ and find significant evolution in their black hole masses characterised by $\\mbh\\propto(1+z)^{3.3\\pm 1.1}$. Comparing this to the observed lack of significant evolution in Eddington ratio ($L/\\ledd\\propto(1+z)^{-0.4\\pm1.1}$) we conclude that luminosity evolution of $L^*$ QSOs is driven primarily by decreasing black hole masses between redshifts 2.5 and 0.5. However, the exact combination of evolution in $\\mbh$ and $L/\\ledd$ is dependent on the slope of the luminosity dependence in the virial mass estimator and any luminosity/redshift dependence in the velocity width of the QSO broad lines as well as the calibrations of the virial mass estimators themselves. Considering this and potential sources of systematic errors in our line width measurements, we find that our data are still consistent with a picture in which both reducing black hole masses and Eddington ratios play an equal role in $L^*$ QSO luminosity evolution as observed in other studies (e.g. Merloni 2004; Heckman et al. 2004).   To extend this work further, detailed analysis of samples with a broader dynamic range will be required, including the analysis of QSO clustering as a function of luminosity.  This has started to be done with small samples (e.g. Adelberger \\& Steidel 2005), but will be extended with new faint QSO surveys such as the 2dF-SDSS LRG and QSO (2SLAQ) Survey (Richards et al. 2005)."
    },
    "0609/astro-ph0609793_arXiv.txt": {
        "abstract": "We explore and discount the hypothesis that the strong, continual H$\\alpha$-emitting T dwarf 2MASS~J12373919+6526148 can be explained as a young, low gravity, very low mass brown dwarf.  The source is already known to have a marginally-fainter absolute magnitude than similar T dwarfs with trigonometric parallax measurements, and has a tangential velocity consistent with old disk kinematics.  Applying the technique of Burgasser, Burrows \\& Kirkpatrick on new near infrared spectroscopy for this source, estimates of its {\\teff}, $\\log{g}$ and metallicity ([M/H]) are obtained.  2M~1237+6526 has a {\\teff} $\\approx$ 800-850~K.  If [M/H] is solar, $\\log{g}$ is as high as $\\sim$5.5 (cgs) and this source is older than 10 Gyr. We find a more plausible scenario to be a modestly subsolar metallicity ([M/H] = -0.2) and moderate $\\log{g}$~$\\sim$~5.0, implying an age older than 2~Gyr and a mass greater than 0.035~M$_{\\sun}$.  The alternative explanation of the unique emission of this source, involving an interacting, close, double degenerate system, should be investigated further.  Indeed, there is some evidence of a {\\teff} $<$ 500~K companion to 2M~1237+6526 on the basis of a possible $Spitzer~IRAC$ [3.6]--[4.5] color excess.  This excess may, however, be caused by a subsolar metallicity. ",
        "introduction": "It is well known that the frequency and relative luminosity of H$\\alpha$ as a likely measure of chromospheric activity drops rapidly beyond spectral types M7~V (Kirkpatrick et al. 2000, Gizis et al. 2000, Mohanty \\& Basri 2003, West et al. 2004).  Very few T dwarfs have exhibited this emission at all, and generally at a very low level compared to active mid-M dwarfs, as measured by the ratio of L$_{H\\alpha}$ / L$_{bol}$ (Burgasser et al.\\ 2003).  Occasionally, a vigorous flare is observed in a late M or L dwarf, causing the ratio to rocket upwards by orders of magnitude.  For the most part, however, magnetic pressures built up by the likely dynamo inside these rapidly-rotating objects are not expressed at the surface, due to the inability of interior plasmas to penetrate the high resistivities of the predominantly neutral atmospheres (Mohanty \\& Basri 2003).  Two exceptions to this behavior have been found, exhibiting in all observations levels of H$\\alpha$ emission 1-2 orders of magnitude higher in L$_{H\\alpha}$ / L$_{bol}$ than any counterparts of similar spectral type.  These are the M9.5 dwarf PC~0025+0047 (Schneider et al. 1991) and the T6.5 dwarf 2MASS~J12373919+6526148 (hereafter 2M~1237+6526, Burgasser et al.\\ 1999).  Continual H$\\alpha$ activity of the former source has been observed for over a decade (Mould et al. 1994; Mart{\\'{\\i}}n, Basri, \\& Zapatero Osorio 1999).  That the T dwarf seems to show similar characteristics, though over a shorter time interval so far, has suggested that emission processes in the two objects may be related.  Mart{\\'{\\i}}n et al (1999) made the case that the strong emission lines, variable amount of optical veiling, and a claimed detection of the Li~6707\\AA\\ doublet indicated that PC~0025+0447 is a substellar object less massive than 0.06M$_{\\bigodot}$ and younger than $\\sim$1~Gyr. Their detection of weak K~I and Na~I resonance doublets indicated a low surface gravity.  They argued, however, that the lack of an infrared excess and the fact that the H$\\alpha$ profile is not similar to classical T Tauris suggests that chromospheric activity rather than an active accretion disk is the likely source of the intense emission.  If the same scenario applies to 2M~1237+6526, this object could be a very low mass, young, very active brown dwarf.  An age of 1~Gyr or less implies a maximum mass of 30~M$_J$ (Burrows et al.\\ 2001) for an effective temperature {\\teff} $\\approx$ 800~K based on the Golimowski et al.\\ (2004) scale (see their Fig.~6).  We have explored an alternative hypothesis for this source: that it is an interacting brown dwarf or double degenerate binary (Burgasser et al. 2000a, 2002a).  Some possible predictions in the former paper were tested weakly in the latter paper.  Monitoring of the J-band flux failed to detect significant variability at the $\\pm$0.025 mag level for periods of hours.  Three separate measurements of the H$\\alpha$ line made at different times with the Keck/LRIS spectrograph between 1999 and 2001 (over a span of 1.6 yr) were inconclusive in establishing or ruling out variability in the radial velocity.  Thus, no evidence was uncovered in favor of the interacting binary hypothesis.  On the other hand, it has not been ruled out.  In this paper, we explore the first hypothesis: that 2M~1237+6526 is young, has low gravity, and a very active chromosphere.  We begin by noting that there is already evidence against this in the literature (\\S~2).  We present additional evidence based on infrared spectrophotometry in \\S~3.  Temperature and gravity sensitive infrared flux indices are compared with atmospheric models to estimate T$_{eff}$, $\\log{g}$ and the atmospheric abundances ([M/H]) for the emission line object, as well as age and mass, in \\S~4. Constraints on the properties of the hypothesized companion to 2M~1237+6526 are discussed in \\S~5, based in part on an apparent mid-infrared color excess detected by Spitzer observations.  Results are summarized in \\S~6. ",
        "conclusions": "On the basis of our spectral analysis, as well as the object's luminosity and kinematics, we conclude that 2M~1237+6526 is likely to be a rather old, massive and slightly metal-poor brown dwarf.  Youthful activity, which is cited as the source for the emission properties of PC~0025+0047, cannot explain this T dwarf's persistent H$\\alpha$ emission.  Other possibilities, including accretion from a close companion, are worth further investigation.  The ``tests'' of the double degenerate hypothesis involving a low rate of transfer have not been conclusive. The problem with searching for possible radial velocity variations was described in the Introduction.  The absence of large infrared variations sought in the Burgasser et al. (2002a) study rules out only a limited phase space for the interacting binary model.  The $IRAC$ photometry suggests a [3.6]--[4.5] color excess which could result from the presence of an unresolved cool companion, but the excess may also be attributable to a lower metallicity.  Any companion would have to be appreciably lower in mass, and probably at least 300~K cooler. Generally speaking, as model atmospheres are improved to predict more accurate 4.5~$\\mu$m fluxes, including the dependence on metallicity, searching for an excess in this band may result in the detection of unresolved companions later in type than any currently-known T dwarf."
    },
    "0609/astro-ph0609720_arXiv.txt": {
        "abstract": "We investigate the importance of the far red wing of the Lyman $\\,\\alpha$ line of hydrogen in the atmospheres of cool white dwarfs of pure hydrogen composition. We find that this absorption process dominates all important sources of opacity in the blue part of the optical spectrum of these stars. Our successful fits to the spectra of cool DA/DC white dwarfs indicate that the far red wing of the $\\rm Ly \\, \\alpha$ line is the source of opacity that had been missing in the models. The observed sequence of cool white dwarfs in color-color diagrams is very well reproduced by our new pure hydrogen atmosphere models, suggesting that the atmospheric composition of the coolest DC white dwarfs must be revisited. ",
        "introduction": "The detection and analysis of large samples of cool white dwarfs over the past decade has led to a significant improvement in our knowledge of the chemical composition and evolution of their atmospheres \\citep{Har06, Kilic06, BLR01, BRL97}. The difficulties in interpreting the spectra of some very cool white dwarfs thought to posses helium-rich atmospheres with small amounts of H \\citep{Bergeron05,BL02,OP01} are most probably related to the extreme physical conditions found in their atmospheres \\citep{Kowalski06a}. On the other hand, the physical conditions in hydrogen-rich atmospheres are much less extreme and the current models should be much more reliable. Nonetheless, models systematically overestimate the flux for $\\lambda \\wig<5000 \\,$\\AA\\ of hydrogen-rich white dwarfs of $T_{\\rm eff}\\rm <6000 \\, K$ \\citep{BLR01,Bergeron01,BRL97}. \\citet{BRL97} suggested that the mechanism responsible for this phenomenon could be the pseudo-continuum bound-free opacity from hydrogen atoms perturbed by collisions in a relatively dense medium. However, a realistic calculation of this opacity source shows that it is too weak to correct the observed blue flux excess in the models \\citep{Kowalski06b}.  In this paper we consider the possibility that the missing opacity is the red wing of the pressure broadened hydrogen Lyman$\\, \\alpha$ line.  This absorption mechanism results from the perturbations of hydrogen atoms by their interaction with other particles, primarily $\\rm H$ and $\\rm H_2$. Such perturbations result in the lowering of the $\\rm Ly \\, \\alpha$ transition energy and the possibility of a bound-bound transition for photon energies that are smaller than that of an isolated hydrogen atom ($E_{\\rm 12}^{\\rm 0}\\rm =10.2 \\, eV$). We present a new, semi-classical calculation of this process based on accurate quantum mechanical calculations and a model that has no free parameters. The inclusion of broadening by collisions with $\\rm H_2$ is the key element of this calculation. The resulting pure hydrogen models successfully reproduce the spectral energy distributions (SED) and colors of very cool DA and DC white dwarfs.  \\begin{figure}[t] \\epsscale{0.7} \\plotone{f1.eps} \\figcaption[f1.ps]{ The interaction energy curves as a function of the inter-particle collision distance for the ground state (lower curves) and the first excited Rydberg state (upper curves) of $\\rm H-H_2$ (solid) and $\\rm H-H$ (dotted) dimers. For $\\rm H-H_2$ the energy curves shown are for a collision angle of $\\rm 90^{o}$, where the collision angle is defined between the line connecting the $\\rm H$ atom to the center of $\\rm H_2$ and the molecular axis. \\label{F1}} \\end{figure} ",
        "conclusions": "The existence of an unidentified absorption mechanism in cool hydrogen white dwarfs atmospheres was reported a decade ago \\citep{BRL97}. The interpretation of this missing opacity as the pseudo-continuum absorption from hydrogen atoms has been recently shown to be incorrect \\citep{Kowalski06b}. On the other hand, the red wing of the $\\rm Ly \\, \\alpha$ line opacity from hydrogen could provide the required absorption \\citep{Wolff02}. We present a new calculation of the extreme pressure-broadening of the $\\rm Ly \\, \\alpha$ line by both $\\rm H$ and $\\rm H_2$. When included in our new pure hydrogen atmosphere models, we obtained an excellent agreement with the UV/optical/near-IR spectral energy distribution (SED) of the cool DA white dwarf BPM 4729 and we successfully fitted the B through K spectral energy distributions of stars with hydrogen-rich atmospheres, as determined by \\citet{BRL97} and \\citet{BLR01}. The inclusion of broadening by collisions with H$_2$ is essential to reproduce the SED of very cool hydrogen white dwarfs.  In color-color diagrams, the new pure hydrogen models follow the observed sequences of cool white dwarfs very well, while an improved physical description of dense helium moves the cool end of the helium sequence to the red, near the locus of black body colors. This suggest that detailed fitting of the SED of the coolest DC white dwarfs will result in a number of them, and perhaps most, having hydrogen-rich composition rather than helium-rich.  An example is provided by our fit to WD$2054-050$.  This has direct implications for the spectral evolution of very cool white dwarfs and for the physical mechanisms responsible for their atmospheric composition.  This research was supported by the United States Department of Energy under contract W-7405-ENG-36. We thank P. Bergeron for illuminating discussions.  \\small"
    },
    "0609/astro-ph0609499_arXiv.txt": {
        "abstract": "We have analysed the parameters characterising the mass and velocity distributions of two samples of relaxed elliptical-like-objects (ELOs) identified, at $z=0$, in a set of self-consistent hydrodynamical simulations operating in the context of a concordance cosmological model. ELOs have a prominent, non-rotating, dynamically relaxed stellar spheroidal component, with very low cold gas content, and sizes of no more than $\\sim $ 10 - 40  kpc (ELO or baryonic object scale), embedded in a massive halo of dark matter typically ten  times larger in size (halo scale). They have  also an extended halo of hot diffuse gas. The parameters characterising the mass, size and velocity dispersion both at the baryonic object  and at the halo scales have been measured in the ELOs of each sample. At the halo scale they have been found to satisfy virial relations; at the scale of the baryonic object the (logarithms of the) ELO stellar masses, projected stellar half-mass radii, and stellar central l.o.s. velocity dispersions define a flattened ellipsoid close to a plane (the intrinsic dynamical plane, IDP), tilted relative to the virial one, whose observational manifestation is the observed FP. Otherwise,  IDPs are  not homogeneously populated, but ELOs, as well as elliptical (E) galaxies in the FP, occupy only a particular region defined by the range of their masses. The ELO samples have been found to show systematic trends with the mass scale in both, the relative content and the relative distributions of the baryonic and the dark mass ELO components, so that homology is broken in the spatial mass distribution (resulting in the IDP tilt), but ELOs are still a two-parameter family where the two parameters are correlated (causing its non-homogeneous population). The physical origin of these trends presumably lies in the systematic decrease, with increasing ELO mass, of the relative amount of dissipation experienced by the baryonic mass component along ELO stellar mass assembly. ELOs also show kinematical segregation, but it does not appreciably change with the mass scale.  The non-homogeneous population of IDPs explains the role played by $M_{\\rm vir}$ to determine the correlations among intrinsic parameters. In this paper  we also show that the central stellar line-of-sight velocity dispersion of ELOs, $\\sigma_{\\rm los, 0}^{\\rm star}$, is a fair empirical estimator of $M_{\\rm vir}$, and this explains the central role played by $\\sigma_{\\rm los,0}$ at determining the observational correlations. ",
        "introduction": "\\label{Intro}  Understanding how local galaxies of  different Hubble types we observe to-day have formed is one of the most challenging open problems in cosmology. Among the different galaxy families, elliptical (Es) are the easiest to study and those that  show the most precise regularities in their empirical properties, some times in the form of tight correlations among their observable parameters. The interest of these regularities lies in that they could encode a lot of relevant informations on the physical processes underlying E formation and evolution.  The Sloan digital sky survey (SDSS, York et al. 2000) has substantially improved the statistics on  E samples. The sample selected by Bernardi et al. (2003a), using morphological and spectral criteria, contains 9000 Es  to date in the redshift range $0.01 \\le z \\le 0.3$ and in every environment from voids to groups to rich clusters. Analyses of their structural and dynamical parameters have shown that the distributions of their luminosities $L$, radii at half projected light $R_{\\rm e}^{\\rm light}$, and central line-of-sight  velocity dispersions $\\sigma_{\\rm los, 0}$ (Bernardi et al. 2003b, 2003c), are approximately gaussian at any $z$. Moreover, a maximum likelihood analysis indicates that the pairs of parameters $\\sigma_{\\rm los, 0}$---$L$ and $R_{\\rm e}^{\\rm light}$---$L$,  or their combinations, such as the mass-to-luminosity ratio within the effective radii  $M_{\\rm e}/L$ and $L$ (where $M_{\\rm e}$ is the dynamical mass defined as $M_{\\rm e} = 2 R_{\\rm e}^{\\rm light} \\sigma_{\\rm los, 0}^{2}/G$), show correlations consistent  with those previously established in literature, obtained from individual galaxy spectra of smaller samples, such as the Faber-Jackson relation (1976); the $D_n$---$\\sigma_{\\rm los, 0}$ relation (Dressler et al. 1987); and the surface brightness ---$R_{\\rm e}$ relation (Kormendy 1977, Kormendy \\& Djorjovski 1989), among others. Furthermore,  early-type galaxies in the SDSS have been found to have roughly constant stellar-mass-to-light ratios (Kauffmann et al. 2003a, 2003b; Padmanabhan et al. 2004).  The correlations involving two variables out of the three needed to fully characterise the structure and dynamics of an E galaxy, are projections of the so-called Fundamental Plane relation (FP, Djorgovski \\& Davis 1987; Dressler et al. 1987a; Faber et al. 1987; Kormendy \\& Djorgovski 1989), involving the three or some or their combinations. The FP relation can be written as  \\begin{equation} \\log_{10} R_{\\rm e}^{\\rm light} = a \\log_{10} \\sigma_{\\rm los, 0} + b \\log_{10} <I^{\\rm light}>_{\\rm e} + c. \\label{FP} \\end{equation}  where $<I^{\\rm light}>_{\\rm e}$ is the mean surface brightness within $R_{\\rm e}^{\\rm light}$. The values of the FP coefficients for the SDSS E  sample are $a \\simeq 1.5$, similar in the four SDSS bands, $b \\simeq -0.77$, and $c \\simeq -8.7$ (see their exact values in   Bernardi et al. 2003c, Table 2) with a small scatter. These SDSS results confirm previous ones, either in the optical (Lucey, Bower \\& Ellis 1991; de Carvalho \\& Djorgovski 1992; Bender, Burstein \\& Faber 1992; J{\\o}rgensen et al. 1993; Prugniel \\& Simien 1996; J{\\o}rgensen et al. 1996) or in the near-IR wavelengths (Recillas-Cruz et al. 1990, 1991; Pahre, Djorgovski \\& de Carvalho 1995; Mobasher et al. 1999), even if the published values of $a$ show larger values in the $K$-band than  at shorter wavelengths (see, for example, Pahre, de Carvalho \\& Djorgovski 1998).  The existence of the FP and its small scatter has the important implication that it provides us with a strong constraint when studying elliptical galaxy formation and evolution (Bender, Burstein \\& Faber 1993; Guzm\\'an, Lucey \\& Bower 1993; Renzini \\& Ciotti 1993). The physical origin of the FP is not yet clear, but it must be a consequence of the physical processes responsible for galaxy assembly. These processes built up early type galaxies as dynamically hot systems whose configuration in phase space are close to equilibrium. Taking an elliptical galaxy as a system in equilibrium, the virial theorem  \\begin{equation} M_{\\rm vir} = c_{\\rm F} (\\sigma_{3, h}^{\\rm tot})^{2} r_{\\rm e, h}^{\\rm tot}/ G , \\label{Virial} \\end{equation}  where  $M_{\\rm vir}$ is its virial mass, $\\sigma_{\\rm 3, h}^{\\rm tot}$ is the average 3-dimensional velocity dispersion of the whole elliptical, including both dark and baryonic matter, $r_{\\rm e, h}^{\\rm tot}$ the dynamical half-radius or radius enclosing half the total mass of the system, and $c_{\\rm F}$ a form factor of order unity, would imply a relation similar to Eq. \\ref{FP} with $a=2$ and $b=-1$ (known as the virial plane, Faber et al. 1987), should the dynamical mass-to-light ratios, $M_{\\rm vir}/L$, and the mass structure coefficients  \\begin{equation} c_{\\rm M}^{\\rm vir} = {G M_{\\rm vir} \\over 3 \\sigma_{\\rm los, 0}^{2} R_{\\rm e}^{\\rm light}}, \\label{CMvir} \\end{equation}  be independent of the E luminosity or mass scale. The observational results described above mean that the FP is {\\it tilted} relative to the virial plane, and, consequently that either $M_{\\rm vir}/L$ or $c_{\\rm M}^{\\rm vir}$, or both, do depend on the E luminosity.  Different authors interpret the tilt of the FP relative to the virial relation as caused by different misassumptions that we comment briefly (note that we can write $M_{\\rm vir}/L =  M^{\\rm star}/L \\times M_{\\rm vir}/M^{\\rm star}$, where $M^{\\rm star}$ is the stellar mass of the elliptical galaxy).  1.1) A first possibility is that the tilt is due to systematic changes of stellar age and metallicity with galaxy mass, or, even, to changes of the slope of the stellar initial mass function (hereafter, IMF) with galaxy mass, resulting in systematic changes in the {\\it stellar}-mass-to-light ratios, $M^{\\rm star}/L$, with mass or luminosity (Zepf \\& Silk 1996; Pahre at al. 1998; Mobasher et al. 1999). But these effects  could explain at most only $\\sim$ one third of the $\\beta \\neq 0$  value in the $B$-band (Tinsley 1978; Dressler et al. 1987; Prugniel \\& Simien 1996; see also Renzini \\& Ciotti 1993; Trujillo, Burkert \\& Bell 2004). Furthermore,  early-type galaxies in the SDSS have been found to have roughly constant stellar-mass-to-light ratios (Kauffmann et al. 2003a, 2003b). Anyhow, the presence of a tilt in the $K$-band FP, where population effects are no important, indicates that it is very difficult that the tilt is caused by stellar physics processes alone, as Bender et al. (1992), Renzini \\& Ciotti (1993), Guzm\\'an et al. (1993), Pahre et al. (1998), among other authors, have suggested.  1.2) A second possibility is that $M_{\\rm vir}/L$ changes systematically with the mass scale because the total dark-to-visible mass ratio, $M_{\\rm vir}/M^{\\rm star}$ changes (see, for example, Renzini \\& Ciotti 1993; Pahre et al. 1998; Ciotti, Lanzoni \\& Renzini 1996; Padmanabhan et al. 2004).  Otherwise, a dependence of $c_{\\rm M}^{\\rm vir}$ on the mass scale could be caused by systematic differences in:  2.1) the dark  versus bright matter spatial distribution,  2.2) the kinematical segregation, the rotational  support and/or velocity dispersion anisotropy in the stellar component  (dynamical non-homology),  2.3) systematic  projection or other geometrical effects.  Taking into account these effects in the FP tilt demands modelling the galaxy mass and velocity three-dimensional distributions and comparing the outputs with high quality data.  Bender et al. (1992) considered  effects iii) and iv); Ciotti et al. (1996) explore ii) - iv) and conclude that an systematic increase in the dark matter content with mass, or differences in its distribution, as well as a dependence of the S\\'ersic (1968) shape parameter for the luminosity profiles with mass, may by themselves formally produce the tilt; Padmanabhan et al. (2004) find evidence of effect ii) in SDSS data. Other authors have also shown that allowing for broken homology, either dynamical (Busarello et al. 1997), in the luminosity profiles (Trujillo et al. 2004), or both (Prugniel \\& Simien 1997; Graham \\& Colless 1997; Pahre et al. 1998), brings the observed FP closer to Eq. (3).   Observational methods suffer from some drawbacks to study in deph the  physical origin of the FP tilt. For example, a drawback is the impossibility to get accurate measurements of the elliptical three-dimensional  mass distributions (either dark, stellar or gaseous) and  another is that  the intrinsic 3D velocity distribution of galaxies is  severely limited by projection. Only the line-of-sight velocity distributions can be inferred from galaxy spectra. And, so, the interpretation  of observational data is not always straightforward. To complement the informations provided by data and circumvent these drawbacks, analytical modelling is largely used in literature (Kronawitter et al. 2000; Gerhard et al. 2001; Romanowsky \\& Kochanek 2001; Borriello et al. 2003; Padmanabhan et al. 2004; Mamon \\& Lockas 2005a, 2005b). These methods give very interesting insights into mass and velocity distributions, as well as the physical processes causing them, but are somewhat limited by symmetry  considerations and other necessary  simplifying hypotheses. These difficulties and limitations could be circumvented should we have at our disposal complete  informations on the phase-space of the galaxy constituents. This is not possible through observations, but can be attained, at least in a virtual sense, through numerical simulations.   Capelato, de Carvalho \\& Carlberg (1995)  first addressed the origin of the FP through numerical simulations. By analyzing the remnants of the dissipationless mergers of two equal-mass one-component King models, and varying their relative orbital energy and angular momentum, they show that the mergers of objects in the FP produces a new objects in the FP. This result was extended by Dantas et al. (2003), who used one- and two-component Hernquist models as progenitors, Goz\\'alez-Garc{\\'\\i}a \\& van Albada (2003), based on Jaffe (1983) models and by Boylan-Kolchin et al. (2005), who used Hernquist+NFW models. Nipoti, Londrillo \\& Ciotti (2003) show, in turn, that the FP is well reproduced by dissipationless hierarchical equal-mass  merging of one- and two-component galaxy models, and by accretion with substantial angular momentum, with the merging zeroth-order generation placed at the FP itself. They also found that both the Faber-Jackson and the Kormendy relations are not reproduced by the simulations, and conclude that dissipation must be a basic ingredient in elliptical formation. In agreement with this conclusion, Dantas et al. (2002) and Dantas et al. (2003) have shown that the end products of  dissipationless collapse generally do not follow a FP-like correlation. Bekki (1998) first considered the role of dissipation in elliptical formation through pre-prepared simulations. He adopts the merger hypothesis (i.e.., ellipticals form by the mergers of two equal-mass gas-rich spirals) and he focuses on the role of the timescale for star formation in determining the structural and kinematical properties of the merger remnants. He concludes that the slope of the FP reflects the difference in the amount of dissipation the merger end products have experienced according with their luminosity (or mass). Recently, Robertson et al. (2006) have confirmed this conclusion on the role of dissipative dynamics to shape the FP, again  through  pre-prepared mergers of disk galaxies.  In this paper we go a step further and we analyse the tilt of the FP in samples of virtual ellipticals formed in a cosmological context through self-consistent hydrodynamical simulations. This numerical approach provide a convenient method to work out the clues of regularity and systematics of elliptical galaxies, and to find out their links with the processes involved in galaxy assembly in a cosmological context. The point important for our present purposes is that they {\\it directly} provide  with complete 6-dimensional phase-space informations on each constituent particle sampling a given galaxy-like object formed in the simulation, that is, they give directly the mass and velocity  distributions of dark matter, gas and stars of each object.  Taking advantage of these possibilities, we have analysed ten self-consistent hydrodynamical simulations run in the framework of a flat $\\Lambda$CDM cosmological model, characterised by cosmological parameters consistent with their last determinations (Spergel et al. 2006). Galaxy-like objects of different morphologies form in these simulations at $z=0$: irregulars, disc-like objects, S0-like objects and elliptical-like objects (hereafter, ELOs). ELOs have been identified as those objects having a prominent dynamically relaxed, roughly non-rotating stellar spheroidal component, with no extended disks and very low cold gas content; the  stellar component has typical sizes of no more than $\\sim $ 10 - 40  kpc, it dominates the mass density at these scales (hereafter, {\\it ELO or baryonic object  scale}) and it is embedded in a halo of dark matter typically ten  times larger in size (hereafter, {\\it halo scale}). In a forthcoming paper (O\\~norbe et al., in preparation) we report on an analysis of the mass and velocity distributions of the different ELO components: dark matter, stars, cold gas and hot gas. In this paper we focus on the quantitative characterisation of these mass and velocity distributions through their corresponding parameters, both at the ELO and at the halo scales. At the baryonic object scale, to characterise the structural and dynamical properties of ELOs, we will describe their three dimensional  distributions of mass and velocity through the three intrinsic (that is, three-dimensional) parameters (the stellar mass at the baryonic object  scale, $M_{\\rm bo}^{\\rm star}$, the stellar  half-mass radii at the baryonic object  scale, $r_{\\rm e, bo}^{\\rm star}$, defined as those radii enclosing half the $M_{\\rm bo}^{\\rm star}$ mass, and the mean square velocity for stars, $\\sigma_{\\rm 3, bo}^{\\rm star}$) whose observational projected counterparts (the luminosity $L$, effective projected size  $R_{\\rm e}^{\\rm light}$, and the stellar central l.o.s. velocity dispersion, $\\sigma_{\\rm los, 0}$) enter the definition of the observed FP (Eq. \\ref{FP}). To help the reader, in Table~\\ref{tabnombres} we give a list of the parameter names and symbols and in Table~\\ref{tabnombres2} a list of the profiles and ratios. These intrinsic parameters have been measured on ELOs, and their correlations have been looked for, and more specifically, the lower--dimensional regions (i.e., {\\it dynamical planes}) they fill in the three dimensional space of the $M_{\\rm bo}^{\\rm star}$, $r_{\\rm e, bo}^{\\rm star}$ and $\\sigma_{\\rm 3, bo}^{\\rm star}$ parameters, because the observational manifestation of these dynamical planes is the FP relation. These informations, combined with that on  the mass, size and velocity dispersion parameters at the halo scale, allows us to test whether or not the $c_{\\rm M}^{\\rm vir}$ coefficients and the $M_{\\rm vir}/M^{\\rm star}$ ratios do  systematically depend on the mass scale, so that the tilt and the scatter of the observed FP can be explained in terms of the  regularities in the structural and dynamical properties of ELOs formed in self-consistent hydrodynamical simulations.   Fully-consistent gravo-hydrodynamical simulations as a method to study E assembly has already proven to be useful. An analysis  of ELO  structural and kinematical properties that can be constrained from observations (i.e., stellar masses,  projected half-mass  radii, central line-of-sight velocity dispersions), has shown that they have counterparts in the local Universe as far as these properties are concerned (S\\'aiz 2003; S\\'aiz, Dom\\'{\\i}nguez-Tenreiro \\& Serna 2004, hereafter SDTS04), including the FP relation and some clues about its physical origin (O\\~norbe et al. 2005), and its lack of dynamical evolution (Dom\\'{\\i}nguez-Tenreiro et al. 2006, hereafter DTal06). Also, ELO stellar populations have age distributions showing similar trends as those inferred from observations (Dom\\'{\\i}nguez-Tenreiro, S\\'aiz \\& Serna 2004, hereafter DSS04).  The paper is organised as follows: in $\\S$2 we briefly describe the simulations, the ELO samples and their generic properties. The ELO size and mass scales and their relations are analysed in $\\S$3. $\\S$4 is devoted to kinematics and in $\\S$5 we report on the intrinsic dynamical plane of ELOs and its comparison with the observed Fundamental Plane, and we analyse its physical origin. Finally, in $\\S$6 we summarise our results and discuss them in the context of theoretical results on halo structure and dissipation of the gaseous component. ",
        "conclusions": "\\subsection{Summary} \\label{summary} We present an analysis of the sample of ELOs formed in ten different cosmological  simulations, run within the same global flat $\\Lambda$ cosmological model, roughly consistent with observations. The normalisation parameter has been taken slightly high, $\\sigma_8 = 1.18$, as compared with the average fluctuations of 2dFGRS or SDSS  galaxies to mimic an active region of the Universe. Newton laws and hydrodynamical equations have been integrated in this context, with a standard cooling algorithm and a star formation parameterization through a Kennicutt-Schmidt-like law, containing our ignorance about its details at sub-kpc scales. No further hypotheses to model  the assembly processes have been made. Individual galaxy-like objects naturally appear as an output of the simulations, so that the physical processes underlying mass assembly can be studied. Five out of the ten simulations (the SF-A type simulations) share the SF parameters and differ in the seed  used to build up the initial conditions. To test the role of SF parameterisation, the same initial conditions have been run with different SF parameters making SF  more difficult, contributing another set of five simulations  (the SF-B type simulations). ELOs have been identified in the simulations as those galaxy-like objects  at $z = 0$ having a prominent, non-rotating dynamically relaxed spheroidal   component made out of stars, with no extended discs and very low gas content. These stellar component is embedded in  a dark matter halo that contributes an important fraction of the mass at distances from the ELO centre larger than $\\sim 10 - 15$ kpc. No ELOs with stellar masses below 3.8 $\\times$ 10$^{10}$ M$_{\\odot}$ or virial masses below 3.7 $\\times$ 10$^{11}$ M$_{\\odot}$ have been found that met the selection  criteria (see Kauffmann et al. 2003b for a similar result in SDSS galaxies and Dekel \\& Birnboim 2006, and Cattaneo et al. 2006 for a possible physical explanation). ELOs have also an extended halo of hot, diffuse gas. Stellar and dark matter particles constitute a dynamically hot component with an important velocity dispersion, and, except in the very central regions, a positive anisotropy.  The informations about position and velocity distributions of the ELO   particles of different kinds (dark matter, stars, cold gas, hot gas) provided by the simulations, allows a detailed study of their  intrinsical three dimensional mass and  velocity distributions, as well as a measure of the parameters characterising their structure and dynamics. In a forthcoming paper, we report on the three dimensional mass density, circular velocity and velocity dispersion  profiles, as well as the projected stellar mass density profiles and the LOS velocity dispersion profiles. In this paper we focus on a  parameter analysis to quantify some of the previous results.  Mass, size and velocity dispersion scales for their different components have been measured in the ELO samples, both at the scale of their halo and at the scale of the baryonic object (a few tens of kiloparsecs). At the {\\it halo scale}, the masses of both cold gas and stars, $M_{\\rm h}^{\\rm cb}$ and $M_{\\rm h}^{\\rm star}$, respectively, have been found to be tightly correlated with the halo total mass, $M_{\\rm vir}$, with the ratios $M_{\\rm h}^{\\rm cb}/M_{\\rm vir}$ and $M_{\\rm h}^{\\rm star}/M_{\\rm vir}$ decreasing as $M_{\\rm vir}$ increases (that is, massive objects miss cold baryons within $r_{\\rm vir}$ when compared with less massive ELOs), presumably because gas gets more difficulties to cool and fall as $M_{\\rm vir}$ increases. The overall half-mass radii, $r_{\\rm e, h}^{\\rm tot}$ shows also a very tight correlation with $M_{\\rm vir}$. Half-mass radii for the cold baryon or stellar mass distributions have a more complex behaviour, as in these cases gas heating in  shocks and energy losses due gas cooling are in competition to determine these distributions.  A very important result we have found when analysing ELOs at the {\\it scale of the baryonic object}, is that $M_{\\rm vir}$ plays an important role to determine the ELO structure also below a few tens of kiloparsecs scales. In fact, both the masses of cold baryons $M_{\\rm bo}^{\\rm cb}$ (i.e., those baryons that have reached the central regions of the configuration), and of stars $M_{\\rm bo}^{\\rm star}$, show a good correlation with $M_{\\rm vir}$, and, moreover, the  $M_{\\rm bo}^{\\rm cb}/M_{\\rm vir}$ and $M_{\\rm bo}^{\\rm star}/M_{\\rm vir}$ ratios (i.e., the relative content of cold baryons or stars versus total mass) decrease as $M_{\\rm vir}$ increases. This is the same qualitative  behaviour shown by these ratios observationally in the SDSS data, and, also, by ELOs at the halo scale. The dependence of $M_{\\rm bo}^{\\rm cb}$ or $M_{\\rm bo}^{\\rm star}$ on the SF parametrisation is only very slight, with SF-A type ELOs having slightly more stars than their SF-B type counterparts. The half-mass radii for cold baryon and stellar masses, $r_{\\rm e, bo}^{\\rm cb}$ and $r_{\\rm e, bo}^{\\rm star}$, show also a good correlation with $M_{\\rm vir}$, but now the values of the SF parameters also play a role, because their change implies a change in the time interval during which gas cooling is turned on, and this changes  the ELO  stellar mass distribution, i.e., its lenghthscale, so that ELO compactness increases from SF-A to SF-B type simulations. Another important result is that, regardless of the SF parametrizations used in this work, the relative distributions of the stellar and dark mass components in ELOs show a systematic trend measured through the $c_{\\rm rd} \\equiv r_{\\rm e, h}^{\\rm tot} / r_{\\rm e, bo}^{\\rm star}$ ratios, with stars relatively more concentrated as $M_{\\rm vir}$ decreases (i.e., a quantification of the spatial  homology breaking). Note that to compare with observational data, the relevant parameters are the {\\it projected} half-mass radii, $R_{\\rm e, bo}^{\\rm star}$. We have checked that they show an excellent correlation with the corresponding three dimensional half-mass radii, with the $c_{\\rm rp} \\equiv r_{\\rm e, bo}^{\\rm star} / R_{\\rm e, bo}^{\\rm star}$ ratios showing no significant dependence on the ELO mass scale.  Concerning kinematics, a useful characterisation of the ELO velocity dispersion is the central stellar line-of-sight velocity dispersion, $\\sigma_{\\rm los, 0}^{\\rm star}$, whose observational counterpart can be measured from elliptical spectra. A very important result is the very tight correlation we have found between $M_{\\rm vir}$ and $\\sigma_{\\rm los, 0}^{\\rm star}$, confirming that the observationally measurable $\\sigma_{\\rm los, 0}$ is a fair virial mass estimator. In addition, $\\sigma_{\\rm los, 0}^{\\rm star}$ is closely related to the mean square velocity of both, the whole elliptical at the halo scale (including the dark matter), $\\sigma_{\\rm 3, h}^{\\rm tot}$, and the stellar component of the central object, $\\sigma_{\\rm 3, bo}^{\\rm star}$. We have also found that  the $c_{\\rm vd} \\equiv (\\sigma_{\\rm 3, h}^{\\rm tot} / \\sigma_{\\rm 3, bo}^{\\rm star})^2$ or the $c_{\\rm vpc} \\equiv (\\sigma_{\\rm 3, bo}^{\\rm star})^2 / 3 (\\sigma_{\\rm los, 0}^{\\rm star})^2$ ratios are roughly independent of the ELO mass scale. And so, ELOs do not show dynamically broken homology, even if their stellar and dark components are kinematically segregated (i.e., $c_{\\rm vd} \\ne 1$). This could lead to inaccurate determinations of the total mass of ellipticals when using stellar kinematics.  A very important result is that, irrespective of the SF parameterisation, the (logarithms of the) ELO stellar  masses $M_{\\rm bo}^{\\rm star}$, stellar half-mass radii  $r_{\\rm e, bo}^{\\rm star}$, and stellar mean square velocity of the central object $\\sigma_{\\rm 3, bo}^{\\rm star}$, define {\\it intrinsic dynamical}  planes (IDPs). These planes are tilted relative to the virial plane and the tilt does not significantly depend on the SF parameterisation, but the zero point does depend. Otherwise, the intrinsic dynamical plane is not homogeneously populated, but ELOs, as well as E galaxies in the FP (Guzm\\'an et al. 1993), occupy only a particular region defined by the range of their masses.    \\subsection{Testing possible resolution effects} \\label{resoleff} To make sure that the results we report in this paper are not unstable under resolution changes, a control simulation with 128$^3$ dark matter and 128$^3$ baryonic particles, a gravitational softening of $\\epsilon = 1.15$ kpc and the other parameters as in SF-A type simulations (the S128 simulation), has been run. The results of its analysis have been compared with those of a 2 $\\times 64^3$ simulation (the S64 simulation), whose initial conditions have been built up by randomly choosing 1 out of 8 particles in the S128 initial conditions, so that every ELO in S128 has a counterpart in the lower resolution simulation and conversely. Due to the very high CPU time requirements for S128, the comparison has been made at $z=1$. The results of this comparison are very satisfactory, as Figure~\\ref{ComCoc} illustrates (compare with Figure~\\ref{CMvirMob4.5}).  \\begin{figure} \\begin{center} \\includegraphics[width=.45\\textwidth]{figure17.eps} \\caption {The $M_{\\rm h}^{\\rm cb}/M_{\\rm bo}^{\\rm cb}$ ratios as a function of the ELO mass scale for three ELOs identified at $z=1$ in S128 (red) and their counterparts in S64 (black). } \\label{ComCoc} \\end{center} \\end{figure}  \\subsection{The dimensionality of ELO and elliptical samples in parameter space} \\label{dimensiomality} The intrinsic dynamical planes and their occupations (see \\S ~\\ref{summary}) reflect the fact that dark matter haloes are a two-parameter family (for example, the virial mass and the energy content or the concentration; see, for example, Hernquist 1990; Navarro, Frenk \\& White 1995, 1996; Manrique et al. 2003; Navarro et al. 2004) where the two parameters are correlated  (see, for example, Bullock et al. 2001; Wechsler et al. 2002;  Manrique et al. 2003). Adding gas implies that heating and cooling processes also play a role at determining the mass and velocity distributions, and, more particularly, the length scales. However, as explained above, we have found that both, the relative content and the relative distributions of the dark and baryonic mass components show systematic trends with the ELO mass scale, that can be written as power-laws of the form $ M_{\\rm vir}/M_{\\rm bo}^{\\rm star} = A_{\\rm vir} (M_{\\rm bo}^{\\rm star})^{\\beta_{\\rm vir}}$ and $ r_{\\rm e, h}^{\\rm tot} / r_{\\rm e, bo}^{\\rm star} = A_{\\rm rd} (M_{\\rm bo}^{\\rm star})^{\\beta_{\\rm rd}}$.  A first consequence of the regularity of the trends with the mass scale found in this paper is that no new parameters are added relative to the dark matter halo family, so that  the baryonic objects are also a two-parameter family, and ELO structural and dynamical parameters define also a plane. A second consequence is that the plane is tilted relative to the halo plane (i.e., the virial plane) because $\\beta_{\\rm rd} - \\beta_{\\rm vir} \\neq 0$. Finally, the plane is not homogeneously populated because of the mass-concentration halo correlation, that at the scale of the baryonic objects appears for example as a mass---size correlation. This explains the role played by $M_{\\rm vir}$ to determine the intrinsic three dimensional correlations. In this paper  we also show that $\\sigma_{\\rm los, 0}^{\\rm star}$ is a fair empirical estimator of $M_{\\rm vir}$, and this explains the central role played by $\\sigma_{\\rm los, 0}$ at determining the observational correlations.   The fundamental plane shown by real elliptical samples is the observationally manifestation of the IDPs when using projected parameters $R_{\\rm e, bo}^{\\rm star}$, $\\sigma_{\\rm los, 0}^{\\rm star}$ and luminosity variables instead of stellar masses $M_{\\rm bo}^{\\rm star}$. We have taken advantage of the constancy of the stellar-mass-to-light ratios of ellipticals in the SDSS (Kauffmann et al. 2003a, 2003b) to put the elliptical sample of Bernardi et al. (2003b, 2003c) in the same projected variables we can measure in our virtual ellipticals. We have found that the FPs shown by the two ELO samples are consistent with that shown by the SDSS elliptical sample in the same variables, with no further need for any relevant contribution from stellar population effects to explain the observed tilt. These effects could, however, have contributed to the scatter of the observed FP, as the IDPs have been found to be thinner than the observed FP.  \\subsection{The physical origin of the tilt in a cosmological context} \\label{physorigtilt} We now turn to discuss  the physical origin of the trends given by  the power laws $ M_{\\rm vir}/M_{\\rm bo}^{\\rm star} = A_{\\rm vir} (M_{\\rm bo}^{\\rm star})^{\\beta_{\\rm vir}}$ and $ r_{\\rm e, h}^{\\rm tot} / r_{\\rm e, bo}^{\\rm star} = A_{\\rm rd} (M_{\\rm bo}^{\\rm star})^{\\beta_{\\rm rd}}$. As explained in Section 2, the simulations provide us with clues on the physical processes involved in elliptical formation (see also DSS04 and DTal06). They also  indicate that the dynamical plane appears at an early violent phase as a consequence of ELO assembly out of gaseous material, with cooling  and on short timescales, and it is preserved  during a later, slower phase, where dissipationless merging plays an important role in stellar mass assembly (see more details in DTal06). Our simulations show  that the physical origin of the trends above lie in the systematic decrease, with increasing ELO mass, of the relative amount of dissipation experienced by the baryonic mass component along ELO stellar mass assembly (DTal06; O\\~norbe et al., in preparation). This possibility had been suggested by Bender et al. (1992), Guzm\\'an et al. (1993) and Ciotti et al. (1996). Bekki (1998) first addressed it numerically in the framework of the merger hypothesis for elliptical formation through pre-prepared simulations of dissipative mergers of disk galaxies, where the rapidity of the star formation in mergers is controlled by a free efficiency parameter $C_{\\rm SF}$. He shows that the star formation rate history of galaxies determine the differences in dissipative dynamics, so that to explain the slope of the FP he {\\it needs to assume} that more luminous galaxies are formed by galaxy mergers with a shorter timescale for gas transformation into stars. Recently, Robertson et al. (2006) have confirmed these findings on the importance of dissipation to explain the FP tilt.  In this paper we go an step further and analyze the FP of virtual ellipticals formed in a  cosmological context, where individual galaxy-like objects naturally appear as an output of the simulations. Our results essentially include previous ones and add important new informations. First, our results on the role of dissipation to produce the tilt of the FP essentially agree with those obtained through dissipative pre-prepared mergers, but it is important to note that, moreover, more massive objects produced in the simulations {\\it do have} older  means  and narrower spreads in their stellar age distributions than less massive ones (see details DSS04); this naturally appears in the simulations and need not be considered as an additional assumption. Second, the preservation of the FP in the slow phase of mass aggregation in our simulations also agrees with previous work based on dissipationless simulations of pre-prepared mergers (Capelato et al. 1995; Dantas et al. 2003; Goz\\'alez-Garc{\\'\\i}a \\& van Albada 2003;  Boylan-Kolchin et al. 2005; Nipoti, Londrillo \\& Ciotti 2003). Moreover, elliptical properties recently inferred from observations (for example, the appearence of blue cores,  Menanteau et al. 2004, and the increase of the stellar mass contributed by the elliptical population since higher $z$, Bell et al. 2004; Conselice, Blackburne, {\\&} Papovich 2005; Faber et al. 2005; see more details in DTal06) can also be explained in our simulations.  \\subsection{Conclusions} \\label{conclusion}  We conclude that the simulations provide an unified scenario where most current observations on ellipticals can be interrelated. In particular, it explains the most important results relative to the physical origin of their FP relation (i.e. the FP tilt is due dissipative dynamics, and disipationless merging in the slow growth phase of mass assembly does not change the tilt). It is worth mentioning that this scenario shares some characteristics of previously proposed scenarios, but it has also significant differences, mainly that most stars in elliptical galaxies form out of cold gas that had never been shock heated at the halo virial temperature and then formed a disk, as the conventional recipe for galaxy formation propounds (see discussion in Keres et al. 2005 and references therein). The scenario for elliptical formation emerging from our simulations has the advantage that it results from simple physical laws acting on initial conditions that are realizations of power spectra consistent with observations of CMB anisotropies.      This work was partially supported by the MCyT (Spain) through grants AYA-0973, AYA-07468-C03-02 and AYA-07468-C03-03 from the PNAyA, and also by the regional government of Madrid through the ASTROCAM Astrophysics network. We  thank the Centro de Computaci\\'on Cient\\'{\\i}fica (UAM, Spain) for computing facilities. AS  thanks  FEDER financial support from Eropean Union."
    },
    "0609/astro-ph0609450_arXiv.txt": {
        "abstract": "{Before the launch of the \\emph{Swift} satellite, the majority of the gamma-ray burst (GRB) afterglows for which Ly$\\alpha$ was redshifted into the observable spectrum showed evidence for a damped Ly$\\alpha$ absorber. This small sample indicated that GRBs explode either in galaxies, or regions within them, having high neutral hydrogen column densities.} {To increase the spectroscopic sample of GRBs with $z>2$ and hence establish the $N(\\ion{H}{i})$ distribution along GRB lines-of-sight.} {We have obtained six $z > 2$ GRB afterglow spectra and fitted the Ly$\\alpha$ absorption line in each case to determine $N(\\ion{H}{i})$. This has been complemented with 12 other \\emph{Swift} $N(\\ion{H}{i})$ values from the literature.} {We show that the peak of the GRB $N(\\ion{H}{i})$ distribution is qualitatively consistent with a model where GRBs originate in Galactic-like molecular clouds. However, a systematic difference, in particular an excess of low column-density systems compared to the predictions, indicates that selection effects and conditions within the cloud (e.g. strong ionization) influence the observed $N(\\ion{H}{i})$ range. We also report the discovery of Ly$\\alpha$ emission from the GRB\\,060714 host, corresponding to a star-formation rate of approximately 0.8\\,$M_{\\odot}\\,\\textrm{yr}^{-1}$. Finally, we present accurate redshifts of the six bursts: $z = 3.240 \\pm 0.001$ (GRB\\,050319), \\mbox{$z = 2.198 \\pm 0.002$} (GRB\\,050922C), $z = 3.221 \\pm 0.001$ (GRB\\,060526), $z = 3.425 \\pm 0.002$ (GRB\\,060707), $z = 2.711 \\pm 0.001$ (GRB\\,060714) and $z = 3.686 \\pm 0.002$ (GRB\\,060906).} {} ",
        "introduction": "In just 18 months, the \\emph{Swift} satellite (Gehrels et al. \\cite{swift}) has already contributed considerably to the progress of gamma-ray burst (GRB) science. It detects roughly two GRBs/week and transmits their accurate localizations (error radius frequently less than 5\\arcsec) to the ground within minutes for rapid follow-up observations. The GRB redshift distribution found by \\emph{Swift} is very different from that of the pre-\\emph{Swift} sample, being skewed to much higher redshifts (Berger et al. \\cite{bergerSwift}; Jakobsson et al. \\cite{palli_swift,palli_swift2}; Daigne et al. \\cite{daigne}; Le \\& Dermer \\cite{le}). Approximately 70$\\%$ of the \\emph{Swift} bursts are found to be located at $z > 2$, while the corresponding fraction was only 20$\\%$ for pre-\\emph{Swift} bursts. \\par It is thus much more common for the Ly$\\alpha$ line in \\emph{Swift} bursts to be redshifted redward of the atmospheric cutoff in optical spectra. This provides us with the opportunity to investigate the GRB host neutral hydrogen column density distribution. In particular we can compare it to the damped Ly$\\alpha$ absorbers (DLAs) seen in absorption against QSO spectra (see Wolfe et al. \\cite{wolfe} for a recent review), defined as displaying $N(\\ion{H}{i}) \\geq 2 \\times 10^{20}$\\,cm$^{-2}$ $(\\log N(\\ion{H}{i}) \\geq 20.3)$. Most of the neutral gas in the Universe in the redshift interval $0 < z < 5$ is in DLAs, providing the fuel for star formation at these epochs. Given that long-duration GRBs are known to have massive stellar progenitors (e.g. Hjorth et al. \\cite{jens}; Malesani et al. \\cite{daniele}; Fruchter et al. \\cite{andy}), GRB-DLAs provide valuable information on the sites of active star formation in the high-redshift Universe. In the pre-\\emph{Swift} era, only seven GRB $N(\\ion{H}{i})$ column density measurements were obtained, with six being classified as DLAs (Vreeswijk et al. \\cite{paul} and references therein). \\par In this Letter we present optical spectroscopy of six GRBs, focusing on the measurement of the $N(\\ion{H}{i})$ column density. We also report the detection of Ly$\\alpha$ emission from one of the bursts (GRB\\,060714). We then discuss how the observed $N(\\ion{H}{i})$ column density distribution from \\emph{Swift} bursts can be understood in terms of the GRB environment. ",
        "conclusions": "\\begin{table} \\caption[]{A list of all \\emph{Swift} GRBs with $z > 2$ known to date (1 October 2006). $N(\\ion{H}{i})$ is derived from optical spectroscopy. References are given in order for the redshift and the $\\ion{H}{i}$ column density (information is not currently available for bursts marked with ---). (1) This work; (2) Watson et al. (\\cite{watson06}); (3) Berger et al. (\\cite{berger505}); (4) Berger et al. (\\cite{berger603}); (5) Starling et al. (\\cite{starling05}); (6) Ledoux et al. (\\cite{ledoux05}); (7) Ledoux et al. (in prep); (8) Kawai et al. (\\cite{kawai06}); (9) Foley et al. (\\cite{foley05}); (10) Quimby et al. (\\cite{quimby05}); (11) Piranomonte et al. (\\cite{pira05}); (12) Prochaska et al. (\\cite{prochaska06}); (13) J. X. Prochaska (private communication); (14) Fynbo et al. (\\cite{fynbo06}); (15) Cucchiara et al. (\\cite{cucchiara06}); (16) Cucchiara et al. (\\cite{cucchiara_poster}); (17) Berger et al. (\\cite{berger223}); (18) Price (\\cite{price06}); (19) Cenko et al. (\\cite{cenko06}); (20) Ferrero et al. (\\cite{ferrero}); (21) Ledoux et al. (\\cite{ledoux06}); (22) Smette et al. (in prep); (23) Rol et al. (\\cite{evert}); (24) D'Elia et al. (\\cite{delia}); (25) Fugazza et al. (in prep); (26) Fynbo et al. (\\cite{5.6}).}  \\label{sample.tab} \\centering \\setlength{\\arrayrulewidth}{0.8pt}   % \\begin{tabular}{@{}llcl@{}} \\hline \\hline \\vspace{-2 mm} \\\\ GRB & \\hspace{1.6 mm} $z$ & $\\log N(\\ion{H}{i})$ & References \\\\ \\hline \\vspace{-2 mm} \\\\ 050319  & 3.24 & $20.9 \\pm 0.2$ & (1) (1) \\\\ 050401  & 2.90 & $22.6 \\pm 0.3$ & (2) (2) \\\\ 050505  & 4.27 & $22.1 \\pm 0.1$ & (3) (3) \\\\ 050603  & 2.82 & ---  & (4) \\\\ 050730  & 3.97 & $22.1 \\pm 0.1$ & (5) (5) \\\\ 050820A & 2.61 & $21.1 \\pm 0.1$ & (6) (7) \\\\ 050904  & 6.30 & $21.3$         & (8) (8) \\\\ 050908  & 3.34 & $19.2$         & (9) (9) \\\\ 050922C & 2.20 & $21.6 \\pm 0.1$ & (1) (1) \\\\ 051109  & 2.35 & ---  & (10) \\\\ 060115  & 3.53 & ---  & (11) \\\\ 060124  & 2.30 & $<$20.3 \\hspace*{1.2 mm} & (12) (13) \\\\ 060206  & 4.05 & $20.9 \\pm 0.1$ & (14) (14) \\\\ 060210  & 3.91 & $21.7 \\pm 0.2$ & (15) (16) \\\\ 060223  & 4.41 & ---  & (17) \\\\ 060510B & 4.94 & ---  & (18) \\\\ 060522  & 5.11 & $20.5 \\pm 0.5$ & (19) (19) \\\\ 060526  & 3.22 & $20.0 \\pm 0.2$ & (1) (1) \\\\ 060605  & 3.71 & ---  & (20) \\\\ 060607  & 3.08 & $<$19.5 \\hspace*{1.2 mm} & (21) (22) \\\\ 060707  & 3.43 & $21.0 \\pm 0.2$ & (1) (1) \\\\ 060714  & 2.71 & $21.8 \\pm 0.1$ & (1) (1) \\\\ 060906  & 3.69 & $21.9 \\pm 0.1$ & (1) (1) \\\\ 060908  & 2.43 & ---  & (23) \\\\ 060926  & 3.21 & $22.7 \\pm 0.1$ & (24) (25) \\\\ 060927  & 5.6  & ---  & (26) \\\\ \\hline \\end{tabular} \\end{table}  In Table~\\ref{sample.tab} we have compiled all available $\\ion{H}{i}$ column density measurements for $z > 2$ \\emph{Swift} bursts (currently displaying a median redshift of 3.4). The bottom panel of Fig.~\\ref{histo.fig} shows a comparison between these $\\ion{H}{i}$ column densities and those of the pre-\\emph{Swift} sample (Vreeswijk et al. \\cite{paul}). The former currently has a median $\\log N(\\ion{H}{i})$ of 21.6, a bit higher than the latter (21.3). However, the small size of the pre-\\emph{Swift} sample does not allow us to reject, with any degree of confidence, the null hypothesis that the two samples are drawn from the same distribution, i.e. a two-sample Kolmogorov-Smirnov (KS) test results in a 90\\% significance. The fraction of DLAs is also similar in both samples ($\\sim$80$\\%$), supporting the conclusion that GRB absorption systems show exceptionally high column densities of gas when compared to DLA systems observed in the lines-of-sight to QSOs (Vreeswijk et al. \\cite{paul}). Indeed, a two-sample KS test indicates there is a less than $10^{-7}$ probability that GRB-DLAs and QSO-DLAs (Prochaska et al. \\cite{QSO-DLA}) are drawn from the same population. This result presumably reflects the fact that GRBs occur in star-forming regions within their hosts, whilst QSOs select more random lines-of-sight through intervening galaxies. \\par In the bottom panel of Fig.~\\ref{histo.fig} we also compare the observed GRB $N(\\ion{H}{i})$ distribution to the expected column density distribution for bursts in Galactic-like molecular clouds (Reichart \\& Price \\cite{reichart}, hereafter RP02). The model (solid, non-filled histogram) is mass-weighted and corrected for geometrical effects, i.e. the clouds being centrally condensed, and the GRB location within a cloud (not behind it). Qualitatively, especially in terms of the location of the peak of the distribution, the match is relatively good. However, a KS test shows that the model and observed $N(\\ion{H}{i})$ distributions are inconsistent with being drawn from the same population at the 3$\\sigma$ level, with a clear overabundance of low-$N(\\ion{H}{i})$ detections compared to the model. A possibility is that GRB host galaxies tend to have lower $N(\\ion{H}{i})$ clouds than the Milky Way (MW) and/or GRBs are more likely to occur in lower $N(\\ion{H}{i})$ clouds. Some evidence for the former exists (Rosolowsky \\cite{MCs}), with the mass distribution of clouds in the LMC (which is more similar to GRB hosts than the MW) having a marginally steeper distribution than the inner disk of the MW. \\par A few other effects can give rise to the excessive detections of $\\log N(\\ion{H}{i}) \\leq 21.0$. The GRB progenitor and nearby massive stars may (partially) ionize their local environments; there is clearly a trend for GRBs with the smallest $N(\\ion{H}{i})$ to exhibit very weak low-ionization lines (e.g. \\ion{Si}{ii} and \\ion{C}{ii}) while displaying strong absorption of e.g. \\ion{Si}{iv} and \\ion{C}{iv} (GRB\\,021004: M\\o ller et al. \\cite{palle}; GRB\\,050908: Prochaska et al. \\cite{x_908}; GRB\\,060124: Prochaska et al. \\cite{prochaska06}). This may support the high-ionization scenario for the low-$N(\\ion{H}{i})$ systems, although implicit is the assumption that low- and high-ionization species trace the same phase. An alternative scenario is that some GRBs are formed by massive runaway stars (e.g. Hammer et al. \\cite{hammer}); these would explode in regions of relatively low $N(\\ion{H}{i})$. Finally, it has been suggested (e.g. White et al. \\cite{white}; Sugitani et al. \\cite{sugitani}; Miao et al. \\cite{miao}) that a significant part of star formation in molecular clouds takes place at their outer edges. The triggering of this star formation also blows the cloud open, resulting in these stars having (largely) unimpeded view outside the cloud and presumably a low $N(\\ion{H}{i})$ column. \\par The drop-off in detections above $\\log N(\\ion{H}{i}) \\geq 22.0$ could be due to selection effects, i.e. the sample might be incomplete due to high extinction. GRBs most likely do not burn through all of the dust in their molecular clouds (e.g. fig.~4 in RP02), so some high-$N(\\ion{H}{i})$ systems are likely dimmed by dust and consequently missed in the sample in Table~\\ref{sample.tab} (see also Savaglio et al. \\cite{sandra}; Vergani et al. \\cite{vergani}). In the top panel of Fig.~\\ref{histo.fig} we have plotted the corresponding afterglow $R$-band magnitudes (UV restframe) interpolated to a restframe epoch of 12\\,hr and redshifted to \\mbox{$z = 3$}. The majority of the bursts have a well-sampled light curve, allowing us to determine the magnitudes via spline interpolation (see appendix A in Kann et al. \\cite{alex} for a detailed description of the method used). For the most recent bursts with a relatively poor data sampling, we have applied extrapolation and adopted large conservative error bars. There is only a tentative evidence of an anti-correlation between the OA flux and higher $N(\\ion{H}{i})$ at the 2$\\sigma$ level as derived from Spearman's rank correlation. It should be noted that intrinsic GRB $R$-band afterglow light curves show evidence for a bimodal luminosity distribution (Kann et al. \\cite{alex}; Liang \\& Zhang \\cite{liang}; Nardini et al. \\cite{nardini}), that increases the scatter in the top panel of Fig.~\\ref{histo.fig}. Excluding the upper limits, there is indeed an indication of bimodality in the data. \\par \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics[bb=16 0 500 663,clip]{6405fig3.eps}} \\caption{\\emph{Bottom:} Histogram of the \\ion{H}{i} column density measured in GRBs for which the redshift was large enough to detect Ly$\\alpha$. Bursts detected by \\emph{Swift} are filled while the pre-\\emph{Swift} sample is hashed (Vreeswijk et al. \\cite{paul}). The current range of known GRB \\ion{H}{i} systems is $19.2 \\leq \\log N(\\ion{H}{i}) \\leq 22.7$. For comparison, the QSO-DLAs (Prochaska et al. \\cite{QSO-DLA}) are overplotted as the dashed histogram (normalized with respect to the GRB observations). The solid, non-filled histogram shows the GRB $N(\\ion{H}{i})$ prediction by RP02. \\emph{Top:} The corresponding afterglow $R$-band magnitudes, corrected for Galactic extinction (Schlegel et al. \\cite{schlegel}), interpolated to a common epoch of 12\\,hr (in the source restframe) and redshifted to $z = 3$. GRBs\\,060210 and 060522 are omitted due to the lack of information on their spectral energy distribution; hence corrections could not be made to account for Ly$\\alpha$ absorption. Where the uncertainty on the \\ion{H}{i} column is not reported in the literature, a value of 0.2 has been plotted.} \\label{histo.fig} \\end{figure} \\par \\par The characteristics of the dust-to-gas ratio in the clouds are discernible from the range of more than three orders of magnitude in the \\ion{H}{i} column densities (lower panel of Fig.~\\ref{histo.fig}), suggesting a similar range in the dust column (e.g. Predehl \\& Schmitt \\cite{predehl}). If dust were entirely responsible for all of the observed variation in the $R$-band (about five magnitudes; upper panel of Fig.~\\ref{histo.fig}), the $A_R$ range would, at most, be of that order, limiting it to a maximum of $A_R \\approx 5$\\,mag. At the Galactic dust-to-gas ratio, the observed $A_R$ (UV restframe) corresponding to the maximum $N(\\ion{H}{i})$ value is around 80 magnitudes (e.g. Pei \\cite{pei}; Diplas \\& Savage \\cite{diplas}), implying that the $A_R/N(\\ion{H}{i})$ ratio in these absorbers must be at most 7\\% of the Galactic value (see also Hjorth et al. \\cite{jens_124}). If some of the variation in the flux is due to other factors, the ratio drops even further. \\par Given that it is thought that most of the star formation in the Universe occurs in molecular clouds, it seems logical that long-duration GRBs originate in such regions (Fig.~\\ref{histo.fig}). However, there is a systematic difference between the model and observations that presumably can be clarified as a result of a combination of selection effects (dust extinction) and intense ionization in the GRB environment. Combining the $N(\\ion{H}{i})$ values derived from optical spectroscopy with the metal column densities from soft X-ray absorption would conceivably shed further light on these issues. A detailed analysis of this is presented in a subsequent paper (Watson et al. \\cite{darach_NH})."
    },
    "0609/astro-ph0609385_arXiv.txt": {
        "abstract": "{We present a  \\ch-LETGS observation of the Seyfert~1 galaxy \\object{\\mrk}. This observation was carried out simultaneously with \\hst\\ and FUSE, in the context of a multiwavelength study of this source. The \\ch\\ pointings were spread over ten days for a total exposure time of $\\sim$360 ks. The maximal continuum flux variation is of the order of 30\\%. The spectrum of \\mrk\\ shows evidence of broad emission features, especially at the wavelength of the \\ovii\\ triplet. We quantitatively explore the possibility that this emission is produced in the broad line region (BLR). We modeled the broad UV emission lines seen in the FUSE and HST-STIS spectra following the ``locally optimally emitting cloud\" approach. This method considers the emission from BLR as arising from ``clouds\" with a wide range of densities and distance from the source. We find that the X-ray lines luminosity derived from the best fit BLR model can match the X-ray features, suggesting that the gas producing the UV lines is sufficient to account also for the X-ray emission. The spectrum is absorbed by ionized gas whose total column density is $\\sim5\\times10^{20}$\\,cm$^{-2}$. The absorption spectrum can be modeled by two distinct gas components (log$\\xi\\sim 0.47$ and $2.49$, respectively) both showing a significant outflow velocity. However, the data allow also the presence of intermediate ionization components. The distribution of the column densities of such extra components as a function of the ionization parameter is not consistent with a continuous, power law-like, absorber, suggesting a complex structure for the gas outflow for \\mrk. ",
        "introduction": "Active galactic nuclei (AGN) are believed to be powered by accretion into a super massive black hole (BH). Through the identification of narrow spectral features in the energy range $E<$ 2\\,keV, the basic phenomenology of the absorbing ionized gas out-flowing from the vicinity of the BH has been established: approximately 60\\% of the Seyfert~1 galaxies display such an outflow \\citep{crenshaw99}. Objects with an UV absorber always display an X-ray absorption component \\citep{crenshaw99,kriss02}. Recent results show that the absorber has to be a multi-ionization component gas and some of the components of the UV and X-ray outflow may be part of the same material \\citep{gabel05a}. Only in the X-ray band higher ionized components, producing mostly He-like and H-like ions of C, N, O, and up to H-like iron, can be detected. These high ionization components have been reported to have a larger column density \\citep[e.g., ][]{turner,netzer03} and sometimes a higher outflow velocity \\citep[e.g.,][]{steen05}. The physical structure of the absorbing systems related to the BH activity is not understood. Independent measurement of the physical structure of the out-flowing ionized gas led to incompatible results. The density structure may be continuous as a function of the ionization level of the gas (NGC~5548, Kaastra et al. 2002, Steenbrugge et al. 2003). Alternatively, the absorber is clumped in discrete components, in pressure equilibrium surrounded by a hotter and tenuous phase \\citep{kriss95,kriss01}. The spectrum of \\object{NGC~3783} seems to be well described by this model \\citep[][]{k03,netzer03}. Finally, there are conflicting results on the spatial extent and the distance from the ionizing source of such an absorber. \\citet{behar03}, based on the \\xmm\\ observation of NGC~3783, proposed that the out-flowing X-ray absorber has the same spatial distribution as the narrow line region (NLR) mapped in UV, as confirmed by the UV spectral analysis of this source \\citep{gabel05a}. This picture would be analogous to the geometry drawn for the Seyfert~2 galaxies (e.g. \\object{NGC~1068}, \\cite{kink02}; \\object{Mrk~3}, \\cite{sako00}). Conversely, a \\ch-LETGS and HETGS study on NGC~5548 suggested that the outflow is more collimated with a small opening angle \\citep{steen05}. The location of the warm absorbers may be as far as the NLR, or closer to the source, at $\\sim$pc scale, where a molecular torus should exist \\citep[see][ and references therein]{blustin} or, finally, the ionized gas outflow may be generated by accretion disk instabilities \\citep[][]{proga}. The X-ray line emission spectrum turned out to be an important component in Seyfert~1 galaxies. Recombination lines, narrow and variable on long time scales are generally believed to arise in the NLR \\citep[e.g., ][]{pounds}. Moreover, broad emission features, several \\AA\\ wide, have recently been detected in the X-ray spectra of some Seyfert galaxies. Whether some of them are due to curvatures of the primary continuum hiding absorption \\citep[][]{lee}, or are relativistically broadened lines \\citep{branduardi,sako,ogle04}, is still unclear. The modeling of other detected broad excesses as broad lines or blends of lines, in the soft X-ray spectrum, points to emission from the BLR \\citep{steen05,ogle04}. A deeper understanding of the physics behind such emission features is also important for a correct interpretation of the entire X-ray spectrum.  \\mrk\\ is a Seyfert~1.5 galaxy, located at redshift $z$\\,=\\,0.0305 \\citep[][]{scott04}. The source spectrum is affected by a relatively low Galactic absorption \\citep[$N_{\\rm H}=1.64\\times 10^{20}$~cm$^{-2}$, ][]{elvis89}. \\mrk\\ has been extensively studied by low resolution instruments \\citep[e.g.,][]{1995ApJ...447..121W,2001ApJ...550..261W} mainly by virtue of the 6.4 keV iron emission line, which is characterized by a broad profile and a variable flux. The first systematic study on the soft X-ray energy band was carried out in 2002 using \\ch-HETGS, simultaneously with FUSE and HST-STIS \\citep{scott04}. At that epoch both UV and X-ray observations show that the source was at the lower end of its historical flux range \\citep{2001ApJ...550..261W}. The $2-10$ keV flux was $\\sim1.2\\times10^{-11}\\ {\\rm erg\\ cm^{-2} s^{-1}}$ and the 1000\\,\\AA\\ flux was $\\sim0.13\\times10^{-11}\\ {\\rm erg\\ cm^{-2} s^{-1}}$\\AA$^{-1}$. While monitored by FUSE, the source became 7.5 times dimmer from 1999 to 2002 and a similar behavior was recorded in the optical \\citep{bachev04}. In the 2002 multiwavelength observation, a relatively short exposure time combined with a low flux state conspired against the detection of any significant absorption in the HETGS spectrum. The ionized absorber in \\mrk\\ has indeed quite a low column density in the X-ray domain (${\\rm N_H}\\sim 5\\times 10^{20} {\\rm cm}^{-2}$, this study), as compared for example to NGC~3783 \\citep[$\\sim3\\times10^{22}$ cm$^{-2}$,][]{kaspi01} or NGC~5548 \\citep[$\\sim5\\times10^{21}$ cm$^{-2}$,][]{kaastra02}.\\\\ A large scale project has more recently focused on \\mrk. The observations were carried out using HST-STIS, FUSE and {\\it Chandra}-LETGS simultaneously. The methodology in studying the UV absorption is described in \\citet[][hereinafter G05]{gabel05}, while \\citet{arav05} focuses on the modeling of the structure of such absorber. For the LETGS data, the theoretical background for the line production from a meta-stable level of \\ov\\ is described in \\citet{jelle04}, while in the present paper we present the total modeling of the X-ray spectrum.\\\\ The X-ray flux we measure with \\ch-LETGS is $\\sim2.18\\times10^{-11}\\ {\\rm erg\\ cm^{-2} s^{-1}}$ in the $2-10$\\,keV range, while the UV flux at 1000\\,\\AA\\ is almost a factor of ten higher ($\\sim1.23\\times10^{-11}\\  {\\rm erg\\ cm^{-2} s^{-1}}$\\AA$^{-1}$, G05) than the 2002 UV observation. The flux history in the optical V band, recorded by the 0.9-m University of Nebraska telescope, shows that the source underwent a large flux rise during the $\\sim$2 months preceding the present campaign (Gaskell et al. 2006, in prep.).\\\\  The paper is organized as follows. Sect.~2 is devoted to the analysis of the data and the modeling of the emission and absorption spectral features of \\mrk. In Sect.~3 we discuss our results and in Sect.~4 we present our conclusions. The cosmological parameters used are: H$_0$=70 km/s/Mpc, $\\Omega_{\\rm m}$=0.3, and $\\Omega_{\\Lambda}$=0.7. The quoted errors refer to 84\\% confidence ($\\Delta\\chi^2=2$ for one interesting parameter), unless otherwise stated. ",
        "conclusions": "We have presented the analysis and modeling of the data of \\mrk, observed for 360\\,ks by \\ch-LETGS. For the first time, we have extended to the X-ray band the ``locally optimally cloud\" model, first proposed by \\citet{baldwin95} to describe the BLR emission lines seen in the UV. This has been achieved by fitting first the luminosity of the UV lines, measured simultaneously by FUSE and HST-STIS (G05) and deriving, from the best fit, the X-ray luminosity.\\\\ In agreement with the LOC model predictions, the distribution of the radial distance of the BLR ``clouds\" follow a power law with index $\\gamma=-1.02\\pm0.14$, assuming a density distribution with index $\\beta=-1$. We find that the inferred X-ray line luminosity well describe the X-ray spectral shape. If, independently, we fit the X-ray broad emission features with Gaussian profiles, the results are, within the errors, consistent with the UV modeling. The most evident X-ray broad emission feature, the \\ovii\\ triplet, which exceeds the continuum of about 30\\%, is well explained by the LOC model. The profiles of other important X-ray lines of the H-like and He-like C, N, O ions are less clear in the spectrum and cannot be studied in more detail. It would be important to quantitatively measure the contribution of such ions, as they may be more sensitive to the physical conditions than \\ovii, which is steadily produced for a wide range of gas densities and distances from the ionizing source. There are not significant spectral residuals in excesses after the inclusion of the X-ray broad lines. This suggests first that the gas of the BLR emitting the bulk of the UV emission is also sufficient to explain the X-ray emission spectrum. Second, that the contribution of lines arising instead from the innermost region of the accretion disks, and thus relativistically broadened, is negligible in the case of \\mrk.\\\\ In the radial distribution of the line luminosity, the luminosity--weighted radius for each ion considered defines a wide region of roughly 90\\,ld for the UV--X-ray BLR (67\\,ld if only the UV ions are considered). This is larger than predicted by reverberation mapping studies \\citep{stirpe}. However, our estimate could be lowered by a factor up to three noting that the BLR size is sensitive to the long term variations of the continuum flux and that at the epoch of the present campaign the optical flux was at the higher end of a gradual rise that lasted nearly two months.\\\\ A proper modeling of the broad emission lines in the LETGS spectrum helps in the analysis of the absorbed spectrum. We find the signature of at least two gas components which can be unmistakably associated with a warm absorber. However, there is some evidence that suggests that this may be only a partial description of the absorption. First, absorption lines do not show any significant variation following the modest changes of the central source flux and, second, the two absorption components cannot be in pressure equilibrium. This evidence, although not decisive, does not strongly support the hypothesis of a compact absorber \\citep{k05}. We investigated the possibility that the absorbing gas is instead distributed over a wide range of ionization parameters \\citep{steen05}. We find that neither of the two models perfectly matches our data. A smooth distribution of the column densities is accepted, but $N_{\\rm H}$ does not monotonically increase as a function of $\\xi$ (following for example a power law), but it is rather significantly bent.\\\\ From the exploratory study of the possible \\ov\\ absorption line to meta-stable level in the \\mrk\\ X-ray spectrum, \\cite{jelle04} infer a density for the absorbing gas of the order of $10^{14}$\\,cm$^{-3}$. In the context of our analysis, we find that an intermediate ionization gas component ($\\log\\xi\\sim 1$), with a column density roughly consistent with our continuous-distribution model, could be consistent to produce the \\ov$^{\\star}$ line. This would imply an upper limit to the distance from the central source of $3\\times 10^{15}$\\,cm for the gas. Although a distance close to the BL emitting region for the warm absorber is not new in the literature \\citep[e.g., ][]{nicastro99,netzer02}, in the case of \\mrk, a global model explaining the coexistence of the complex warm absorber structure and the BLR cannot be easily depicted."
    },
    "0609/astro-ph0609317_arXiv.txt": {
        "abstract": "The ARGO-YBJ experiment is almost completely installed at the YangBaJing Cosmic Ray Laboratory (4300 m a.s.l., Tibet, P.R. China). The lower energy limit of the detector (E $\\sim$ 1 GeV) is reached with the scaler mode, i.e., recording the single particle rate at fixed time intervals. In this technique, due to its high altitude location and large area ($\\sim$ 6700 m$^2$), this experiment is the most sensitive among all present and past ground-based detectors. In the energy range under investigation, signals due to local (e.g. solar GLEs) and cosmological (e.g. GRBs) phenomena are expected as significant enhancements of the counting rate over the background. Results on the search for GRBs in coincidence with satellite detections are presented. ",
        "introduction": "\\label{intro}  Gamma Ray Bursts (GRBs) have been deeply studied in the keV -- MeV energy range, but only little information in the GeV range has been provided in the past decade by EGRET measurements. Only 3 bursts have been detected at energies $>$1 GeV, with one photon reaching 18 GeV \\cite{egret}. The study of the high energy emission of GRBs could provide extremely useful data able to constrain the emission models and the value of the ambient parameters. At these high energies the detection from space is hampered by the very low fluxes, requiring large collection areas. From ground the last generation Cherenkov telescope MAGIC \\cite{magic}, designed also to point at the detected GRB direction in a very short time, is still making efforts to lower the threshold at energies $\\leq$100 GeV. Its duty cycle and field of view are however very small: a wide field detector, able to cover simultaneously and continuously a significant ($\\sim$1 steradian) fraction of the sky, is thus necessary. With such a detector the Milagro Collaboration reported evidence for TeV emission from GRB 970417a with a significance slightly greater than 3$\\sigma$ \\cite{milagrito}. The study of the high energy spectrum of GRBs is perhaps the strongest motivation for an all-sky VHE detector. The study of transient phenomena can be successfully performed at energies down to 1 GeV by air showers arrays working in \"single particle mode\" \\cite{vernetto}, i.e., counting all the particles hitting the individual detectors during fixed time intervals. The observation of an excess in coincidence with a GRB detected by satellites would be an unambiguous signature of the nature of the signal. Both the sensitivity and the energy threshold improve with larger detection areas and higher observation levels, making air shower detectors at very high altitude the most suitable.  The ARGO-YBJ experiment, located at the YangBaJing Cosmic Ray Laboratory (4300 m a.s.l.) with a detection area of $\\sim$ 6700 $m^2$, is an air shower array exploiting the full coverage approach at very high altitude, with the aim of studying the cosmic radiation with a low energy threshold. In this paper we present results on the search for GRBs in coincidence with satellite detections performed in the December 2004 - May 2006 period with the ARGO-YBJ experiment. ",
        "conclusions": "A search for VHE emission from GRBs has been performed with an increasing detector area of the ARGO-YBJ experiment. A total of 16 satellite-triggered GRBs in the field of view ($\\theta\\leq 40^{\\circ}$) of ARGO-YBJ in the Dec. 2004 - May 2006 period has been analyzed. No significant emission was detected and typical fluence upper limits of $\\approx 10^{-4}$ erg cm$^{-2}$ in the 1 -- 100 GeV energy range were obtained using the measured counting rates and GRB parameters determined by the satellite observations. We expect to increase the sensitivity by a factor $\\sim$2 converting the secondary photons with a 0.5 cm thick layer of lead."
    },
    "0609/astro-ph0609121_arXiv.txt": {
        "abstract": "{} {We present \\emph{empirical} color transformations between the Sloan Digital Sky Survey (SDSS) \\emph{ugriz} photometry and Johnson-Cousins \\emph{UBVRI} system and Becker's \\emph{RGU} system, respectively. Owing to the magnitude of data that is becoming available in the SDSS photometric system it is particularly important to be able to convert between this new system and traditional photometric systems. Unlike earlier published transformations we based our calculations on stars actually measured by the SDSS with the SDSS 2.5-m telescope.  The photometric database of the SDSS provides in a sense a single-epoch set of 'tertiary standards' covering more than one quarter of the sky. Our transformations should facilitate their use to easily and reliably derive the corresponding approximate Johnson-Cousins or \\emph{RGU} magnitudes.  } {The SDSS survey covers a number of areas that were previously established as standard fields in the Johnson-Cousins system, in particular, fields established by Landolt and by Stetson. We used these overlapping fields to create well-photometered star samples on which our calculated transformations are based.  For the \\emph{RGU} photometry we used fields observed in the framework of the new Basel high-latitude field star survey.} {We calculated \\emph{empirical} color transformations between SDSS photometry and Johnson-Cousins \\emph{UBVRI} and Becker's \\emph{RGU} system.  For all transformations we found linear relations to be sufficient. Furthermore we showed that the transformations between the Johnson-Cousins and the SDSS system have a slight dependence on metallicity.} {} ",
        "introduction": "The Sloan Digital Sky Survey (SDSS) is the largest photometric and spectroscopic survey in the optical wavelength range.  The SDSS is mapping one quarter of the entire sky and is measuring positions and magnitudes for over 100 million celestial objects \\citep{York2000, Grebel2001}. The SDSS photometric system was specifically designed for the survey \\citep{Fukugita1996, Smith2002}, but many observatories worldwide have since purchased SDSS filters as well. The five SDSS \\emph{ugriz} filters  are a modified Thuan-Gunn system~\\citep{Thuan1976}.  Their most prominent feature are the wide passbands, which cover an effective wavelength range of $\\sim 380$ nm to $\\sim 920$ nm, and the lack of overlap between neighboring passbands.  These properties ensure a high efficiency for faint object measurements, independent spectral information in each band, and coverage of essentially the entire optical wavelength range accessible from the ground.  The SDSS provides a photometric catalog of unprecedented depth, homogeneity, and quality. Owing to the magnitude of the data that are becoming publicly available in this new photometric system \\citep{Stoughton2002, Abazajian2003, Abazajian2004, Abazajian2005, Adelman-McCarthy2006}, it is important to have well-calibrated transformation relations between this system and traditional photometric systems such as the Johnson-Cousins system \\citep{Johnson1953, Cousins1976}.  It is easy to imagine situations where one wishes to know, e.g., the V-band magnitude of a star of interest that happens to be in the SDSS database and for which no other photometry is available. Moreover, the large area coverage of the SDSS and the high quality of its drift-scan photometry make it also suitable as a source of \\textquoteleft tertiary standards\\textquoteright, although there is no information on variability for the majority of the SDSS objects.  This disadvantage is compensated by the large number of photometered sources even within a small patch of the sky.  Thus a few variable objects will merely appear as outliers and will not have a major effect on a photometric transformation.  Also, if observations in a traditional photometric system are being obtained during a non-photometric night, existing transformable SDSS photometry of the same field will prove very useful for at least an approximate photometric calibration \\citep{Koch2004a, Koch2004b}.  Other large-area photometric survey catalogs are already being used in this manner; for instance, \\citet{Udalski1998, Udalski2000, Udalski2002, Epchtein1999, Zaritsky2002, Zaritsky2004, Skrutskie2006}.  Prior to the start of SDSS observations, \\citet{Fukugita1996} derived theoretical transformation relations between the Johnson-Cousins system and the SDSS \\emph{u$'$g$'$r$'$i$'$z$'$} system.  The primes refer to the filter-detector combination envisioned to be used at the $20''$ photometric monitoring telescope at Apache Point Observatory. This auxiliary telescope observes SDSS photometric standards while the science observations are done in drift-scan mode with the actual SDSS survey telescope, a dedicated 2.5-m telescope at the same site (see, e.g., \\citet{York2000, Stoughton2002, Gunn2006}). \\citet{Fukugita1996} calculated synthetic magnitudes using the planned filter-detector combinations and spectral energy distributions of stars from the \\citet{Gunn1983} spectrophotometric atlas.  These synthetic magnitudes were then used to determine photometric transformations.  \\citet{Smith2002} calculated transformations between the SDSS \\emph{u$'$g$'$r$'$i$'$z$'$} system and the Johnson--Cousins photometric system based on actual observations in \\emph{u$'$g$'$r$'$i$'$z$'$} filters.  In this context the primes refer to the SDSS filter-detector combination used at the 1.0-m telescope at the US Naval Observatory (USNO), Flagstaff Station.  These observations were used to set up a system of 158 bright standard stars that define the \\emph{u$'$g$'$r$'$i$'$z$'$} system and to derive the above mentioned transformation equations.  \\citet{Smith2002} point out that there are ``small but significant'' differences between the USNO SDSS filters and the SDSS filters used at the 2.2-m telescope at Apache Point Observatory, leading to expected systematic differences between the USNO magnitudes and the final SDSS magnitudes.  Similarly, \\citet{Karaali2005} used observations obtained in \\emph{u$'$g$'$r$'$} filters at the Isaac Newton Telescope (INT) at La Palma, Spain.  These filters were designed to reproduce the SDSS system. These data were complemented by \\citet{Landolt1992} \\emph{UBV} standard star photometry and used to calculate transformations between the INT SDSS \\emph{u$'$g$'$r$'$} filter-detector combination and standard Johnson--Cousins photometry.  \\citet{Karaali2005} presented for the first time transformation equations depending on two colors. Since the INT filters and detector differ from the actual filter-detector combinations used by the SDSS, again systematic deviations are to be expected.  Direct empirical transformations between SDSS point-source photometry from the ``early data release'' (EDR; \\citet{Stoughton2002}) and certain Johnson-Cousins filters were calculated by \\citet{Odenkirchen2001} and by \\citet{Rave2003}.  \\citet{Rave2003} based their transformations on resolved stellar photometry of the Draco dwarf spheroidal galaxy obtained with the SDSS as well as with various other telescopes in different variants of the Johnson-Cousins system.  The SDSS EDR photometry was then superseded by later SDSS data releases \\citep{Abazajian2003, Abazajian2004, Abazajian2005, Adelman-McCarthy2006}.  \\begin{figure} \\centering \\includegraphics[width=88mm]{6082fig1.eps} \\caption{Comparison of the (normalized) filter curves of the Johnson-Cousins \\emph{UBVRI} passbands (upper panel, the SDSS \\emph{ugriz} passbands (middle panel), and Becker's \\emph{RGU} photometry in the lower panel.} \\label{fig:filters} \\end{figure}  A recent \\emph{theoretical} transformation was carried out by \\citet{Girardi2004}.  These authors used the actual SDSS passbands and CCD sensitivity curves of the SDSS camera \\citep{Gunn1998} employed at the 2.5-m SDSS telescope at Apache Point Observatory  to transform theoretical stellar isochrones into the SDSS \\emph{ugriz} system.  The third photometric system that we are considering here is the \\emph{RGU} system by Becker \\citep{Becker1946a}, a broad-band photometric system that was initially introduced with the goal of measuring stellar parameters and Galactic structure.  No earlier transformations between SDSS photometry and \\emph{RGU} photometry have been published. For transformation relations between \\emph{RGU} and \\emph{UBV} see \\citet{Steinlin1968a, Steinlin1968b, Bell1972, Buser1978a, Buser1978b}. While \\emph{RGU} photometry is not widely used, it plays a continuing role due to the existence of the ``new Basel high-latitude field star survey'' \\citep{Buser1985, Buser1998}.  The goal of this single-epoch photographic survey is to map the Galactic density structure and the metallicity distribution in the various components of the Milky Way.  For this purpose, 14 intermediate and high-latitude fields were observed, in which three-color \\emph{RGU} photometry of over 18\\,000 stars was measured. The limiting magnitude of this Galactic survey is approximately $G\\sim19$ or $V\\sim18$.  This survey may seem to be obsolescent by modern standards due to its small number of stars, the shallow exposures, the limited photographic resolution, and the availability of much deeper, more homogeneous, and much larger CCD surveys.  On the other hand, this photographic survey has played an important role in uncovering the structure of our Galaxy (e.g., \\cite{Buser1999}). Transformation relations between SDSS and \\emph{RGU} photometry help to evaluate the earlier findings with respect to modern explorations of the structure of the Milky Way (e.g., \\citet{Chen2001, Juric2005}). The Basel survey provides star counts (or space densities) also for distances relatively close to the Sun, whereas there is a gap at these distances in studies carried out using SDSS data due to the saturation limit at $r \\sim 14$.  Space densities for stars with brighter apparent magnitudes can be combined with Hipparchos parallaxes, thus providing reliable Galactic model parameters. Moreover, the metallicity sensitivity of the \\emph{RGU} system (e.g., \\cite{Buser1990}) adds continuing value to the Basel high-latitude field star survey.  Finally, this survey comprises a number of important Galactic sight lines that are not included in the currently available wide-field CCD surveys, and the analysis of these sight lines is continuing.  Here we present \\emph{empirical} color transformations between the SDSS \\emph{ugriz} photometric system using point-source photometry obtained with the SDSS 2.5-m telescope at Apache Point Observatory, and the Johnson-Cousins \\emph{UBVRI} system as well as the \\emph{RGU} system, respectively. These empirical colors have the advantage of not having to rely on synthetic colors or on passband-detector combinations of other telescopes.  Hence, they provide the most direct transformations possible.  In Sect. 2 we present the sources of our star samples and the criteria applied to choose the stars. Section 3 contains our calculations and the resulting transformation equations. In Sect. 4 we discuss our results and some possible future developments. ",
        "conclusions": "We presented empirical color transformations for the conversion of the SDSS \\emph{ugriz} photometric system into the Johnson-Cousins \\emph{UBVRI} system and into Becker's \\emph{RGU} system, respectively.  These are the first transformations between the SDSS and the Becker system, whereas several earlier transformations between variants of the SDSS and of the Johnson-Cousins system have been published.  For the Johnson-Cousins conversions, the novelty of our approach lies in the use of actual SDSS data obtained with the 2.5-m SDSS telescope at Apache Point Observatory and the use of well-defined Johnson-Cousins standard stars taken from the lists of \\citet{Landolt1992} and \\citet{Stetson2000}. For all transformations linear relations were found to be sufficient, with a slope change in transformations involving $(V-R)$.  A comparison with previously published transformations shows that they qualitatively reproduce our transformations, but that they show systematic differences that may amount to $\\sim 0.1$ mag. This is due to the fact that the earlier transformations were either done with data that used other filter--detector--telescope combinations than that of the actual 2.5-m SDSS telescope and therefore differ intrinsically from the actual SDSS data, or that the data are based on early versions of the SDSS photometry catalogs before the SDSS photometry was in its final form, or that non-standard Johnson-Cousins data were used.  For conversions between the SDSS and the Johnson-Cousins systems, we tested the transformations for a possible metallicity dependence. Some of the Stetson standard stars lie in a field centered on the Draco dwarf spheroidal galaxy or on metal-poor ($\\sim -2.0$ dex) globular clusters.  Most of the stars in these fields are Population II stars and have much lower metallicities than the majority of the stars in Landolt's equatorial standard fields, where Population I stars dominate.  The latter sample necessarily comprises stars with a range of ages and abundances, while our Population II sample is essentially limited to very old and very metal-poor stars.  The calculated coefficients for the metal-poor stars result in slightly different slopes as compared to the more metal-rich stars, which affect very blue and very red stars most.  Although our metallicity differentiation is quite crude, there does seem to be a slight dependence of the transformations on the metallicity.  A more accurate separation into metal-rich and metal-poor stars -- ideally aided by spectroscopically measured stellar abundances -- is needed to fully evaluate the magnitude of this trend."
    },
    "0609/astro-ph0609298_arXiv.txt": {
        "abstract": "Theoretical studies predict that Trojans are likely a frequent byproduct of planet formation and evolution.  We present a novel method of detecting Trojan companions to transiting extrasolar planets which involves comparing the midtime of eclipse with the time of the stellar reflex velocity null.  We demonstrate that this method offers the potential to detect terrestrial-mass Trojans using existing ground-based observatories.  This method rules out Trojan companions to HD 209458b and HD 149026b more massive than $\\simeq~13~M_{\\oplus}$ and $\\simeq~25~M_{\\oplus}$ at a 99.9\\% confidence level.  Such a Trojan would be dynamically stable, would not yet have been detected by photometric or spectroscopic monitoring, and would be unrecognizable from radial velocity observations alone.  We outline the future prospects for this method, and show that the detection of a ``Hot Trojan'' of any mass would place a significant constraint on theories of orbital migration. ",
        "introduction": "Stable Trojan companions to extrasolar planets may be common.  In our solar system, Mars, Jupiter, and Neptune share their orbit with asteroids orbiting near the stable (L4/L5) Lagrange points that lead/trail the planet by $\\simeq~60^{\\circ}$.  Orbits near the L4/L5 points of the terrestrial planets, Saturn and Uranus are less stable due to perturbations from the other planets (Nesvorny \\& Dones 2002).  Saturn's satellites also include small moons orbiting about the L4/L5 points of Tethys and Dione.  While the mass ratios of the Trojan systems in our solar system are rather extreme ($\\le7\\times10^{-9}$), extrasolar planets may have more massive Trojans.  Theorists have already outlined several mechanisms to form Trojans with mass ratios as large as unity.  For example, Laughlin \\& Chambers (2002) present hydrodynamic simulations of a protoplanetary disk where disk material lingers near the L4 and L5 points of a planet (near the gap-opening threshold). The resulting vortex could trap particles and lead to the accretion of a Trojan {\\em in situ} (Chiang \\& Lithwick 2005).  If disk torques caused the planet to gradually migrate inwards, then the Trojan would migrate with the planet.  Unlike resonant migration in the 2:1 mean motion resonance, the eccentricity or libration amplitude of the Trojan would not be excited by the migration (Laughlin \\& Chambers 2002). Alternatively, a body could be captured into an orbit about the L4/L5 point after a violent event, as has been suggested for the formation of Jupiter and Neptune Trojans in our solar system (Morbidelli et al.\\ 2005).  Capture into a Trojan orbit could also occur due to rapid mass growth of the planet or a collision of two objects near L4/L5 (Chiang \\& Lithwick 2005 and reference therein).  Another possibility is that convergent migration could trap multiple protoplanets into a 1:1 mean motion resonance (Thommes 2005, Cresswell \\& Nelson 2006). In each of these scenarios, the captured bodies could initially have a large libration amplitude or reside on horseshoe-type orbits.  However, if the capture occurred before or during the planet's inward migration, then interactions with either a gaseous or planetesimal disk would damp the libration amplitude.  This mechanism is even capable of causing objects initially on horseshoe orbits to evolve into tadpole orbits and then small amplitude libration near the L4/L5 fixed point.  Such behavior has been found in numerical simulations of multiple planet systems interacting with either a gaseous or planetesimal disk (Cresswell \\& Nelson 2006; Ford \\& Chiang 2006).  Trojans of Jupiter and Neptune have provided clues about our solar system's history (Michtchenko, Beauge \\& Roig 2001; Kortenkamp, Malhotra \\& Michtchenko 2004; Chiang \\& Lithwick 2005; Morbidelli et al.\\ 2005).  Similarly, the detection of extrasolar Trojans would be useful for constraining theories of planet formation. While all the above mechanisms predict that Trojans would survive the migration process, there are alternative models of planet migration that predict Trojans would not survive. For example, while planet formation models generally agree that planets should form on nearly circular orbits, it is possible that gravitational perturbations by other planets or a binary companion could excite sizable eccentricities.  One possible formation mechanism for short-period giant planets is that a planet acquires a large eccentricity (e.g., due to strong planet-planet scattering, secular perturbations from a binary companion, or being tidal captured) and comes so close to the star that tidal dissipation circularizes the orbit at a small semimajor axis (Rasio \\& Ford 1996, Wu \\& Murray 2003, Gaudi 2003, Ford \\& Rasio 2006).  Detecting a Trojan companion to a short-period planet would present a serious challenge for these mechanisms for forming ``hot Jupiters'' and would imply that the planet in such a system was formed via migration through a dissipative disk. Thus, searching for extrasolar Trojans can test models of planet formation. Here, we present a method for detecting Trojan companions to extrasolar planets by combining RV and photometric observations of transiting extrasolar planets.  We refer to all bodies librating about the L4/L5 fixed point of a planet as ``Trojans'' and focus our attention on Trojans that are significantly less massive than the currently known planet and not currently recognizable from radial velocity (RV) observations alone. ",
        "conclusions": "In principle, Trojans could be detected via their radial velocity, astrometric, transit, or transit timing signatures. If a Trojan is sufficiently massive and has a sufficiently large libration amplitude, then it could be detected from the deviations from a Keplerian perturbation to the stellar radial velocity or astrometric signal caused by a single planet.  Laughlin \\& Chambers (2002) have shown that two comparable mass planets occupying a 1:1 mean motion resonance would typically have strong planet-planet gravitational interactions on a secular timescale. However, these signatures may not be unique: a reanalysis of the RV observations of HD 128311 and HD 82943 have shown that both of the current data sets are consistent with a pair of planets in a 1:1 mean motion resonance (Gozdziewski \\& Konacki 2006), as well as the originally published orbital solutions involving higher-order mean motion resonances.  Trojans may also be detectable if they transit their parent star. Photometric or spectroscopic monitoring of stars with transiting planets (particularly at times offset from the planet transit by $\\sim~P$/6) may reveal the Trojan transit via the decrease in stellar flux or anomalous RV excursions due to the Rossiter-McLaughlin effect (Gaudi \\& Winn 2006).  For planets discovered via a photometric transit search, there will typically be observations at epochs useful for searching for large Trojans.  Unfortunately, a Trojan might not transit its parent star if it has a significant inclination (e.g., Morbidelli et al.~ 2005).  Since the libration period can be large, long-term monitoring would be required to ensure detection. Currently, the most stringent photometric constraints on Trojan companions to HD 209458b come from the continuous photometry of the system for 14 days by the MOST satellite (Rowe et al.\\ 2006).   When heavily binned into $\\simeq~2$~hr intervals, this photometry has a fractional uncertainty of $\\simeq~3\\times10^{-4}$ (Rowe et al.\\ 2006). Assuming an average density equal to that of Earth, this corresponds to a 3-$\\sigma$ detection threshold of $48M_{\\oplus}$.  It is not clear whether the data reduction techniques used in their analysis of HD 209458 might subtract part of the signal due to a Trojan (Rowe et al.\\ 2006).  Regardless, if a Trojan had a vertical libration amplitude greater than $\\simeq~9^{\\circ}$, it would not always transit the star.  Since the putative libration period of $\\simeq 53$~d (Murray \\& Dermott 2000) is significantly longer than the duration of the MOST observations, it could have missed even a much larger Trojan.  In a sense, our method is most similar and complementary to the recently-proposed method of searching for gravitational perturbations due to low mass planets using transit timing (Holman \\& Murray 2005; Agol et al.\\ 2005).  In contrast to the transit timing method, the unique geometry of Trojan orbits results in a nearly constant perturbation (assuming small amplitude libration about L4/L5) that recurs at {\\em every} transit.  In principle, it is not necessary to make precise measurements of the time of many transits to search for a complex pattern of perturbations.  Thus, our method can be practically applied to transiting planets with long periods.  Nevertheless, multiple transits should be observed to ensure that the same offset is observed and avoid potential confusion with perturbations from a more distant planet (Holman \\& Murray 2005; Agol et al.\\ 2005).  In principle, our technique could be applied to search for terrestrial-mass Trojans of giant planets orbiting in the habitable zone of their stars (Ji et al.\\ 2005; Schwarz et al.\\ 2005).  While present search techniques are strongly biased towards finding transiting planets at short orbital periods, future space missions (e.g., Corot, Kepler) offer the prospect of finding transiting planets in the habitable zone of their stars, particularly for low mass stars where the habitable zone can be $\\simeq~0.015$~AU away from the star."
    },
    "0609/astro-ph0609251_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:intro}  Globular clusters (GCs) are well known as useful probes of the formation histories of galaxies, especially at early times. {\\it Kinematics} of globular cluster systems (GCSs) brings an extra dimension to these probes, helping firstly to {\\it distinguish GC subsystems} by their distinctive kinematical signatures. The need for such information is highlighted by recent doubts about bimodalities in GC metallicity distributions. Kinematical parameters can also help identify {\\it correlations between GCs and field stars} within galaxies, and so determine which objects had a common origin. The present-day motions of GCs can provide {\\it signatures of their formation and evolution,} as different physical processes will lead to different orbital properties. Last but not least, the residence of GCs at very large radii around galaxies makes them excellent {\\it halo mass tracers}.  The informative capacity of GCS kinematics can be illustrated by the well-known case of the Milky Way, where there are at least two subpopulations (the bulge/disk and halo GCs). The cornerstone characteristics of these subpopulations are their spatial distributions, metallicities, and kinematics (e.g. \\cite{romanow:zinncote,romanow:pritzl05}): without all three clues, the distinction would be much less clear. One may make further decompositions into ``thin/thick'' disk GCs, and ``young/old'' halo GCs---which is a level of detail difficult to achieve in external galaxies. Note that these nomenclatures are not arbitrary, but connect to the picture that GC and field star formation are associated. GCs have also proved valuable for measuring the mass of the Milky Way halo (e.g. \\cite{romanow:wilkinson99,romanow:sakamoto03}). ",
        "conclusions": ""
    },
    "0609/astro-ph0609067_arXiv.txt": {
        "abstract": "In this paper and a companion paper, we report on a wide-field imaging survey of the globular cluster (GC) populations around M87 carried out with Suprime-Cam on the 8.2m Subaru telescope. Here we describe the observations, data reduction, and data analysis and present luminosity functions of GC populations around M87 and NGC 4552, another luminous Virgo elliptical in our survey field. The imaging data were taken in the $B$, $V$, and $I$ bands with a sky coverage of $2^{\\circ} \\times 0_{\\cdot}^{\\circ}5$ extending from the M87 centre out to $\\sim$ 0.5 Mpc. GC candidates were selected by applying a colour criterion on the $B-V$ and $V-I$ diagram to unresolved objects, which greatly reduces contamination. The data from control fields taken with Subaru/Suprime-Cam were also analyzed for subtraction of contamination in the GC sample. These control field data are compatible with those in the M87 field in terms of the filter set ($BVI$), limiting magnitudes, and image quality, which minimizes the possibility of introducing any systematic errors into the subtractive correction. We investigate GC luminosity functions (GCLFs) at distances $\\leq 10^{\\prime}$ ($\\lesssim$ 45 kpc) from the host galaxy centre in detail. By fitting Gaussians to the GCLFs, the $V$-band turnover magnitude ($V_{\\rm TO}$) is estimated to be $23.62 \\pm 0.06$ mag and $23.56 \\pm 0.20$ mag for the GC population in M87 and NGC 4552, respectively. The GCLF is found to be a function of GC colour; $V_{\\rm TO}$ of the red GC subpopulation ($V-I > 1.1$) is fainter than that of the blue GC subpopulation ($V-I \\leq 1.1$) in both M87 and NGC 4552, as expected if the colour differences are primarily due to a metallicity effect and the mass functions of the two subpopulations are similar. The radial dependence of the GCLF is also investigated for the GC population in M87. The GCLF of each subpopulation at $1^{\\prime} \\leq R \\leq 5^{\\prime}$ is compared with that at $5^{\\prime} \\leq R \\leq 10^{\\prime}$ but no significant trend with distance is found in the shape of the GCLF. We also estimate GC specific frequencies ($S_N$) for M87 and NGC 4552. The $S_N$ of the M87 GC population is estimated to be 12.5 $\\pm$ 0.8 within 25$^{\\prime}$. The $S_N$ value of the NGC 4552 GC population is estimated to be 5.0 $\\pm$ 0.6 within $10^{\\prime}$. ",
        "introduction": "Globular clusters (GCs) are homogeneous stellar systems containing stars with a single age and metallicity, which are in principle simpler to interpret than photometric and spectroscopic observations of the integrated stellar light of a galaxy. GCs are therefore considered to be powerful probes with which to understand the star formation and chemical enrichment history of their host galaxy. One of the basic findings from observations about GC populations in luminous galaxies is that while thousands of GCs are associated with luminous elliptical galaxies, a significantly smaller number of GCs exist around spiral galaxies with similar luminosities (e.g., Harris 1991; Barmby 2003). This indicates that the specific frequency of GCs ($S_N$), which is considered to be related to the relative efficiency of GC formation and/or survival compared to galactic halo/bulge stars, depends on galaxy morphology. In fact, $S_N$ has also been suggested to be correlated with local galaxy density, with galaxies in denser environments having larger $S_N$ values (West 1993). The fact that this trend appears to exist even when the sample of galaxies is restricted to ellipticals suggests that GC formation efficiency is more physically linked with galaxy environment. One possibility to explain this observation is biased GC formation in galaxies inhabiting denser environments (West 1993; Blakeslee 1999; McLaughlin 1999). Alternatively, a substantial number of GCs in a luminous galaxy may have an external origin; GCs could be captured from other galaxies through galaxy interactions and accrete onto luminous galaxies in clusters, which would then enhance $S_N$ values as observed.  Recent studies of GCs in the central regions of luminous ellipticals conducted with the {\\it Hubble Space Telescope} (HST) have revealed that many luminous ellipticals have bimodal or multimodal colour distributions of GCs (Gebhardt \\& Kissler-Patig 1999; Larsen et al. 2001; Brodie et al. 2005). It is found that the mean colours of both red (metal-rich) and blue (metal-poor) GC subpopulations are correlated with the host galaxy luminosities and colours (Larsen et al. 2001; Strader, Brodie \\& Forbes 2004; Strader et al. 2005; Peng et al. 2006), which will be important constraints on the proposed scenarios for the formation and evolution of GC population such as multiphase collapse scenario (Forbes, Brodie \\& Grillmair 1997), merger scenario (Ashman \\& Zepf 1992), hiearchical merging scenario (Beasley et al.  2002), and accretion scenario (C\\^{o}t\\'{e}, Marzke \\& West 1998). The small field of view of the HST, however, does not allow one to collect GC populations in the outer halo of a galaxy and investigate their spatial structures, which will also be key pieces of the puzzle (e.g., Moore et al. 2005). Much wider-field studies (e.g., out to $\\sim$ 100 kpc from the galaxy centre) of GC populations therefore need to be performed using ground-based data (Rhode \\& Zepf 2001, 2004; Dirsch et al. 2003; Bassino et al.  2006). Studying GC populations at large distances from the host galaxy is of great importance because the outer halo of the host galaxy, and even the intergalactic space, are presumed to be large reservoirs of blue and metal-poor GCs in the accretion scenario.  \\setlength{\\tabcolsep}{1.5mm} \\begin{table*} \\centering \\begin{minipage}{150mm} \\begin{tabular}{lccccccccc} \\hline\\hline \\multicolumn{1}{c}{Field ID}     & \\multicolumn{3}{c}{Field centre} & \\multicolumn{3}{c}{Integration time ($m_{\\rm lim}$)} & \\multicolumn{3}{c}{Seeing size} \\\\[2pt] & & & & \\multicolumn{3}{c}{[sec (mag)]}  & \\multicolumn{3}{c}{[arcsec]} \\\\ \\multicolumn{1}{c}{(1)} & \\multicolumn{3}{c}{(2)} & \\multicolumn{3}{c}{(3)} & \\multicolumn{3}{c}{(4)} \\\\ \\hline & $\\alpha$(J2000) & $\\delta$(J2000) & $b$(J2000)      & $B$ & $V$ & $I$ & $B$ & $V$ & $I$ \\\\  Field 1       & $12^{h}31^{m}18^{s}_{\\cdot}4$ & $12^{\\circ}29^{\\prime}13^{\\prime\\prime}$ & 74$^{\\circ}_{\\cdot}$6 & 3680 (25.6) & 1350 (25.1) & 3480 (24.6) & 1.8 & 1.0 & 1.0 \\\\ Field 2       & $12^{h}33^{m}42^{s}_{\\cdot}8$ & $12^{\\circ}27^{\\prime}37^{\\prime\\prime}$ & 74$^{\\circ}_{\\cdot}$8 & 2640 (25.9) & 1350 (25.2) & 3690 (24.7) & 1.2 & 1.0 & 1.0 \\\\ Field 3       & $12^{h}36^{m}08^{s}_{\\cdot}8$ & $12^{\\circ}24^{\\prime}35^{\\prime\\prime}$ & 74$^{\\circ}_{\\cdot}$9 & 1800 (25.7) & 1350 (25.2) & 2640 (24.5) & 1.5 & 1.0 & 1.1 \\\\ Field 4       & $12^{h}38^{m}33^{s}_{\\cdot}2$ & $12^{\\circ}24^{\\prime}35^{\\prime\\prime}$ & 75$^{\\circ}_{\\cdot}$0 & 2280 (26.1) & 1350 (25.2) & 4380 (24.8) & 1.1 & 1.0 & 1.0 \\\\ \\hline HDF-N         & $12^{h}36^{m}46^{s}_{\\cdot}7$ & $62^{\\circ}11^{\\prime}50^{\\prime\\prime}$ & 54$^{\\circ}_{\\cdot}$8 & 6000 (26.6) & 4800 (25.5) & 4200 (25.2) & 0.8 & 1.1 & 0.9 \\\\ Lockman Hole  & $10^{h}35^{m}55^{s}_{\\cdot}2$ & $57^{\\circ}42^{\\prime}18^{\\prime\\prime}$ & 51$^{\\circ}_{\\cdot}$3 & 6000 (26.9) & 4800 (26.0) & 3600 (24.8) & 1.0 & 1.1 & 1.3 \\\\ \\hline \\end{tabular}  \\caption{Observation log. Data for HDF-N and Lockman Hole are retrieved from SMOKA. Col. (3): The integration times were calculated including data taken under non-photometric conditions on the first night (hence the interpretation is not straightforward). Number in the parentheses indicates 50 \\% completeness to point sources estimated with artificial star test. Col. (4): The seeing sizes are estimated in the stacked images.} \\label{basic}  \\end{minipage} \\end{table*}  \\begin{figure*} \\begin{center} \\includegraphics[height=11cm,keepaspectratio,clip]{ntamurafig01.ps} \\caption{Locations and IDs of the observed fields. Each box indicates one SCam field of view. The background is a DSS image. North is up and East is left.}  \\label{deff} \\end{center} \\end{figure*}  In this paper and a companion paper (Tamura et al. 2006; Paper II hereafter), we report a wide-field imaging survey of the GC populations around M87 conducted with Suprime-Cam on the Subaru telescope. Several moderately wide-field studies of GCs surrounding M87 have already been carried out using photometry (Strom et al. 1981; McLaughlin, Harris, \\& Harris 1994; Harris, Harris, \\& McLaughrin 1998; Hanes et al. 2001) and spectroscopy (Cohen \\& Rizhov 1997; Cohen, Blakeslee \\& Rizhov 1998; Kissler-Patig \\& Gebhardt 1998; C\\^{o}t\\'{e} et al. 2001), but all of these studies explored only the regions $\\lesssim 10^{\\prime}$ ($\\lesssim$ 45 kpc or 5 $r_e$) from M87. In contrast, the area of our survey is approximately $2^{\\circ} \\times 0_{\\cdot}^{\\circ}5$ (560 kpc $\\times$ 140 kpc) extending from the centre of M87 out to $\\sim$ 0.5 Mpc, which is the widest survey yet undertaken of the GC populations in luminous galaxies. In this paper, we focus on describing the observations, data reduction (\\S~\\ref{obsredcal}) and data analyses such as selection and photometry of GC candidates, incompleteness correction, and subtraction of foreground and background contamination in the GC candidates (\\S~\\ref{analyses}). We derive GC luminosity functions around M87 and NGC 4552, another Virgo luminous elliptical galaxy in our survey field, and estimate the global GC specific frequencies of these luminous ellipticals in \\S~\\ref{resultsanddiscussions}. We investigate colour distributions and spatial distributions of GC candidates in Paper II. We adopt distances of 16.1 Mpc (distance modulus of 31.03) to M87, and 15.4 Mpc (distance modulus of 30.93) to NGC 4552, based on measurements using the surface brightness fluctuation method (Tonry et al. 2001). An angular scale of 1$^{\\prime}$ corresponds to 4.7 kpc and 4.5 kpc at the distance of M87 and NGC 4552, respectively. ",
        "conclusions": "\\label{resultsanddiscussions}  \\subsection{Globular Cluster Luminosity Function}  In Fig. \\ref{gclf_cor}, GCLFs obtained within annuli centered on M87 are presented. The dotted lines indicate raw GCLFs without incompleteness correction or control field subtraction, and the thin black lines show GCLFs after the incompleteness is corrected for (but no subtractive correction for contamination is performed yet). The red (thick black) lines show GCLFs after the control field subtraction is also performed based on the HDF-N (LH) field, respectively. In Fig. \\ref{gclf}, the GCLFs after the incompleteness correction and field subtraction using the HDF-N data are divided into red ($V-I > 1.1$) and blue ($V-I \\leq 1.1$) GC subpopulations. The black line shows the GCLF for the total GC population (blue $+$ red). The boundary colour of the blue and red GCs corresponds approximately to the middle of the peak colours in the bimodal colour distributions (see Paper II). Note that there is a local maximum of GC number in the $70^{\\prime} < r < 80^{\\prime}$ bin due to the GC population of NGC 4552 at a distance of $\\sim$ 75$^{\\prime}$ from M87. GCLFs in the outermost regions tend to have some negative bins due to the field subtraction. In both of these figures, the shaded region indicates the magnitude range where the completeness is lower than 50 \\%. Note that the completeness is calculated by considering not only the simple detection completeness on the $V$ band image but also the selection efficiency of unresolved objects based on the \\verb|CLASS_STAR| index and the completeness in $B$ and $I$ bands at the corresponding magnitudes depending on the GC colours. We investigate and discuss the spatial distributions of GC populations in detail in Paper II.  A Gaussian fitted to the GCLF obtained using only GCs at distances smaller than $10^{\\prime}$ ($\\sim$ 45 kpc) from the host galaxy is plotted in Fig. \\ref{gauss} (we fit a Gaussian to the binned data). The fainter part of the GCLF in the shaded region where the completeness is lower than 50 \\% is not used in this fitting process. The turnover magnitude ($V_{\\rm TO}$) and dispersion ($\\sigma$) of the GCLF are then estimated to be $V_{\\rm TO} = 23.62 \\pm 0.06$ mag and $\\sigma = 1.50 \\pm 0.04$ mag for the GC population in M87. For the NGC 4552 GCs, $V_{\\rm TO} = 23.56 \\pm 0.20$ and $\\sigma = 1.34 \\pm 0.12$ mag. In Fig. \\ref{gclfcol}, the GCLFs of the red and blue GC subpopulations are shown with fitted Gaussians. Again the shaded region indicates the magnitude range where the completeness calculated by using all the GCs (i.e., blue $+$ red) is lower than 50 \\%. This limiting magnitude does not change a lot if only the red or blue GC subpopulation is considered since our B- and I-band images are deep enough not to miss a significant number of red or blue globular clusters detected on the V-band image. The fitted Gaussians suggest that the GCLF depends on subpopulation; the $V_{\\rm TO}$ of the red GC subpopulation is $\\sim 0.5$ mag and 0.4 mag fainter than that of the blue one for M87 and NGC 4552, respectively, although the results for the NGC 4552 GCs are less significant due to the large errors. Larsen et al. (2001) present similar results using the HST/WFPC2 data, although their $V_{\\rm TO}$ values tend to be brighter than those from other studies. This difference of $V_{\\rm TO}$ between the GC subpopulations is perhaps because of a metallicity difference (Ashman, Conti \\& Zepf 1995; Elson \\& Santiago 1996; Jord\\'{a}n et al. 2002). $V_{\\rm TO}$ and $\\sigma$ of the GCLFs are summarized in Table \\ref{gclfparams}.  \\begin{figure*} \\begin{center} \\includegraphics[height=14cm,angle=-90,keepaspectratio]{ntamurafig16.ps} \\end{center} \\caption{The GCLF of all GCs, red GCs, and blue GCs is indicated by a black solid line, dotted line, and grey solid line, respectively. Incompleteness was corrected and contamination was subtracted based on the HDF-N data. Gaussians fitted to the GCLFs are also plotted.} \\label{gclfcol} \\end{figure*}  \\begin{figure*} \\begin{center} \\includegraphics[height=14cm,angle=-90,keepaspectratio]{ntamurafig17.ps} \\end{center} \\caption{Same as Fig. \\ref{gclfcol}, but the GCLFs in the region of $1^{\\prime} \\leq R \\leq 5^{\\prime}$ around M87 are compared with those in the region of $5^{\\prime} \\leq R \\leq 10^{\\prime}$.} \\label{gclfinout} \\end{figure*}  The GCs within $10^{\\prime}$ of M87 are further divided into two samples at a boundary of $5^{\\prime}$ and the GCLFs of all GCs, red GCs, and blue GCs in the inner and outer regions are presented in Fig. \\ref{gclfinout}. Gaussians are also fitted to these GCLFs and the $V_{\\rm TO}$ and $\\sigma$ are summarized in Table \\ref{gclfparams}. This indicates that the shape of GCLF is not significantly different between the inner region ($1^{\\prime} \\leq R \\leq 5^{\\prime}$) and the outer region ($5^{\\prime} \\leq R \\leq 10^{\\prime}$) for red GCs or blue GCs. Furthermore, the GCLF shape for all GCs is consistent with that in the core region ($\\leq 1^{\\prime}$) obtained by Kundu et al. (1999) using HST/WFPC2: $V_{\\rm TO} = 23.67 \\pm 0.07$ mag ($\\sigma = 1.39 \\pm 0.06$). This suggests that the GCLF shape of the M87 GCs is not a strong function of distance from the host galaxy. We note that the fainter part of the GCLF for all GCs tends to be more deficient at the larger distance; this is due to the lower contribution of the red GC subpopulation in the outer region. Any radial dependence of the GCLF is unclear for the GC population around NGC 4552 because the number of GCs is not substantial enough especially in the outer region ($5^{\\prime} \\leq R \\leq 10^{\\prime}$). Nevertheless, the $V_{\\rm TO}$ of the GCLF obtained at $1^{\\prime} \\leq R \\leq 10^{\\prime}$ is consistent with that in the core region obtained by Kundu \\& Whitmore (2001) using HST/WFPC2 who found $V_{\\rm TO} = 23.47 \\pm 0.14$ mag with a fixed value of $\\sigma = 1.3$, again suggesting that the GCLF shape does not depend significantly on distance from the host galaxy.  \\subsection{GC Specific Frequency}  Using the $V_{\\rm TO}$ and $\\sigma$ of the Gaussians fitted to the GCLFs presented in Fig. \\ref{gauss}, we can estimate the total number of GCs associated with M87 or NGC 4552 and calculate a GC specific frequency ($S_N$). We first integrate a GCLF obtained in an annulus ($1^{\\prime}$ width) centered on the host galaxy down to the magnitude where the completeness becomes 50 \\%. This number is then multiplied by a correction factor to include the fainter GCs based on the Gaussian GCLF. This calculation is repeated out to a certain distance from the galaxy centre. We note that it is ideal to fit a Gaussian to a GCLF obtained in each annulus but this is difficult in practice, especially at large distances because the number of GCs is not substantial and statistical errors are significant. We therefore use the fitted Gaussian in the inner region ($\\leq 10^{\\prime}$) of M87 or NGC 4552 independently of the distance from the host galaxy. In order to estimate the number of GCs at distances smaller than $\\leq 1^{\\prime}$, where the brightness of host galaxy halo light exceeds the linearity regime of the CCD, the radial profile of GC surface density outside the core region is fitted with a de Vaucouleurs law profile (see Paper II for details) and the fitted formula is extrapolated towards the galaxy centre. The estimated total number of M87 GCs is $12000 \\pm 800$ within $25^{\\prime}$ from the galaxy centre (cf. $13200 \\pm 1500$ by HHM). For NGC 4552, the total number of GCs is estimated to be $1400 \\pm 170$ within $10^{\\prime}$. Note that the fraction of GCs at $R \\leq 1^{\\prime}$ estimated from the extrapolation is 8.5 \\% for M87 GCs and 17 \\% for NGC 4552 GCs. If we adopt $M_V = -22.46$ mag as the $V$-band absolute magnitude of M87, which is also adopted by HHM (a difference in adopted distance modulus of 0.03 mag is corrected), the $S_N$ value of M87 is calculated to be $12.5 \\pm 0.8$, which is only slightly smaller than $14.1 \\pm 1.6$ obtained by HHM. The $V$-band luminosity of NGC 4552 is obtained from our data by fitting a de Vaucouleurs law to the $V$-band surface brightness profile and integrating it out to the same distance (10$^{\\prime}$); $M_V = -21.12$ mag. The $S_N$ value of NGC 4552 is then estimated to be $5.0 \\pm 0.6$.\\footnote{We calculated the total number of GCs and $S_N$ within 10$^{\\prime}$ for NGC 4552 to avoid possible contributions of M87 GCs and intergalactic GCs outside of this radius (see Paper II). If we fit a de Vaucouleurs law to the GC surface density profile within 10$^{\\prime}$ from the NGC 4552 centre and integrate it out to 25$^{\\prime}$ as done for M87 GCs, the number of GCs and $S_N$ are estimated to be $2000 \\pm 660$ and $6.2 \\pm 2.1$, respectively.}"
    },
    "0609/astro-ph0609584_arXiv.txt": {
        "abstract": "We trace the assembly history of red galaxies since $z=1$, by measuring their evolving space density with the $B$-band luminosity function. Our sample of 39599 red galaxies, selected from $6.96~{\\rm deg}^2$ of imaging from the NOAO Deep Wide-Field and {\\it Spitzer} IRAC Shallow surveys, is an order of magnitude larger, in size and volume, than comparable samples in the literature. We measure a higher space density of $z\\sim 0.9$ red galaxies than some of the recent literature, in part because we account for the faint yet significant galaxy flux which falls outside of our photometric aperture. The $B$-band luminosity density of red galaxies, which effectively measures the evolution of $\\sim L^*$ galaxies, increases by only $36\\pm 13\\%$ from $z=0$ to $z=1$. If red galaxy stellar populations have faded by  $\\simeq 1.24$ $B$-band magnitudes since $z=1$, the stellar mass contained within the red galaxy population has roughly doubled over the past $8~{\\rm Gyr}$.  This is consistent with star-forming galaxies being transformed into $\\lesssim L^*$ red galaxies after a decline in their star formation rates. In contrast, the evolution of $\\simeq 4L^*$ red galaxies differs only slightly from a model with negligible $z<1$ star formation and no galaxy mergers. If this model approximates the luminosity evolution of red galaxy stellar populations, then $\\simeq 80\\%$ of the stellar mass contained within today's $4L^*$ red galaxies was already in place at $z=0.7$. While red galaxy mergers have been observed, such mergers do not produce rapid growth of $4L^*$ red galaxy stellar masses between $z=1$ and the present day. ",
        "introduction": "\\label{sec:intro}  Red galaxies contain the majority of the stellar mass at low redshift \\markcite{hog02}({Hogg} {et~al.} 2002), so understanding their formation and assembly is one of the key goals of both observational and theoretical extragalactic astronomy. The stellar populations of these galaxies are dominated by an old component, with little ongoing star formation \\markcite{tin68}(e.g., {Tinsley} 1968). Observationally, it is unclear if this star formation occurred in situ at $z>1$, or if these stars formed in lower mass galaxies which were later assembled via galaxy mergers. Simulations of red galaxy evolution predict varying assembly and star formation histories, with some papers concluding that assembly takes place primarily at high redshift \\markcite{mez03,naa05}(e.g., {Meza} {et~al.} 2003; {Naab} {et~al.} 2005) while others suggest it continues today \\markcite{del06}(e.g., {De Lucia} {et~al.} 2006).  The color-magnitude diagram of low redshift galaxies provides several important clues about the assembly of red galaxies. At low redshift there is a tight locus of red galaxies in color-magnitude space \\markcite{bow92,hog04,mci05}(e.g., {Bower}, {Lucey}, \\& {Ellis} 1992; {Hogg} {et~al.} 2004; {McIntosh} {et~al.} 2005), with the most luminous red galaxies being slightly redder than less luminous red galaxies. This indicates that red galaxies of a given luminosity and metallicity have similar star formation histories, at least over the last few ${\\rm Gyr}$. A red galaxy locus is also observed within the $z\\sim 1$ cluster  and field galaxy populations \\markcite{van00,bla03,bel04,wil05}({van Dokkum} {et~al.} 2000; {Blakeslee} {et~al.} 2003; {Bell} {et~al.} 2004; {Willmer} {et~al.} 2005). The absence of very massive star-forming galaxies at low redshifts indicates that the most massive red galaxies are either formed at high redshift or they are assembled via mergers at $z<1$. As the colors of red galaxies are a function of luminosity, mergers of red galaxies should produce galaxies slightly blueward of the color-magnitude relation, and thus increase the scatter within the relation.  Mergers of red galaxies, perhaps without significant merger-triggered star formation (dry mergers), have been observed at low redshift \\markcite{lau88,van05}(e.g., {Lauer} 1988; {van Dokkum} 2005), but it is unclear how large a role they play in galaxy formation. While there have been valiant attempts to measure the galaxy merger rate with redshift \\markcite{lef00,bel06}(e.g., {Le F{\\`e}vre} {et~al.} 2000; {Bell} {et~al.} 2006), there is debate about the selection function of merger candidates and the duration of observable galaxy mergers. At low redshift, measured rates of red galaxy stellar mass growth via mergers span from $\\simeq 1\\%$ per ${\\rm Gyr}$ \\markcite{mas06}({Masjedi} {et~al.} 2006) to $\\simeq 10\\%$ per ${\\rm Gyr}$ \\markcite{van05}({van Dokkum} 2005).  The space density of red galaxies with redshift is a relatively direct measure of the galaxy assembly history, and has fewer uncertainties than measurements of the galaxy merger rate. At the redshifts where galaxy assembly takes place, the space density of galaxies as a function of stellar mass will evolve with redshift. While conceptually simple, it is difficult to measure the galaxy space density accurately, as $z=1$ $L^*$ red galaxies are optically faint and their strong spatial clustering \\markcite{bro03}(e.g., {Brown} {et~al.} 2003) results in significant cosmic variance. Despite these difficulties, several groups have used measurements of the space density of galaxies with redshift to constrain red galaxy evolution and assembly \\markcite{lil95,lin99,che03,bel04,bun05,wil05,wak06}(e.g., {Lilly} {et~al.} 1995; {Lin} {et~al.} 1999; {Chen} {et~al.} 2003; {Bell} {et~al.} 2004; {Bundy} {et~al.} 2005; {Willmer} {et~al.} 2005; {Wake} {et~al.} 2006). \\markcite{fab05}{Faber} {et~al.} (2005) provides a useful summary of prior studies, and describes their varying conclusions. The prior literature provides several plausible scenarios for $z<1$ red galaxy evolution, including passive evolution, assembly via dry mergers, formation from fading blue galaxies, or a combination of the above.  This is the first in a series of papers which will discuss the assembly and evolution of $z<1$ red galaxies detected by the multiwavelength imaging surveys of the entire \\bootes~ field. Other papers in this series will measure the stellar mass function, spatial clustering, and AGN content of red galaxies.  In this paper we present the evolving $B$-band luminosity function of $0.2<z<1.0$ red galaxies, and discuss the resulting implications for red galaxy evolution and assembly. While the rest-frame $B$-band is not ideal for tracing stellar mass, it has the advantages of being within the well studied optical wavelength range and remaining within the observed $B_WRI$ passbands of the NOAO Deep Wide-Field Survey (NDWFS) for redshifts of $z<1$.  The structure of this paper is as follows. We provide an overview of the NDWFS and {\\it Spitzer} IRAC Shallow Survey in \\S\\ref{sec:surveys}. We discuss our catalogs, photometry, photometric redshifts, and galaxy rest-frame properties in \\S\\ref{sec:catalogs}. In \\S\\ref{sec:sample} we describe the selection of our red galaxy sample. In \\S\\ref{sec:lf} we present rest-frame $B$-band luminosity functions of red galaxies and compare our measurements with simple galaxy evolution models. In \\S\\ref{sec:litcomp} we compare our results with recent red galaxy surveys, and discuss discrepancies between various studies. Our principal results and conclusions are summarized in \\S\\ref{sec:sum}. Throughout this paper we use Vega based magnitudes and adopt a cosmology of $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$, $H_0=70~{\\rm km}~{\\rm s}^{-1}~{\\rm Mpc}^{-1}$, $h=0.7$, and $\\sigma_8=0.8$. ",
        "conclusions": "\\label{sec:sum}  We have measured the evolution and assembly of red galaxies over the last $8~{\\rm Gyr}$, using the evolving  $B$-band luminosity function of $0.2<z<1.0$ red galaxies. Our sample of 39599 $0.2<z<1.0$ red galaxies, selected from $6.96~{\\rm deg}^2$ of imaging, is an order of magnitude larger in size and volume than comparable $z\\sim 0.9$ red galaxy samples. Accurate photometric redshifts were determined using the empirical ANNz photometric redshift code and photometry from both the NOAO Deep Wide-Field and {\\it Spitzer} IRAC shallow surveys. The accuracy of these redshifts has been verified with spectroscopy, and is comparable to the best broad-band photometric redshifts in the literature.  We find the color-magnitude and size-luminosity relations of red galaxies evolve with redshift, and this evolution can be approximated by a \\markcite{bru03}{Bruzual} \\& {Charlot} (2003) $\\tau=0.6~{\\rm Gyr}$ stellar population synthesis model. The size-luminosity relation predicts a significant fraction of the galaxy flux is at radii of several arcseconds, even at $z\\sim 0.9$. Luminosity functions which do not account for this underestimate the space density of red galaxies at high redshift, and overestimate the assembly rate of red galaxies at $z<1$.  We find the luminosity density, $j_B$, of red galaxies increases by $36\\pm13\\%$ from $z=0$ to $z=1$. In contrast, the $B$-band luminosity of a $\\tau=0.6~{\\rm Gyr}$ stellar population model increases by 213\\% over the same redshift range. If red galaxy stellar populations fade by $1.24$ magnitudes from $z=1$ to $z=0$, the stellar mass contained within red galaxies has approximately doubled over the past $8~{\\rm Gyr}$. Blue galaxies are being transformed into $\\lesssim L^*$ red galaxies at $z<1$, after a steady decline or rapid truncation of their star formation.  The evolution of the most luminous red galaxies at $z<1$ cannot be dominated by fading blue galaxies, as blue galaxies with large stellar masses are exceptionally rare at $z<1$. We have measured the evolution of $4L^*$ red galaxies using $M_B(10^{-3.5})$, the absolute magnitude corresponding to a fixed space density of $10^{-3.5}~h^3~{\\rm Mpc^{-3}}~{\\rm mag^{-1}}$. Our measurements of $M_B(10^{-3.5})$, along with those we derived from the literature, show that the stellar masses of the most luminous red galaxies evolve slowly at $z<1$. A \\markcite{bru03}{Bruzual} \\& {Charlot} (2003) $\\tau=0.6~{\\rm Gyr}$ model, with little stellar mass evolution and normalized to the 2dFGRS, only overestimates $M_B(10^{-3.5})$ by $\\simeq 0.21$ magnitudes at $z=0.7$. We therefore conclude that $\\simeq 80\\%$ of the stellar mass contained within today's $M_B-5{\\rm log}h=-21$ red galaxies was already in place within these galaxies at $z=0.7$. While red galaxy mergers have been reported in the prior literature, such mergers do not rapidly increase the stellar masses of $4L^*$ red galaxies between $z=1$ and the present day."
    },
    "0609/hep-ph0609144_arXiv.txt": {
        "abstract": "The model of Natural Inflation is examined in light of recent 3-year data from the Wilkinson Microwave Anisotropy Probe and shown to provide a good fit.  The inflaton potential is naturally flat due to shift symmetries, and in the simplest version takes the form $V(\\phi) = \\Lambda^4 [1 \\pm \\cos(N\\phi/f)]$.  The model agrees with WMAP3 measurements as long as $f > 0.7 \\mpl$ (where $\\mpl = 1.22 \\times 10^{19}$GeV) and $\\Lambda \\sim \\mgut$.  The running of the scalar spectral index is shown to be small -- an order of magnitude below the sensitivity of WMAP3. The location of the field in the potential when perturbations on observable scales are produced is examined; for $f>5\\mpl$, the relevant part of the  potential is indistinguishable from a quadratic, yet has the advantage that the required flatness is well-motivated.  Depending on the value of $f$, the model falls into the large field ($f \\ge 1.5 \\mpl$) or small field ($f<1.5 \\mpl$) classification scheme that has been applied to inflation models. Natural inflation provides a good fit to WMAP3 data. ",
        "introduction": "Introduction}  \\begin{figure} \\insertfig{rnplane} \\caption[$r-n_s$ plane] { Natural inflation predictions and WMAP3 constraints in the $r$-$\\ns$ plane.  (Solid/blue) lines running from approximately the lower left to upper right are predictions for constant $N$ and varying $f$, where $N$ is the number of e-foldings prior to the end of inflation at which current modes of scale $k = 0.002$ Mpc$^{-1}$ were generated and $f$ is the width of the potential. The remaining (dashed/red) lines are for constant $f$ and varying $N$. The (light blue) band corresponds to the values of $N$ for standard post-inflation cosmology with $(\\textrm{1 GeV})^4 < \\rhoRH < V_\\textrm{end}$.  Filled (nearly vertical) regions are the parameter spaces allowed by WMAP3 at 68\\% and 95\\% C.L.'s (error contours taken from Ref.~\\cite{Kinney:2006qm}). Natural inflation is consistent with the WMAP3 data for $f \\gae 0.7\\mpl$ and essentially all likely values of $N$. } \\label{fig:rnplane} \\end{figure}  Inflation was proposed \\cite{Guth:1980zm,Kazanas:1980tx, Starobinsky:1980te,Sato:1981ds,Sato:1980yn} to solve several cosmological puzzles: an early period of accelerated expansion explains the homogeneity, isotropy, and flatness of the universe, as well as the lack of relic monopoles.  While inflation results in an approximately homogeneous universe, inflation models also predict small inhomogeneities.  Observations of inhomogeneities via the cosmic microwave background (CMB) anisotropies and structure formation are now providing tests of inflation models.  The release of three years of data from the Wilkinson Microwave Anisotropy Probe (WMAP3) \\cite{Spergel:2006hy} satellite have generated a great deal of excitement.  First, generic predictions of inflation match the observations: the universe has a critical density ($\\Omega=1$), the density perturbation spectrum is nearly scale invariant, and superhorizon fluctuations are evident. Second, current data is beginning to differentiate between inflationary models and already rules some of them out \\cite{Spergel:2006hy,Alabidi:2006qa, Peiris:2006ug,Easther:2006tv,Seljak:2006bg,Kinney:2006qm,Martin:2006rs, Peiris:2006sj}. (For example, quartic potentials  and generic tree-level hybrid models do not provide a good match to the data.) It is the purpose of this paper to illustrate that the model known as Natural Inflation is an excellent  match to current data.  Inflation models predict two types of perturbations, scalar and tensor, which result in density and gravitational wave fluctuations, respectively.  Each is typically characterized by a fluctuation amplitude ($\\Pscalarrt$ for scalar and $\\Ptensorrt$ for tensor, with the latter usually given in terms of the ratio $r \\equiv \\Ptensor/\\Pscalar$) and a spectral index ($\\ns$ for scalar and $\\nt$ for tensor) describing the mild scale dependence of the fluctuation amplitude.  The amplitude $\\Pscalarrt$ is normalized by the height of the inflationary potential.  The inflationary consistency condition $r = -8 \\nt$ further reduces the number of free parameters to two, leaving experimental limits on $\\ns$ and $r$ as the primary means of distinguishing among inflation models.  Hence, predictions of models are presented as plots in the $r$-$\\ns$ plane.  Most inflation models suffer from a potential drawback: to match various observational constraints, notably CMB anisotropy measurements as well as the requirement of sufficient inflation, the height of the inflaton potential must be of a much smaller scale than that of the width, by many orders of magnitude (\\ie, the potential must be very flat).  This requirement of two very different mass scales is what is  known as the ``fine-tuning'' problem in inflation, since very precise couplings are required in the theory to prevent radiative corrections from bringing the two mass scales back to the same level. The natural inflation model (NI) uses shift symmetries to generate a flat potential, protected from radiative corrections, in a natural way \\cite{Freese:1990rb}.  In this regard, NI is one of the best motivated inflation models.  One of the major results of the paper is shown in \\reffig{rnplane}. The predictions of NI are plotted in the $r$-$\\ns$ plane for various parameters: the width $f$ of the potential and number of e-foldings $N$ before the end of inflation at which present day fluctuation modes of scale $k=0.002$ Mpc$^{-1}$ were produced.  $N$ depends upon the post-inflationary universe and is $\\sim$50-60.  Also shown in the figure are the observational constraints from WMAP's recent 3-year data, which provides some of the tightest constraints on inflationary models to date \\cite{Spergel:2006hy}.  The primary result is that NI, for $f \\gae 0.7\\mpl$, is consistent with current observational constraints.  In this paper we take $\\mpl = 1.22 \\times 10^{19}$ GeV.  Our result extends upon a previous analysis of NI \\cite{Freese:2004un} that was based upon WMAP's first year data \\cite{Spergel:2003cb}. Earlier analyses \\cite{Adams:1992bn,Moroi:2000jr} have placed observational constraints on this model using COBE data \\cite{Smoot:1992td}.  Other papers have more recently considered NI in light of the WMAP1 and WMAP3 data \\cite{Alabidi:2005qi,Alabidi:2006qa}.  This paper emphasizes two further results as well. First, we investigate the running of the spectral index in natural inflation, \\ie\\ the dependence of $\\ns$ on scale, and find that it is small: two orders of magnitude smaller than the sensitivity of WMAP3 and below the sensitivity of any planned experiment.  Second, we find how far down the potential the field is at the time structure is produced, and find that for $f > 5 \\mpl$ the relevant part of the potential is indistinguishable from a quadratic potential.  (Still, the naturalness motivation for NI renders it a superior model to a quadratic potential as the latter typically lacks an explanation for its flatness).  We will begin by discussing the model of natural inflation in \\refsec{NI}: the motivation, the potential, the evolution of the inflaton field, and relating pre- and post-inflation scales.  In \\refsec{Fluctuations}, we will examine the scalar and tensor perturbations predicted by NI and compare them with the WMAP 3-year data. In \\refsec{Running}, we will address the running of the spectral index.  In \\refsec{Potential}, we will examine the location on the potential at which the observable e-folds of inflation take place and examine where NI falls in the small field/large field/hybrid model categorization scheme.  We conclude in \\refsec{Conclusion}. ",
        "conclusions": "Conclusion}  Remarkable advances in cosmology have taken place in the past decade thanks to Cosmic Microwave Background experiments.  The release of the 3 year data set by the Wilkinson Microwave Anisotropy Probe is leading to exciting times for inflationary cosmology.  Not only are generic predictions of inflation confirmed (though there are still outstanding theoretical issues), but indeed individual inflation models are beginning to be tested.  Currently the natural inflation model, which is extremely well-motivated on theoretical grounds of naturalness, is a good fit to existing data.  In this paper, we showed that for potential width $f > 0.7 \\mpl$ and height $\\Lambda \\sim \\mgut$ the model is in good agreement with WMAP3 data.  Natural inflation predicts very little running, an order of magnitude lower than the sensitivity of WMAP. The location of the field in the potential while perturbations on observable scales are produced was shown to depend on the width $f$. Even for values $f>5 \\mpl$ where the relevant parts of the potential are indistinguishable from quadratic, natural inflation provides a framework free of fine-tuning for the required potential.  There has been some confusion in the literature as to whether natural inflation should be characterized as a `small-field' or `large-field' model.  In \\reffig{rnplane2} we demonstrated that either categorization is possible, depending on the value of $f$, and that both are in agreement with data.  Natural inflation makes definite predictions for tensor modes, as shown in \\reffig{rnplane}. Of particular significance is that current observational constraints place a \\textit{lower limit} on the tensor/scalar ratio for Natural Inflation of order $10^{-3}$, a value which is within range of proposed future high-precision cosmological measurements \\cite{Kinney:1998md,Friedman:2006zt}. Therefore Natural Inflation represents a model which is both well-motivated and testable."
    },
    "0609/astro-ph0609701_arXiv.txt": {
        "abstract": "We present the first tests of a new method, the Correlated Component Analysis (CCA) based on second-order statistics, to estimate the mixing matrix, a key ingredient to separate astrophysical foregrounds superimposed to the Cosmic Microwave Background (CMB). In the present application, the mixing matrix is parameterized in terms of the spectral indices of Galactic synchrotron and thermal dust emissions, while the free-free spectral index is prescribed by basic physics, and is thus assumed to be known. We consider simulated observations of the microwave sky with angular resolution and white stationary noise at the nominal levels for the {\\sc Planck} satellite, and realistic foreground emissions, with a position dependent synchrotron spectral index. We work with two sets of {\\sc Planck} frequency channels: the low frequency set, from 30 to 143 GHz, complemented with the Haslam 408 MHz map, and the high frequency set, from 217 to 545 GHz. The concentration of intense free-free emission on the Galactic plane introduces a steep dependence of the spectral index of the global Galactic emission with Galactic latitude, close to the Galactic equator. This feature makes difficult for the CCA to recover the synchrotron spectral index in this region, given the limited angular resolution of {\\sc Planck}, especially at low frequencies. A cut of a narrow strip around the Galactic equator ($|b|<3^\\circ$), however, allows us to overcome this problem. We show that, once this strip is removed, the CCA allows an effective foreground subtraction, with residual uncertainties inducing a minor contribution to errors on the recovered CMB power spectrum. ",
        "introduction": "The Cosmic Microwave Background (CMB) is by far the most powerful cosmological probe. The power spectra of its temperature and polarization anisotropies encode detailed information on the key cosmological parameters. Tremendous experimental efforts, and especially the currently flying \\emph{Wilkinson Microwave Anisotropy Probe} (WMAP, Bennett et al. 2003a, Spergel et al. 2006) and the forthcoming {\\sc Planck} satellite (Tauber 2004; Lamarre et al. 2003; Mandolesi, Morgante, \\& Villa 2003) will substantially advance the sensitivity and resolution of maps of the microwave sky. Correspondingly, efficient methods to extract all the available information from the measured signal need to be implemented, exploiting the advances of the data analysis science. One of the most challenging tasks is the separation of the astrophysical components superimposed on the CMB, usually denominated ``foregrounds''.  On angular scales larger than about $30'$ the dominant foregrounds in the relevant spectral region are diffuse emissions from our own Galaxy (De Zotti et al. 1999). Synchrotron (Haslam et al. 1982) and free-free (Haffner, Reynolds, \\& Tufte 1999, Finkbeiner 2003) emissions dominate below $\\simeq 60-80$ GHz (Bennett et al. 2003b), while at higher frequencies thermal dust (Schlegel, Finkbeiner, Davies 1998; Finkbeiner, Schlegel, Davies 1999) takes over. On smaller angular scales foreground fluctuations are dominated by several populations of extra-galactic sources, with different spectral behaviour: radio sources, dusty galaxies and the Sunyaev-Zel'dovich effect from galaxy clusters.  A big deal of work has been recently dedicated to develop algorithms performing component separation, based on different ideas and techniques from signal processing science. Several algorithms, referred to as ``non-blind'', assume a perfect knowledge of the frequency dependence of sources. The most widely used techniques exploiting this approach are Wiener Filtering (WF; Tegmark \\& Efstathiou 1996; Bouchet, Prunet, \\& Sethi 1999) and the Maximum Entropy Method (MEM; Hobson et al. 1998; Barreiro et al. 2004; Maisinger et al. 2004; Stolyarov et al. 2005). On the other hand, the emission spectra are generally poorly known. This fact has motivated ``blind'' approaches, which do not make any assumption on the spectral shape. Particularly promising are methods exploiting Independent Component Analysis (ICA) techniques (Amari \\& Chichocki 1998; Hyv\\\"arinen 1999), relying on the statistical independence of the different components. These methods proved to be very effective in extracting the CMB, which is independent of all other sources (Baccigalupi et al. 2000, 2004; Maino et al. 2002, 2003; Stivoli et al. 2006; Patanchon et al. 2005). However, Galactic emissions are tightly correlated to each other, so that the assumption of mutual independence breaks down. Still following the ICA approach, Belouchrani et al. (1997) dropped statistical independence for mere uncorrelation, and recovered the missing information by exploiting the physically plausible assumption of a significant spatial autocorrelation of the individual sources. Bedini et al. (2005) also exploited second-order statistics and spatial autocorrelation but, at variance with Belouchrani et al. (1997), they gave up the assumption of mutually uncorrelated sources. By exploiting a parametric knowledge of the mixing matrix, they succeeded in estimating both the mixing matrix and the relevant cross-covariances.  Here we present the first tests on this method, referred to as \\emph{correlated component analysis} (CCA), on a data set as realistic as possible. Also, we work either on the whole sky or on spherical sky patches, rather than on small plane patches as in Bedini et al. (2005).  The paper is organized as follows. In Section 2 and 3 we describe the basic aspects of CCA, and its implementation. In Section 4 we illustate the tests on simulated skies. In Section 5 and 6 we estimate the precision of the mixing matrix and CMB power spectrum recovery. In Section 7 we draw our conclusions. ",
        "conclusions": "The tests described here demonstrate that the Correlated Component Analysis (CCA) method (Bedini et al. 2005), applied to simulated data with {\\sc Planck} specifications, is a promising tool to estimate the mixing matrix parameterizing the frequency scaling of different astrophysical signals in cosmic microwave background (CMB) observations. For the low frequency {\\sc Planck} channels, from 30 to 143 GHz, and with the angular resolution of $33'$, the most accurate estimates are obtained in the latitude range $[-30^\\circ,+30^\\circ]$, where the foregrounds we aim to recover are relatively strong.  However, the strong concentration of intense free-free emission in a narrow strip (a few degrees wide) around the Galactic equator introduces a large gradient of the global spectral index of Galactic radio emission (synchrotron plus free-free) that prevents an accurate recovery of the synchrotron spectral index with the low resolution of the LF channels. The problem is cured cutting out the strip at $|b|< 3^\\circ$. Having done that, both the synchrotron and the dust spectral index are recovered with a mean absolute error of 0.05. At higher latitudes, the mean error on the synchrotron index increases to 0.08, while that on the dust index is unchanged.  The situation is even better with the high frequency {\\sc Planck} channels, from 217 to 545 GHz, thanks to the better spatial resolution ($5'$) and to the fact that the only relevant foreground, at least on intermediate angular scales, is Galactic dust.  We have shown that, exploiting non-optimized component separation techniques, such as Wiener filter and pseudo-inverse filter, such errors allow us to estimate the CMB power spectrum  with an uncertainty of the order of $1\\%$ on the angular scales constrained by our patch size.  We conclude that the CCA can be a promising independent way to probe the spectral behaviour of the main foregrounds affecting the CMB. This second order statistical approach may allow us to increase our knowledge on foregrounds and to perform component separation with traditional non-blind methods with a minimum number of priors."
    },
    "0609/astro-ph0609537_arXiv.txt": {
        "abstract": "We analyse those objects in the Brightest 55 sample of clusters of galaxies which have a short central cooling time and a central temperature drop.  Such clusters are likely to require some form of heating.  Where clear radio bubbles are observed in these clusters, their energy injection is compared to the X-ray cooling rate.  Of the 20 clusters requiring heating, at least 14 have clear bubbles, implying a duty cycle for the bubbling activity of at least 70 per cent.  The average distance out to which the bubbles can offset the X-ray cooling, $r_{\\rm heat}$, is given by $r_{\\rm heat}/r_{\\rm cool}=0.86\\pm0.11$ where $r_{\\rm cool}$ is defined as the radius as which the radiative cooling time is $3\\gyr$. 10 out of 16 clusters have $r_{\\rm heat}/r_{\\rm cool}\\gtrsim1$, but there is a large range in values.  The clusters which require heating but show {\\it no} clear bubbles were combined with those clusters which have a radio core to form a second sub-sample. Using $r_{\\rm heat}=0.86r_{\\rm cool}$ we calculate the size of an average bubble expected in these clusters.  In five cases (3C129.1, A2063, A2204, A3112 and A3391) the radio morphology is bi-lobed and its extent similar to the expected bubble sizes.  A comparison between the actual bubble size and the maximum expected if they were to offset the X-ray cooling exactly, $R_{\\rm max}$, shows a peak at $R_{\\rm bubble}\\sim 0.7 R_{\\rm max}$ with a tail extending to larger $R_{\\rm bubble}/R_{\\rm max}$.  The offset from the expected value of $R_{\\rm bubble}\\sim R_{\\rm max}$ may indicate the presence of a non-thermal component in the innermost ICM of most clusters, with a pressure comparable to the thermal pressure. ",
        "introduction": "Since the discovery of ``holes'' in the X-ray emission at the centre of the Perseus Cluster \\citep{Bohringer}, and following the launch of \\emph{Chandra}, many more depressions in the intra-cluster medium (ICM) of low redshift clusters have been found (e.g. Hydra A, \\citealp{McNamaraHydra00}; A2052, \\citealp{Blanton01}; A2199, \\citealp{JohnstoneA2199}; Centaurus, \\citealp{SandersCent02}).  Recent compilations are given in \\citet{DunnFabian05, DunnFabian04, Birzan04}.  Such holes have been observed to anti-correlate spectacularly with the radio emission from the active galactic nucleus (AGN) at the centres of these clusters.  Their morphology, particularly in the closest clusters, has led to the interpretation that these are bubbles of relativistic gas blown by the AGN into the thermal ICM. This relativistic gas is less dense than the ICM and so the bubbles are expected to detach from the core and rise up buoyantly through the cluster, e.g. Perseus \\citep{Churazov00,ACF_Halpha_PER03}.  The older, detached, bubbles tend not to have $\\ghz$ radio emission associated with them and have been termed ``Ghost'' bubbles.  The X-ray emission of the ICM naturally leads to the conclusion that the plasma should be cooling.  To maintain pressure support, the gas is expected to flow on to the central galaxy as a ``cooling flow.''  The subsequent increase in density would lead to a ``cooling-catastrophe'', with extremely rapid cooling in the cluster centre.  However, with the high spatial and spectral resolution of \\emph{Chandra} and \\emph{XMM-Newton} little of the expected X-ray emitting cool gas has been found (\\citealp{Peterson_2003}, see \\citealp{PetersonFabian2006} for a review).  Many mechanisms have been proposed by which the cool gas could be heated, including, for example, thermal conduction \\citep{KimNarayan03,VoigtFabian04}, but this appears not to work for clusters below $5 \\kev$.  A majority (71 per cent) of ``cooling core''  clusters harbour radio sources \\citep{Burns90}, and a similar fraction of clusters which require heating (likely to be a cooling core) harbour clear bubbles \\citep{DunnFabian05}.  The action of creating the bubbles at the centre of the cluster by the AGN is a favoured method of injecting energy into the central regions of the cluster and so prevent the ICM from cooling.  This process sets up sound/pressure waves in the ICM, the dissipation of which requires the ICM to be viscous \\citep{ACF_Halpha_PER03,ACF_Per_Mega_06}.  The viscous dissipation of the pressure waves allows the energy from the bubble creation to be dissipated far from the cluster centre, as gentle, continuous and distributed heating as required \\citep{VoigtFabian04}. \\citet{Birzan04} surveyed 80 clusters in the {\\it Chandra} archive without discriminating between cooling and non-cooling clusters and found 16 that contained bubbles.  This was interpreted as clusters having AGN/bubble activity only 20 per cent of the time.  In this work we take a sample of clusters which forms a subset of an almost complete sample and investigate the bubble population therein.  As the parent sample is almost complete we are able to investigate prevalence of cooling core clusters, duty cycles of bubble activity and radio sources within the general cluster population.  The sample selection is described in Section \\ref{selection}, and the data preparation and reduction in Section \\ref{data_prep}. The subset is itself then split into two: those clusters which are expected to require some form of heating and harbour clear bubbles; and a combination of those which require some form of heating but do not harbour bubbles and those clusters which have a central radio source. The analysis of those clusters with clear bubbles and of those without is discussed in Sections \\ref{bubble_cluster} and \\ref{nobubbles} respectively.  The overall implications of this work is discussed in Section \\ref{discussion} with future investigations outlined in Section \\ref{furtherwork}. ",
        "conclusions": "We have investigated a sub-sample of the Brightest 55 cluster sample. At least 36 per cent of clusters have a cool core, and of these at least 70 per cent harbour clear bubbles. First we analysed those clusters which have a short ($<3\\gyr$) central cooling time and a central temperature drop of a factor of two.  These clusters are expected to have the most rapid cooling in the centre and so require some form of heating.  At least 36 per cent of clusters have a cool core, and of these at least 70 per cent harbour clear bubbles, implying a duty cycle also of at least 70 per cent.  Up to 90 percent could have some form of bubbles.  The energy provided in the bubbles was compared to the X-ray cooling within the cooling radius (for $r_{\\rm cool}$ at $t_{\\rm cool}=3\\gyr$).   The mean distance out to which the bubble power can offset the X-ray cooling is $r_{\\rm heat}/r_{\\rm cool}\\geq0.86\\pm0.11$, with a large spread. In most clusters (10 out of 16) the AGNs energy input to the central regions of the cluster is sufficient to offset the X-ray cooling.  The bubble energy in the remainder does not offset most of the X-ray cooling.  In some cases this is the result of catching these bubbles at a very young age (e.g. A2199).  Also analysed were those clusters in the B55 sample which have a central radio source.  The fraction of clusters which have a central radio source is at least 50 per cent. These were combined with the clusters from the previous subset which did not have clear bubbles.  For clusters without clear bubbles we have used the mean value of $r_{\\rm heat}/r_{\\rm cool}$ to estimate bubble sizes from the X-ray cooling.  In five clusters (3C129.1, A2063, A2204, A3312 \\& A3391) the radio sources have bi-lobed morphologies, whose sizes are similar from the expected sizes.  Only in two of these five clusters is the X-ray observation of sufficient depth that any interaction between the radio and ICM could be expected to be seen. It is not clear why no clear bubbles are seen, especially in A2204.  As a result bubbles may be more common than previously thought.  We investigated the ratio of the actual bubble size to that of their maximum size if they were to offset all the X-ray cooling, and find the peak occurs at $R_{\\rm av}/R_{\\rm max}\\sim 0.7$ with a tail extending to higher $R_{\\rm av}/R_{\\rm max}$.  Either we are biased to imaging young bubbles or some assumption is incorrect.   A likely explanation is that a non-thermal component comparable to the thermal component is present in the innermost regions of most clusters."
    },
    "0609/astro-ph0609471_arXiv.txt": {
        "abstract": "As part of our ongoing effort to measure exoplanet sizes and transit times with greater accuracy, we present $I$ band observations of two transits of OGLE-TR-111b. The photometry has an accuracy of 0.15-0.20\\% and a cadence of 1--2~minutes. We derive a planetary radius of \\rp~and a stellar radius of \\rs. The uncertainties are dominated by errors in the photometry, rather than by systematic errors arising from uncertainties in the limb darkening function or the stellar mass. Both the stellar radius and the planetary radius are in agreement with theoretical expectations. The transit times are accurate to within 30 seconds, and allow us to refine the estimate of the mean orbital period: $4.0144479\\pm 0.0000041$~days. ",
        "introduction": "The aim of the Transit Light Curve (TLC) project is to gather accurate photometry during the transits of exoplanets across the disks of their parent stars. The immediate scientific harvest is the improved accuracy with which the basic system parameters are known. High-accuracy, high-cadence transit photometry allows for the determination of the stellar radius, the planetary radius, and the orbital inclination, for an assumed value of the stellar mass (see, e.g., Seager \\& Mall\\'{e}n-Ornelas~2003 for a discussion of the theory, and Brown et al.~2001 for the most famous example of a transit light curve). These parameters are interesting in their own right, and are important in the interpretation of other measurements such as reflected light (Rowe et al.~2006), thermal emission (Charbonneau et al.~2005, Deming et al.~2005), atmospheric absorption (Charbonneau et al.~2002, Vidal-Madjar et al.~2003), and the Rossiter-McLaughlin effect (Queloz et al.~2000, Winn et al.~2005, Wolf et al.~2006).  In the longer term, repeated observations of transits will reduce the uncertainties in the system parameters (and especially the orbital period) still further. More interestingly, it may be possible to detect additional transiting objects, satellites, rings, or even reflected light, by combining the data from many individual light curves. In addition, the existence of hitherto-undetected planets and satellites might be betrayed by periodic patterns in the measured times of mid-transit or variations in the orbital inclination (Miralda-Escud\\'{e}~2002, Holman \\& Murray 2005, Agol et al.~2005).  The TLC project is currently in its initial phase, in which we are observing all of the known transiting exoplanets, refining the estimates of each system's parameters and assessing the feasibility of continued long-term monitoring. We have previously reported on observations of the exoplanet XO-1b (Holman et al.~2006). In this paper, we present TLC results for OGLE-TR-111b.  The periodic 2\\% dimming events of the star OGLE-TR-111 were discovered by Udalski et al.~(2002) in a survey for transiting planets within a rich star field in Carina. Spectroscopic follow-up by Pont et al.~(2004) revealed a periodic Doppler shift, confirming that the dimming events were caused by the transits of a Jovian planet. Santos et al.~(2006) obtained optical spectra with a higher signal-to-noise ratio to study the properties of the host star, which is an early K dwarf with roughly solar metallicity.  The orbital period is just over 4~days, which is the longest period among the 5 exoplanets identified in the OGLE survey, but which is typical of the periods of the ``hot Jupiters'' that have been discovered in abundance in radial-velocity surveys.  For this reason, Pont et al.~(2004) referred to OGLE-TR-111 as the ``missing link'' between the OGLE survey and the radial-velocity surveys.  Initially, it was thought that the two surveys were yielding discrepant results because of the shorter periods of the OGLE objects, but this is now understood as a selection effect (Pont et al.~2005, Gaudi et al.~2005, Gould et al.~2006).  This paper is organized as follows. We describe the observations in the next section, and the photometric procedure in \\S~3. In \\S~4, we describe the techniques we used to estimate the physical and orbital parameters. In \\S~5 we give the results, and the final section is a brief summary. ",
        "conclusions": "Through observations of two closely spaced transits of the exoplanet OGLE-TR-111b, we have improved upon the estimates of the system parameters. The improvement comes not only from an overall increase in the signal-to-noise ratio, but also from the ability to resolve the ingress and egress and thereby determine the stellar radius photometrically. Our results confirm the previous estimates of the stellar and planetary radii, and are subject to a smaller systematic error. We have also provided a more precise transit ephemeris. All of these results will help with future observations and interpretations of this system."
    },
    "0609/astro-ph0609192_arXiv.txt": {
        "abstract": "Doubly-degenerate binary systems consisting of two white dwarfs both composed of carbon and oxygen and close enough that mass is transferred from the less massive to the more massive are possible progenitors of type~Ia supernovae.  If the mass transfer rate is slow enough that the accreting white dwarf can reach a mass of $1.38\\,\\rm M_{\\odot}$ then it can ignite carbon degenerately at its centre.  This can lead to a thermonuclear runaway and thence a supernova explosion. However if the accretion rate is too high the outer layers of the white dwarf heat up too much and carbon ignites there non-degenerately.  A series of mild carbon flashes can then propagate inwards and convert the carbon to neon relatively gently.  There is no thermonuclear runaway and no supernova.  We examine the critical accretion rate at which ignition switches from the centre to the surface for a variety of white dwarfs and find it to be about two fifths of the Eddington rate.  In a real binary star the mass transfer rate falls off as mass transfer proceeds and the system widens.  Even if the initial transfer rate is high enough for carbon to ignite at the outer edge, if such rapid accretion were to persist, we find that it can extinguish if the rate drops sufficiently quickly.  The interior of the white dwarf remains carbon rich and, if sufficient mass can still be transferred from the companion, it can eventually ignite degenerately at the centre.  The primary white dwarf must be about $1.1\\,M_\\odot$ or above and the companion about $0.3\\,M_\\odot$. Though white dwarfs of such low mass are expected to be pure helium we note that a star of initial mass $2.5\\,M_\\odot$ has a CO core of about $0.3\\,M_\\odot$ when it begins to ascend the asymptotic giant branch. Alternatively if the accretion rate can be limited to a maximum of $0.46$ of the Eddington rate then a $1.1\\,M_\\odot$ white dwarf accretes sufficiently slowly to explode from a companion white dwarf of any large enough mass. ",
        "introduction": "Type~Ia supernovae (SNe~Ia) are the brightest objects in normal galaxies. They appear to be standard candles and hence are useful cosmological measures of the Universe \\citep{Baa}. They are also a major source of iron \\citep{tout2001} which tends to be trapped in the neutron star remnants of other supernovae.  However, their progenitors remain uncertain (see for example \\citet{tout2005}) and this leaves their standard nature questionable particularly in the light of the chemical evolution of the Universe.  Type~I supernovae have no hydrogen lines and type~Ia are further distinguished by their prominent silicon lines.  They are almost certainly exploding carbon and oxygen (CO) white dwarfs.  Their available nuclear energy exceeds their binding energy so that the whole star can be destroyed in a thermonuclear runaway.  Most of the material approaches nuclear statistical equilibrium and about $0.6\\,\\rm M_{\\odot}$ of $^{56}\\rm Ni$ is expelled \\citep{ropke2005}. The radioactive decay of $^{56}$Ni to $^{56}$Fe via $^{56}$Co powers the supernova and the decay signature has been clearly observed \\citep{branch1998}.  It is fairly certain that SNe~Ia are detonated by mass accretion on to the CO white dwarf.  As the mass of a cold white dwarf increases towards the Chandrasekhar mass, the gravitational collapse heats the material and fusion ignites in the degenerate core. Typically CO white dwarfs explode at about $1.38\\, \\rm M_{\\odot}$.  Accreting white dwarfs have been known for some time as the engines of cataclysmic variables, the source of novae and dwarf novae \\citep{warner1995}, and so were the first candidates to be considered. However if the accreting material is hydrogen-rich, accumulation of a layer of only $10^{-5}-10^{-3}\\,M_\\odot$ of cold material leads to degenerate ignition of hydrogen burning sufficiently violent to eject most, if not all of or more than, the accreted layer in the well known nova outbursts of cataclysmic variables.  The white dwarf mass does not significantly increase and ignition of its interior is usually avoided.  However if the accretion rate is high, $\\dot M > 10^{-7}\\,M_\\odot\\,\\rm yr^{-1}$, compressional heating of the surface layers raises the degeneracy and hydrogen can burn relatively gently as it is accreted, bypassing novae explosions \\citep{paczynski1978}, allowing the white dwarf mass to grow.  Though, if it is not much larger than this, $\\dot M > 3\\times 10^{-7}\\,M_\\odot\\,\\rm yr^{-1}$, hydrogen cannot burn fast enough so that accreted material builds up a giant-like envelope around the core and burning shell which eventually leads to more drastic interaction with the companion and probably the end of the mass transfer episode. Rates in the narrow range for steady burning are found only when the companion is in the short-lived phase of thermal-timescale expansion as it evolves from the end of the main sequence to the base of the giant branch.  Super-soft X-ray sources \\citep{kahabka1997} are probably in such a state but cannot be expected to remain in it for very long (a SNe~Ia rate of only about one millionth of the observed rate form by this channel in the population synthesis of \\cite{hurley2002}) and white dwarf masses almost never increase sufficiently to explode as SNe~Ia.  If He-rich material is accreted instead about $0.1 - 0.15\\,M_\\odot$ of degenerate material can accumulate before ignition \\citep{nomoto1982} and an accretion rate above $3\\times 10^{-8}\\,M_\\odot\\,\\rm yr^{-1}$ can lead to steady burning \\citep{saio1987}. \\par The currently popular model overcomes these problems by postulating that, when the mass-transfer rate exceeds that allowed for steady burning, only just the right fraction of the mass transferred is actually accreted by the white dwarf.  A viable mechanism for this is a strong wind from the accretion disc that expels material from the system before it reaches the white dwarf \\citep{hachisu1996}.  Alternatively the white dwarf might indeed swell up to giant dimensions but the resulting common-envelope evolution could be sufficiently efficient that the small amount of excess material can be ejected without the cores spiralling in.  This is quite consistent with the findings of \\citet{nelemans2005} and \\citet{nelemans2000} that such efficiency necessary for at least one phase of common-envelope evolution in the formation of close double white dwarf systems.  Helium-accreting CO white dwarfs were once popular when it was thought that a thermonuclear runaway in the accreted helium layer could set off the CO core in an edge-lit detonation \\citep{woosley1994}.  However 2--D models indicate that central ignition is unlikely and light curves and spectra do not fit well so they are no longer considered viable progenitors \\citep{branch1995}.  In this work we consider the alternative that many complications can be avoided altogether if the white dwarf can accrete material of similar composition to itself, carbon and oxygen.  This can be achieved in a double degenerate model in which two CO white dwarfs merge \\citep{Webbink1984}.  Both form as the cores of asymptotic giant branch stars and, after one or two phases of mass transfer, one or both of which involves substantial orbital shrinkage through common-envelope evolution, they are brought close enough together that angular momentum losses owing to gravitational radiation lead to Roche lobe overflow.  A major difficulty with this model lies in the fact that high accretion rates also heat the white dwarf.  This heating is due mainly to the gravitational compression of the white-dwarf material as the mass increases.  The actual accretion luminosity liberated by the material falling down the potential well on to the white-dwarf surface has a negligible heating effect in comparison and is in any case mostly liberated in an accretion disc or boundary layer and radiated away.  A white dwarf can lose heat by conduction to the surface and radiation or by interior neutrino loss processes such as electron-positron pair production and annihilation. The material at the surface undergoes the most rapid compression as it becomes degenerate. At slow rates energy released in these surface layers can escape fast enough but at high rates the surface layers reach temperatures required for fusion to begin. True ignition occurs when energy production by carbon fusion exceeds that lost in neutrinos, a situation that is rapidly followed by the onset of convection but at velocities that are not sufficient to carry away igniting packets before thermonuclear runaway ensues. \\cite{nomoto1985} calculated that carbon ignites near the surface at constant accretion rates in excess of about one fifth of the Eddington rate (see section~\\ref{secedd}). Ignition at the surface raises the degeneracy so that burning of carbon to neon and sodium proceeds relatively gently.  Surface burning triggers the gentle ignition of a deeper shell and carbon is thus burnt successively throughout the white dwarf. Once the central carbon has been exhausted the core can't ignite until hot enough for neon burning by photodisintegration or oxygen fusion but by this point electron capture by magnesium has taken the collapse beyond the point at which the available nuclear energy can explode the white dwarf. It simply collapses quietly to a neutron star releasing energy in neutrinos. Thus for a SN~Ia a thermonuclear runaway must begin in the core before carbon ignites sufficiently at the surface.  Because white dwarfs expand as they lose mass, unless the Roche lobe of the mass donating companion expands even faster the very process of mass transfer causes the mass-losing (lower mass) white dwarf to overfill its Roche lobe yet more and mass transfer accelerates to dynamical rates.  Only when the mass ratio $q=M_{\\rm loser}/M_{\\rm accretor}<0.628$ can this positive feedback be avoided.  It is not known what the final outcome of dynamical mass transfer is but it is likely that much of the loser is lost to the interstellar medium and that the remainder is accreted at very high rates.  This already excludes the majority (about four-fifths in the population synthesis models of \\cite{hurley2002}) of merging CO white dwarf pairs which tend to have component masses between $0.5$ and~$1.1\\rm M_{\\odot}$. In the remaining cases, where dynamical mass transfer is avoided, the initial rate is still in excess of \\cite{nomoto1985}'s limit and so merging CO white dwarfs are not currently favoured SNe~Ia progenitors.  Various mechanisms might force merging white dwarfs to spin rapidly. \\cite{piersanti2003a,piersanti2003b} modelled the effects of rapid spin on carbon ignition in CO white dwarfs.  They found that compressional heating was not much affected but that the rate of diffusion of heat to the centre was slowed.  This generally reduces the maximum rate of accretion that allows central carbon ignition.  On the other hand spinning near breakup prevents accretion but also distorts the white dwarf so that it spins down by its own gravitational wave emission, while accreting just enough mass to keep it spinning close to break up.  This leads to a natural limit to the accretion rate of about $4\\times 10^{-7}\\,M_\\odot\\,\\rm yr^{-1}$, low enough to allow central carbon ignition.  In this work we only consider slowly rotating white dwarfs because we are interested in dynamically stable mass transfer.  For instance, when a $0.3\\,M_\\odot$ white dwarf begins mass transfer to a $1.1\\,M_\\odot$ white dwarf (as in section~\\ref{secsup}) the period of the binary is $146\\,$s, well above the break-up spin of the $1.1\\,M_\\odot$ white dwarf of $5.6\\,$s. It might however be argued that accretion from the inner edge of a Keplerian accretion disc would still spin up the massive white dwarf before it reached its ignition mass.  However this is avoided if the white dwarf has a magnetic field strong enough to disrupt the accretion disc as in the polar and intermediate-polar cataclysmic variables \\citep{warner1995}.  Indeed more massive white dwarfs are thought to form with stronger magnetic fields \\citep{wickramasinghe2005}.  We thus expect the accreting white dwarfs which interest us here to be relatively slowly rotating and so can be modelled without rotation.  Here we reconsider the limit on the accretion rate.  After describing the details of our evolution code and how we identify carbon ignition (section~\\ref{seccode}) we first work with constant rates (section~\\ref{seccon}) for direct comparison with \\cite{nomoto1985}. With modern computer power we can calculate on a much finer grid of initial masses and accretion rates.  Secondly, we consider accretion at a constant fraction of the Eddington rate (section~\\ref{secedd}) and find somewhat higher limits than \\cite{nomoto1985}.  Finally we consider true binary evolution in which the rate is initially high but falls off as the donor white dwarf loses mass and the system widens (section~\\ref{secbin}).  In section~\\ref{secsup} we explore the range of initial binary star parameters for which we find central carbon ignition and thence a supernovae and in section~\\ref{secconc} we present our conclusions and consider whether realistic progenitor systems might exist. ",
        "conclusions": "\\label{secconc}  Using detailed models of accreting CO white dwarfs we have critically investigated what rates allow central degenerate carbon ignition as opposed to ignition near the surface.  To compare with earlier work we first considered accretion at constant rates and then at fractions of the Eddington-limited rate.  Our results are qualitatively similar to those of \\cite{nomoto1985} but we find central ignition at rates up to twice what they did or two-fifths of $\\dot M_{\\rm EDD}$.  This can be attributed to their relatively course grid and updated stellar physics. We have also considered real double white dwarf binaries with mass transfer by Roche lobe overflow. Variation in the initial temperature and composition of the accreting white dwarf have only a small effect. We note that with fully conservative mass transfer a $1.1\\,M_\\odot$ CO white dwarf with a companion CO white dwarf of $0.3\\,M_\\odot$ can reach the conditions required for a SN~Ia even though the accretion rate initially rises to a maximum of almost $1.7\\times 10^{-5}\\, \\rm M_{\\odot}yr^{-1}$ and carbon burning in the non-degenerate layers reaches a luminosity of nearly $21\\, \\rm L_{\\odot}$.  The mass transfer rate and hence the luminosity fall off sufficiently quickly with the mass-transfer rate that the burning is extinguished before it runs away and burns the entire white dwarf.  Though standard single star evolution with reasonable mass loss only produces CO white dwarfs from about $0.55\\,M_{\\odot}$ to $1.2\\,M_{\\odot}$ we note that a star of initial mass $2.5\\,M_{\\odot}$ has a CO core of $0.28\\,M_\\odot$ at the start of the asymptotic giant branch, immediately after convective core helium burning.  We also recall that one or two phases of common envelope evolution are required to bring the two white dwarfs close enough for gravitational radiation to operate and that it would be easy to strip such a star's hydrogen rich envelope to leave a naked helium star of about $0.55\\,M_\\odot$.  Such a star is very small and would not undergo further interaction until gravitational radiation sets in but it is also very luminous (about $200\\,L_\\odot$) and maintains this luminosity for a long time (about $5\\times 10^8\\,$yr owing to helium burning.  This could well drive off the helium rich material, in a similar way to Wolf-Rayet stars, and expose the required low-mass CO core.  \\cite{hamann1995} suggest a mass-loss rate from naked helium stars of as much as $10^{-11.95}(L/L_\\odot)^{1.5}$ which is considerably more than would be needed here.  Undoubtedly the parameter space leading to these conditions is small. In the population synthesis of \\cite{hurley2002} only one in five hundred of the merging CO white dwarfs would actually make it to central carbon ignition under these conditions and thence a supernova rate of only about a thousandth of that observed.  However much of the physics, including that of common envelope evolution and mass loss from naked helium stars is highly uncertain, as are the initial mass functions, mass ratio distributions and separation distributions of binary stars so any population synthesis model is still very uncertain.  We note also that interactions between stars in dense stellar environments can significantly increase the number of merging white dwarfs.  \\cite{shara2002} find a factor of fifteen increase in open clusters.  On the other hand if the mass accretion rate can be limited to $0.46$ of $\\dot M_{\\rm Edd}$ then a $1.1\\,M_\\odot$ white dwarf accreting from a companion of any mass above $0.28\\,M_\\odot$ can make it to a supernova.  This would indeed be the case if a disc wind similar to that postulated by \\citet{hachisu1996} operated in the double-degenerate case in a similar way to the single-degenerate model.  We conclude that double CO white dwarfs remain as viable progenitors of SNe~Ia as any others currently proposed but are unlikely to be responsible for the majority of SNe~Ia.  In particular, if the accretion luminosity were somehow limited to less than 46~per cent of the Eddington rate in the early stages, then stable mass transfer could account for an important fraction of observed Sne~Ia."
    },
    "0609/astro-ph0609186_arXiv.txt": {
        "abstract": "{} {The competition between the formation and destruction of coronene clusters under interstellar conditions is investigated theoretically.} {The unimolecular nucleation of  neutral clusters is simulated with an atomic model combining an explicit classical force field and a quantum tight-binding approach. Evaporation rates are calculated in the framework of the phase space theory and are inserted in an infrared emission model and compared with the growth rate constants.} {It is found that, in interstellar conditions, most collisions lead to cluster growth. The time evolution of small clusters (containing up to 312 carbon atoms) was specifically investigated under the physical conditions of the northern photodissociation region of NGC~7023. These clusters are found to be thermally photoevaporated much faster than they are reformed, thus providing an interpretation for the lowest limit of the interstellar cluster size distribution inferred from observations. The effects of ionizing the clusters and density heterogeneities are also considered. Based on our results, the possibility that PAH clusters could be formed in PDRs is critically discussed.} ",
        "introduction": "Many astronomical objects show a distinct set of emission bands in the mid-infrared range, known collectively as the  unidentified infrared emission features. The most intense of these features, the Aromatic Infrared Bands (AIBs) falling at 3.3, 6.2, ``7.7'', 8.6, 11.3 and 12.7~$\\mu$m, are the signatures of aromatic CC and CH bonds. They are observed systematically from different regions of the interstellar medium (ISM) irradiated by UV photons. The carriers of these bands have been identified as Polycyclic Aromatic Hydrocarbons (PAHs) some twenty years ago by \\cite{leger84} and \\cite{allamandola85}.  \\cite{Boulanger90} and \\cite{Bernard93} analysed the emission in the IRAS photometric bands measured for several molecular clouds, and suggested that free-flying PAHs are produced by photoevaporation of larger grains. The observation of a mid-infrared continuum in the reflection nebula Ced 201 has been attributed by \\cite{Cesarsky00} to Very Small carbonaceous Grains (VSGs) producing the AIB carriers. More recently, in their study of the two reflection nebulae NGC 7023 and $\\rho$ Ophiucus-SR3, \\cite{rapacioli05a} showed that PAH molecules are produced by the destruction of small carbonaceous grains inside molecular clouds, these grains being interpreted as PAH clusters. A minimal size of 400 carbon atoms per cluster was inferred from the analysis of astronomical observations \\citep{rapacioli05a}. A plausible scenario is that these grains undergo photo-evaporation at the edge of clouds, thus leading to isolated PAHs. These observations motivate theoretical and experimental studies on the physical properties of PAH clusters, such as structure, stability or aggregation.  Small clusters of PAH molecules have been studied by several authors in the past, both experimentally \\citep{benharash,song,piuzzi,miller84} and theoretically \\citep{vandewaal1,gonzalez99,marzec,piuzzi,grimme}. Most of these studies have focused on rather small species,  usually containing no more than a few tens of carbon atoms, far below the limit of 400 carbon atoms per VSG as suggested by \\cite{rapacioli05a}. Studies on larger systems have been initiated only lately, motivated by the astrophysical context. Clusters of coronene molecules containing up to 20 units have been produced in a gas aggregation source \\citep{brechignac05, schmidt06}. Interestingly, the authors showed that the dissociation of coronene clusters induced by multi-photon excitation of a near-UV laser proceeds via the ejection of van der Waals-bonded coronene units.   In a previous work \\citep{rapacioli05b}, we investigated the stable structures of PAH clusters containing up to 800 carbon atoms (32 coronene molecules). The considered PAHs ranged from pyrene (C$_{16}$H$_{10}$) to circumcoronene (C$_{54}$H$_{18}$). The most stable cluster structures generally result from stacking the PAH molecules, first in a one-dimensional pattern at small cluster sizes. Above some critical size (depending on the PAH itself), structures more stable than the single stack are found, consisting of shorter stacks lying next to each other, and growing as two-, eventually three-dimensional structures. In large PAH clusters, the multiple ways of arranging the short stacks lead to a competition between stable structures.  Our present interest lies in the formation and destruction of PAH clusters in the ISM, and particularly the competition between these two processes. The coronene molecule (C$_{24}$H$_{12}$) provides a convenient example of large PAHs, especially from the computational point of view, and is used  in the present work as a prototype for interstellar PAHs. We wish to address here the following questions: \\begin{itemize} \\item[(i)]  What are the astrophysical conditions that favour the nucleation of PAH clusters? \\item[(ii)] How much energy is required for thermal evaporation of these clusters? \\item[(iii)] How do the rates of formation and destruction compare for these clusters under interstellar conditions? \\end{itemize}  In this paper, we focus on small neutral coronene clusters containing up to 312 carbon atoms, or 13 molecules. In particular, we attempt to offer a physical explanation for the existence of the observed lower limit to the interstellar PAH cluster size distribution of 400 carbon atoms per cluster \\citep{rapacioli05a}.   Addressing these questions requires a variety of theoretical methods to be used. Nucleation processes are explicitly simulated with an extended version of the tight-binding (TB) model initially developed by \\cite{parneix} for single PAHs. Concerning the thermal stability of coronene clusters, the rates of unimolecular evaporation as well as the photoemission rates in the infrared are obtained and compared. The evaporation rates are obtained in the statistical framework of the phase space theory \\citep{bowers77}. The competition between evaporation and IR emission is calculated using a master equation modeling based on microcanonical statistics \\citep{joblin02}. In the light of these calculations, we discuss several chemical scenarios to interpret the observations of the reflection nebulae \\citep{rapacioli05a}.  The paper is organized as follows. In the next section, we describe the atomic model and discuss the numerical simulation of coronene clusters growth by unimolecular nucleation. In section \\ref{sec:relax}, the evaporation rates of coronene clusters are calculated. In section \\ref{sec:compet}, we study the competition between evaporation and IR emission. In section \\ref{sec:photophys}, the photophysics of a coronene cluster in a radiation field is simulated. In section \\ref{sec:balance}, we study the competition between formation and destruction of coronene clusters under interstellar conditions. We also highlight the astrophysical implications of our results for possible models of chemical evolution in the photodissociation regions. In particular, the possible roles of ionization and density heterogeneities in the PDR in the existence and stability of PAH clusters are discussed. Some conluding remarks are given in Sec.~\\ref{sec:con}. ",
        "conclusions": "\\label{sec:con}  The formation and physico-chemical evolution of PAHs is a question of great interest in astrochemistry. \\cite{rapacioli05a} presented strong evidence that  PAH clusters are good candidates for small carbonaceous grains. These authors suggested that those grains are indeed PAH clusters, and a minimum size of 400 carbon atoms per cluster was inferred in NGC~7023. In this paper, the competition between the formation and destruction processes of such PAH clusters in the neutral state was investigated theoretically under realistic astrophysical conditions. Coronene clusters have been considered here as representative prototypes of PAH clusters.  The growth of neutral PAH clusters by nucleation of individual molecules was simulated using a model combining an explicit atomic force field to describe the intermolecular interactions, along with a quantum tight-binding approach for the intramolecular interactions. It was found that, under interstellar conditions, most of the collisions lead to cluster growth. This was interpreted as due to the large number of intramolecular modes that contribute to absorbing the excess collision energy.  Evaporation was treated in the framework of phase space theory, supplemented with molecular dynamics simulations. The clusters (C$_{24}$H$_{12}$)$_{4}$ and (C$_{24}$H$_{12}$)$_{13}$ were investigated, and a good agreement between the simulation results and the predictions of the statistical theory was found for the distribution of kinetic energy released after evaporation. The calculated evaporation rates were then introduced in an infrared emission model based on microcanonical statistics. This model allowed the simulation of these clusters in a specified interstellar radiation field.  When applied to the reflection nebula NGC~7023, our model indicates that (C$_{24}$H$_{12}$)$_{4}$ and (C$_{24}$H$_{12}$)$_{13}$ are both photoevaporated much faster than they can be reformed. In the case of (C$_{24}$H$_{12}$)$_{13}$, this evaporation is possible if the UV absorption rate is faster than the IR emission rate, so that at least $\\sim$3~eV of internal energy are stored in the cluster before absorption of a VUV photon. These results provide a plausible explanation for the absence of small VSGs observed in this PDR. They also suggest that PAH clusters are formed inside the molecular cloud and transported in the PDR by dynamical effects.  The present investigation is a first step towards the treatment of the complex dynamics of nucleation, thermal and photo fragmentations of PAH clusters in the interstellar medium. Our work could be extended by limiting the approximations done in the calculations of aggregation rates (impact parameter and initial energy distribution) and of evaporation rates (mainly the angular momentum). A next step could be to study larger clusters, especially those beyond the 400 carbon atom limit. It would be important to determine whether such large clusters could survive inside a reflection nebula like NGC~7023. Unfortunately, the computational efforts for simulating the evaporation of large clusters by molecular dynamics, to calculate their vibrational densities of states with Monte Carlo simulations, or to apply the IR/evaporation model will all tend to become extremely costly, at least one order of magnitude higher than in the present study. Further approximations will then become unavoidable for a proper extension of the present work.  Two important factors, namely ionization and density heterogeneities, might enhance the stability of PAH clusters in PDRs. These factors have been qualitatively discussed here. Only singly-charged cationic clusters are expected to be more stable than neutrals. Conversely, multicharged species and anionic clusters should be less stable due to electrostatic repulsion and electronic photodetachment processes, respectively. For the formation problem, it would be very useful to build a complete model to describe singly charged cationic assemblies of PAH molecules.  To be chemically relevant, such a model would need to include two main effects, that are not important for neutral clusters. Firstly, the charge will lead to $1/r^4$ polarization contributions, or even many-body implicit interactions between induced dipoles at higher order. Secondly, the model should account for charge delocalization, the missing electron being able to hop between neighbouring molecules. In homogeneous molecular systems, the charge is not restricted to be carried by a single PAH, but may extend over several units, typically between two and four \\citep{piuzzi}. This phenomenon is known as charge resonance. Obviously, polarization and charge delocalization effects are not unrelated, but their many-body and quantum characters make them significantly harder to describe than the pairwise repulsion-dispersion-electrostatic interactions in the presently investigated neutral clusters.  Thus, for cationic PAH clusters a more complex atomic model is necessary to take into account electron tranfer. A model along those lines was developed by \\cite{piuzzi}  and \\cite{Bouvier02} in order to investigate the structure of small clusters of aromatic molecules. The extension of this approach to large clusters and/or large PAH molecules will be computationally demanding, considering the extensive sampling required in the study of unimolecular evaporation. In the context of nucleation, the coupling between intermolecular and intramolecular modes will further burden the application of this model. With respect to neutral clusters, the theoretical modeling of cationic PAH clusters is therefore expected to be more involved. Efforts in this direction will be invaluable to improve our understanding of the physical and chemical evolution of photodissociation regions."
    },
    "0609/astro-ph0609465_arXiv.txt": {
        "abstract": "Binary population synthesis calculations and associated predictions, especially event rates, are known to depend on a significant number of input model parameters with different degrees of sensitivity. At the same time, for systems with relatively low formation rates, such simulations are heavily computationally demanding and therefore the needed explorations of the high-dimensional parameter space require major -- often prohibitive -- computational resources. In the present study, to better understand several key  event rates involving binary evolution and binaries with two compact objects in Milky Way-like galaxies and to provide ways of reducing the computational costs of complete parameter space explorations: (i)  we perform a methodical parameter study of the  \\emph{StarTrack} population synthesis code ; and (ii) we develop a formalism and methodology for the derivation of {\\em multi-dimensional fits} for event rates.   We significantly generalize our earlier study, and we  focus on ways of thoroughly assessing the accuracy of the fits.   We anticipate that the efficient tools developed here can be applied in lieu of large-scale population calculations and will facilitate the exploration of the dependence of rate predictions on a wide range binary evolution parameters. Such explorations can then allow  the derivation of constraints on these  parameters, given empirical rate constraints and accounting for fitting errors. Here we describe in detail the principles and practice behind constructing these fits, estimating their accuracy, and comparing them with observations in a manner that accounts for their errors. ",
        "introduction": "Models for binary stellar evolution and population syntheses are necessary to provide quantitative theoretical predictions for the relative likelihood of assorted events involving the evolution of binary stars.  The resulting predictions are particularly critical when no empirical estimates exist for topics of immediate astrophysical interest, such as mergers of  double-compact object (DCO) through the emission of gravitational waves.  The most practical and widely applied binary population synthesis codes available in the community --  such as the \\emph{StarTrack}  code described in \\citet{StarTrack} and significantly updated in \\citet{StarTrack2}; the BSE code, described in \\cite{2002MNRAS.329..897H}; the SeBa code described in \\cite{1996AA...309..179P}; and the StarFaster code described in \\cite{1998ApJ...496..333F}  --  rely on a large set of fairly simple parameterized rules to characterize many complex and often ill-understood physical processes.  Unfortunately, current binary population synthesis codes are greatly and often forbiddingly computationally demanding (depending  on their level of sophistication): even with substantial simplifications, exploring the entire parameter space relevant to the simulations is beyond present-day computational capability.  However,  observational information can provide us with constraints that help us improve our understanding of massive binary evolution. For example, pulsar searches continue to discover and refine observations of isolated pulsars and new binary pulsar systems; e.g., see \\citet{lrr-2005-7} for a review. Specifically, the samples of binary pulsars with neutron star and relatively massive white dwarf companions have been used for a statistical derivation of empirical rate estimates for their formation (most recently see \\citet{Chunglee-nsns-1}, \\citet{Chunglee-wdns-1},  \\citet{Chunglee-wdns-proceedings} and references therein). Additionally, many ground-based gravitational wave detectors now operating at or near design sensitivity (i.e., LIGO, GEO, TAMA) are designed to detect the late stages of double compact object (DCO) inspiral and merger.  Based only on early-stage data, these instruments have already provided conservative upper limits to certain DCO merger rates \\citep[see,e.g.][]{LIGOS2nsns,LIGOS2macho}.  With LIGO now very close to design sensitivity, a year of LIGO  data could definitively exclude (and even possibly confirm) the most  optimistic theoretical predictions for BH-BH merger rates \\citep[see, e.g.][for a range of BH-BH merger rates arising from binary evolution in Milky Way-like galaxies]{PSutility}.  Thus, gravitational-wave based upper limits (and, eventually, detections)  will shortly provide constraints on theoretical models of DCO formation.  Faced with the availability of empirical rate constraints, and yet at first unable to quantitatively impose them on population synthesis predictions because of the extremely limited exploration of the multi-dimenional parameter space, \\citet{PSutility} realized that any single unambiguous population synthesis  prediction could be sampled loosely and then fit over the most sensitive population synthesis parameters. In the same study, we also presented a technique to accelerate the synthesis code used (\\emph{StarTrack}) to study a  single target subpopulation, which we called ``partitioning'': we used  experience gained from prior simulations to reject binary parameters highly unlikely to produce the current event of interest. \\citet{PSconstraints} first applied these early fits to allow a direct comparison between  \\emph{StarTrack}-produced population synthesis predictions and the observed formation rate for  NS-NS binaries. Though only a  small fraction ($2\\%$) of \\emph{StarTrack} models appeared consistent with the constraints, conceptual challenges with seven-dimensional visualization prevented \\citet{PSutility} from clearly describing the constraint-satisfying region. A forthcoming paper, \\citet{PSmoreconstraints} will significantly extend that earlier preliminary analysis, adding significantly more observational constraints as well as a clearer investigation of the constraint-satisfying models.   In this paper, we extend and generalize the analysis of \\citet{PSutility}, and we present a  thorough discussion of our much-updated and vastly larger population synthesis archive (\\S\\ref{sec:ps:archive}) and particularly of the fitting methods we employ to extract predictions and assess their uncertainties and quality (\\S\\ref{sec:fit}). Because several imporant  and yet not immediately obvious pitfalls must be identified avoided  when constructing, testing, and applying fits, in this paper we provide a thorough and pedagogical discussion of our fitting method  and philosophy. For example, we explain how systematic fitting errors that were ignored in \\cite{PSconstraints} can limit our ability to constrain the  merger rates in the Milky Way. In a companion paper  \\citep{PSmoreconstraints} we explore the astrophysical applications of these fits, emphasizing their use in deriving empirical constraints on population syntheses and rate predictions.  Though we develop the fitting formulation specifically for the results of the \\emph{StarTrack} code, the methods described in this paper can be applied to any family of population synthesis simulations.  For this reason, we survey the physics underlying  the \\emph{StarTrack} model only in our companion paper \\cite{PSmoreconstraints}, in which these fits are used to  cover the parameter space and apply empirical rate constraints from supernovae and binary pulsars. ",
        "conclusions": "To  develop a more comprehensive understanding of population synthesis predictions  and   to allow those theoretical predictions to be systematically compared with observations of the end products of high-mass single and binary stellar evolution, we have fit eight predictions from the \\emph{StarTrack} code over seven of its input parameters. These fits are available on requests sent to the first author.  In a companion paper, \\citet{PSmoreconstraints} apply these fits along with estimates of their systematic errors to discover robust constraints on the seven parameters that enter into population synthesis.  Additionally, we have demonstrated that in analagous model problems the constraint-satisfying region defined by using these fits can, under appropriate conditions, very accurately trace the underlying constraint-satisfying region. Finally, we have presented a thorough diagnostic formalism, including a large list of diagnostic quantities and tests ($I,J,\\sigma_{E,DF,DD},\\chi^2,\\eta, r_+$, etc.), which can be applied to studies of fits to large archives of \\emph{any} (individual) population synthesis code."
    },
    "0609/astro-ph0609534.txt": {
        "abstract": "The sensitivity of a search for sources of TeV neutrinos can be improved by grouping potential sources together into generic classes in a procedure that is known as  source stacking. In this paper, we define catalogs of Active Galactic Nuclei (AGN) and use them to perform a source stacking analysis. The grouping of AGN into classes is done in two steps: first, AGN classes are defined, then, sources to be stacked are selected assuming that a potential neutrino flux is linearly correlated with the photon luminosity in a certain energy band (radio, IR, optical, keV, GeV, TeV). Lacking any secure detailed knowledge on neutrino production in AGN, this correlation is motivated by  hadronic AGN models, as briefly reviewed in this paper.  The source stacking search for neutrinos from generic AGN classes is illustrated using the data collected by the AMANDA-II high energy neutrino detector during the year 2000. No significant excess for any of the suggested groups was found. ",
        "introduction": "With AMANDA-II, different searches for high energy neutrino point sources have been performed~\\cite{tonioPS,paoloPS,4yrPS}, showing that neither the all sky search for hot spots in the neutrino sky nor the search at predefined and astrophysically motivated source positions has led to a discovery of a statistically significant neutrino signal. This however does not exclude the possibility that the signal of a superposition of generically equal sources, each contributing below the individual significance threshold, sums up to a significant signal for that specific source type. In case of no detection, the stacked signal may be used to obtain a limit on neutrino emission from that generic source type. Source stacking methods have been applied in gamma-ray astronomy, e.g., in searches for non-blazar AGN classes at GeV photon energies \\cite{cillis1,cillis2}.  AGN belong to the most promising potential neutrino sources. Since the detailed mechanism of neutrino production in AGN is still unknown, in this paper a neutrino-production-model free attempt to classify the sources according to their geometrical properties and electromagnetic emission is introduced. Based on this classification, we develop an AGN stacking analysis to search for TeV neutrinos. %, selecting candidate sources according to phenomenological %principles.   %The aim of source stacking methods is to obtain greater sensitivity to %emissions from classes of %sources by analyzing the cumulative signal of a sample of %several sources instead of single candidate sources. %Due to larger statistics, the %relative background fluctuations decrease. Hence, a smaller signal %contribution from each source is necessary to obtain a statistically %significant detection. If, in a source list ordered by luminosity, the %signal contribution of each source decreases slower than the relative %background fluctuations, source stacking generally provides a gain in %sensitivity. %Source stacking methods have been applied in gamma-ray astronomy, e.g., in %searches for non-blazar AGN classes at GeV photon energies %\\cite{cillis1,cillis2}. %In this paper, we develop an AGN stacking analysis to search for TeV %neutrinos, selecting candidate sources according to phenomenological %principles. %Since no TeV neutrino point sources have been detected up to now %\\cite{paoloPS}, the classification of TeV neutrino source candidates has to %rely on photon data.  In Sec.~\\ref{classi}, we describe in detail the adopted systematic classification of AGN based on an axisymmetric model \\cite{urry}. Possible neutrino production scenarios in AGN are reviewed in Sec.~\\ref{theory}. Pion production and subsequent decay results in correlated neutrino and photon production. Neutrino fluxes from these reactions are expected to be of the same order of magnitude as the gamma-ray fluxes. However, photons interact in the source environment possibly leading to a different photon spectrum than the neutrino spectrum. In Sec.~\\ref{samples}, we define catalogs of AGN classes based on our classification, and select from them interesting neutrino candidate sources. Hence, we order sources according to their photon flux at different energies. In Sec.~\\ref{opti}, a procedure is presented to determine the optimum number of sources from each catalog to be included in the stacking analysis. This procedure is applied to determine source samples to be analyzed with the AMANDA neutrino telescope, located at the South Pole~\\cite{AMANDA}. Thus selected sources in AMANDA's field of view are in the Northern sky. Assuming hypothetical distributions of the source strength as a function of the ordered source number we determine under which conditions a gain in sensitivity is obtained with the source stacking method with respect to other methods \\cite{tonioPS,kirsten_icrc}. In Sec.~\\ref{signal}, we illustrate our method by evaluating the signal from the resulting source samples using a data set collected by AMANDA in the year 2000 \\cite{tonioPS}. ",
        "conclusions": "We have performed a systematic classification of AGN. This is the basis of a source stacking analysis of TeV neutrinos with the AMANDA neutrino telescope. Neutrinos are assumed to be produced in optically thick jets or accretion disks. The optical depths of these sites and consequently the average energy of photons leaving the source are unknown. We therefore select the sources according to the photon flux at different energies. Most of the available photon data were not collected at the same time as the considered AMANDA data set. Thus, the time variability of AGN was not considered in this analysis.  With the hypothesis that the TeV neutrino flux is  proportional to the corresponding photon flux, we have optimized the number of sources to be included.  For most source classes,  an optimum of $8-12$ sources to be stacked has been found. For these, the source stacking analysis is more sensitive than the point source analysis of a single source. Exceptions have been found for FR-I and FR-II radio galaxies. While FR-I galaxies are dominated by the single source M87, there are too many FR-II galaxies with a similar flux in order to converge at a reasonably small number of sources to be stacked.  For the data of the year 2000, no significant deviation from the background expectation was found. In a next step, the stacking method will be applied to several years of AMANDA data. Optimizing the method to an angular resolution of less than $1^\\circ$ and better background rejection, it will be adapted to the $1$ km$^3$ IceCube array \\cite{IceCube}. Starting with the selection of southern sky sources, it is also applicable to other neutrino telescopes \\cite{BAIKAL,ANTARES,NESTOR,NEMO,KM3}.   \\newpage \\appendix"
    },
    "0609/astro-ph0609653_arXiv.txt": {
        "abstract": "Because of the similarity of the primary star of HS~1136+6646 to the planetary nebula central star BE Ursae Majoris, we did wide field imaging of the former with an H$\\alpha$ filter.  No nebulosity was detected.  On the other hand, the point spread function of the star appeared extended.  A partially-resolved red component is detected in the image with the five-band Sloan Digital Sky Survey.  Most importantly, a companion is easily resolved in the \\emph{HST} acquisition image for the published \\emph{STIS} observation.  A companion to the pre-cataclysmic binary is present at a separation of 1.349'' at position angle 54.4$^o$.  Evidence indicates that it is likely of K spectral type.  We cannot demonstrate conclusively that this component has common proper motion with the close binary.  However, the similar apparent $z$ magnitudes and spectral types of HS~1136+6646B and the resolved component make it likely that we have in reality a hierarchial triple system.  In any case, the presence of this component needs to be taken into account in future ground-based studies. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609379_arXiv.txt": {
        "abstract": "}[2]{{\\footnotesize\\begin{center}ABSTRACT\\end{center} \\vspace{1mm}\\par#1\\par \\noindent {~}{\\it #2}}}  \\newcommand{\\TabCap}[2]{\\begin{center}\\parbox[t]{#1}{\\begin{center} \\small {\\spaceskip 2pt plus 1pt minus 1pt T a b l e} \\refstepcounter{table}\\thetable \\\\[2mm] \\footnotesize #2 \\end{center}}\\end{center}}  \\newcommand{\\TableSep}[2]{\\begin{table}[p]\\vspace{#1} \\TabCap{#2}\\end{table}}  \\newcommand{\\FigCap}[1]{\\footnotesize\\par\\noindent Fig.\\  % \\refstepcounter{figure}\\thefigure. #1\\par}  \\newcommand{\\TableFont}{\\footnotesize} \\newcommand{\\TableFontIt}{\\ttit} \\newcommand{\\SetTableFont}[1]{\\renewcommand{\\TableFont}{#1}}  \\newcommand{\\MakeTable}[4]{\\begin{table}[htb]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newcommand{\\MakeTableSep}[4]{\\begin{table}[p]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newenvironment{references}% { \\footnotesize \\frenchspacing \\renewcommand{\\thesection}{} \\renewcommand{\\in}{{\\rm in }} \\renewcommand{\\AA}{Astron.\\ Astrophys.} \\newcommand{\\AAS}{Astron.~Astrophys.~Suppl.~Ser.} \\newcommand{\\ApJ}{Astrophys.\\ J.} \\newcommand{\\ApJS}{Astrophys.\\ J.~Suppl.~Ser.} \\newcommand{\\ApJL}{Astrophys.\\ J.~Letters} \\newcommand{\\AJ}{Astron.\\ J.} \\newcommand{\\IBVS}{IBVS} \\newcommand{\\PASP}{P.A.S.P.} \\newcommand{\\Acta}{Acta Astron.} \\newcommand{\\MNRAS}{MNRAS} \\renewcommand{\\and}{{\\rm and }} {We obtained $BV$ time series photometry for 12 variable stars from the field of the old open cluster NGC~2243. The sample includes 3 newly identified detached/semi-detached binaries. There are now four detached eclipsing binaries which are likely members of the cluster. Determination of the absolute parameters of the components would provide a valuable check on evolutionary models of low-mass stars. An accurate ephemeris and orbital period analysis are presented for the previously-known detached binary NV~CMa. We also provide ephemerides for 7 other periodic variables. We show that 3 contact binaries are likely members of the cluster. } {Stars:  binaries: eclipsing, binaries -- stars: individual: NV CMa  -- open clusters and associations: individual: NGC~2243} ",
        "introduction": "Detached eclipsing binaries (DEBs)  at the moment provide the most accurate determinations of the absolute stellar parameters of mass and radius (Andersen 1991, 1998). Although recent developments in high-precision astrometry,  combined with radial velocity measurements, hold great promise in this respect (eg. Muterspaugh et al. 2005), DEBs can be observed at much larger distances than visual binaries and, moreover, they allow a more direct determination of stellar radii and temperatures. Systems of early spectral type can be used as standard candles as far as in the Local Group Galaxies (Paczy\\'nski 1997; Clausen 2004; Ribas et al. 2005). DEBs located in open and globular clusters provide a particularly interesting check on stellar evolution theory (Paczy\\'nski 1997).  NGC~2243 is an intermediate age open cluster located in the direction of Galactic anticenter ($l=239^{\\circ}.5$, $b=-18^{\\circ}.0$) in a region of low reddening. An early photometric study of the cluster by Hawarden (1975) was followed by CCD-based studies by Bonifazi (1990) and Bergbusch et al. (1991), and more recently by Anthony-Twarog et al. (2005).  The last authors used $uvbyCaH{\\beta}$ photometry to derive reddening $E(B-V)=0.055\\pm 0.004$, metallicity ${\\rm [Fe/H]}=-0.57\\pm 0.03$, age of $3.8\\pm 0.2$~Gyr, and an apparent distance modulus $(m-M)_{V}=13.15\\pm 0.1$.  Spectroscopic studies conducted by Gratton \\& Contarini (1994; only two giants observed) and  Friel et al. (2002) gave ${\\rm [Fe/H]}=-0.48\\pm 0.15$ and ${\\rm [Fe/H]}=-0.49\\pm 0.05$, respectively.  The CCD-based color magnitude diagrams of NGC~2243 (see references above) exhibit a well marked sequence of `photometric binaries', formed by stars located about 0.7 mag above the main sequence defined by the `single' stars. This indicates that the cluster possesses a rich population of binary stars with mass ratio close to unity.  The field of NGC~2243 was surveyed for variable stars by Kaluzny et al. (1996; hereafter KKM).  The sample of 6 detected objects included 3 DEBs, 2 contact binaries and a background RR Lyr star. Recently, Anthony-Twarog et al. (2005) identified 4 new candidate variables in the cluster field.  The present study resulted from an effort aimed mainly at obtaining a good quality light curve for the eclipsing binary V1=NV~CMa detected by KKM. The system is located near the main-sequence turnoff on the cluster color-magnitude diagram. Its light curve exhibits two eclipses of similar depth with $\\delta V\\approx 0.6$~mag. The flatness of the light curve outside of eclipses together with an orbital period of $P=1.19$~days indicates that the binary has a detached configuration and hence deserves more detailed study.  As a first step toward this goal we present $BV$ light curves of V1=NV~CMa. We  also report some results for other objects, including the detection of new variables in the central field of the cluster. ",
        "conclusions": "We have expanded from 6 to 14 the sample of variable stars known in the field of the open cluster NGC~2243. Of primary interest are 5 detached or semi-detached binaries  which are likely members of the cluster. Detailed analysis of these objects can provide a direct measure of the distance to NGC~2243 based on the surface brightness method. Four out of five binaries are composed of stars located in different parts of the cluster main sequence --- assuming that all 4 binaries are members of NGC 2243. Determination of absolute parameters for up to 8 stars of the same age, metallicity and distance would provide a powerful test of evolutionary models of low mass stars.  Unfortunately, for the moment orbital periods are known with confidence only for one detached and one semi-detached binary: V1 and V7.  We have also obtained almost complete light curves  for these two systems. A dedicated extended photometric survey will  be necessary to determine periods and to obtain good coverage of eclipses for the remaining 3 detached eclipsing binaries.  In addition, spectroscopic data suitable for accurate determination of systemic velocities of all the binaries are needed to confidently resolve the question of their membership status.  (It should be noted that stars in old open clusters exhibit small dispersions of radial velocities. For example Mathieu et al (1990) quote a value of 0.5~km/s for M67 whose age is comparable to that of NGC 2243. Hence, measurement of radial velocities provides a powerful discrimination between field stars and cluster members.)  There are three certain and one likely W~UMa type variables in the central area of the cluster. The surveyed sample included about 1000 stars and is dominated by cluster members, as can be seen in Fig. 7. Hence, the relative frequency of occurrence of contact binaries in NGC~2243 is about 0.3-0.4\\%. This is comparable to the frequency of 0.2\\%  observed for binary dwarfs in the Solar neighborhood (Rucinski 2002).    \\Acknow{ JK, WK \\& GS were supported by grants 1~P03D~001~28 and 76/E-60/SPB/MSN/P-03/DWM 35/2005-2007 from the Ministry of Science and Information Society Technologies, Poland. IBT was supported by NSF grant AST-0507325. }"
    },
    "0609/astro-ph0609523_arXiv.txt": {
        "abstract": "I  review the  uncertainties  in two  observational local  constraints  of the  Galactic  disc  chemical  evolution :  the metallicity distribution of  long-lived dwarfs and the age-metallicity relation. Analysing most recent data, it is shown first that the observed metallicity distribution at solar galactocentric radius, designed with standard methods, is more fit to a closed-box model than to the infall metallicity distribution. We argue that this is due to the specific contribution of the thick disc population, which has been overlooked in both the derivation of the observed metallicity distribution and in standard chemical evolution models. Although this agreement disqualifies the metallicity distribution as the best supportive (indirect) evidence for infall, we argue that the evolution must be more complex than described by either the closed-box or standard infall models.    It is then shown that recent determinations of the age-metallicity distribution from large Str\\\"omgren photometric surveys  are dominated by  noise resulting  from systematic  biases in metallicities  and  effective  temperatures.   These biases are  evaluated and a new age-metallicity distribution  is obtained, where particularities of  the previous determinations  are phased out. The new age-metallicity relation shows a mean increase limited to about a factor of 2 in Z over the disc age. It is shown  that below 3 Gyrs, the dispersion in  metallicity is about  0.1 dex, which, given  the observational uncertainties  in  the  derived metallicities, is compatible with the small cosmic dispersion measured on the ISM and meteoritic  presolar dust grains.  A population that is progressively  older  and  more  metal-rich arises  at  a  metallicity greater  than that of the  Hyades, to  reach [Fe/H]$\\approx$+0.5  dex  at ages greater  than  5  Gyrs.  We suggest that this is best explained by radial migration.  A symmetrical  widening of  the metallicity interval towards lower values is seen at about the same age, which is attributed to a similar  cause. Finally, the  new  derived  ages  are  sufficiently  consistent  that  an age-metallicity relation  within the  thick disc is  confirmed.  These new features  altogether draw a  picture of the chemical  evolution in the solar neighbourhood where  dynamical effects and complexity in the age-metallicity distribution dominate,  rather than a generalised high dispersion at all ages. ",
        "introduction": "The metallicity distribution of long-lived dwarfs and the age-metallicity relation are two often-quoted constraints of the chemical evolution of the Galactic disc: the latter constraint is considered to be weak, whereas the former one is considered to be a strong, or even the 'strongest' constraint. Both conclusions are usually perceived as definitive\\footnote{In a recent article on modeling of the chemical evolution of the Milky Way, the authors (Romano et al. 2005.) argue that 'the uncertainties in the data have become really small' to focus on theoretical uncertainties. A similar conclusion is reached by Nordstr\\\"om et al. (2004) on the 'G-dwarf problem'.} observational results. On the first point,  this is illustrated by the vast majority of papers published in the last decade on the chemical evolution of the Galactic disc, which is described  similarly with prolonged infall controlling its progressive inside-out formation. The main  motivation for adopting infall being the capacity to reduce the number of metal poor stars born at early ages, providing a solution to the so-called 'G-dwarf' problem. How robust is the observational basis - the dwarf metallicity distribution - of such uniformity in models ? The unbiased metallicity distribution of dwarfs is usually obtained by weighting the metallicity bins of local samples using a relation between metallicity and  the vertical velocity dispersion ($\\sigma_W$) as measured on these same samples. Due to the small statistics in the lower metallicity bins and the uncertainties that plague the photometric metallicity scales, this process leads to large uncertainties in the level of corrections that are applied, as can be testified by the variety of [Fe/H]-$\\sigma_W$ relations found in the literature (see Fig. 3). The recent release of large catalogues of solar neighbourhood stars (Nordstr\\\"om et al. 2004, Valenti \\& Fischer 2005) containing radial velocities should significantly improve the situation. The level that this relation should reach at [Fe/H]$<$-0.4 dex is usually not questioned (with the exception of Sommer-Larsen (1991)), however it is shown here that the inclusion of the thick disc stars introduces significant modifications to the diagnostic of the 'G dwarf problem'. After having questioned the metallicity scales in Haywood (2001, 2002), we point here to the ambiguity caused by the thick disc population, which, if included in local samples, makes the observed distribution compliant with the predictions of the simple closed-box model. This suggests that some caution should be exercised as long as most of the  evidence for prolonged infall is limited to the so-called 'G dwarf problem'.   Since Edvardsson et al. (1993) the age-metallicity distribution (hereafter AMD) is often commented as being highly dispersed at all ages, providing a poor constraint to chemical evolution models of the Galaxy. Two recent studies, based on large photometric datasets, have claimed a similar result (Feltzing, Holmberg \\& Hurley 2001; Nordstr\\\"om et al. 2004). However, and though it has received little attention (but see Pont \\& Eyer 2004), these recent AMDs illustrate a rather  atypical chemical evolution, where the youngest  stars ($<$3 Gyrs) show the largest range of metallicity,  stars between 5 and 10 Gyrs show negative evolution (i.e. mean decreasing metallicity with age), and the lowest metallicity  stars ([Fe/H]$<$-0.5) are more often 'young' (age $<$ 5 Gyrs) than old. Although such non-conventional evolution is not excluded (one can think of an enhanced infall episode increasing the dilution of metals in the ISM), confirmation is needed. Moreover, the large dispersion at all ages is difficult to reconcile with the very small dispersion that is measured either in the local ISM (Cartledge et al. 2006), on meteoritic pre-solar grains (Nittler 2005), or abundance ratios, which all suggest that the gas from which stars are born is very homogenenous at all times in the disc, the thick disc, and probably the halo (Arnone et al. 2005)  In the following section, we present a discussion of recent metallicity distributions and scale height corrections, while the third section focuses on the biases in  the age-metallicity relation. Section 4 presents a derivation of a corrected AMD, which although not optimised, is quite different from the previously cited studies. We conclude in section 5. ",
        "conclusions": "We reviewed the metallicity distribution of stars in the solar neighbourhood published in the recent years  and find that  recent spectroscopic datasets show good agreement with our own finding (Haywood 2001),  with a peak at [Fe/H]=-0.05. Using scale height corrections based on a $\\sigma_W$-[Fe/H] relation derived from a sample of solar neighbourhood stars,  we arrive at the conclusion that there is no deficit of metal poor stars relative to the  closed box model. The low metallicity part ([Fe/H]$<$-0.5 dex) of our relation being dominated by thick disc stars, this conclusion is valid only if the thick disc stars are genuine disc stars, not of extragalactic origin.  It is also in agreement with expectations from simple arguments of Galactic structure. Standard chemical evolution models usually acknowledge the existence of a thick disc as a genuine Galactic population, but its contribution to the local solar radius metallicity distribution is often underevaluated. On the other hand, looking at the metallicity distributions recently published, we find the contribution of stars having the metallicity of the thick disc to be generally incompatible with simple Galactic structure constraints.  Despite the recent detailed analysis of large photometric Str\\\"omgren data sets, it appears that published AMD  are essentially reflecting noise in the determination of atmospheric parameters and derived ages. Significant improvement should come from reducing systematic biases in the determination of metallicity and effective temperatures from photometric indices. A first step in this direction is made here, and shows a limited increase in the mean metallicity of the thin disc, superposed on another, steeper relation in the thick disc. This scheme confirms that the disc has endured most of its chemical evolution during the thick disc phase, and has remained (chemically) unchanged since then.  In the last 3 Gyrs, the measured dispersion in metallicity is 0.10-0.11 dex, compatible with error measurements, and implying a small scatter of the young disc population. After 3 Gyrs, the progressive appearance of  a metal rich population reaching [Fe/H]$\\approx$+0.5 dex at ages greater than 5 Gyrs is established. It is suggested that these super metal-rich stars are the contaminant objects resulting from radial migration  from the inner disc. A similar spread at lower metallicity is seen at approximately the same age, giving stars reaching [Fe/H]$\\approx$-0.7 at 5 Gyrs, attributable to an equivalent migration from the outer disc. These characteristics are consistent with the (B-V, [Fe/H]) distribution presented in section 3, which, although independent of any age determination, is a source of information on the AMD. Note also that the metallicity scale of the  (B-V, [Fe/H]) distribution is independent of the Str\\\"omgren metallicity scale used in the derivation of ages. If this general picture is correct, we expect that the metallicity distribution at age greater than 3 Gyrs is a superposition of a narrow (about 0.05 dex or less) distribution of stars born at solar radius and a broader distribution of stars that haved moved to the solar radius. We note that the metallicity distribution of Fuhrmann (2004, figs. 49 and 50) presents such characteristics, with clearly a superposition of 2 gaussians with dispersions of about 0.05 and 0.25 dex centred on solar metallicity.   Altogether, these elements point to a picture  where chemical and dynamical evolution lead to structured, if complex,  patterns of age and metallicity, but not a general, structureless, dispersion. It is apparent that, in the level of abundance reached by the Galaxy prior to the thin disc formation as well as in the shape of the local metallicity distribution, the thick disc must have played a central role. This is now discussed.  \\subsection{The thick disc : Is there a continuity argument ? }   Is there a continuity argument to invoke when discussing the origin of the thick disc ? Inspection of  $\\alpha$ elements ratios $vs$ metallicity plots (see Pritzl, Venn \\& Irwin (2005), Reddy, Lambert, \\& Allende Prieto (2006) for example)  suggests the following : (1)  The halo presents no evidence of having been enriched by SNeIa (even when allowing for a generous thresold in galactocentric rotational velocity at about +100-120 km/s, see Fig. 2 in Pritzl et al, 2005). There have been suggestions that low [$\\alpha$/Fe] implied contamination by SNeIa. However, most recent studies (Arnone et al. (2005) and references therein) indicate that  halo stars seems to have a higher level of homogeneity in element ratios [X/Fe] than previously thought ($<$ 0.06 dex), with levels of [$\\alpha$/Fe] compatible with no contamination by SNIa. (2)  while there is a clearly visible increase of the [$\\alpha$/Fe] ratio towards low metallicity stars of the thin disc ([$\\alpha$/Fe]=0.1--0.2 dex at about [Fe/H]$\\approx$-0.5 -0.7 dex, see Reddy et al. (2006), Fig. 12), there is no evidence that the thin disc has escaped contamination by SNIa. (3) in the solar neighbourhood at least, the thick disc seems to be the {\\it only} population that shows both enrichment phases by SNeII and SNeII+SNeIa, suggesting that the thick disc could be a transition population between the halo and thin disc. The suggestion is made stronger if one realises that an extragalactic thick disc would require that an accreted satellite  have just the right ($\\alpha$/Fe, Fe/H) pattern to meet the halo on one side and the thin disc on the other. As a matter of fact, there is no detected chemical or kinematical discontinuity between the 'rotating halo' and thick disc (see Gratton et al. 2003) so far. If the small dispersion of abundance ratios measured on halo stars is confirmed, then the continuity between enrichment levels in the different elements between the halo and the thick disc should become a stringent constraint for the origin of this last population. For example, Nissen et al. (1994) measured [Mg/Fe]=0.41 with a rms dispersion of $<$0.06 dex on halo stars with -3.2$<$[Fe/H]$<$-1.8, while Reddy et al. (2006 and references therein) measured thick disc stars to have a similar rms dispersion and [Mg/Fe] comparable with that of Nissen et al. (1994) at [Fe/H]$<$-0.8dex. Confirmation on a larger scale is needed, but it suggests that the 'rotating halo' and the thick disc may be the same population.  There is however an observed discontinuity between the thin and thick discs, with the  ($\\alpha$/Fe, Fe/H) sequence of the two being almost parallel. It has been proposed that since the most metal-rich stars in the thick disc have a higher metallicity than the most metal-poor stars in the thin disc, there has been a dilution of metals by an infall episode in the temporal  gap between the two populations (Bensby et al. 2005; Reddy et al. 2006). The AMD of the previous section shows that metal-poor disc stars are not specifically old for their metallicity, a characteristic that is not expected if it is assumed that the first thin disc stars had formed after the infall episode at [Fe/H]$\\approx$-0.6 dex. Moreover, the AMD of Fig. 13 shows that the mean metallicity of old disc stars is nearer to [Fe/H]=-0.2 dex than -0.6 dex, a fact which does not fit well with the dilution scenario above. A perhaps more convenient explanation is that metal-poor thin disc stars found in local samples have been moved to the solar neighbourhood by radial migration from outside the solar  circle. This scenario, suggested by our conclusions that the AMD has been shaped by radial migration, is also well in agreement with the results presented by Carney et al. (2005) and Yong, Carney \\& de Almeida (2005). These authors present evidences, from open clusters, field giants and cepheids, that the mean disc metallicity decreases to (-0.4, -0.6) dex at R$_{GC}$=10-11 kpc. An even more interesting clue found by these authors is that the [$\\alpha$/Fe] ratios of these stars reach about +0.1 to +0.2 dex. This is just the characteristics of the most metal-poor thin disc stars in the solar neighbourhood sample, which also show this slight increase in  [$\\alpha$/Fe]. Finally, we note that Carney et al. (2005)  and Yong et al. (2005) find that the gradient in metallicity and [$\\alpha$/Fe] flattens beyond  R$_{GC}$=10-11 kpc, which may explain why the metal-poor thin disc stars found in local samples are limited to about  [Fe/H]$\\approx$-0.6 dex. Taking these considerations  into account, it might then be expected that the thin disc starts forming its stars (at solar galactocentric radius) not at -0.6 dex, but at a higher metallicity which Fig. 13 suggest might be around -0.2 dex, or above. This is also consistent with the ([$\\alpha$/Fe], [Fe/H]) plot in Reddy et al. (2006), which, showing excellently the separation between the thin and thick discs, also shows that the tip of the metal rich thick disc is around (-0.3,-0.2) dex.   This leaves the possibility that there is a gap in  [$\\alpha$/Fe] between the thin disc and the thick disc (Pagel 2001) of the order of +0.05 dex, consistent with a temporal gap between the two populations. A jump in  [$\\alpha$/Fe] with no significant variation of the general metal content ($<$0.1 dex) does not require exchange of gas (inflow or outflow) but simply that star formation ceased for a period of time sufficiently long that the ISM was enriched in iron by SNeIa for [$\\alpha$/Fe] to decrease by 0.05 dex.   \\subsection{Life without the 'G-dwarf problem'}    In view of what has just been said about the thick disc, how do we interpret the fit between the dwarf metallicity distribution and the closed-box model ? We cannot eliminate the option that the thick disc is an accreted population, since there is by now no clear-cut evidence against this solution. In this eventuality however, we note that the metallicity distribution would then have to be limited to 'pure' thin disc, solar radius stars. Limiting the sample to objects born within a restricted range around this radius would probably make the metallicity much thinner than the one derived in Haywood (2001) - probably less than 0.1 dex dispersion. Casuso \\& Beckman (2004) have explicitely considered the thick disc should not be part of the metallicity distribution,  although not formally considering the thick disc as a population of extragalactic origin. Their conclusion that infall should be an increasing function of time, is not surprising.    In the case that the thick disc is a genuine Galactic population, the disc (thin + thick) metallicity distribution is not in marked disagreement with the closed-box model. The problem is reminiscent of the discussion about the MDF (metallicity distribution function) of the halo. According to predictions of hierarchical clustering in $\\Lambda$CDM cosmologies, the stellar halo is supposed to be an archetypal open system, formed from many independent units. However, the observed MDF of the halo is shown to be fitted reasonably well by a simple box model distribution (Oey 2003) over a large range of metallicities. More generally, it is apparent that the closed-box model gives a good fit to the MDF of spheroids for most of the metallicity interval. There is often a problem in fitting the metal poor tail of these distributions, but this concerns a small fraction of the stars, and is not comparable with the local disc 'G-dwarf problem'. The question however remains : why does the closed box model provides an honest fit to these systems ? The question for the disc MDF is more or less the same. The exact profile of the low metallicity tail of the thick disc is unknown, but globally, the simple model provides a fair fit to the observed distribution. Because discs are now conceived as open systems that build up continuously from accreted gas, this result is not expected. We emphasized that strictly speaking, our result only means that the allegated mismatch between the observed MDF and the simple model cannot be used as an argument for infall models. It does not imply infall is not a crucial ingredient in models, although it strongly suggests that standard infall models should be revised. We notice that first results from chemo-dynamical models (Brook et al. 2005) are encouraging. In particular, Brook et al. (2005) have come quite close to reproducing the characteristics of the distinct ($\\alpha$-element) chemical signatures of the thin and thick discs, with the thick disc formed during a merging phase of  gas-rich  'building block', while the thin disc forms inside-out, at a more quiescent epoch, from continuously accreting gas. The interesting point is given by inspection of the metallicity distribution of the (thin + thick) disc generated by their model (Martel et al., 2005, their Fig. 3), which shows it is quite similar to a closed box model distribution. This suggests that although the local data are not conclusive indication for prolonged infall, as classically proposed, a metallicity distribution with the simple model characteristics is neither an indication of a closed system\\footnote{Ironically though, it may be noted that at least 2 hypotheses  of the simple  model are little disputed, and one fact is hard to prove wrong : that the IMF is constant with time, the interstellar medium is well mixed, and, with the present results, that the solar neighbourhood metallicity ditribution is like a closed box model distribution.}. Said differently,  the simple model may have lost its paradigmatic strength.    \\begin{figure} \\includegraphics[width=9cm]{fig12ab.eps} \\caption{(a) (B-V,[Fe/H]) for a subsample of the GCS catalogue, selected as described in the text. (b) (B-V,[Fe/H]) for these same stars with metallicity derived as described in the text. } \\label{revisedbvfeh} \\end{figure}   \\begin{figure} \\includegraphics[width=9cm]{fig13abcd.eps} \\caption{AMD for various age and metallicity determinations. The horizontal line at [Fe/H]=0.2 dex is given to ease the comparison. The thick curve on the first 3 plots is the dispersion in metallicity calculated over 300 points. The solid stars in all plots are thick disc stars according to their kinematics (see text). {\\bf (a)} The AMD  according to the subsample of the GCS catalogue, for stars selected as detailed in the text. Ages and metallicities are all from the CGS catalogue. {\\bf(b)} The AMD for the same stars, with the atmospheric parameters from the GCS catalogue, but with ages derived from the Yi et al. (2003) isochrones. The $\\alpha$-element content was set to 0 for all stars. {\\bf(c)} AMD for the same sample, but derived from the new [Fe/H] (see text),  effective temperatures from the GCS corrected as described in the text, and using the isochrones of Yi et al. (2003). The $\\alpha$-element content was set to 0 for all stars. The open circles  are the new ages derived for the Hyades stars. The thin lower curve is the mean metallicity calculated over 100 points. {\\bf(d)} Same as (c), except that the ages of thick disc stars (solid stars) have been derived taking into account the $\\alpha$-element content according to the value listed in Soubiran \\& Girard (2005). Open stars are objects with [Fe/H]$<$-0.4 from Soubiran \\& Girard (2005) flagged as disc stars according to their kinematics and $\\alpha$-element content. } \\label{revisedamr} \\end{figure}"
    },
    "0609/astro-ph0609009_arXiv.txt": {
        "abstract": "Many quasi-simultaneous optical observations of 17 blazars were obtained from the previous published papers over last 19 years in order to investigate the spectral slope variability and understand the radiation mechanism of blazars. The long period dereddened optical spectral slopes were calculated in this paper. Our analysis upon the average spectral slope distribution suggests that the spectra of Flat Spectrum Radio Quasars (FSRQs) and High energy peaked BL Lac objects (HBLs) are probably deformed by other emission components. The average spectral slopes of Low energy peaked BL Lac objects(LBLs), which scatter around 1.5, show a good accordance with Synchrotron Self-Compton (SSC) loss-dominated model. We present and discuss the variability between spectral slope and optical luminosity. The spectra  of all HBLs and LBLs get flatter when they turn brighter, while for FSRQs this trend does not exist or may be in a reverse situation. This phenomenon may imply that there is a thermal contribution to the optical spectrum for FSRQs. For the FSRQ 1156+295, there is a hint that the slope gets flatter at both the brightest and faintest states. Our result shows that three subclasses locate in different regions in the pattern of slope variability indicator versus the average spectral slope. Furthermore we also delineate that relativistic jet mechanism is supported by the significant correlation between the optical Doppler factor and the average spectral slope. ",
        "introduction": "Blazars are a subset of Active Galactic Nuclei (AGNs), which are compact, flat spectrum radio sources with highly variable and polarized nonthermal continuum emission ranging from radio up to X-ray and often to $\\gamma$-ray frequency (e.g. \\citealt{fugmann,bregman,sillanpaa,maraschi,urry}). Their observed properties  are mainly considered to originate from a relativistic beaming jet \\citep{rees}, which is possibly powered and accelerated by a rotating and accreting supermassive black hole. There are two classes of blazars, BL Lac objects (BL Lacs) and FSRQs, the former one has a featureless optical continuum, while the latter has a great many strong and broad emission lines. According to the peak frequency of synchrotron emission relative to the peak flux, BL Lacs are classified as LBLs, whose peak is located in infrared or optical band, and HBLs, whose peak is located in the ultraviolet or X-ray frequency \\citep{padovani,giommi1999}. For HBLs, the broadband spectral index $\\alpha_{RX}$ is less than or equal to 0.75 \\citep{urry} or 0.80 \\citep{sambruna}, while for LBLs, it is greater than 0.75 or 0.80.  The studies of the spectral energy distribution (SED) of blazars are very important for the understanding of the physical radiation mechanism and to constrain the parameters involved in physical model, and the studies are crucial in analyzing individual emission components. Synchrotron inverse Compton emission model predicts that the spectrum gets harder as the source turns brighter. Many investigators found that the variation amplitude of blazars at high frequencies was larger than that at low frequencies, namely, the spectrum became flatter when the flux increased, while the spectrum became steeper when the flux decreased (e.g. \\citealt{racine,gear,massano}). But \\citet{ghosh} suggested that it might not be always correct (e.g. \\citealt{ramirez}). \\citet{damicis} reported the optical spectral slope variability of 8 BL Lacs. They found that the spectral slope became flatter when the source turned brighter for all 8 objects. Four objects showed marked correlation between spectral slopes and \\emph{R} magnitudes. \\citet{trevese} analyzed the spectral slope variability of 42 PG quasars. They concluded that the spectral variability must be intrinsic of the nuclear component. Their numerical simulation showed that hot spots on the disk were able to represent the observed spectral variability, while single host galaxy contribution or single changes of accretion rate was insufficient to explain the slope variability. \\citet{vagnetti} analyzed the data from \\citet{damicis} under a simple synchrotron mode and the thermal bump (TB) mode. They showed that the spectral variability, even restricted to the optical band, could be used to set limits on the relative contribution of the synchrotron component and the thermal component. They deduced that the spectral variability of BL Lacs differed from the variability of quasars. The 8 BL Lacs they studied were all LBLs. \\citet{ramirez} presented that the spectrum of a FSRQ PKS 0736+017 became softer when it turned brighter. The correlation between spectral slope and flux was strong. Its variability relation between slope and flux was reverse to the result that was mentioned above for LBLs.  In this paper we show and discuss the optical spectral slope variability of 17 blazars, which include 9 LBLs, 4 HBLs and 4 FSRQs. In section 2 we describe the data and data reduction process. The average optical spectral slope distribution is analyzed in section 3, the spectral slope variability is given in section 4, and conclusions are presented in the last section. ",
        "conclusions": "  \\begin{enumerate} \\item All the average spectral slopes of our 9 LBLs are greater than unity and scatter around 1.5, this is a good support to the SSC-loss-dominated model \\citep{chiang}. \\item The optical spectra of all our HBLs and LBLs become flatter when the sources turn brighter, the average linear regression correlation coefficient of HBLs is higher than that of LBLs. The relation between spectral slope and \\emph{R} magnitude of FSRQs suggests that the emission of FSRQs probably has ``thermal bump\" contribution and the thermal component must be considered. \\item The slope distribution and variability for HBLs indicate that their spectra appear to be deformed by other components, which presumably come from the thermal accretion disk or originate from different regions of the jet. \\item Different subclasses of blazars have the tendency to locate in different regions in the pattern of the slope variability indicator versus average spectral slope. \\item There is a significant correlation (correlation coefficient is 0.79) between the average optical spectral slope and optical Doppler factor, which suggests that Doppler beaming effect of relativistic jet is the main mechanism for blazars. \\end{enumerate}"
    },
    "0609/hep-ph0609199_arXiv.txt": {
        "abstract": "We study the effect of a uniform magnetic field on the dynamics of axions. In particular, we show that the Peccei-Quinn symmetry is explicitly broken by the presence of an external magnetic field. This breaking is induced by the non-conservation of the magnetic helicity and generates an electromagnetic contribution to the axion mass. We compute the magnetic axion mass in one loop approximation, with no restriction on the intensity of the magnetic field, and including thermal effects. ",
        "introduction": "Axions are the pseudoscalar particles predicted by the Peccei and Quinn (PQ) mechanism~\\cite{PQ}, a quite elegant attempt to explain the smallness of the CP-violation induced by the QCD Lagrangian (the so-called strong CP problem; for a review see, e.g., Ref.~\\cite{Kim}). They emerge as Pseudo-Goldstone Boson modes associated with the global axial PQ-symmetry $U(1)_{\\rm PQ}$, spontaneously broken at the PQ-scale $f_a$. This last is also known as the PQ- or axion-constant, and characterizes all the axion phenomenology. More specifically the axion mass is given by the relation~\\cite{Kim} \\begin{equation} \\label{m0} \\mq /\\eV \\simeq 6.2 \\times 10^{6} \\GeV / f_a, \\end{equation} (here QCD means that this is induced by interaction with gluons) while its interactions with fermions are measured by $g_{i}\\sim m_i/f_a$, where $m_i$ represents the fermion mass (e.g.\\ $m_e,m_N,...$ for electrons, nucleons, etc.). In addition, though chargeless, axions interact with photons through the electromagnetic anomaly: \\begin{equation} \\label{agamma} \\mc L_{a\\gamma} = \\frac14 \\, g_{a\\gamma} \\, a \\, F \\widetilde F, \\end{equation} where $F$ is the electromagnetic field, $\\widetilde F^{\\mu\\nu} = (1/2) \\, \\varepsilon^{\\mu \\nu \\alpha \\beta} F_{\\alpha \\beta}$ its dual, and $g_{a\\gamma} = \\aem \\xi / (2\\pi f_a)$. Here, $\\aem$ is the electromagnetic fine structure constant and $\\xi$ an order one, model dependent constant. Presently, the PQ-constant's allowed range is rather narrow~\\cite{Raff}: terrestrial experiments together with astrophysical and cosmological considerations have in fact excluded all the values of $f_a$ up to $10^{9} \\, \\GeV$, \\footnote{In the case of the hadronic axion \\cite{KSVZ}, a small window around $f_a \\sim 10^6$GeV can be also permitted (see also Ref.~\\cite{hadronic1}).} and above $10^{12} \\, \\GeV$. \\footnote{Attempts to relax the above bounds on the PQ-constant are discussed in Ref.~\\cite{Maurizio}, whereas possibilities of different cosmological bounds are discussed in Ref.~\\cite{Dav86}.}  Because of the axion-photon interaction term (\\ref{agamma}), the phenomenology of axions can be influenced by the presence of an external electromagnetic field. Probably, the most discussed result of this well studied interaction is the axion-photon conversion in an external magnetic field (see, e.g., Ref.~\\cite{Raff88}), which presently seems the best way to detect axions by terrestrial experiments~\\cite{Sikivie}.  Another interesting consequence of interaction (\\ref{agamma}) was discussed in Ref.~\\cite{Mikheev} where it was shown that axions in a strong magnetic field undergo a finite mass shift. However, in that computation of the mass shift the temperature effects were not taken into account. This makes the result in Ref.~\\cite{Mikheev}, though correct, not directly applicable to the axion phenomenology. In fact, since the axion-photon coupling is very small, the induced mass can be relevant for the axion phenomenology only in the case of very intense magnetic fields, which normally exist only at high temperatures. In this paper, we reconsider this effect accounting also for the thermal contribution. As we will show, the result given in Ref.~\\cite{Mikheev} is a very good approximation only for temperatures well below the electron mass $T \\ll m_e$. For higher temperatures and sufficiently intense magnetic fields the thermal corrections are relevant.  Interestingly, the resulting mass shift induced by the magnetic field is independent of the gluon dynamics and, therefore, does not vanish if the (QCD induced) axion mass (\\ref{m0}) is set to zero. In other words, the external electromagnetic field itself is responsible for the breaking of the PQ-symmetry, playing a role similar to the gluon fields. We will discuss this point in Section 2 of this paper. We will show that the magnetic contribution is indeed due to a topological effect (just as for gluons) which is expected when the axion dynamics take place in an external magnetic field. Ultimately, we will associate this effect with the generation of magnetic helicity, whose important role in axion cosmology has already been pointed out (see Ref.~\\cite{noi2}). This property makes the magnetic mass shift a true mass term. Therefore, the terminology ``mass shift'' is perhaps inappropriate. In the following, we will refer to this contribution as the {\\it magnetic induced axion mass} $\\mb$. It contributes to the total axion mass as \\begin{equation} m^2 = \\mq^2 + \\mb^2, \\end{equation} where $\\mq$ is given in (\\ref{m0}).  Because of this property, our result applies not only to axions but to any pseudoscalar particle which interacts with photons as in Eq.~(\\ref{agamma}), independent of its interaction with gluons. This is the case with arions~\\cite{Anselm}, (almost) massless pseudoscalar particles which do not interact with gluons at all. \\footnote{Since arions do not interact with gluons, they are not affected by the color anomaly. Therefore, they do not have a QCD-induced mass term as in Eq.~(\\ref{m0}). The phenomenology of arions in an external magnetic field has been reconsidered recently as a tentative to analyze the problem of supernova dimming in terms of photon-arion oscillations~\\cite{Csaki} (see also Ref.~\\cite{Mirizzi}).} On the other hand, concerning the axion phenomenology, we expect the contribution $\\mb$ to be especially relevant at temperatures above $\\sim 1 \\GeV$, where the gluon contribution to the axion mass is strongly suppressed~\\cite{GPY}, or in regions of space and time where a very strong magnetic field is present. A more detailed analysis of the phenomenological implications of the magnetic induced axion mass is in progress.  The paper is organized as follows: In Section~2 we comment on the origin of the magnetic induced mass term. We then compute the magnetic induced axion mass in Section~3, discussing the zero temperature limit as well as the finite temperature result. Finally, in Section~4 we summarize and give our conclusions. ",
        "conclusions": "The phenomenology of the axion in an external magnetic field has been deeply studied in the last few decades. Nowadays the relevance of an in depth understanding of the effects of axion interaction in an external magnetic field to astrophysics and cosmology is clear. In this paper we have further investigated this problem computing the axion mass induced by an external magnetic field. The origin of this phenomenon is related to the generation of magnetic helicity, induced by the axion itself. We computed the resulting induced mass in one loop approximation, accounting for the thermal corrections, and with no restriction on the intensity of the magnetic field. Our result is given in Eqs.~(\\ref{dm}) and (\\ref{Pbars}) (see also Figs.~4 and 5), and indicates that thermal effects are not relevant for temperatures significantly below the electron mass. For $T \\ll m_e$ our result is indeed in agreement with the existent literature. However, for higher temperatures, the axion mass is strongly influenced by temperature effects if the external magnetic field is much stronger than the Schwinger value, $B \\gg B_c \\sim 10^{13} \\G$. In this case, thermal effects give rise to an increment of the axion mass with respect to the zero temperature limit, which is proportional to the magnetic field: $\\mb^2(T,B)/\\mb^2(0,B) \\simeq (\\aem/3\\pi) \\, (B/B_c)$.  It is clear that the value of the magnetic field that could induce a sizable change of the axion dispersion relation is very high, and certainly well above the possibility of the present terrestrial experiments. However, in an astrophysical or cosmological contest the magnetic induced mass could play an important role. This topic will be the object of future investigations.  \\vspace*{0.5cm}"
    },
    "0609/astro-ph0609073_arXiv.txt": {
        "abstract": "Modern astrometry is based on angular measurements at the micro-arcsecond level. At this accuracy a fully general relativistic treatment of the data reduction is required. This paper concludes a series of articles dedicated to the problem of relativistic light propagation, presenting the final microarcsecond version of a relativistic astrometric model which enable us to trace back the light path to its emitting source throughout the non-stationary gravity field of the moving bodies in the Solar System. The previous model is used as test-bed for numerical comparisons to the present one. Here we also test different versions of the computer code implementing the model at different levels of complexity to start exploring the best trade-off between numerical efficiency and the $\\mu\\textrm{as}$ accuracy needed to be reached. ",
        "introduction": "Modern space technology will soon provide stellar positioning with micro-arcsecond accuracy ($\\mu$as). At this level one has to take into account the general relativistic effects on light propagation arising from metric perturbations due not only to the bulk mass but also to the rotational and translational motion of the bodies of the Solar System, and to their multipole structure (see \\citealp{2002PhRvD..65f4025K,2003AJ....125.1580K,2004CQGra..21.4463L} and references therein). Our aim is to develop a Relativistic Astrometric MODel (RAMOD) which enabled us to deduce, to the accuracy of one $\\mu$as, the astrometric parameters of a star in our Galaxy from observations taken by modern space-born astrometric satellites like Gaia \\citep{2005tdug.conf.....T}, which are fully consistent with the precepts of General Relativity.  In this paper we present an astrometric model which contains an extension to the dynamical case, {\\em i.e.} with the inclusion of the $1/c^{3}$ terms, of our previous model which was only accurate to $1/c^{2}$ \\citep{2004ApJ...607..580D}. We term it as RAMOD3 since it was intended as the successor of two previous models (\\citealt{1998AAp...332.1133D,2001AAp...373..336D}). Following the same scheme, we shall refer to the model described here as RAMOD4.  The inclusion of terms of the order of $1/c^{3}$ corresponds to an accuracy of $0.1~\\mu$as, at least one order of magnitude better than the expected precision of the Gaia measurements. Here we require that the Solar System is isolated and source of a weak gravitational field. These two conditions imply that the velocities of the gravitational sources within the Solar System are very small compared to the velocity of light, typically of the order of $10\\,$km s$^{-1}$. Under these conditions we select a coordinate system $(x^{i},x^{0}\\equiv ct)_{(i=1,\\,2,\\,3)}$ such that the background geometry has a post-Minkowskian form: \\begin{equation} g_{\\alpha\\beta}=\\eta_{\\alpha\\beta}+h_{\\alpha\\beta}+O\\left(h^{2}\\right) \\label{eq:metric} \\end{equation} where $\\eta_{\\alpha\\beta}$ is the Minkowski metric and $h_{\\alpha\\beta}$ are small perturbations which describe effects generated by the bodies of the Solar System. These perturbations are \\textit{small} in the sense that $|h_{\\alpha\\beta}|\\ll1$, their spatial variations are of the order of $|h_{\\alpha\\beta}|$ while their time variations are of the order of $|h_{\\alpha\\beta}|/c$; here and in what follows Greek indeces run from $0$ to $3$. Clearly metric form (\\ref{eq:metric}) is preserved under coordinate gauge transformations of the order of $h$. The $h_{\\alpha\\beta}$'s contain terms of the order of at least $1/c^{2}$, hence we shall keep our approximation to first order in $h$. To the order of $1/c^{3}$ the time dependence of the background metric cannot be ignored therefore the time-like vector field $\\boldsymbol{\\eta}=\\boldsymbol{\\partial}_0$ tangent to the coordinate time axis, will not in general be a Killing field (namely an isometry for the space-time) unless one moves to far distances from the Solar System where the metric tends to be Minkowski's. The components of the vector field $\\boldsymbol{\\eta}$ are $\\eta^{\\mu}=\\delta_{0}^{\\mu},\\,\\eta_{\\mu}=g_{0\\mu}$ hence we easily deduce that the congruence $C_{\\boldsymbol{\\eta}}$, namely the family of curves having the vector field $\\boldsymbol{\\eta}$ as tangent field% \\footnote{These curves are also termed integral curves of the vector field $\\boldsymbol{\\eta}$.}, will have a non zero vorticity. From its definition \\begin{equation} \\omega_{\\alpha\\beta}=P(\\eta)_{\\alpha}{}^{\\sigma}P(\\eta)_{\\beta}{}^{\\rho}\\nabla_{[\\rho}\\eta_{\\sigma]}, \\end{equation} where $P(\\eta)_{\\lambda}^{\\mu}=\\delta_{\\lambda}^{\\mu}+\\eta_{\\lambda}\\eta^{\\mu}$ is the operator which projects orthogonally to $\\boldsymbol{\\eta}$, $\\nabla_{\\rho}$ the covariant $\\rho-$derivative relative to the given metric and square brackets mean antisymmetrization, we have that the only non vanishing components are, to the lowest order: \\begin{equation} \\omega_{ij}=\\partial_{[j}h_{i]0}+O\\left(1/c^{4}\\right),\\quad i,j=1,\\,2,\\,3. \\label{eq:vort} \\end{equation} We remind that $h_{0i}$ and $\\partial_{j}h_{0i}$ are $\\sim O(1/c^{3})$. Condition (\\ref{eq:vort}) implies that the surfaces $t=\\mathrm{constant}$ are not orthogonal to the integral curves of $\\boldsymbol{\\eta}$ at least in and nearby the Solar System \\citep{2004ApJ...607..580D}. Nevertheless, because the time lines with tangent field $\\boldsymbol{\\eta}$ are asymptotically Killing and vorticity free then the slices $S(t):\\; t=\\mathrm{constant}$ allow for a non ambiguous $3+1$ splitting of spacetime with a coordinate representation such that space-like coordinates are fixed within each slice. We fix the origin of the space-like coordinates at the barycenter of the Solar System and assume that the spatial coordinate axes point to distant sources chosen in such a way to assure that the system is kinematically non rotating. Adopting the IAU prescriptions \\citep{2000IAU-res....B1.3}, this system is termed Barycentric Celestial Reference System (BCRS) and will be our main system since all coordinate tensorial components will be relative to it. The non stationarity of the spacetime makes the construction of a relativistic astrometric model much less straightforward than for the static case considered in \\citet{2004ApJ...607..580D}.  In section~\\ref{sec:light-trajectory} these complications are handled to define the geometrical environment for light propagation and data analysis. The spacetime metric is not diagonal since terms as $h_{0i}$ are different than zero; they are mainly generated by the velocity of the metric sources relative to the given BCRS. Then the gravitational potential at each point of the light trajectory depends on the sources in the Solar System at the appropriate retarded position as specified in section~\\ref{sec:retarded-time}. This will have consequences when fixing the observables and the boundary conditions for the differential equations of the light rays in section~\\ref{sec:boundary-conditions}. Section~\\ref{sec:Tests} shows how this model was numerically tested and presents the results of these tests.  In what follows we shall use geometrized units such that $c=1=G$, $G$ being the gravitational constant; with these units a mass $\\mathcal{M}$, expressed in kilograms, has the dimension of a length according to the relation $M=G\\mathcal{M}/c^{2}$; similarly the time coordinate $x^{0}=ct$ will be simply written as $t$ and the spatial velocities are in units of $c$. For sake of clarity, the velocity of light $c$ will appear explicitly only when we specify the order of magnitude of terms under discussion. ",
        "conclusions": "Aim of modern relativistic astrometry is to produce a three dimensional rendering of our Galaxy with an accuracy of one $\\mu\\mathrm{as}$ in the measurements of angles. In order to reach this goal one needs an appropriate algorithm to trace back to their emission space-time points the light signals which are intercepted by a suitable observational device.  In the near future, one of such devices will be the astrometric satellite Gaia which will orbit the Sun from nearby the Sun-Earth outer Lagrangian point $L_{2}$. The expected accuracy of the satellite observations requires that one should treat the light propagation through the Solar System in a general relativistic framework and consider the general relativistic effects induced by the bodies of the Solar System up to the order of $1/c^{3}$.  Our strategy was to construct a series of models with increasing generality and complexity, exploiting each of them as a test-bed for the more advanced one. Our last model accurate to $1/c^{3}$ is RAMOD4. As previously discussed, we first compared RAMOD4 to its less accurate predecessor RAMOD3, which was accurate to $1/c^{2}$. With a suitable handling of the items considered in the model, we are able to highlight the relative importance of the individual sources of relativistic perturbations and judge about their physical relevance.  Our first conclusion is that RAMOD4 behaves as expected with respect to RAMOD3, that is, the differences between the two models are at the expected level of $\\sim10~\\mu\\mathrm{as}$. So this work successfully completes the series of models from a theoretical point of view and represents a launching platform to deduce with a completely numerical algorithm the astrometric parameters of a celestial object from a well-defined set of relativistically measured quantities.  This model, however, is part of a broader project whose completion requires its implementation into a feasible and efficient structure for the data reduction of the Gaia mission. In this sense, the practical application of this model requires also to investigate the specific ways of implementing it with respect to the accuracy goals of the mission. For this reason, then, we tried to single out the different contribution which come out together from the motion of the perturbing bodies, and evaluate them separately. We found that the contribution to the total deflection coming from the velocity-induced terms of the metric is generally smaller than that of the reduced distance, but not negligible at the $0.1~\\mu\\mathrm{as}$ level. The former one can in fact reach a maximum amount of $\\sim1.8\\:\\mu\\mathrm{as}$ while the latter goes up to $\\sim4.5\\:\\mu\\mathrm{as}$. Both these cases, and the greatest differences with RAMOD3, come out when the photon is grazing Jupiter, while the Sun contributes up to less than $0.3~\\mu\\mathrm{as}$ that is one order of magnitude lower.   Given the unprecedented accuracy reachable by Gaia-like missions, several new relativistic tests related to the light propagation in a non-stationary gravitational field could be carried out. From this point of view it is advisable to create a cross-check framework of models to validate the relativistic outputs induced by the modeling itself at the microarcsecond level. In fact, our Relativistic Astrometric Model naturally confronts itself with the works of Kopeikin, Sch\\\"afer and Mashhoon \\citep{1999PhRvD..60f124002K,2002PhRvD..65f4025K} and \\citet{2003AJ....125.1580K}. While the latter stems from the general formulation of the former conforming to its basic tenets, our model responds to an inherently different strategy. As far as the astrometric problem is concerned, in the Kopeikin-Sch\\\"afer and Klioner approach the light ray is reconstructed as a sum of terms which allows for a direct evaluation of the individual relativistic effects induced by the gravitating bodies that the light ray is expected to encounter on its way. Each relativistic effect enters as a perturbation of the light ray trajectory and is treated to a sufficiently high order (not less than $1/c^3$) to cope with the accuracy expected from Gaia's observations.  Our model instead aims to determine a full solution for the light trajectory which naturally includes, in a curved space-time, all the individual effects; the latters are somewhat hidden in the covariant formalism of our approach and directly contribute to the solutions of our master equations.  Evidently any specific effect one is interested to explore can be independently deduced from our formalism as a branch output, once we adapt the master equation to a required specific case.  Although our model and Klioner's are of comparable accuracy as it can be deduced from individual tests, the overall structure of Klioner's model makes it ready to handle the reduction of a large number of observational data in a computational efficient way; nevertheless we are confident that our model will soon reach the same versatility and operate as an essential tool of verification. However, the general covariant formulation of our model, including the complete relativistic treatment of the satellite attitude (\\citealp{2003CQGra..20.2251B,2003CQGra..20.4695B}), represents a well defined framework where any desired advancement in the light tracing problem using astrometric data is contemplated."
    },
    "0609/astro-ph0609044.txt": {
        "abstract": "\\noindent We present \\zebra{}, the Zurich Extragalactic Bayesian Redshift Analyzer. \\\\ The  current version of  \\zebra{} combines and extends several of the classical approaches to produce accurate photometric redshifts down to faint magnitudes. In particular, \\zebra{} uses the template-fitting approach to produce  Maximum Likelihood and  Bayesian redshift estimates based on: \\\\ $(1.)$ An automatic iterative technique to correct the original set of galaxy templates to best represent the SEDs of real galaxies at different redshifts;  \\\\ $(2.)$ A training set of spectroscopic redshifts for a small fraction of the photometric sample to improve the robustness of the photometric redshift estimates; and \\\\ $(3.)$ An iterative technique for  Bayesian redshift estimates, which extracts  the full two-dimensional  redshift  and template probability function for each galaxy. \\\\  We demonstrate the performance of  \\zebra{} by applying it to a sample of  866  $I_{AB}\\le 22.5$ COSMOS galaxies with  available $u*$, $B$, $V$, $g'$, $r'$, $i'$, $z'$ and $K_s$ photometry and zCOSMOS spectroscopic redshifts in the range $0 <  z < 1.3$. Adopting a 5-$\\sigma$-clipping that excludes $\\le10$  galaxies, both the Maximum Likelihood and Bayesian  \\zebra{}  estimates for this sample have an accuracy $\\sigma_{\\Delta{}z/(1+z)}$ smaller than 0.03. Similar accuracies are recovered using mock galaxies.\\\\   \\zebra{}  is made available to the public at:  {\\it http://www.exp-astro.phys.ethz.ch/ZEBRA.} ",
        "introduction": "\\label{sec:Introduction}  Current imaging surveys of faint high redshift galaxies such as, e.g., COSMOS \\citep{scoville2006}, already return millions of galaxies with magnitudes well beyond the current observational spectroscopic limits. As spectroscopic redshifts for such large distant galaxy samples will thus remain practically unobtainable in the foreseeable future,  photometric redshifts of increasing accuracy will have to be constructed in order to properly exploit the wealth of information, as a function of cosmic epoch,  that is potentially extractable from state-of-the-art and future large imaging surveys.  The importance of estimating accurate redshifts from multi-wavelength medium- and broad-band photometry for large galaxy samples is reflected in the extensive efforts that have been devoted to improving algorithms and methodologies to increase the accuracy of the photometric  estimates (see, e.g.,  \\citealt{sawicki1997}; \\citealt{1998astro.ph..9347Y}; \\citealt{arnouts1999}; \\citealt{2000ApJ...536..571B} (BPZ); \\citealt{2000AJ....120.1588B}; \\citealt{firth2003}; \\citealt{benitez2004}; \\citealt{2006ApJS..162...20B}; \\citealt{ilbert2006}, and references therein). These works are based on a few basic principles, namely:  $(i)$ $\\chi^2$ minimization of the difference between a model-galaxy Spectral Energy Distribution (SED; the model SEDs are hereafter referred to as {\\it templates}) and the observed galaxy photometry; $(ii)$  Neural network approaches that rely on the availability of a small sample of spectroscopic redshifts to find a functional dependence between photometric data and redshifts; $(iii) $ Hybrid approaches that perform  standard $\\chi^2$ minimization while using  a small spectroscopic training sample to optimize the initial set of galaxy templates; $(iv)$ Bayesian methods which use additional information provided by a prior to obtain final photometric redshift estimates.  Motivated by the scientific returns of deriving accurate photometric redshifts for large numbers of faint COSMOS  galaxies, we have developed \\zebra{}, the {\\it Zurich Extragalactic Bayesian Redshift Analyzer}.  In this paper we describe the current version of \\zebra{}, which combines and extends several of the above-mentioned approaches to produce accurate photometric redshifts down to faint magnitudes. More specifically, the paper is structured as follows:  Section \\ref{sect:zebraIntro}: {\\it About \\zebra{}}, specifies the input requirements and the output of \\zebra{}.  Section \\ref{sect:zebraPrinciples}: {\\it The principles of \\zebra{}}, describes  the general design and  methodological details of  the code. A flow chart indicating the architectural structure of  \\zebra{} is shown in Figure  \\ref{fig:DesignZEBRA}. Basically, \\zebra{} produces two separate estimates for the photometric redshifts of individual galaxies: A Maximum Likelihood (ML) estimate,  and a Bayesian (BY) estimate. These achieve a high accuracy by combining together some novel features with several of the approaches that have been published in the  literature. In particular, \\zebra{}:  \\begin{description} \\item{-} Uses a novel automatic iterative technique to correct an original set of galaxy templates to best represent the SEDs of real galaxies at different redshifts. These template corrections depend  on the accuracies and systematic errors in the absolute photometric calibrations; therefore,  prior to performing the individual template corrections, \\zebra{} automatically removes systematic calibration errors in the input photometric catalogs. The template corrections substantially reduce  the photometric redshift inaccuracies that are generated by galaxy-template mismatches; \\item{-}  Can be fed with a training set of spectroscopic redshifts for a small fraction of the photometric sample, to improve the robustness of the photometric redshift estimates; \\item{-}  Adopts  an  iterative technique for  Bayesian photometric redshift estimates that extracts  the full two-dimensional  redshift  {\\it \\underline{and}}  template likelihood function for each galaxy. \\end{description}  Section \\ref{sect:application}: {\\it The $1^{st}$ application of \\zebra{}}, demonstrates the performance of  \\zebra{} by comparing our photometric redshifts estimates for a sample of $866$   $I_{HST,AB}<22.5$  ACS-selected COSMOS galaxies with high-quality zCOSMOS  spectroscopic redshifts $z_{spec}\\le1.3$ \\citep{lil06}.  Based on the currently available passbands and  photometric accuracies,  both the  ML and BY  \\zebra{} photo-$z$'s  for COSMOS galaxies have a   $5\\sigma-$clipped accuracy of $\\Delta{}z /(1+z)=0.027$ over the entire redshift range (with $\\sim1\\%$ outliers).  Section \\ref{sect:Conclusions}: {\\it  Concluding remarks}, briefly comments on the first applications of the COSMOS \\zebra{} photometric redshifts, and lists  the  developments which we are already implementing in the  next version of  \\zebra{}.  The three Appendices introduce the notation and conventions that we use throughout the paper (Appendix \\ref{sect:definitions}), present in detail the explicit mathematical formulation of \\zebra{}'s algorithms (Appendix \\ref{sect:tempCorrection}), and demonstrate the \\zebra{} performance on a Mock catalogue produced for the COSMOS survey (courtesy of Manfred Kitzbichler; Appendix \\ref{sect:mock}).  \\zebra{} is made available to the general public at {\\it http://www.exp-astro.phys.ethz.ch/ZEBRA}\\footnote{The \\zebra{} website is currently under construction.}.  The use of \\zebra{} should please be acknowledged with an explicit reference  to this paper in the bibliographic list of any resulting publication. ",
        "conclusions": "\\label{sect:Conclusions}  A more thorough comparison of the \\zebra{}  ML and BY photo-$z$'s with the currently available zCOSMOS redshifts is given by \\cite{lil06}. Furthermore, the \\zebra{} ML and BY photometric redshift estimates are compared by \\cite{mobasher2006} with photo-$z$ estimates derived for the same galaxies with independent codes (these are either public codes, e.g. BPZ or have been developed by other teams within the COSMOS collaboration).  The \\zebra{} ML photometric redshifts estimates for the COSMOS sample studied in this paper have already been used to derive the evolution up to $z\\sim1$ of the luminosity functions for morphologically-classified early-, disk- and irregular-type galaxies (according to the classification scheme of {\\it ZEST}, the Zurich Estimator of Structural Types; \\citealt{scarlata2006}), to study the evolution up to similar redshifts of the number density of intermediate-size and large disk galaxies \\citep{sargent2006} and to study the evolution of the luminosity function of elliptical galaxy progenitors \\citep{scarlata2006a}.  The current version (1.0) of \\zebra{} is being packaged with a user-friendly web-interface at the URL provided in the abstract. The \\zebra{}  website will be constantly updated to provide the newest improved versions of the \\zebra{} code, and the associated documentation describing in detail the implemented changes. In the meanwhile, \\zebra{} is being upgraded with several new modules, including $(1)$ A module that incorporates dust absorption and reddening, according to several user-specifiable extinction corrections; $(2)$ An improved treatment of AGN's; and $(3)$ A module that uses several synthetic template models and a large choice of self-consistent star formation and metal enrichment schemes to estimate stellar masses, average ages and metallicities (and their uncertainties)."
    },
    "0609/astro-ph0609245_arXiv.txt": {
        "abstract": "CH Cygni is a symbiotic star consisting of an M giant and an accreting white dwarf, which is known to be a highly variable X-ray source with a complex, two-component, spectra.  Here we report on two {\\sl Suzaku\\/} observations of CH Cyg, taken in 2006 January and May, during which the system was seen to be in a soft X-ray bright, hard X-ray faint state. Based on the extraordinary strength of the 6.4 keV fluorescent Fe K$\\alpha$ line, we show that the hard X-rays observed with {\\sl Suzaku\\/} are dominated by scattering. ",
        "introduction": "A symbiotic star is a binary composed of a red giant and a hot blue component, often an accreting white dwarf.  Many symbiotic stars have been detected as soft X-ray sources using {\\sl ROSAT\\/} \\citep{MWJ1997}.  The {\\sl ROSAT\\/} detections can usually be classified into two types.  One is of the supersoft type, which can be understood as photospheric emission from the hot white dwarf.  The second type is optically thin plasma emission with temperatures of order 10$^7$K, which \\authorcite{MWJ1997} attributed to the colliding winds. That is, a strong, fast ($\\sim$1000 km\\,s$^{-1}$) wind from the vicinity of the accreting white dwarf, which, in many cases, are known to be present through studies of UV and optical emission lines, is colliding with the dense, low-velocity wind from the M giant, shock-heating the gas to X-ray emitting temperatures.  More recently, a subset of symbiotic stars have been shown to be a strong source of hard ($>$10 keV) X-rays in the {\\sl INTEGRAL\\/} Galactic plane scan and in the {\\sl Swift\\/} BAT survey (\\cite{Kea2006} and references therein), which requires the presence of yet another type of X-ray emission in these symbiotic stars.  CH Cygni is an oft-observed, yet poorly-understood symbiotic star, at a distance of 245$\\pm$50 pc according to the {\\sl Hipparcos\\/} parallax measurements \\citep{Pea1997}.  It has produced radio \\citep{TSM1986} and optical \\citep{S1987} jets, which appear to be ejected in the plane of the sky as seen from the Earth, implying an accretion disk and/or binary with an inclination angle of nearly 90$^\\circ$.  \\citet{Hea1993} analyzed a large body of infrared spectra and proposed CH Cyg to be a triple system, consisting of a 756 day period inner binary in an 14.5 yr orbit around a third component.  However, \\citet{Mea1996} cautioned against the triple system interpretation on several grounds. In particular, the orbital solution of \\authorcite{Hea1993} requires a low inclination angle, contrary to the indications from the jet data. In addition, such a low inclination then results in the unrealistically low mass (0.2M$_\\odot$) for the white dwarf.  To add to the confusion, \\citet{Sea1996} claimed that CH Cyg was a triple system as proposed by \\citet{Hea1993} but both the inner and the outer binaries were eclipsing, thus requiring high inclination angles for both.  More recently, the similarity between the 756 day period of CH Cyg with non-radial pulsations in quite a few semi-regular variables has become apparent; therefore, the 756 day period of CH Cyg is likely to be a non-radial pulsation mode rather than an orbital period (\\cite{Sea2006} and references therein).  CH Cyg has been known to be a variable X-ray source for many years; the early X-ray data from {\\sl HEAO-1\\/} to {\\sl EXOSAT\\/} are summarized by \\citet{LT1987}.  It was detected with {\\sl ROSAT\\/} showing the colliding wind-type spectrum and an additional hard component \\citep{MWJ1997}.  However, it was the {\\sl ASCA\\/} observation that revealed the dramatic two-component X-ray spectrum \\citep{EIM1997}. There is a relatively unabsorbed soft component, dominant below 2 keV, that can be modeled as two-temperature (kT$\\sim$0.2, 0.7 keV) optically thin thermal plasma emission.  Above 2 keV, there is a separate, highly absorbed hotter component (kT$\\sim$7 keV).  \\citet{EIM1997} interpreted the latter as due to accretion onto a $>$0.44 M$_\\odot$ white dwarf, and the former as either from the M giant or from the jets.  \\citet{W2001} proposed an alternative interpretation of a single spectral component seen through an ionized wind of the M giant.  A {\\sl Chandra\\/} HETG observation was performed to decide between the two models, but CH Cyg was found to be much fainter than during the {\\sl ASCA\\/} observation, making the grating data of little use.  Instead, \\citet{GS2004} analyzed the 0th order data and discovered a weak, spatially extended emission from the jet, in addition to the much brighter, spatially unresolved, component.  Here we report on two {\\sl Suzaku\\/} observations of CH Cyg performed in 2006.  We have also re-analyzed the {\\sl ASCA\\/} data and the {\\sl Chandra\\/} zeroth order data, and present the results in a uniform manner. ",
        "conclusions": "Our two {\\sl Suzaku\\/} observations of CH Cyg have caught the system in a soft X-ray bright, hard X-ray dim state.  Based on the Fe K$\\alpha$ spectra, we believe we currently observe only the scattered hard X-ray component.  The decline in the hard X-ray flux is only an apparent decline, and the intrinsic X-ray luminosity of CH Cyg may still be as high as during the {\\sl ASCA\\/} observation, though the actual value is unknowable.  A symbiotic star such as CH Cyg is filled with possible scattering media, but we favor material above the accretion disk as the scattering site. As for the object that blocks our direct view of the hard X-ray emission, we believe it is more likely to be an edge-on accretion disk than the M giant mass donor.  The hard X-rays presumably originate from accretion onto the white dwarf, but further progress would require high-quality X-ray spectroscopy when the direct view is not blocked.  \\subsection*{Acknowledgements}  We thank Drs. M. Corcoran and S. Drake for useful discussion on colliding winds and M giant coronae.  \\newpage"
    },
    "0609/astro-ph0609135_arXiv.txt": {
        "abstract": "We report the first detection of gas-phase CO$_2$ emission in the star-forming region Cepheus A East, obtained by spectral line mapping of the $\\nu_2$ bending mode at 14.98 $\\mu$m with the Infrared Spectrograph (IRS) instrument onboard the {\\it Spitzer Space Telescope}. The gaseous CO$_2$ emission covers a region about 35$''$ $\\times$ 25$''$ in extent, and results from radiative pumping by 15 $\\mu$m continuum photons emanating predominantly from the HW2 protostellar region. The gaseous CO$_2$ exhibits a temperature distribution ranging from 50 K to 200 K. A correlation between the gas-phase CO$_2$ distribution and that of H$_2$ S(2), a tracer of shock activity, indicates that the CO$_2$ molecules originate in a cool post-shock gas component associated with the outflow powered by HW2. The presence of CO$_2$ ice absorption features at 15.20 $\\mu$m toward this region and the lack of correlation between the IR continuum emission and the CO$_2$ gas emission distribution further suggest that the gaseous CO$_2$ molecules are mainly sputtered off grain mantles -- by the passage of slow non-dissociative shocks with velocities of 15-30 km s$^{-1}$ -- rather than sublimated through grain heating. ",
        "introduction": "Cepheus A is a well-known site of star formation that has been studied extensively. Previous results have shown that it contains a series of heavily embedded far-infrared and radio-continuum sources, one of which dominates the luminosity of the entire region (HW2 with 2.5$\\times$10$^{4}$ $L_{\\odot}$; Hughes \\& Wouterloot 1984; Evans et al. 1981). Ground- (e.g. Goetz et al. 1998) and space-based observations (e.g. Hartigan et al. 1996; Wright et al. 1996; van den Ancker et al. 2000) revealed that the multipolar outflow existing in this region consists of complex structures of shock-excited atomic and molecular gas. H$_2$ emissions and line emissions from ionized species were detected along high-velocity (HV) and extra-high-velocity (EHV) outflows, indicating the presence of both dissociative (J-type) and non-dissociative (C-type) shocks, likely resulting from successive episodes of activity (Narayanan \\& Walker 1996). The protostellar object HW2 was determined to be the dominant powering source of the EHV outflow, while the source of the HV jet is still under debate (Goetz et al. 1998; Hiriart et al. 2004). The cavities carved into the surrounding molecular gas by the quadrupolar outflow are clearly seen in the NH$_3$ maps obtained by Torrelles et al. (1993).   In this {\\it Letter}, we focus on new observations of Cepheus A East obtained with the Infrared Spectrograph (IRS) onboard the {\\it Spitzer Space Telescope} ({\\it SST}).  The benefit of the {\\it SST} data lies in the much greater sensitivity and the significantly better spatial resolution offered by the IRS compared to the {\\it Infrared Space Observatory (ISO)}. Such advances allow for the detection and spectral mapping of variations in the molecular and atomic line intensities on scales of a few arcsec instead of the previous few tens of arcsec or few arcmin scales, thus enabling our understanding of shock physics and chemistry at much finer scales. Detections of H$_2$ S(1) through S(7), C$_2$H$_2$, [\\ion{Ne}{2}], [\\ion{Ne}{3}], [\\ion{S}{1}], [\\ion{S}{3}] and [\\ion{Fe}{2}] emissions as well as absorption from CO$_2$ and H$_2$O ices will be discussed in a future paper. The present work will focus on the first detection of gas-phase CO$_2$ emission toward Cepheus A East. Section 2 describes the observations and data analysis. Sections 3 \\& 4 discuss our new results as well as the spatial distributions of gas-phase CO$_2$ and H$_2$ S(2) in the context of shock chemistry and outflow activity. ",
        "conclusions": "The presence of CO$_2$ ice absorption features at most spatial positions in the gaseous CO$_2$-emitting region, their progressive weakening along the EHV outflow axis, as well as the localization of the gaseous CO$_2$ emission into a cool post-shock component, all pose the question of the origin of the gas-phase CO$_2$ molecules in Cepheus A East. Was CO$_2$ injected into the gas-phase by sublimation from CO$_2$-rich dust mantles after substantial grain heating or by grain sputtering caused by the passage of slow non-dissociative shocks known to exist in this active region?  We searched for evidence of a correlation between the dust continuum emission, a tracer of grain heating, and the gaseous CO$_2$ distribution in the CO$_2$-emitting region by comparing the CO$_2$ column density measurements with measurements of the dust continuum emission at 7.3 $\\mu$m, 14.8 $\\mu$m and 18.0 $\\mu$m. We found no obvious relation between $N$(CO$_2$) and the dust continuum emission at these wavelengths. Furthermore, we determined that the dust temperature is never greater than the CO$_2$ sublimation temperature ($\\sim$ 90 K) at the locations of the CO$_2$ column peaks. While grain heating likely occurs in the active region close to HW2, CO$_2$-rich grain mantle sublimation is not the dominant source of gaseous CO$_2$ in Cepheus A East. Sputtering off CO$_2$-rich ice mantles is, therefore, the most likely scenario in the present case. Our measurements of the CO$_2$ ice column density range from $\\sim$ 10$^{17}$ to $\\sim$ few $\\times$ 10$^{18}$ cm$^{-2}$ over the region exhibiting CO$_2$ gas emission, indicating that only a few percent of the material has passed through mantle-destroying shocks.  To derive the abundance of gaseous CO$_2$, an estimate of $N$(H$_2$) in the CO$_2$ gas component is needed. Because the gas-phase CO$_2$ molecules arise from a post-shock component much cooler than that containing the H$_2$ gas we detected, we cannot directly measure the total column density of the gas at $T=$ 50-200 K from our data. Our abundance estimate will, hence, rely on previous indirect measures. G\\'omez et al. (1999) derived $N$(H$_2$) $=$ 1.5 $\\times$ 10$^{22}$ cm$^{-2}$ in the outflow using HCO$^+$ emission measurements; a similar estimate was obtained by van den Ancker (2000) using far-infrared CO emission at the {\\it ISO}/Long Wavelength Spectrometer angular resolution (75$''$ beam size). Adopting this H$_2$ column density yields an average gas-phase CO$_2$ abundance of few $\\times$10$^{-7}$.  This value should be regarded as a lower limit since the $N$(CO$_2$) might be underestimated if the true distance from HW2 is greater than the projected distance we assumed. A small fraction of the gaseous CO$_2$ might arise from CO$_2$ ice sublimation since some CO$_2$ ice profiles, especially close to HW2, show substructures characteristic of grain heating (Gerakines et al. 1999). A fraction of the gas-phase CO$_2$ emission might also come from quiescent gas associated with Cep-A2 (knot at [R.A., Dec]= [7$''$,18$''$] in Fig.~1). While these additional sources of gaseous CO$_2$ certainly contribute to the scatter seen in the right panel of Fig.~3, none of these sources can account for the correlation we observe between $N$(CO$_2$) and $N$(H$_2$) S(2) over such a large spatial extent.  The low gas-phase CO$_2$ abundances detected by {\\it ISO} ($\\sim$10$^{-7}$ instead of the expected few$\\times$10$^{-6}$) toward active star-forming regions with large CO$_2$ ices reservoirs (e.g. van Dishoeck et al. 1996) fostered theoretical efforts to investigate the effects that shocks might have on CO$_2$ chemistry. Charnley \\& Kaufman (2000) showed that for shocks faster than a critical shock speed, which is a function of preshock gas density, CO$_2$ molecules sputtered off grain mantles are expected to be efficiently destroyed in the shock through reactions with H and H$_2$. On the other hand, a shock speed higher than 15 km s$^{-1}$ is necessary to allow efficient sputtering.  Comparison of our results with the Charnley \\& Kaufman (2000) shock-model predictions indicates that slow C-type shocks with speeds of at least $\\sim$ 15 km s$^{-1}$ interacted with the ambient gas over the extent of the CO$_2$-emitting region (see Fig~1). Adopting a pre-shock density of $n_H \\sim$ 10$^{5}$ cm$^{-3}$ (Goetz et al. 1998), our abundance limit further indicates that the shock speed is no greater than 30 km s$^{-1}$ since, at this density, higher shock speeds would efficiently remove CO$_2$ from the gas phase, mainly through reactions with hydrogen atoms (see Fig.~3 of Charnley \\& Kaufman 2000). The CO$_2$ emission might arise at the flanks of the bow shock modeled by Froebrich et al. (2002). The average gas phase CO$_2$ abundance derived above (few $\\times$10$^{-7}$) is consistent with a scenario in which a few percent of the grain material has been subject to shocks fast enough to sputter CO$_2$ ice mantles (with CO$_2$(ice)/H$_2$ $\\sim$ few $\\times$10$^{-5}$)."
    },
    "0609/astro-ph0609629_arXiv.txt": {
        "abstract": "Using the Millennium N--body simulation we explore how the shape and  angular momentum of galaxy dark matter haloes  surrounding the largest cosmological voids are oriented. We find that the major and intermediate axes of the haloes tend to lie parallel to the surface of the voids, whereas the minor axis points preferentially in the radial direction. We have quantified the strength of these alignments at different radial distances from the void centres. The effect of these orientations is still detected at distances as large as 2.2 R$_{void}$ from the void centre. Taking a subsample of haloes expected to contain disc--dominated galaxies at their centres we detect, at the 99.9\\% confidence level,  a signal that the angular momentum of those haloes tends to lie parallel to the surface of the voids. Contrary to the alignments of the inertia axes, this signal is only detected in shells at the void surface (1$<$R$<$1.07 R$_{void}$) and disappears at larger distances. This signal, together with the similar alignment observed using real spiral galaxies (Trujillo, Carretero \\& Patiri 2006), strongly supports the prediction of the Tidal Torque theory  that both dark matter haloes and baryonic  matter have acquired, conjointly, their angular momentum before the moment of turnaround. ",
        "introduction": "The angular momentum  of a galaxy  plays a central role in determining its evolution and final type. However, understanding the origin and properties of the galactic angular momentum  has been one of the key problems in astrophysics in the last five decades. The current 'standard' theory for the origin of the angular momentum, within the cosmological framework of  hierarchical structure formation, is the tidal--torque theory (hereafter TTT). This theory is based on early ideas from Hoyle (1951) that suggested that the angular momentum of a galaxy arises from the tidal field of neighbouring galaxies. This idea was further developed and quantified by Peebles (1969), Doroshkevich (1970) and White (1984).    The TTT suggests that most of the angular momentum is gained gradually by the protohaloes during the linear regime of the growth of density fluctuations, due to  tidal torques from neighbouring fluctuations. Angular momentum grows linearly with time at this early epoch and saturates when the halo decouples from the expanding background at the moment of turnaround. It is assumed that during this phase the baryonic component follows the dark matter distribution and consequently gains a similar specific angular momentum to that of the halo. A subsequent large collapse factor of the baryonic matter to the centre of the halo will explain the centrifugally supported nature of the galactic discs (Fall \\& Efstathiou 1980).    To first order, the angular momentum of the haloes results from the misalignment between the principal axes of the inertia momentum tensor (I$_{ij}$) of the matter being torqued and the principal axes of the shear or tidal tensor (T$_{ij}$=-$\\partial^2\\phi/\\partial x_i\\partial x_j$) generated by neighbouring density fluctuations. The leading term of the torque is given by L$_i$$\\propto$T$_{jk}$(I$_{jj}$-I$_{kk}$), where $i$, $j$, and $k$ are cyclic permutations of 1, 2 and 3. Several aspects of this picture have been confirmed by a number of studies of the angular momentum properties of dark matter haloes both analytically  (Heavens \\& Peacock 1988; Catelan \\& Theuns 1996) and in N--body simulations (Barnes \\& Efstathiou 1987; Sugerman, Summers \\& Kamionkowski 2000; Lee \\& Pen 2000; Porciani, Dekel \\& Hoffman 2002a,b).    If $\\bf{I}$ and $\\bf{T}$ were uncorrelated (as frequently has been assumed under the argument that the former is local and the latter is dominated by the distribution of the large--scale structure) the direction of the angular momentum in the linear regime should be aligned with the intermediate axis of $\\bf{I}$ (the direction that maximizes the difference between I$_{jj}$ and I$_{kk}$). However, Porciani et al. (2002b) have found in their simulations that there is a strong correlation between the $\\bf{I}$ and $\\bf{T}$ tensors, in the sense that their minor, major and intermediate principal axes tend to be aligned at the protohalo stage. The strong correlation between $\\bf{I}$ and $\\bf{T}$ should produce an angular momentum vector of the haloes that is perpendicular to the minor axis of the sheet  they are embedded in (i.e.  perpendicular to the direction of the maximum compression of the large--scale structure at that point). However, this last correlation, at least at redshift zero, is expected to be very weak (or totally erased) by non--linear effects (such as exchange of angular momentum between haloes) at late times.   From the observational point of view, Trujillo, Carretero \\& Patiri (2006) have found that the rotation axes of spiral galaxies located on the shells of the largest cosmic voids lie preferentially on the void surface (in agreement with the expectation for the angular momentum of haloes given in the previous picture). The observational relation could be explained then as a consequence of the spin of the baryonic matter still retaining memory of the angular momentum properties of the haloes at the moment of  turnaround (Navarro, Abadi \\& Steinmetz 2004). It is key, consequently, to explore whether the signal is also found in the haloes of cosmological N--body simulations when we mimic the observational technique. If so, this will strongly support our current understanding of how haloes and baryonic matter have acquired, conjointly, their angular momentum.   Together with the orientation of the angular momentum, the alignment of the shape of the galaxy dark matter haloes (M$<$10$^{13}h^{-1}$M$_{\\sun}$) with their surrounding large--scale structure can have important observational consequences. In fact, the shapes of dark matter haloes can affect the coherence of tidal streams (Sackett 1999), can be related to galactic warps (Ostriker \\& Binney 1989; Debattista \\& Sellwood 1999; L\\'opez--Corredoira et al. 2002) or can affect the distribution of the orbits of satellite galaxies (Holmberg 1969; Zaritsky et al. 1997; Sales \\& Lambas 2004; Agustsson \\& Brainerd 2006; Yang et al. 2006). In addition, infall of material into the haloes is not    isotropic but  is expected to be through the filaments where the haloes are embedded. Consequently, the orientation of the dark matter haloes within these structures can affect the characteristics of the galaxy properties previously mentioned.    The alignment of massive (group and cluster) haloes (M$>$10$^{13}h^{-1}$M$_{\\sun}$) with their surrounding large--scale structure has been explored in detail (Splinter et al. 1997; Onuora \\& Thomas 2000; Faltenbacher et al. 2002; Kasun \\& Evrard 2005; Hopkins et al. 2005; Basilakos et al. 2006). These works indicate that the major axes of neighbouring galaxy clusters are aligned. A result that is in agreement with the \"Binggeli (1982) effect\". The origin of these alignments is still under debate and could be associated to infall of material (van Haarlem \\& van de Weygaert 1993) and/or tidal fields (Bond et al. 1996). Due to the lack of resolution in previous simulations, it is only now that an exploration of the alignment of the galaxy dark matter haloes has become possible (Bailin \\& Steinmetz 2005, hereafter BS05; Altay, Colberg \\& Croft 2006, hereafter ACC06; Patiri et al. 2006b, hereafter PA06; Arag\\'on--Calvo et al. 2006). These works suggest that the major axis of the haloes lies along the filaments. The quantification of the strength of these alignments is key to studies of strong and weak lensing where the intrinsic distribution and alignment of galaxy shapes plays an important role in interpreting the signal (see e.g. Heavens, Refregier \\& Heymans 2000; Croft \\& Metzler 2000; Heymans et al. 2006).    The aim of this paper is to characterise the alignment of both the shape and angular momentum of galaxy dark matter haloes with their surrounding large--scale structure to an unprecedent statistical level using the Millennium simulation (Springel et al. 2005).  In particular, we will focus our attention on haloes surrounding the largest cosmological voids. Contrary to filaments (which are strongly affected by redshift--space distortion), large cosmological voids are a feature easy to characterise from the observational point of view (Trujillo et al. 2006). In addition, another important advantage of the void scheme is that (because of the radial growth of the voids) the vector joining the centre of the void with the galaxy (halo) is a good approximation to the direction of the maximum compression of the large--scale structure at that point.  This naturally generates a framework for exploring the alignments of the shape and angular momentum of the haloes with their surrounding matter distribution. Consequently, our work mimics the observational framework to provide an easy interpretation of current and future observations. The large volume sampled by the Millennium simulation combined with the excellent spatial resolution allows us to conduct this project. This  unprecendent statistical power is absolutely crucial if we want to explore signals expected to be very weak like the alignment of  halo angular momentum and the large--scale structure.    This paper is structured as follows. Section 2 provides a description of the Millennium simulation itself, the void and halo samples used, and details of how the shapes and spins were measured. Section 3 describes our results and we discuss our findings in Section 4.        \\section[]{Numerical simulation}    \\subsection[]{N--body simulation}   The main simulation we have used for this study is the {\\it Millennium Simulation} of Springel et al (2005). This employs $2160^3$ dark matter particles each of mass $8.6\\times 10^8 h^{-1}{\\rm M}_\\odot$ within a comoving box of side $500h^{-1}{\\rm Mpc}$. This simulation was performed in a $\\Lambda$CDM universe with cosmological parameters: $\\Omega_\\Lambda = 0.75, \\Omega_M = 0.25, \\Omega_b = 0.045, h = 0.73, n=1$, and $\\sigma_8 = 0.9$, where the Hubble constant is characterised as $100h {\\rm km s^{-1} Mpc^{-1}}.$ These cosmological parameters are consistent with recent combined analyses from {\\it WMAP} data (Spergel et al 2003) and the 2dF galaxy redshift survey (Colless et al 2001), although the value for $\\sigma_8$ is a little higher than would perhaps have been desirable in retrospect. The spatial resolution is $5h^{-1}{\\rm kpc}$ everywhere inside the simulation volume.    \\subsection[]{Void and halo samples}    To explore the effect of the large--scale structure on the orientation of the dark matter haloes, we use a similar technique to the one adopted in analysing the orientation of disc galaxies  in real observations (Trujillo et al. 2006). To this aim we need, first, to find and characterise the radius of the large voids in the simulation. To speed up the process of finding the voids, we used as an initial guess for the position of the void centres the position of the most underdense particles in the box. To do this, for every dark matter particle in the Millennium Simulation we have estimated the local density by using a standard (Monaghan \\& Lattanzio 1985; Hernquist \\& Katz 1989) smoothing kernel averaging over 32 neighbours. We produced a list of the most underdense particles and then radially sorted surrounding particles in distance from these points. In this way the radius of a volume underdense by a factor of 10, centred on the most underdense particle, was calculated. All particles with densities less than 0.035 of the cosmic mean and further than $2h^{-1}{\\rm Mpc}$ from an even less dense particle were tried as prospective centres. After cleaning this catalogue  by removing  those positions that lay within a larger void and limiting the size to be larger than $5h^{-1}{\\rm Mpc}$ in radius, this technique produced 3024 potential void centre candidates. We estimate that this catalogue of potential void centres is greater than $99\\%$ complete, given that we tried over 200,000 random centres but only found 2 additional voids in the last 50,000 (i.e. voids with central densities close to our limit of 0.035).    Once a list of initial positions for searching for voids in the simulation had been created, we used these positions to search for the maximum spheres that are empty of haloes with mass larger than a given value. To conduct this search we follow  an algorithm  based on a modification of the one presented in Patiri et al. (2006a). The haloes were identified using a minimal spanning tree that links together particles with density exceeding 900 times the background density (Thomas et al. 1998). This is to focus on the core properties. In our case, we have used all haloes with masses larger than 8.6$\\times$10$^{11} h^{-1}$M$_{\\sun}$ (i.e. those haloes with more than 1000 particles). Using this mass cut we are selecting haloes that contain galaxies with stellar masses similar to the ones used in the observational data. To characterise the final position of the void centres and their radii, we populate a sphere of radius R=5 h$^{-1}$ Mpc, centered on each initial position, with 2000 random points. For every point in this sphere, we estimate the position of the closest 4 haloes lying in geometrically \"independent\" octants. Then we built the sphere defined by these 4 haloes. This is repeated for all the 2000 random points. As a characterization (position and radius) of the void we choose the biggest empty spherical region generated in the previous step.  It is important to stress that the position of the void defined in this way normally does not match the position of the initial guess.      To match the observations we select only those voids whose radius (as measured by the largest sphere that is empty of haloes greater than a given mass) is larger than 10 $h^{-1}$ Mpc. This cut produces 2932 voids in our box with a median radius of 14 $h^{-1}$ Mpc. To explore the alignments of the haloes we have concentrated only on haloes with masses 8.6$\\times$10$^{11} h^{-1}$M$_{\\sun}$ $<$M$<$ 8.6$\\times$10$^{12} h^{-1}$M$_{\\sun}$(i.e. we restrict the sample to galaxy--sized haloes) and located within shells beyond the surface of the voids.    In addition, we have also created a subsample of haloes which contain a disc--dominated galaxy at their centre. These haloes are expected to have had a relatively 'quiet' life and have suffered smaller non--linear effects after turnaround such that their spin properties should remain matched to those of the baryonic matter. To select these haloes we have used the semianalytic galaxy catalogue of Croton et al. (2006). We have created the subsample  by selecting the brightest dominant galaxy within 200 $h^{-1}$ kpc of our halo centre. We then select those haloes where the semianalytic code indicates that there is a galaxy brighter than M$_{K}$$<$-23 mag that has a bulge--to--total (B/T) ratio 0$<$B/T$<$0.4 (i.e. a Milky--Way like disc galaxy).   \\subsection[]{Halo shapes and spins}    To characterise the orientation of the haloes we have determined the orientation of their principal axes and their angular momentum vectors. The principal axes of the halo mass distribution are measured by diagonalising the inertia tensor,  \\begin{equation} I_{ij}=\\frac{1}{N_p}\\sum_{k=1}^{N_p}m_kx_{k,i}x_{k,j} \\end{equation}  where the sum is over all the particles in the halo N$_p$, and the coordinates are defined with respect to the centre of  mass of the halo of mass M$_h$. The resulting eigenvalues M$_h$a$^2$/5, M$_h$b$^2$/5, and M$_h$c$^2$/5 are sorted by size,  in descending order. The eigenvectors give the directions of the principal axes.    The angular momentum vector of the halo is given by:  \\begin{equation} \\bold{L}=\\sum_{k}m_{k}\\bold{r_k}\\times\\bold{v_k} \\end{equation}  where the sum is again over all the particles in the halo and $\\bold{r_k}$ and $\\bold{v_k}$ are the position and velocity of each particle relative to the centre of mass of the halo.   The uncertainty in the position of the inertia axes and of the angular momentum has been evaluated by comparing these quantities in the same haloes in two different runs that differ by a factor of 20 in resolution.  We found that we can measure the orientation of the angular momentum vector with an uncertainty (as provided by the full width half maximum of the ditribution) of $\\sim$14$^\\circ$ when the number of particles in the (low resolution) halo is larger than 1000. The uncertainties in the inertia axes are 13$^\\circ$ for the semimajor and semiminor axes and 20$^\\circ$ for the intermediate axis.  Once the angular momentum and the inertia vectors are evaluated, we estimate the cosine of the angle between those vectors and the vector joining the centre of the void with the centre of the halo $\\bold{R}$:   \\begin{equation} \\mu=\\cos\\theta=\\big(\\frac{\\bold{R} \\cdot \\bold{V}}{|\\bold{R}||\\bold{V}|}\\big) \\end{equation} ",
        "conclusions": "In this paper, we have shown that when haloes are selected in order to contain a Milky--Way like disc galaxy at their centres the angular momentum of the dark matter halo is oriented preferentially parallel to the surface of the voids. Observationally, the same alignment is detected using the baryonic matter (Trujillo et al. 2006). These two pieces of information are in agreement with the TTT prediction that both the dark and the baryonic matter component have conjointly acquired  their angular momentum before the moment of the turnaround. Interestingly, the signal detected in the real observation c=0.7$^{+0.1}_{-0.2}$ is higher than the one found in the simulations c=0.151$\\pm$0.046. This is to be expected taking into account that the signal in the dark matter haloes should be erased by non--linear  effects such as exchange of angular momentum between the haloes. Future work, consequently, should explore the strength of the alignment of the haloes at the moment of  turnaround. At that early epoch the strength of the signal should be as strong as the one measured using the baryonic component. Porciani et al. (2002b) shows hints that this should be the case by comparing the relation between the halo spin and the linear shear tensor at different redshifts from z=50 to z=0.    We have compared our work with previous analysis of the alignment of the inertia axes and angular momentum with their surrounding large--scale structure using different simulations. We concentrate first on the alignment of the angular momentum: using the void framework, neither Heymans et al. (2006) nor PA06 have found a signal of the alignment of the angular momentum within the shell of the voids. Their results are easily understood taking into account that no preselection of the haloes was done in either of these works and that they explored the signal in just one  shell of width 4 $h^{-1}$ Mpc. In fact, if we mimic their analysis we find c=-0.014$\\pm$0.012 (p=0.996$\\pm$0.005), in agreement with their findings. In the same shell, selecting those haloes with 0$<$B/T$<$0.45 produces c=0.030$\\pm$0.021 (p=1.012$\\pm$0.008). This is a factor of 5 weaker than the signal we find in the closest shell to the void surface. As we have seen in this work, the signal is only significant at the void surface, consequently, taking a wide shell can mask it. In the particular case of Heymans et al. (2006) their mass resolution limit could be another source of concern, since it is a factor of $\\sim$10 worse than the present simulation.    Comparison with other work is more complicated since the analysis of the alignments is done using a different scheme. BS05 measure the alignment of the angular momentum along filaments. Interestingly, they found that the angular momentum of galaxy mass haloes shows a weak tendency  to point along filaments, while those of group and cluster mass haloes show a very strong tendency to point perpendicular to the filaments. The significance and strength of their signal in terms of the c parameter is, however, not quantified. Consequently, the agreement between ours and their work can be done only qualitatively. Porciani et al. (2002b) explored the alignment of massive haloes (M$>$10$^{13}h^{-1}$M$_{\\sun}$) between their $final$ spin and the $initial$ shear tensor at the halo position. They found c=0.07$\\pm$0.04. The mass regime explored by these authors and their comparison between an initial and a final property of the haloes prevents us making a direct comparison between their and our measurements of the c parameter. Finally, in a recent paper, Arag\\'on--Calvo et al. (2006), using galaxy mass haloes find that the strength of alignment of their spins in walls is c=0.13$\\pm$0.02. This is in excellent agreement with the value reported in this paper.    If we focus  our attention on the alignment of the inertia axes, we find that our results are in good agreement with those of BS05, ACC06 and PA06. All these authors find that the tendency of the minor axis to lie perpendicular to large--scale filaments is the strongest of the alignments. It is interesting to note that the strength of these alignments seems to be dependent on the mass of the haloes, being stronger for the most massive (cluster--sized) ones. BS05 and ACC06 suggest that the different strength could be related to fact that most massive haloes receive a larger infall of matter from filaments. This could also help to explain the tendency of the angular momentum of the most massive haloes to be perpendicular to the filaments. In this sense, the cluster mass haloes would acquire most of their current angular momentum from major mergers along the filaments, whereas the angular momentum of the galaxy mass haloes will still have memory of the initial tidal fields.    The alignment of the haloes with their local large--scale structure is not only of interest to constrain models of galaxy formation, it could also be relevant to explain other observational features. For example, the tendency of satellite galaxies to avoid orbits that are coplanar with their host spiral galaxies (know as the ''Holmberg effect'') found in observations (Holmberg 1969; Zaritsky et al. 1997; Sales \\& Lambas 2004, but see Agustsson \\& Brainerd 2006) and in simulations (Zentner et al. 2005; Libeskind et al. 2006). This could be due to  the preferential accretion of satellites along filaments, that we have seen are preferentially aligned with the major axis of the host halo.   Finally, it is worth  pointing out the potential importance of the alignments we have discussed here to strong and weak lensing studies. In particular, these alignments could contribute to the weak lensing signal producing a shear--ellipticity correlation (Hirata \\& Seljak 2004). The degree of contamination that these alignments will produce in the weak lensing surveys should be explored in future work."
    },
    "0609/astro-ph0609303_arXiv.txt": {
        "abstract": "We describe the general problem of estimating black hole masses of AGN by calculating the conditional probability distribution of $M_{BH}$ given some set of observables. Special attention is given to the case where one uses the AGN continuum luminosity and emission line widths to estimate $M_{BH}$, and we outline how to set up the conditional probability distribution of $M_{BH}$ given the observed luminosity, line width, and redshift. We show how to combine the broad line estimates of $M_{BH}$ with information from an intrinsic correlation between $M_{BH}$ and $L$, and from the intrinsic distribution of $M_{BH}$, in a manner that improves the estimates of $M_{BH}$. Simulation was used to assess how the distribution of $M_{BH}$ inferred from the broad line mass estimates differs from the intrinsic distribution, and we find that this can lead to an inferred distribution that is too broad. We use these results and a sample of 25 sources that have recent reverberation mapping estimates of AGN black hole masses to investigate the effectiveness of using the C IV emission line to estimate $M_{BH}$ and to indirectly probe the C IV region size--luminosity ($R$--$L$) relationship. A linear regression of $\\log L_{\\lambda} (1549\\AA)$ on $\\log M_{BH}$ found that $L_{1549} \\propto M_{BH}^{1.17 \\pm 0.22}$. A linear regression also found that $M_{BH} \\propto L_{1549}^{0.41 \\pm 0.07} FWHM_{CIV}^2$, implying a C IV $R$--$L$ relationship of the form $R_{CIV} \\propto L_{1549}^{0.41 \\pm 0.07}$. Including the C IV line $FWHM$ resulted in a reduction of a factor of $\\sim 1/3$ in the error in the estimates of $M_{BH}$ over simply using the continuum luminosity, statistically justifying its use. We estimated $M_{BH}$ from both C IV and H$\\beta$ for a sample of 100 sources, including new spectra of 29 quasars. We find that the two emission lines give consistent estimates if one assumes $R \\propto L^{1/2}_{UV}$ for both lines. ",
        "introduction": "\\label{s-intro}  It is widely accepted that the extraordinary activity associated with quasars involves accretion onto a supermassive black hole (SMBH). Furthermore, the evidence that almost all massive galaxies host SMBHs has become quite convincing. It has been found that SMBH mass is correlated with the host galaxy's bulge luminosity \\citep[e.g.,][]{korm95,mag98,mclure01,marc03} as well as the stellar velocity dispersion \\citep[e.g.,][]{gebh00a,merr01,mclure02,trem02}. Because luminous quasars have been observed to reside in massive early-type galaxies \\citep{mclure99,kuk01,mcleod01,nolan01,perc01,dunlop03}, this implies that the evolution of spheroidal galaxies and quasars is intricately tied together \\citep[e.g.,][]{silk98,haehn00,adams01,merr04,dimatt05}. Therefore, understanding the cosmic evolution of SMBHs is an important task of modern astronomy.  Reverberation mapping \\citep{bland82,peter93} is often used to estimate SMBH mass, $M_{BH}$, in Type 1 active galaxies \\citep{wand99,kaspi00,peter04}. One of the principal advantages of this method is that it does not require high spatial resolution, but rather relies on the time lag between the continuum and emission line variability. Under the assumption that the broad line region (BLR) is in Keplerian motion, the time lag is combined with the line width to give an estimate of SMBH mass. This, in principal, makes it immediately applicable for both faint and distant quasars. Although there are many potential systematic uncertainties in the technique \\citep{krolik01}, there has been good agreement between SMBH masses inferred from reverberation mapping and those inferred from the stellar velocity dispersion \\citep{gebh00b,ferr01,nelson04,onken04}.  Unfortunately, reverberation mapping requires long-term intensive monitoring, which is not practical for large samples. In addition, the long time scales for variability in bright high redshift sources make reverberation mapping for these sources unfeasable. Fortunately, a correlation has been found that relates the BLR size and monochromatic continuum luminosity \\citep[the $R$--$L$ relationship, ][]{kaspi00,kaspi05}, making it possible to estimate virial masses from single-epoch spectra using the H$\\beta$ width \\citep{wand99}. In addition, \\citet{wu04} find a relationship between BLR size and the luminosity of H$\\beta$, but \\citet{woo02a} do not find any correlation between SMBH mass and bolometric luminosity. H$\\beta$ is redshifted into the near-infrared at $z \\sim 0.9$, making it difficult to observe from the ground for large samples. \\citet{vest02} and \\citet{bhmmgii} have argued for the use of C IV and Mg II, respectively, to estimate virial masses from single-epoch spectra, allowing quasar SMBH masses to be estimated at high $z$ from the ground. Many studies have exploited these results and estimated SMBH masses for large samples of quasars \\citep[e.g.,][]{bech03,corb03,warn03,vest04,mclure04a,mclure04b}. \\citet{diet04} have found that the estimates based on H$\\beta$ and C IV agree well for high $z$ quasars, whereas Mg II-based estimates are typically a factor of $\\sim 5$ times lower. In constrast to this, \\citet{shem04} and \\citet{bask05} argue that C IV does not give as accurate of an estimate of SMBH mass as H$\\beta$.  Previous methods that utilize single-epoch spectra have relied on empirical linear relationships that estimate the BLR size for a given source luminosity, and then use this estimate of $R$ in the virial relationship. Although this will give good estimates for $R$, this may give less efficient estimates for SMBH mass. The reason for this is that the luminosity may also be correlated with $M_{BH}$ in a manner independent of the $R$--$L$ relationship, such as through the accretion process. By only using the luminosity as a proxy for $R$ when estimating $M_{BH}$, one is ignoring the additional information of $M_{BH}$ that is contained within $L$. In other words, the standard broad line estimates are based on the distribution of velocities at a given luminosity and $M_{BH}$, as folded through the $R$--$L$ correlation and under the assumption of Keplerian motion. However, the distribution of black hole masses at a given luminosity and line width is not simply obtained by inverting the distribution of velocities at a given luminosity and black hole mass, unless $L$ and $M_{BH}$ are statistically independent.  Recent studies have found evidence for a correlation between luminosity and black hole mass \\citep[e.g.,][]{corb03, netzer03, peter04}, suggesting that black hole mass and luminosity are not statistically independent.  However, in contrast to other studies, \\citet{woo02a} have argued that there is no significant correlation between bolometric luminosity and $M_{BH}$. If a correlation between $M_{BH}$ and $L$ exists, then we can combine the $M_{BH}$--$L$ correlation with the broad line mass estimates to obtain, on average, more accurate estimates of $M_{BH}$ for a given luminosity and line width.  In this work, we outline a formalism that allows one to estimate the probability distribution of an AGN's black hole mass, given some set of observables. We focus on the special case of estimating $M_{BH}$ given some monochromatic luminosity, $L_{\\lambda}$, and the width of an emission line. We also search for other parameters of the C IV line that may contribute additional information of $M_{BH}$, thus decreasing the uncertainty in the estimated $M_{BH}$. Although it is possible to include other quasar properties, such as radio loudness, X-ray loudness, or variability \\citep[e.g.,][]{xie05, oneill05, pessah06}, for simplicity we have chosen to only include parameters that may be measured from a single spectrum containing C IV. This allows $M_{BH}$ to be estimated using only one spectrum, and it is thus not necessary to compile observations from several different spectral regions.  We use the formalism developed here to justify using the C IV line width in estimating $M_{BH}$, and attempt to indirectly infer the C IV $R$--$L$ relationship. We have chosen the C IV line because it is readily observable from the ground over a wide range in redshift ($1.5 \\lesssim z \\lesssim 4.5)$, is less effected by blends with iron and other lines, has shown to give consistent mass estimates with H$\\beta$ \\citep{vest02}, is commonly employed to estimate SMBH mass, and archival UV spectra are available for most of the sources with reverberation-based masses \\citep{peter04}. Despite its common usage, the $R$--$L$ relationship for C IV is mostly unexplored, as there are only seven data points with reliable C IV reverberation mapping data \\citep{peter05}. Often it is assumed that C IV BLR size has the same dependence on luminosity as H$\\beta$ \\citep[e.g.,][]{vest02,netzer03,vest06}. Another possibility is to assume $R \\propto L^{1/2}$, as predicted from simple photoionization theory or if the BLR size is set by the dust sublimation radius \\citep{netz93}; this was done by \\citet{wand99} and \\citet{shields03} for the H$\\beta$ line. \\citet{peter05} performed a linear regression using five AGN with a total of seven data points and find $R \\propto L_{UV}^{0.61 \\pm 0.05}$, similar to that for the Balmer lines\\footnote{This value appears in an erratum to this paper}. Unfortunately, this result is almost entirely dependent on the inclusion of NGC 4395, the least luminous known Seyfert 1 galaxy, since the data points for the other AGN are clustered around $\\lambda L_{\\lambda}(1350\\AA) \\approx 10^{43.75}\\ {\\rm ergs\\ s^{-1}}$. \\citet{vest06} found that using the C IV line to calculate black hole masses assuming $R \\propto L_{UV}^{0.53}$ gave results consistent with masses obtained by reverberation mapping.  The layout of the paper is as follows. In \\S~\\ref{s-bhmass} we describe the general problem of estimating SMBH mass from single-epoch spectra, and in \\S~\\ref{s-sample} we describe our two samples. The first sample is a set of 25 quasars with black holes from reverberation mapping and archival UV spectra, and the second sample is a set of 100 quasars for which we have spectra containing both the H$\\beta$ and C IV emission lines. In \\S~\\ref{s-lineparams} we describe the method we employ to estimate the emission line parameters. In \\S~\\ref{s-varselect} we test if any other UV or C IV parameters contribute useful information about $M_{BH}$.  In \\S~\\ref{s-mlrel} we use the sample with reverberation mapping data to investigate the $M_{BH}$--$L$ relationship, and in \\S~\\ref{s-civrl} we test if including the C IV line $FWHM$ is preferred by this sample and investigate the nature of the C IV $R$--$L$ relationship.  In \\S~\\ref{s-hbeta} we use our larger sample to compare estimates of $M_{BH}$ obtained from both single-epoch H$\\beta$ and C IV. We summarize our results in \\S~\\ref{s-summary}.  In this work we adopt the WMAP best-fit cosmological parameters \\citep[$h=0.71, \\Omega_m=0.27, \\Omega_{\\Lambda}=0.73$,][]{wmap}. We will use the common statistical notation where a point estimate of a quantity is denoted by placing a $\\hat{}$ above it, e.g., $\\hat{M}_{BH}$ would be an estimate of $M_{BH}$.  We will also commonly refer to the bias of an estimate, where the bias of $\\hat{M}_{BH}$ is $Bias = E(\\hat{M}_{BH}) - M_{BH}$. Here, $E(\\hat{M}_{BH})$ is the expectation value of $\\hat{M}_{BH}$. An unbiased estimate of $M_{BH}$ is one such that $E(\\hat{M}_{BH}) = M_{BH}$. ",
        "conclusions": "\\label{s-summary}  In this work we have undertaken a statistical investigation of the best method to estimate SMBH mass based on the single-epoch C IV line and AGN continuum at $1549\\AA$. The main conclusions are : \\begin{enumerate} \\item Estimating AGN black hole masses from a single-epoch spectrum is a considerably different problem than in the reverberation-mapping case. Because of this, one is not estimating $M_{BH}$ directly, but calculating a probability distribution of $M_{BH}$ given the observed luminosity, line width, and redshift, $p(m|l,v,z) \\propto p(v|l,m,z) p(l|m,z) p(m|z)$. \\item Combining the information in $L$ from both an intrinsic $M_{BH}$--$L$ correlation and the $R$--$L$ relationship results in improved black hole mass estimates. However, because of the current systematic uncertainties in the $M_{BH}$--$L$ relationship, estimates based on its inclusion should be viewed with caution. In addition, incorporating information from the intrinsic distribution of $M_{BH}$ also results in more accurate estimates, on average. \\item The distribution of $M_{BH}$ inferred from the broad line mass estimates, $\\hat{m}_{BL}$, is broader than the intrinsic distribution of $M_{BH}$, as $\\hat{m}_{BL}$ are estimates of $m$ contaminated by `measurement error.' In addition, it is necessary to estimate or assume $p(l|m,z)$ in order to estimate a survey's completeness in $M_{BH}$. \\item The best estimates of $M_{BH}$ based on the C IV emission line are obtained using the UV continuum luminosity and $FWHM$. We find evidence that the C IV line shift, centroid, $EW$, and spectral slope of the UV continuum do not contribute any additional information to estimating $M_{BH}$ for a given $L_{1549}$ and $FWHM_{CIV}$. \\item Using the reverberation mapping sample, we find an $M_{BH}$--$L$ relationship of the form \\begin{equation} l = 35.72(\\pm 1.67) + 1.17(\\pm 0.22) m \\nonumber. \\end{equation} The intrinsic scatter about this relationship is $\\sigma_l \\approx 0.61$ dex. Combining mass estimates based on this relationship with the broad line mass estimates results in a reduction in the statistical error of $\\approx 18\\%$. \\item We estimate a C IV $R$--$L$ relationship of the form $R \\propto L_{1549}^{0.41 \\pm 0.07}$. This is consistent with an $R$--$L$ relationship of the form $R \\propto L^{1/2}$ predicted from simple photoionization physics or if the BLR size is set by the dust sublimation radius. The scatter about the C IV $R$--$L$ relationship inferred from the reverberation mapping estimates of $M_{BH}$ is $\\sigma_r \\approx 0.35$ dex. \\item A broad line estimate of $m$ based on the C IV emission line may be calculated as \\begin{equation} \\log \\hat{M}_{BH} / M_{\\odot} = \\hat{m}_{CIV} = -17.82 (\\pm 2.99) + 0.41 (\\pm 0.07) \\log \\lambda L_{\\lambda}(1549\\AA) + 2 \\log FWHM_{CIV}. \\label{mciv} \\end{equation} Here, $\\lambda L_{\\lambda}$ is in units of ${\\rm ergs\\ s^{-1}}$ and $FWHM$ is in units of ${\\rm km\\ s^{-1}}$. The correlation between regression coefficients is $Corr(\\hat{m}^{CIV}_0, \\hat{\\theta}_l) = -0.9996$. The intrinsic scatter in $m$ about Equation (\\ref{mciv}) is $\\approx 0.40$ dex. \\item The C IV- and H$\\beta$-based mass estimates are consistent if one assumes $R \\propto L_{UV}^{1/2}$ for both emission lines. The two emission lines give estimates of $M_{BH}$ with comparable accuracy. \\item We find evidence that the $R$--$L$ relationships for C IV and H$\\beta$ are similar in their dependence on $L_{UV}$. For both, $R \\propto L_{UV}^{1/2}$ is consistent with the available data. Also, the scatter in the C IV $R$--$L$ relationship is a factor $\\lesssim 1.6$ larger. This suggests that the C IV broad line region gas has a larger dispersion in its properties (e.g., density, ionization parameter, inclination) than the H$\\beta$ BLR gas. \\end{enumerate}"
    },
    "0609/astro-ph0609827_arXiv.txt": {
        "abstract": "The origin of the particles in the highest energy end of the cosmic ray spectrum is discussed in the context of the wider problem of the origin of the whole cosmic radiation as observed at the Earth. In particular we focus our attention on the acceleration problem and on the transition from galactic to extragalactic cosmic rays. ",
        "introduction": "There are mainly two reasons for the attention shown by the community for the origin of the highest energy cosmic ray spectrum, say above $10^{19}$ eV: on the one hand most acceleration processes need to be pushed to their extreme (or beyond) in the attempt to reach such energies. On the other hand, even in the presence of numerous nearby sources, the spectrum is expected to show a suppression, known as the GZK feature (Greisen 1966; Zatsepin \\& Kuzmin 1966) above the threshold for photopion production on the photons of the cosmic microwave background (CMB). The suppression should become evident around $\\sim 10^{20}$ eV. Here we review some recent ideas put forward to address both points mentioned above.  The acceleration of cosmic rays to the observed energies is a problem even in supernova remnants, as possible sources of galactic cosmic rays, therefore it is natural to try and learn from that instance. As for the presence of the GZK feature we summarize the results of current experiments and try to address the issue of what future experiments could be expected to observe. However, in order to understand the origin of extragalactic cosmic rays and more specifically of UHECRs, it is crucial to understand which cosmic rays are in fact extragalactic. The two main lines of thought in this respect will be summarized and discussed: in the so-called {\\it ankle scenario} (Hillas 2005; Wibig \\& wolfendale 2005) the transition takes place around $10^{19}$ eV where a steep galactic spectrum encounters the flat spectrum of extragalactic cosmic rays. In the {\\it dip scenario} (Berezinsky et al. 2002, 2005) the transition takes place at energies roughly one order of magnitude lower. The case of a mixed chemical composition of extragalactic cosmic rays (Allard et al. 2005a; Allard et al. 2005b) will also be discussed.  The paper is structured as follows: in \\S \\ref{sec:acceleration} we discuss the problem of acceleration and some recent findings on magnetic field amplification in the acceleration region. In \\S \\ref{sec:trans} we discuss the current ideas on the transition between galactic and extragalactic cosmic rays. In \\S \\ref{sec:GZK} we describe the current status of the observations of the end of the cosmic ray spectrum and of the anisotropies in such energy range. We conclude in \\S \\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  We discussed three of the crucial issues in the investigation of the origin of UHECRs: 1) Particle acceleration in astrophysical sources; 2) the transition from galactic to extragalactic cosmic rays; 3) the spectrum and anisotropies of UHECRs.  The theoretical investigation of particle acceleration at shock fronts is recently experiencing a boost, due to some interesting calculations of the nonlinear effects of particle acceleration (e.g. dynamical reaction of the accelerated particles and self-generation of scattering) and some observational results on the spectra of single chemical components (e.g. measurements from KASCADE). These lines of investigation are leading to some consensus on the fact that protons may be accelerated by sources in the Galaxy up to energies $\\sim 10^{16}$ eV. A rigidity dependent argument would imply that the spectrum of iron extends to energy $\\sim 3\\times 10^{17}$ eV. The flux of galactic cosmic rays should rapidly decay above this energy, so that a transition to extragalactic cosmic rays is expected to take place at energy $\\sim 10^{18}$ eV. This conclusion is at odds with the traditional interpretation of the ankle (Wibig and Wolfendale 2005) as the transition feature, while it appears to agree with the prediction of the dip scenario (Berezinsky et al. 2006). A similar energy for the transition energy is however inferred in the case of a mixed composition of extragalactic cosmic rays, as discussed in (Allard et al. 2005a; Allard et al. 2005b). Whether the dip or the mixed composition scenario best describes the observations will most likely be decided when an accurate measurement of the chemical composition between $10^{17}$ and $10^{19}$ eV will become available.  At larger energies, a crucial issue is represented by the measurement of the spectrum at energies around $10^{20}$ eV, where the GZK suppression has long been searched and missed. AGASA and HiRes spectra are in a mild contradiction which may well be due to statistical fluctuations and a systematic error in the energy determination with the two different techniques used. The Pierre Auger Observatory should soon settle this issue.  At sufficiently high energies the magnetic field of the Galaxy is not expected to bend the trajectories of high energy particles by more than the angular resolution of the operating cosmic ray experiments. This may lead to the identification of small scale anisotropies, flagging the presence of discrete sources of UHECRs. A signal of this type was found by AGASA, though its statistical significance was later questioned (Finley \\& Westerhoff 2004; De Marco, Blasi \\& Olinto 2006b). The level of SSA detected by AGASA would correspond to a source density $\\sim 10^{-5}\\rm Mpc^{-3}$ (Blasi and De Marco 2004). This density would roughly correspond to the density of active galactic nuclei. This has inspired numerous searches for correlations between the directions of arrival of UHECRs and the locations of AGNs. Despite several claims of positive correlations, a definite answer is still missing."
    },
    "0609/astro-ph0609016_arXiv.txt": {
        "abstract": "Using the multiwavelength photometric and spectroscopic data covering the Chandra Deep Field South obtained within the Great Observatories Origins Deep Survey, we investigate the rest-frame UV properties of galaxies to $z\\sim 2.2$, including the evolution of the luminosity function, the luminosity density, star formation rate (SFR) and galaxy morphology. We find a significant brightening ($\\sim 1$~mag) in the rest-frame 2800\\AA~characteristic magnitude ($M^\\ast$) over the redshift range $0.3\\lsim z\\lsim1.7$~and no evolution at higher redshifts. The rest-frame 2800\\AA~luminosity density shows an increase by a factor $\\sim 4$~over the redshift range investigated. We estimate the star formation rate density to $z\\sim$2.2 from the 1500\\AA~and 2800\\AA~luminosities. When no correction for extinction is made, we find that the star formation rate derived from the 2800\\AA~luminosity density is almost factor two higher than that derived from the 1500\\AA~luminosities. Attributing this difference to differential dust extinction, we find that an E(B--V)=0.20 results in the same extinction corrected star formation rate from both 1500\\AA~and 2800\\AA~luminosities. The extinction corrected SFR is a factor $\\sim 6.5$~($\\sim 3.7$) higher than the uncorrected SFR derived from 1500\\AA~(2800\\AA) luminosity. We investigate the morphological composition of our sample by fitting S\\'{e}rsic profiles to the $HST$~ACS galaxy images at a fixed rest-frame wavelength of 2800\\AA~at $0.5\\lsim z\\lsim2.2$. We find that the fraction of apparently bulge-dominated galaxies (S\\'{e}rsic index $n>2.5$) increases from $\\sim$10\\%~at $z\\sim 0.5$~to $\\sim$30\\%~at $z\\sim 2.2$. At the same time, we note that galaxies get bluer at increasing redshift. This suggests a scenario where an increased fraction of the star formation takes place in bulge-dominated systems at high redshift. This could be the evidence that the present day ellipticals are a result of assembly (i.e., mergers) of galaxies at $z\\gsim$1. Finally, we find that galaxy sizes for a luminosity selected sample evolves as $r_h\\propto (1+z)^{-1.1}$~between redshifts $z=2.2$~and $z=1.1$. This is consistent with previous measurements and suggests a similar evolution over the redshift range $0\\lsim z \\lsim$6. ",
        "introduction": "Over the last decade, a large number of studies have measured the evolution of cosmic star formation rate (SFR) in the range $0 < z < 6$, using different diagnostics (Steidel et al. 1995; Madau et al. 1996; Chary \\& Elbaz 2001; Giavalisco et al. 2004a; Schiminovich et al. 2005; Takeuchi et al. 2005; see also Hopkins (2004) for a compilation of the SFR from a large number of multiwavelength surveys). These studies show a rapid increase in the SFR to $z\\sim 1-2$, beyond which it flattens or turns over. While most studies are in broad agreement on the shape of the SFR, the absolute normalization is still uncertain, mainly due to unknown extinction in the UV to optical bands and systematic effects depending on star formation diagnostic, selection of the star-forming population and the adopted flux limit of the surveys used (e.g., Hopkins \\& Beacom 2006). A further question in these investigations concerns the morphology of star forming galaxies and how it evolves with redshift. These studies require complete surveys of star forming galaxies with known redshifts, SFRs and morphologies.  Using the high spatial resolution and sensitivity of the Advance Camera for Surveys (ACS) on-board the Hubble Space Telescope (HST), a number of large, multi-waveband and deep surveys have now been completed, with the two deepest being the Great Observatories Origins Deep Survey (GOODS; Giavalisco et al. 2004b) and the Hubble Ultra-Deep Field (HUDF; Beckwith, S. V. W., et al. 2006, in preparation). The combined wavelength coverage, spatial resolution and depth of these surveys allow accurate measurements of the SFR, morphology and size of these galaxies to faint flux levels ($m_z\\gsim 25$) and high redshifts ($z\\sim 6$).  Recently, the ultraviolet morphological properties of star forming galaxies in the GOODS and HUDF were explored by a number of investigators. Using WFPC2 observations of the GOODS parallel fields at F300W, de Mello et al. (2006) studied rest-frame UV properties of galaxies to $z\\sim 1.5$. They found that their sample included all major morphological types, with compact and peculiar morphologies becoming relatively more abundant at higher redshift ($z\\gsim 0.7$). Furthermore, Lotz et al. (2006) studied the UV morphologies of star-forming galaxies in the GOODS-S at $z\\sim 1.5$~and $z\\sim 4$ and found no significant differences between the galaxy morphologies at these redshifts. Conselice et al. (2004) used the $HST$~ACS observations of GOODS-S to identify populations of luminous diffuse objects and luminous asymmetric objects at $1<z<2$~which they argue are the progenitors of todays normal disk and elliptical galaxies.  Ravindranath et al. (2006) find that among their sample of bright LBGs at $z>2.5$, about 40\\% have light profiles that can be approximated by an exponential profile as seen in disks, while 30\\% have close to r$^{1/4}$ profiles as seen spheroids. However, they note that even these galaxies do show clumpy or faint asymmetric features characterestic of tidal interaction or minor merger. The method of Gini coefficients used by Lotz et al. (2006) is more sensitive to merger-like features. Lotz et al. (2006) find that among their sample of LBGs at $z=4$, 30\\% have relatively undisturbed spheroid-like morphologies, about 10-25\\% are major mergers. About 50\\% are like exponential disks or have minor mergers. Except for small differences arising from the different methods that are employed, the results from Ravindranath et al. (2006) are in good agreement with Lotz et al. (2006).  Finally, the evolution of galaxy sizes has also been investigated using the GOODS and HUDF data sets by Bouwens et al. (2004a, 2004b, 2005) and Ferguson et al. (2004). In this paper, we study UV properties of a complete sample of star-forming galaxies in the GOODS-S at rest-frame 1500\\AA~and 2800\\AA~ wavelengths. The investigation includes a study of the evolution of the luminosity function, the luminosity density, the SFR, morphology, and galaxy size. In \\S 2 we describe our data. This is followed by results on the luminosity functions and luminosity densities (\\S 3), the SFR (\\S 4), and morphology and size (\\S 5). We conclude and summarize our results in \\S 6.  Throughout this paper we use $\\Omega_{\\Lambda}=0.7$, $\\Omega_M=0.3$~ and $H_0=70$ km s$^{-1}$ Mpc$^{-1}$. Magnitudes are in the AB system. ",
        "conclusions": "We use the GOODS CDF-S optical and near-IR observations to derive the evolution of the rest-frame 1500\\AA~and 2800\\AA~luminosity functions, luminosity densities and star formation rates to $z\\sim 2.2$. Taking advantage of the high resolution $HST$~ACS imaging, we also derive the evolution the rest-frame UV morphological properties. We find: \\begin{itemize}  \\item{There is a strong evolution in the UV characteristic magnitude with redshift. We find a brightening in $M^*_{2800}$~by $\\sim 1$~mag between $z\\sim 0.3$~and $z\\sim 1.7$. At both 1500\\AA~and 2800\\AA, we find no significant evolution at $z\\gsim 1.7$.}  \\item{The rest-frame 2800\\AA~luminosity density increases by a factor $\\sim 4$~over the redshift range $0.3<z<1.7$, while at rest-frame 1500\\AA~the increase is a factor $\\sim 2$~between $z\\sim 1.1$~and $z\\sim 1.7$. The evolution flattens in both wave-lengths at $z\\gsim 1.7$.}  \\item{We find that the uncorrected SFR derived from 2800\\AA~luminosity is a factor $\\sim 1.7$~higher than the SFR derived from 1500\\AA~luminosity. Assuming that the difference is due to differential dust extinction, we find that E(B--V)=0.20$\\pm0.05$~ (using a Calzetti attenuation law and a standard Salpeter IMF) can explain the difference in the uncorrected SFRs.}  \\item{The extinction corrected SFR is a factor $\\sim 6.5$ ($\\sim 3.7$) higher than the uncorrected SFR calculated from 1500\\AA~(2800\\AA) luminosity densities. These corrections lead to a $\\sim 50\\%$~higher SFR compared to the SFRs derived from dropout samples which assume a smaller extinction correction (e.g., Giavalisco et al. 2004a).}  \\item{The SFR we derive shows a steep increase out to $z\\sim 1.5-2$. Comparing with the local value from Wyder et al. (2005), we find that the SFR is a factor $\\gsim 5$~higher at $z\\sim 1.5$~compared to $z=0$. We note, however, that our measurements are consistent with a flat rate at $z\\gsim$1.7.}  \\item{We find that the fraction of galaxies with high S\\'{e}rsic index ($n>2.5$), indicating a bulge-like morphology, increases at higher redshift. At $z\\sim 0.5$, we find that $\\sim 10\\%$~of the galaxies have bulge morphology, while the fraction is $\\sim 30\\%$~at $z\\sim 2.2$. At the same time, the mean color of the galaxies gets bluer at higher redshift. This suggests that an increased fraction of the star formation takes place in objects with bulge-like morphology at high redshift. This could be the sign of the formation of todays elliptical population via mergers and infall at redshifts $z\\gsim$1. At least part of the high redshift bulge systems may, however, evolve into the bulges of todays population of old spirals.}  \\item{Investigating the luminosity-size redshift evolution of disk galaxies at rest-frame 2800\\AA~over the redshift range $1.1<z<2.2$, we find that galaxies become larger at lower redshift. Our results are consistent with trends found at both higher and lower redshift and suggests a size evolution $r_h\\propto(1+z)^{-1.1}$ between redshifts $z\\sim 6$~and $z=0$. This is somewhat steeper then the simple scaling law $R_s\\propto H(z)^{-2/3}$, expected in the hierarchical model for a sample selected by mass. This suggests a stellar mass-to-light ratio that decreases at higher redshift.}  \\item{We find a somewhat smaller evolution $r_h\\propto (1+z)^{-0.98}$~ when selecting galaxies by mass (where mass is indirectly derived from the galaxies rest-frame colors and absolute magnitudes). This milder evolution supports the change in stellar mass-to-light ratio with redshift. Compared to measurements in the literature, we find a steeper evolution in the size evolution for the mass selected sample. However, investigations are made at different rest-frame bands and redshift ranges, and are therefore difficult to directly compare.} \\end{itemize}"
    },
    "0609/astro-ph0609220_arXiv.txt": {
        "abstract": "The origin of blue straggler stars (BSS) in globular clusters (GCs) is still not fully understood: they can form from stellar collisions, or through mass transfer in isolated, primordial binaries (PBs). In this paper we use the radial distribution of BSS observed in four GCs (M3, 47~Tuc, NGC~6752 and $\\omega{}$ Cen) to investigate  which formation process prevails. We find that both channels co-exist in all the considered GCs. The fraction of mass-transfer (collisional) BSS with respect to the total number of BSS is around $\\sim 0.4-0.5$ ($\\sim 0.5-0.6$) in M3, 47~Tuc, and NGC~6752. The case of $\\omega$~Cen is peculiar with an underproduction of collisional BSS. The relative lack of collisional BSS in $\\omega$~Cen can be understood if mass segregation has not yet driven to the core a sizeable number of PBs, which dominate stellar collisions through three- and four-body processes. The spatial distribution of BSS provides strong hints to their origin: the BSS in the cluster outskirts form almost exclusively from mass transfer in PBs, whereas the BSS found close to the cluster core most likely have a collisional origin. ",
        "introduction": "Blue straggler stars (BSS) are stars lying above and blue-ward of the turn-off in the color-magnitude diagram (CMD) of a star cluster.  At least two different processes have been proposed to explain their formation (Fusi Pecci et al. 1992; Bailyn 1995; Bailyn and Pinsonneault 1995; Procter Sills, Bailyn \\& Demarque 1995; Sills \\&{} Bailyn 1999; Sills et al.  2000; Hurley et al. 2001). The first, dubbed the {\\it mass-transfer scenario}, suggests that BSS are generated by primordial binaries (hereafter PBs) that evolve mainly in isolation until they start mass-transfer and possibly coalesce (McCrea 1964; Carney et al. 2001).  The second, known as the {\\it collisional scenario}, states that BSS are the product of a merger between two main sequence stars (MSs) in a dynamical interaction that involves a MS-MS collision,  most likely in  a binary-MS encounter (Davies, Benz \\& Hills 1994; Lombardi et al. 2002). The collisional BSS (COL-BSS) differ kinematically from the mass-transfer BSS (MT-BSS), since they are believed to acquire kicks due to dynamical recoil. In both these hypotheses, the resulting BSS have  mass exceeding the turn-off mass of the cluster and are fueled by hydrogen thanks to the mixing of the hydrogen-rich surface layers of the two progenitor stars.   \\begin{figure} \\epsfxsize=8.truecm \\epsfysize=8.truecm \\epsfbox{fig1c.eps} \\label{fig1} \\vspace{-0.5truecm} \\caption {Comparison between the observed surface star density profile of the considered GCs (solid line and full circles) and the adopted multi-mass King model (dashed line). The units on the y-axis are number of stars per arcsec$^2.$ The observed profiles are taken from Ferraro et al. (1997a) for M3, Ferraro et al. (2006a) for $\\omega$~Cen, Mapelli et al. (2004) for 47~Tuc, Ferraro et al. (2003b) for NGC~6752. } \\end{figure}   The two aforementioned scenarios do not necessarily exclude each other and might co-exist within the same star cluster (Ferraro et al. 1993; Davies, Piotto \\& De Angeli 2004; Mapelli et al. 2004).  However, until recently, it has been difficult to estimate the relative importance of the two formation channels in a given globular cluster (GC).  A new approach to solve this problem has been presented by Mapelli et al. (2004, hereafter paper~I), who revised and upgraded an original attempt by Sigurdsson, Davies \\& Bolte (1994). The basic input for the new method is the observation of the shape of the BSS radial distribution in a star cluster. After the pioneering works on M3 (Ferraro et al. 1993, 1997a) and M55 (Zaggia, Piotto \\& Capaccioli 1997), accurate determinations of the spatial distribution of BSS are now being obtained for an increasing sample of GCs (Ferraro et al. 2004 for 47~Tuc; Sabbi et al. 2004 for NGC~6752; Ferraro et al. 2006a for $\\omega$~Cen; Warren, Sandquist \\& Bolte 2006 for M5).  A very interesting result is emerging from these observations: BSS display a clear tendency to follow a bimodal spatial distribution, with a peak in the core, decreasing at intermediate radii and rising again at larger radii.  In paper~I, Mapelli et al. focused on 47~Tuc and showed that the bimodal distribution can be reproduced assuming a  suitable combination of the two proposed mechanisms for forming BSS. In particular, it was suggested (paper~I) that the BSS in the core of 47~Tuc are mainly COL-BSS, while the external BSS are  MT-BSS.  This result agrees with studies of the BSS luminosity function in 47~Tuc (Bailyn \\& Pinsonneault 1995; Sills \\& Bailyn 1999; Sills et al. 2000; Ferraro et al. 2003a; Monkman et al. 2006).   In this paper, we explore two other GCs, whose BSS have a bimodal radial distribution. The aim is to test whether the conclusions drawn for 47~Tuc can be extended to GCs having different mass, concentration, central density, and velocity dispersion with respect to 47~Tuc. In particular, we study whether the position of a BSS in a GC can reliably be regarded as a signature for its origin.  We also consider the case of $\\omega$~Cen, whose radial distribution of BSS is flat. We follow the same procedure of paper~I, carrying out dynamical simulations of BSS evolved in the gravitational potential well of their associated cluster. Details about the simulations are illustrated in Section 2; the sample of the investigated clusters is presented in Section 3. The results for M3, NGC~6752 and 47~Tuc are discussed in Section 4, whilst Section 5 is devoted to the outlier, $\\omega$~Cen. In Section 6 we  discuss our findings in comparison with the findings of Davies et al. (2004). Finally in Section 7 we present our conclusions. ",
        "conclusions": "We have explored whether the radial distribution of BSS observed in our sample of GCs can be reproduced with a model which predicts two different classes of progenitors for the BSS: {\\it (i)} PBs evolved in isolation and {\\it (ii)} binary (or single) stars which underwent collisions in the high density environment of the GC core.  Fig.~2 shows the results obtained for M3, 47~Tuc, NGC~6752 and $\\omega$~Cen.  In all cases, our adopted model provides a statistically acceptable fit to the observations.  The best fit value of the fraction $\\eta{}_{\\rm MT}$ of MT-BSS with respect to the total number of BSS [i.e. $\\eta{}_{\\rm MT}\\equiv{}N_{\\rm MT-BSS}/(N_{\\rm MT-BSS}+N_{\\rm COL-BSS})$, where $N_{\\rm MT-BSS}$ and $N_{\\rm COL-BSS}$ are the number of mass-transfer and collisional BSS, respectively] % is reported in Table 2, as well as the reduced $\\chi^2$ of the best fit solution.    In order to determine the best fit solution, we have created a one-dimensional grid of simulations for each cluster, allowing the only free parameter in our model\\footnote{Other possible parameters (such as the initial velocity of BSS) were found to have only minor impact on the results (see Section~2) and have been set to a fixed value in all the considered runs. Also, the cluster properties are not free parameters in the fitting procedure, because they have been chosen {\\it a priori} (see Section 3) and kept fixed in all simulations of a given cluster.~An~exception was made only for the peculiar case of the velocity dispersion in NGC~6752 (see Section 4.3).}, $\\eta{}_{\\rm MT}$, to vary between 0 and 1 with an initial step of $0.2.$ Then, the grid was iteratively refined down to a step of $0.01$ for the values surrounding the best fit.  The values ${\\tilde{\\chi{}}}^2$ of the reduced $\\chi{}^2$ were derived by comparing each simulation with the observed data and their uncertainties. The range of values for $\\eta{}_{\\rm MT}$ in the second column of Table~2 is calculated at $2\\,{}\\sigma{}$ level,  i.e. by selecting all the simulations for which $\\chi{}^2\\le\\chi{}^2_{\\rm best}+4,$ where $\\chi{}^2_{\\rm best}$ is the value for the best fit solution.  Inspection of Table~2 reveals a clear dichotomy between the cases of 47~Tuc, M3 and NGC~6752 and that of $\\omega$~Cen. The $2\\,{}\\sigma{}$ intervals of confidence for $\\eta{}_{\\rm MT}$ for the first three GCs are in agreement with an almost even distribution of MT-BSS and COL-BSS (a slight predominance of COL-BSS being indicated by the best fit models). Instead, in the best fit model of $\\omega$~Cen only 14\\% of BSS are COL-BSS and the $2\\,{}\\sigma{}$ interval of confidence is compatible with $\\eta{}_{\\rm MT}=1.$   In the following, we will discuss the case of these three clusters separately, whereas $\\omega$~Cen will be examined in Section 5.     \\subsection{M3}  In the case of M3, our findings contrast earlier suggestions by Sigurdsson et al. (1994). They hypothesized that the BSS observed in the peripheral regions of M3 could be COL-BSS born initially in the core and then ejected in the periphery because of their natal kick. Indeed, our dynamical simulations show that COL-BSS are ejected from the entire cluster (if the kick is too high), or sink back to the core in $\\simlt 1$ Gyr, if the kick is low enough to retain them in the cluster potential well. Hence, they cannot account for the peripheral BSS. Instead, MT-BSS have in general much more circular orbits than the COL-BSS and can remain in the peripheral regions of the GC for a time comparable to the typical lifetime of a BSS.  \\subsection{47~Tuc}  As for 47~Tuc, we confirm the main outcome from paper~I: i.e. most of the central BSS of 47~Tuc are born from collisions and all the peripheral BSS originate from mass-transfer in binaries. However, in paper~I it was suggested that only $\\sim25$\\% of the total BSS are MT-BSS. The upward correction of this estimate is due to the fact that we have refined our parameter grid with respect to paper~I. In paper~I we considered only cases with a relatively low fraction of MT-BSS ($\\lesssim{}30$\\%) and one single case with $\\eta{}_{\\rm MT}=1$. For this paper we explored a more complete parameter space, the percentage of MT-BSS (with respect to the total number of BSS) going from 0 to 100\\%. Our present best fit (46\\% MT-BSS and 54\\% COL-BSS) is still consistent (at $2\\,{}\\sigma{}$) with the result of paper~I (25\\% MT-BSS and 75\\% COL-BSS), given the observational uncertainties, especially in the central bins (see Fig.~2).   \\subsection{NGC~6752}  The parameters of the best fit model for NGC~6752 are similar to those inferred for M3 and 47~Tuc.  However, for this cluster there may be a caveat about the central velocity dispersion, whose value is still very uncertain. From the integrated-light spectra of the core, Dubath et al. (1997) derived $\\sigma=4.9^{+2.4}_{-1.4}$ km s$^{-1}$. More recently, from proper motion measurements of a sample of $\\sim1000$ stars with the Wide Field Planetary Camera 2/{\\it Hubble Space Telescope} (WFPC2/HST), Drukier et al. (2003) obtained $\\sigma=12.4\\pm0.5$ km s$^{-1}$. Thus, we ran simulations for each of these two possibilities.  If we assume $t_{\\rm last}=1.5$ Gyr, the best fit model obtained for $\\sigma=4.9$ km s$^{-1}$ (Fig.~3) is not satisfactory, since we cannot reproduce the increasing number of BSS located in the cluster outskirts. In fact, a low velocity dispersion  implies that the dynamical friction time-scale is too short ($\\ll{}1$ Gyr) and PBs rapidly shrink towards the core. If we adopt $\\sigma=12.4$ km s$^{-1}$, the rising trend of the BSS radial distribution in the cluster periphery can be reproduced; but the $\\tilde{\\chi{}}^2$ ($\\sim{}5$) is still unacceptable, because of the disagreement with the data points in the inner regions.  Since none of the two options was satisfying, we ran new simulations for other values of $\\sigma$ (ranging between 4.9 and 12.4 km s$^{-1}$), until we found an acceptable fit to the observed BSS distribution.  That occurs for $\\sigma=7^{+3}_{-1}$ km s$^{-1}$ ($\\tilde{\\chi{}}^2=1.7$, reported in Table~2), marginally consistent with the measurement of Dubath et al. (1997).  However, for this cluster other marginally acceptable solutions for the BSS radial profile exist: e.g. assuming $t_{\\rm last}=4$ Gyr, a fit ($\\tilde{\\chi{}}^2\\lesssim{}3$) can be obtained also for $\\sigma=12.4$ km s$^{-1}$ (open triangles and short-dashed line in Fig.~3).  Considering  this uncertainty and the peculiarities of NGC~6752 (e.g. the unusual density profile), the aforementioned best fit value $\\sigma=7^{+3}_{-1}$ km s$^{-1}$ cannot be taken as an indirect measurement of the central velocity dispersion in the cluster. A new measurement of $\\sigma$ is strongly needed to understand the dynamics and evolution of NGC~6752.   \\begin{figure} \\epsfxsize=8.truecm \\epsfysize=8.truecm \\epsfbox{fig3c.eps} \\label{fig3} \\vspace{-0.5truecm} \\caption {Distribution of BSS normalized to the distribution of HB stars in NGC~6752. The solid line indicates the observations (filled circles; Sabbi et al. 2004). The dotted line (connecting open squares) reports the best fit numerical simulation for $\\sigma=4.9$ km s$^{-1}$; the long-dashed line (connecting  open circles) is for $\\sigma=7$ km s$^{-1}$ (reported also in Fig. ~2); the dot-dashed line (connecting crosses) is for the case $\\sigma=12.4$ km s$^{-1}$. In all these cases $t_{\\rm last}=1.5$ Gyr. The short-dashed line connecting open triangles is obtained for $\\sigma=12.4$ km s$^{-1}$ and $t_{\\rm last}=4$ Gyr.} \\end{figure}  \\subsection{The location of the minimum in the BSS radial distributions}  The radial distribution of BSS (normalized to the number of RGB or HB stars) in M3, 47~Tuc, and NGC~6752 clearly displays a minimum at a distance from the cluster center in the range $5-10\\,{}r_{\\rm c}.$ The area surrounding this minimum was dubbed 'zone of avoidance' (paper~I). We expect that its location roughly corresponds to the radius $r_{\\rm av}$ below which all the PBs of mass $\\simgt m_{\\rm BSS}$ (where $m_{\\rm BSS}$ is the minimum mass for generating BSS via mass transfer) have already sunk towards the cluster center because of dynamical friction. In other words, there is a small probability of observing a MT-BSS located at a distance from the cluster center of the order of $r_{\\rm av}.$ Most likely, the MT-BSS  which were originally located at $r\\simlt r_{\\rm av}$ can now be detected at $r\\simlt r_{\\rm c}$.  The time of dynamical friction $t_{\\rm df}$ for a mass $m_{\\rm BSS}$ located at a radius $r$ from the center of a GC is given by (Binney \\& Tremaine 1987) \\q \\label{eq:eq1} t_{\\rm df}=\\frac{3}{4\\,\\ln{\\Lambda}\\,G^2(2\\,\\pi)^{1/2}}\\, \\frac{\\sigma(r)^3}{m_{\\rm BSS}\\,\\rho(r)}, \\nq where $\\ln{\\Lambda}\\sim10$ is the Coulomb logarithm and $G$ the gravitational constant. $\\rho(r)$ and $\\sigma(r)$ are the cluster density and the velocity dispersion at the radius $r$. We note the dependence of $t_{\\rm df}$ on the density $\\rho(r)$ and especially on the cube of the velocity dispersion. Given the definition of this time-scale, $r_{\\rm av}$ can be calculated setting $t_{\\rm df}=t_{\\rm gc}$, where $t_{\\rm gc}$ is the lifetime of the cluster ($\\sim{}$12 Gyr), and using the simulated stellar density profile and the velocity dispersion distribution of the cluster for inferring $r_{\\rm av}$ from $\\rho(r_{\\rm av})$ and $\\sigma(r_{\\rm av})$. In doing the calculation, we choose $m_{\\rm BSS}=1.2\\,{\\rm M_\\odot}.$  The results (for $r_{\\rm av}$ normalized to $r_{\\rm c}$) are shown in the last column of Table 2 and in Fig.~2 (vertical dotted line). The values of $r_{\\rm av}$ are always consistent with the zone of avoidance observed in the data. This confirms for M3 and NGC~6752 what had been found for 47~Tuc (paper~I).  Third and fourth columns of Table~2 report the values of the fraction ${\\eta{}}_{\\rm MT}$ {\\it  (i)} for the MT-BSS located within the GC core $r_{\\rm c}$ ($\\eta{}_{\\rm MT} (<r_{\\rm c})$), and  {\\it (ii)} for the MT-BSS found at a distance from the cluster center larger than $r_{\\rm av}$ ($\\eta{}_{\\rm MT} (>r_{\\rm av})$).  These findings indicate that the position of a BSS in a GC (at least in those showing a bimodal distribution of the BSS) can be used as a strong indication of the nature of the progenitor. In particular, a BSS which is located outside the zone of avoidance almost certainly is a MT-MSS; on the contrary, a BSS found close to the cluster core is a COL-BSS.  Among the considered clusters,  only $\\omega$ Cen does not show any zone of avoidance. If we calculate the expected radius of avoidance also for this cluster, we find $r_{\\rm av}\\sim{}r_{\\rm c}$. This indicates that in $\\omega$ Cen the dynamical friction is far less efficient than in the other clusters. In $\\omega{}$~Cen the value of ${\\eta{}}_{\\rm MT}(>r_{\\rm av})$ is similar to the other clusters, whereas ${\\eta{}}_{\\rm MT}(<r_{\\rm c})$ is much higher (see Section 5).   We note that very similar values of both $\\eta{}_{\\rm MT}$ and $r_{\\rm av}/r_{\\rm c}$ have been obtained for an intermediate density GC (M3) and for two much denser clusters (47~Tuc and NGC~6752). The analogy between these three clusters is even more evident in Fig.~4, where radial distances are normalized to the avoidance radius, $r_{\\rm av}$. The radial distributions of BSS in M3, 47~Tuc and NGC~6752 nearly superimpose.  This also indicates that $r_{\\rm av}$ is a crucial parameter in describing the bimodal distribution of BSS.     { \\begin{figure}{ \\center{{ \\epsfxsize=8.truecm \\epsfysize=8.truecm \\epsfbox{fig4c.eps} }}\\label{fig4} \\footnotesize \\caption{ Observed distribution of BSS normalized to the distribution of HB (47~Tuc, NGC~6752) or RGB stars (M3). On the x-axis we plot the distance from the center normalized to $r_{\\rm av}$. Crosses indicate the observational data for 47~Tuc, open squares for M3 and open triangles for NGC~6752. The observational points are connected by a dashed line for M3, a dotted line for 47~Tuc and a solid line for NGC~6752.  }  } \\end{figure}     \\begin{table} \\begin{center} \\tabcolsep 4pt \\leavevmode \\caption{Mass Transfer vs Collisional Scenario} \\begin{tabular}{lccccc} \\hline \\hline & ${\\eta{}}_{\\rm MT}$ & ${\\eta{}}_{\\rm MT}(<r_{\\rm c})$ & ${\\eta{}}_{\\rm MT}(>r_{\\rm av})$ & $\\tilde{\\chi}^2$ & $r_{\\rm av}/r_{\\rm c}$\\\\ \\hline NGC~104       & 0.46$^{+0.06}_{-0.22}$ & 0$^{+0.002}_{-0}$  & 0.95$^{+0.02}_{-0.10}$    &  0.5  &  9.0\\vspace{0.1cm}\\\\ NGC~5272      & 0.41$^{+0.16}_{-0.20}$ & 0$^{+0.002}_{-0}$  & 0.91$^{+0.02}_{-0.01}$     &  2.2  &  4.9\\vspace{0.1cm}\\\\ NGC~6752      & 0.41$^{+0.12}_{-0.04}$ & 0.02$^{+0.03}_{-0.02}$ & 0.97$^{+0.01}_{-0.05}$  &  1.7  & 12.9\\vspace{0.1cm}\\\\ NGC~5139      & 0.86$^{+0.14}_{-0.08}$ & 0.64$^{+0.36}_{-0.04}$  & 0.90$^{+0.10}_{-0.002}$ &  0.8  &  1.05\\vspace{0.1cm}\\\\ \\noalign{\\vspace{0.1cm}} \\hline \\end{tabular} \\end{center} \\footnotesize { The second column reports the fraction $\\eta{}_{\\rm MT}$ of MT-BSS with respect to the total number of BSS (the errors are at 2$\\,{}\\sigma{}$, i.e. referred to $\\chi{}^2=\\chi{}_{\\rm best}^2+4$). The third (fourth) column reports the fraction of MT-BSS inside $r_{\\rm c}$ (outside  $r_{\\rm av}$) with respect to the total number of BSS inside $r_{\\rm c}$ (outside  $r_{\\rm av}$). The fifth column lists the reduced $\\chi^2$ obtained for the best fit models. The rightmost column reports $r_{\\rm av}$ normalized to $r_{\\rm c}.$ } \\end{table}    We have exploited the recent determination of the radial distribution of BSS in four GCs, in order to investigate which mechanism of BSS formation prevails in these stellar systems.  Our conclusion is that the two main formation paths proposed so far,  i.e. mass-transfer in PBs and merging of MS stars due to collisions in the cluster core, must {\\it co-exist and have similar efficiency both in a low density cluster (M3) and in much denser clusters, like 47~Tuc and NGC~6752.}  In particular, in M3, 47~Tuc, and NGC~6752 the COL-BSS sum to $\\sim 50-60\\%$ of the total and mostly reside in the central region of the cluster. The MT-BSS are slightly less abundant than the COL-BSS, but populate all the GC.  The density of BSS reaches a minimum in a so-called zone of avoidance, which separates the portion of the GC mostly occupied by COL-BSS from the cluster outskirts, where the MT-BSS dominate.  The location of the zone of avoidance is explained by accounting for the effects of the dynamical friction on the PBs which were massive enough for generating the observed BSS.  The picture described above can also be applied to $\\omega$~Cen; but in this case the lack of a central peak in the BSS radial distribution requires that the large majority of the BSS derive from PBs. The very low rate of production of COL-BSS could be in turn attributed to the fact that mass segregation has not yet driven a sizeable number of PBs to the central region of the cluster to produce BSS.  A very interesting further development of this research will be to perform a comparison between the location of a significant sample of BSS in a GC and their spectroscopic properties. According to the findings of this work, the position in the GC might represent a strong dynamical clue for the formation mechanism of a given BSS. If it is located outside the zone of avoidance, the BSS almost certainly results from evolution of a PB; if it is harbored in the cluster core, the BSS has most likely a collisional origin.  On the other hand, indication about the origin of the same BSS can be independently obtained from high resolution spectroscopy. Indeed the chemical signature of the MT-BSS formation process has been recently discovered in 47~Tuc (Ferraro et al. 2006b). The acquisition of similar sets of data in clusters with different structural parameters and/or in different regions of the same cluster will provide an unprecedented tool for confirming the scenario presented here and to finally address the BSS formation processes and their complex interplay with the dynamical evolution of the cluster."
    },
    "0609/astro-ph0609150_arXiv.txt": {
        "abstract": "{The microquasar 1E~1740.7$-$2942 is a source located in the direction of the Galactic Center. It has been detected at X-rays, soft gamma-rays, and in the radio band, showing an extended radio component in the form of a double-sided jet. Although no optical counterpart has been found so far for 1E~1740.7$-$2942, its X-ray activity strongly points to a galactic nature.}{We aim to improve our understanding of the hard X-ray and gamma-ray production in the system, exploring whether the jet can emit significantly at high energies under the light of the present knowledge.}{We have modeled the source emission, from radio to gamma-rays, with a cold-matter dominated jet model. INTEGRAL data combined with radio and RXTE data, as well as EGRET and HESS upper-limits, are used to compare the computed and the observed spectra.}{From our modeling, we find out that jet emission cannot explain the high fluxes observed at hard X-rays without violating at the same time the constraints from the radio data, favoring the corona origin of the hard X-rays. Also, 1E~1740.7$-$2942 might be detected by GLAST or AGILE at GeV energies, and by HESS and HESS-II beyond 100 GeV, with the spectral shape likely affected by photon-photon absorption in the disk and corona photon fields.}{} ",
        "introduction": "\\label{intro}   1E~1740.7$-$2942, discovered at X-rays by the {\\it Einstein} satellite (Hertz \\& Grindlay 1984)  and located at less than one degree from the Galactic Center, is considered to be the first microquasar since  the observation of radio jets extending $\\sim 1'$ (Mirabel et~al. 1992). There are hints of correlation between the radio and the strong X-ray emission (Mirabel et al. 1993), and X-ray state changes have been observed with timescales similar to those of other X-ray binaries (Del Santo et~al. 2005 and references therein). Moreover, a 12.7-days periodical modulation of a 3--4\\% in the X-ray emission seems to be produced by the orbital motion of the system  (Smith et al. 2002). This X-ray behavior suggests that the source is galactic in nature. The modulation period and the upper-limits in magnitude obtained in the infrared band  suggest that the stellar companion in 1E~1740.7$-$2942 might be a low-mass star (Mart\\'i et~al. 2000). The behavior of the X-ray spectrum, very similar to those of other galactic black-hole candidates (or BHC) (e.g. Sunyaev et~al. 1991), points to a black-hole as the compact object. The inferred star mass, the strong X-ray radiation and the little X-ray orbital modulation, hinting perhaps to a circular orbit, point to accretion via Roche-lobe overflow. The fact that the companion has not been detected yet is likely due to the dense surrounding medium, with high concentrations of dust and a large hydrogen column density of about $10^{23}$~cm$^{-2}$ (Gallo \\& Fender 2002), which depletes photons from the infrared up to soft X-ray energies. The same authors also detected hints of an extended X-ray halo. This  hydrogen column  density is even higher than the typical value in the direction of the Galactic Center ($\\sim 10^{22}$~cm$^{-2}$; Dickey \\& Lockman 1990), and might imply that this object is embedded within a molecular cloud (Yan \\& Dalgarno 1997). This peculiar location  could also explain the very extended structures detected at different wavelengths around this source, as was pointed out by Mirabel et~al. (1992). Regarding the distance to 1E~1740.7$-$2942, we have adopted in this work the Galactic Center distance, 8~kpc.  Concerning the emission properties, the source shows spectrally flat and faint radio emission coming from the core with flux densities of less than 1~mJy (Fender 2001 and references therein;  $\\sim 10^{29}$~erg~s$^{-1}$ at 8~kpc), as well as optically thin radio emission coming from the radio lobes. In the 10-100 keV band, {\\it RXTE} and  {\\it INTEGRAL} observed the source in August 2003 when it was in the low-hard state finding fluxes of about $10^{-9}$~erg~s$^{-1}$~cm$^{-2}$ ($\\sim 10^{37}$~erg~s$^{-1}$; Del Santo et~al. 2005). Further data have been taken by {\\it INTEGRAL},  which altogether with the previous ones have been reduced and analyzed for this work with the most recent software package (OSA~5.1), obtaining very similar results to those already published. EGRET observed intensively the Galactic Center region in the range $\\sim 0.1-10$~GeV. No source was detected at the position of 1E~1740.7$-$2942, which provides upper-limits for its emission at these energies. At the Galactic Center distance the inferred luminosity upper-limits are about $10^{35}$~erg~s$^{-1}$ (Hartman et al. 1999). Moreover, the flux of the nearby source 3EG~J1744$-$3011 can be taken as an absolute upper-limit\\footnote{The EGRET sensitivity in that region is confuse, as it is seen in Fig.~3 of Hartman et~al. (1999), and the angular resolution is poor all along the Galactic plane.}, being about $10^{-10}$~erg~s$^{-1}$~cm$^{-2}$ ($\\sim 10^{36}$~erg~s$^{-1}$). The Cherenkov telescope HESS, working above  0.1~TeV has not detected 1E~1740.7$-$2942, although significant diffuse emission has been found in the surroundings of the Galactic Center (Aharonian et al. 2006) that could cover intrinsic TeV emission from  this source. It is probably premature to state that this source cannot be detected by HESS at very high energies with longer exposures, or by HESS-II in the near future.  We want to explain the emission of this object in the context of microquasars during the low hard state (Fender et~al. 2003), when emission at radio and higher energies from a compact jet is produced and accretion processes are likely contributing or even dominant at X-rays. Data in the radio band and X-rays/soft gamma-rays (from {\\it RXTE} and {\\it INTEGRAL})\\footnote{{\\it CHANDRA} data from the observations performed by Gallo \\& Fender (2002) is not  considered here due to its non-simultaneity with the observations carried out by {\\it RXTE} and {\\it INTEGRAL}. We note however that at that time the flux at few keV was about one order of magnitude lower than the flux presented here in  similar energy ranges. We note that radio is either non-simultaneous, but only moderate radio flux variations have been found (Mirabel et~al. 1993).}, and upper-limits at high-energy and very high-energy gamma-rays (EGRET and HESS), have been considered. We have applied a detailed jet model to explain the data and to make predictions at high energies, and to give some constraints about the physics of the processes involved in producing the broadband emission. ",
        "conclusions": "There is a wide debate in the astrophysical community about the possible origin of the X-rays in microquasars  since, at the present state of knowledge, the corona and the jet scenarios seem to be roughly consistent in some cases with observations (Markoff et~al. 2005). In 1E~1740.7$-$2942, the hard X-ray emission is difficult to be explained if coming from a jet since it would imply an energetic efficiency in the jets significantly larger than for the corona emission (see Sect~\\ref{res}). Actually, in the context of our model, synchrotron emission from the jet seems to be able to explain the low-hard state X-ray spectrum of the source, but it exceeds largely the observed core radio fluxes. 1E~1740.7$-$2942 appears to show related radio and X-ray variability (Mirabel et al. 1993), although is radio underluminous in the context of the radio-X-ray luminosity correlation found for black hole candidate X-ray binaries and associated with the accretion/ejection activity (Gallo et~al. 2003, Corbel et~al. 2003, Fender et~al. 2003). This might be due to a particularly radiatively efficient corona. We remark that radio variations (Mirabel et~al. 1993) could hardly explain a flux as large as that predicted with the synchrotron jet model  (there remains the possibility that other  jet models, with a much higher electron minimum energy, may be consistent with observations). All this points to the fact that, although accretion and jet phenomena are probably linked, the dominant component at X-rays can be either the jet or the disk/corona depending on the source.  After the discovery of gamma-ray emission from microquasars (Paredes et~al. 2000, Aharonian et al. 2005, Albert et al. 2006), gamma-rays are becoming of great interest as a probe for X-ray binaries. In this context, we remark that jets in general, if sharing similar properties, are likely to emit above X-rays, reaching in some cases very high energies. Our model predicts significant amounts of gamma-rays from 1E~1740.7$-$2942 that might be detected by future space- and ground-based instruments like {\\it GLAST}, {\\it AGILE} or HESS and HESS-II, although photon-photon absorption in the disk and corona fields could prevent these detections, as can be seen comparing  the two cases explored here (see Fig.~\\ref{plot1}; for more precise calculations showing the importance of the discussed effects, see e.g. Akharonian \\& Vardanian 1985, Wu et~al. 1993, Bednarek 1993).  Concerning the large-scale jets, Gallo \\& Fender (2002) stated that the magnetic field should be $\\ge 1$~$\\mu$G, in agreement with our semi-quantitative  results but inferred on the basis of completely different grounds. Our estimate of the magnetic field would imply that no significant IC gamma-rays can be originated in that region, since the synchrotron emission channel would be more efficient than the IC one. It would not prevent high-energy emission to be produced via other mechanisms, like hadronic radiative mechanisms possibly at work within the mentioned molecular cloud  that could surround 1E~1740.7$-$2942 (see, e.g., Bosch-Ramon et al. 2005)."
    },
    "0609/astro-ph0609199_arXiv.txt": {
        "abstract": "Cataclysmic variables undergo periodic nova explosions during which a finite mass of material is expelled on a short timescale. The system widens and, as a result, the mass-transfer rate drops. This state of hibernation may account for the variety of cataclysmic variable types observed in systems of similar mass and period. In the light of recent changes to the theory of nova ignition and magnetic braking we investigate whether hibernation remains a viable mechanism for creating cataclysmic variable diversity. We model the ratio of time spent as dwarf novae (DNe) to nova-like systems (NLs). Above a critical mass-transfer rate the system is NL and below it a DN. The dominant loss of angular momentum is by magnetic braking but the rate is uncertain. It is also uncertain what fraction of the mass accreted is expelled during the novae. We compare the models of the ratios against the period of the system for different magnetic braking rates and different ejected masses with the ratio of the number of observed NLs to DNe. We deduce that a rate of angular momentum loss a factor of ten smaller than that traditionally assumed is necessary if hibernation is to account for the observed ratios. ",
        "introduction": "Cataclysmic variables \\citep{b6} are interacting close binary stars in which a white dwarf is accreting from a companion. In the vast majority of cases the companion is a low-mass main-sequence star or red dwarf and the periods are measured in hours. The red dwarf is transferring hydrogen-rich material to the white dwarf by Roche lobe overflow at a rate, $10^{-11}$ to $10^{-8} M_\\odot\\,\\rm yr^{-1}$. As the hydrogen settles on the surface, the base of the layer becomes degenerate and heats. When enough, $10^{-5}$ to $10^{-3} M_\\odot$, has accumulated it ignites. The degenerate conditions lead to a thermonuclear runaway and a nova explosion. The system brightens by ten or so magnitudes and most of the layer if not all of it together with some of the underlying white dwarf material is blown off.  Though some of this material may fall back on to the white dwarf and some of it may be accreted back by the red dwarf, a significant fraction escapes carrying with it only the intrinsic angular momentum of the white dwarf. The system expands such that $aM=\\rm const$, where $a$ is the separation and $M$ the total mass which has just decreased. If the red dwarf had no atmosphere, that is a distinct boundary just filling its Roche lobe, it would no longer do so and mass transfer would cease. This is the idea of hibernation first documented by \\cite{b4}.  At that time they were motivated by the apparent lack of cataclysmic variables in the solar neighbourhood when compared with theoretical estimates and the frequency of classical novae in M31.  The complete detachment during hibernation would render the majority of cataclysmic variables unobservable at any given time. \\par Because the red dwarf has a slightly extended atmosphere it actually must overfill its Roche lobe in order to maintain an equilibrium mass-transfer rate. The ratio of the atmosphere scale height to the radius of the star is comparable to the ratio of the mass lost from the system to its total mass, so the change in separation is not sufficient to shut off mass transfer altogether but does reduce the rate significantly. This change in rate could be responsible for the diversity of cataclysmic variable types at a given period: dwarf novae that undergo disc instability outbursts require a lower mean mass-transfer rate than nova-likes that do not suffer the disc instability \\citep{Faulk}. However the idea lost favour when it was realised that the nova ejecta would temporarily form a common envelope around the system and consequential angular momentum loss would actually cause lower-period systems to shrink \\citep{livio91}. Recent work by \\cite{Nele} indicated that the common-envelope phase is very poorly understood and that in some cases systems with tenuous envelopes can actually expand. Since we wish to compare the simplest case we neglect any common-envelope phase in this work. \\par A detailed analysis of the period distribution was made by \\cite{shafter92}.  He emphasized that the lack of dwarf novae between 3 and~$4\\,$hr is difficult to reproduce theoretically with any reasonable magnetic braking law.  He speculated that the solution might lie in a better understanding of any of additional angular momentum loss alongside mass transfer, weak but unnoticed magnetic fields that disrupt the inner parts of the disc, correlation between white dwarf mass and orbital period or hibernation.  Among these he deduced that hibernation might prove the most promising and it is this that we explore here in more detail. \\par We investigate the hibernation scenario for cataclysmic variable diversity in the light of two recent changes to the accepted model of cataclysmic variables. These are \\begin{enumerate} \\item the relation between the mass of hydrogen that must be accreted to ignite $M_{\\rm ign}$ and the secular accretion rate \\citep{TandB} \\item the possibility that magnetic braking, that drives the evolution of cataclysmic variables, could be much weaker than previously thought \\citep{And}. \\end{enumerate} We consider hibernation in its simplest form to deduce what combinations of $\\dot J$, $M_{\\rm ign}$ and $M_{\\rm ej}$, the mass actually ejected in the nova, could be consistent with the observed variety of cataclysmic variable types. ",
        "conclusions": "For the full \\cite{Rapp} rate ($\\dot J$ factor of $1$) only the unrealistically high $M_{\\rm ej}$ of $10^{-3}M_\\odot$ follows the observed trend but then the variation of ratio with period is too small.  With $\\dot J$ reduced by a factor of ten all the models show the right trends with the ratio rising to a peak at a period of about $3\\,\\rm hr$ and then falling distinctly by a period of $4.4\\,\\rm hr$. We cannot then distinguish between the different mass ejection models but do find a distinct improvement when $M_{\\rm ej}=M_{\\rm ign}$ when proper account is taken of white dwarf heating by accretion \\citep{TandB}. Hibernation is most effective at reproducing observed trends when our $\\dot J$ factor is about $0.1$. This is ten times smaller than that proposed by \\cite{Rapp} and ten times larger than that proposed by \\cite{And}."
    },
    "0609/astro-ph0609685_arXiv.txt": {
        "abstract": "IGR~J00291+5934 is one of the seven accreting millisecond pulsars (AMPs) discovered so far. We report on the aperiodic timing and color analysis of its X-ray flux, using all the {\\it RXTE} observations of the 2004 outburst. Flat-top noise and two harmonically related quasi-periodic oscillations, all of them at very low frequencies (0.01-0.1~Hz), were present in the power spectra during most of the outburst as well as a very high fractional variability ($\\sim$50~\\%). These properties are atypical not only for AMPs but also for neutron star low-mass X-ray binaries (LMXBs) in general. There are instead some remarkable similarities with the variability observed in black hole systems, reinforcing the connections between these two types of LMXB, as well as some interesting differences. We note finally that the results of this paper are difficult to reconcile with interpretations where any break frequency of power density spectra scales inversely with the mass of the central object at an accuracy sufficient to distinguish between the masses of neutron stars and black holes in LMXBs. ",
        "introduction": "\\label{sec:intro}  Most low-mass X-ray binaries (LMXBs) contain rapidly spinning neutron stars (NSs), which only in a few cases become apparent as millisecond X-ray pulsars. This growing sub-class of accreting millisecond pulsars (AMPs; see \\citealt{Wijnands05} for a review) consists of dim transients that typically go in outburst during a few weeks once every few years (so far the only exception is HETE~J1900.1-2455, still active at the moment of writing, more than a year after its discovery). AMPs are thought to be the evolutionary link between neutron star LMXBs and millisecond radio pulsars, according to the so-called recycling scenario \\citep[][for a review]{BhatHeuv91}. Eight years after the first of the AMPs was discovered \\citep[SAX~J1808.4--3658,][]{Wijnands98} it is still unknown why these systems show pulsations and the rest of neutron star LMXBs do not, although some explanations have been put forward (e.g. screening of the magnetic field, \\citealt{Cumming01}; smearing of the pulsations, \\citealt{Titarchuk02}).  Aperiodic variability in the X-ray lightcurves of black hole and neutron star LMXBs (BH/NS-LMXBs), such as broadband noise or quasi-periodic oscillations (QPOs), offers a powerful tool to study the accretion flow in its inner region, where the bulk of the radiation is thought to originate and where gravitational fields are extreme \\citep{vanderKlis00,vanderKlis06}. Using this tool NS-LMXBs were classified according to their correlated timing and energy spectral behavior into Z and atoll sources \\citep{Hasinger89}. As noted several years ago \\citep[e.g.][]{Klis94}, very similar physical mechanisms are expected to drive the accretion process in BH- and NS-LMXBs, and the difficulties in finding unequivocal signatures of a black hole or a neutron star in the energy or power spectral properties of LMXBs are a natural consequence of this fact. Several attempts to find direct BH signatures have been made based on the spectral properties at high energies (\\citealt{Sunyaev91}; but see also \\citealt{DiSalvo04} or \\citealt{Barret00}) or in the power spectral properties at high frequencies (\\citealt{SunRev00}; but see also \\citealt{Kleinwolt04b}). On the other hand, apart from pulsations and type I X-ray bursts, twin variable-frequency kilohertz quasi-periodic oscillations (twin kHz QPOs) are now considered a NS signature \\citep{vanderKlis06}. NS-LMXBs are also known to be fainter than BH-LMXBs at radio wavelengths \\citep{Migliari06}, and BH-LMXBs are generally fainter in X-rays when they are in quiescence \\citep{Garcia01}.  We present in this work the results of the aperiodic timing and X-ray color analysis of one of the seven known AMPs, IGR~J00291+5934. We show the exceptional behavior of this source in its 2004 outburst, more similar with respect to aperiodic X-ray variability to a BH-LMXB than to the presently known NS-LMXBs: IGR~J00291+5934 was variable at very low frequencies with a very high fractional amplitude. However, previously identified NS signatures were also present in the highest frequencies. In Section~\\ref{sec:review} we briefly review some information on this source and in Sections~\\ref{sec:data} and \\ref{sec:results} we describe our analysis and give its results, respectively. Section~\\ref{sec:discussion} contains our final remarks and conclusions. ",
        "conclusions": "\\label{sec:discussion}  \\subsection{Exceptional island state} \\label{ssec:extreme}  Our analysis of the fastest spinning AMP, IGR~J00291+5934, reveals the exceptional behavior of this source in its 2004 outburst compared to the rest of NS-LMXBs: very strong X-ray variability at very low Fourier frequencies. Its power spectra showed flat-topped noise (i.e. approximately flat below a certain ``break'' frequency and decreasing above that) with extremely low break frequencies ($\\sim$0.04~Hz), as well as very high integrated fractional rms variability ($\\sim$50~\\%). These break frequencies are about two orders of magnitude lower than those observed in Z sources and about one order of magnitude lower than those previously found in extreme island (low-hard) states of atoll sources \\citep[e.g,][]{Barret00,Belloni02}\\footnote{A reanalysis of GS~1826--238 data in \\citet{Barret00} led to a break frequency above 0.1~Hz}. The measured fractional rms amplitude of the variability is the highest found so far for in a NS-LMXB. Both the break frequencies and fractional variability observed in IGR~J00291+5934 are instead similar to those detected in low-hard states of BH-LMXBs (though the frequencies are low even for the lowest BH-LMXB range; see next section for a more detailed comparison). It has been noted \\citep{Straaten05, Linares05} that AMPs hardly ``make it to the soft state'' and instead remain in (low-hard) island or extreme island states. The (2.5--25.0~keV) peak luminosity of IGR~J00291+5934 (between 0.5 and 5.5~\\% of the Eddington luminosity according to the flux measured by \\citealt{Galloway05} and the different distance estimates, see Sec.~\\ref{sec:review}) was similar to that of other transient NS-LMXBs when they show ``canonical'' extreme island states, which seems to indicate that the exceptional state of our source was not purely a consequence of an exceptionally low accretion rate.  The lack of high frequency variability ($\\gtrsim$100~Hz) during the outburst of IGR~J00291+5934 might be related to a combination of a dynamically important magnetic field (needed in AMPs in order to produce the observed pulsations) and a fast spin (spin frequency of $\\sim$600 Hz in the case of IGR~J00291+5934) that could prevent the accretion disk from reaching the innermost regions, where the highest dynamical frequencies are present. This could in turn be related to the strong overall variability that we measured. If, as has been argued \\citep[e.g.][]{Churazov01, Gilfanov03}, the disk sets the variability frequencies and the comptonizing medium (corona, boundary layer, spreading layer, ADAF or similar) sets the variability amplitude, then a disk truncated at a large radius could provide both low dynamical frequencies and strong {\\it fractional} variability, as its flux contribution would be relatively low and hence so would be the decrease that this flux would cause in the relative amplitude of the variability. We note, however, that in standard descriptions \\citep[e.g.][]{Gosh78} the inner disk radius does not depend on the spin rate, and if the so-called propeller mechanism is responsible for clearing out the inner disk annulus this would be in contradiction with the reported spin-up \\citep{Falanga05b, Burderi05}, as the expelled matter would remove angular momentum from the neutron star. The interaction between a rapidly spinning magnetized neutron star and an accretion disk is a complex problem.  Most of the work so far concerns the ``classical'' slow X-ray pulsars \\citep{Lamb73, Gosh78} and only recently the AMP case has received more attention \\citep{Psaltis99c,Rappaport04}.  An anticorrelation in the flat-topped noise of Cyg~X-1 was discovered by \\citet{BH90} linking the rms level of the flat top with the break frequency: the lower the break frequency, the higher is the flat-top (in rms normalization). In Figure ~\\ref{fig:BHWK} (top) we see how IGR~J00291+5934 extends this relation to the lowest frequencies and the highest rms levels, even when comparing with the BH system Cyg~X-1 \\citep[see][for an update of this correlation and comparison with other neutron star systems, which are all at frequencies $>$0.1~Hz, and the BH-LMXB XTE~J1118+480, which is contiguous with but still at slightly higher frequecies than our NS-LMXB IGR~J00291+5934]{Belloni02}.   \\subsection{IGR~J00291+5934 vs. BH-LMXBs} \\label{ssec:NSBH}  The similarities between IGR~J00291+5934 and BH-LMXBs in low-hard states are illustrated in Figure~\\ref{fig:comp}, where a power spectrum of IGR~J00291+5934 is compared to one of XTE~J1118+480 and one of XTE~J1550-564 \\citep[e.g.][respectively]{Revni00, Cui99}. The flat-top noise has very similar power, break frequency and QPOs. There are also obvious yet interesting differences: as previously noted in a study of low states of BH/NS-LMXBs the NS power spectra show more power at high frequencies ($\\gtrsim$10~Hz) than the BH ones \\citep{SunRev00}. In this extreme case, however, the differences between the power spectra of IGR~J00291+5934 and XTE~J1118+480 start at $\\sim$0.2 Hz, where the NS power spectrum rises and the BH one starts to fall (see Fig.~\\ref{fig:comp}). The power spectrum of XTE~J1550-564 has a very similar shape to that of IGR~J00291+5934 and XTE~J1118+480, but shifted towards higher frequency and lower rms. It also shows a prominent bump at $\\sim$1-10~Hz (similar to what is seen in IGR~J00291+5934 at slightly lower frequency) followed by a steep decline above $\\sim$10~Hz (characteristic of BH systems; see \\citealt{Kleinwolt06} for an extensive study and further discussion). In contrast to the ones studied in \\citet{SunRev00} this NS-LMXB shows no significant power above $\\sim$100~Hz.  We found two low-frequency QPOs at $\\sim$0.02 and $\\sim$0.04~Hz, similar to the ones seen in BH- and other NS-LMXBs in low-luminosity states \\citep{Casella05, Straaten05}. Their centroid frequency ratios were 2.06$\\pm$0.06 and 1.91$\\pm$0.06 in A1 and A2 respectively, indicating an harmonic relation. In general this kind of QPO is associated with low-luminosity states and therefore it requires a relatively low mass accretion rate in the inner regions of the disk. Several attemps have been made to investigate the nature of these low-frequency QPOs, e.g. in the context of shot-noise models \\citep{Vikh94} or thermal-viscous instabilities \\citep{Chen94}. In any case our work shows clearly that neither the presence of a solid surface nor a magnetic field affects this phenomenon so that the same physical mechanism must be at work in NS and BH systems in order to produce this common feature: harmonically related QPOs superposed on a break in the flat-top noise at very low frequencies.  \\citet{WK99} found a correlation between the break frequency and the frequency of the ``hump'' present in the power spectra of atoll, Z sources, black hole candidates and AMPs ($L_h$ or sometimes an associated QPO, $L_{LF}$). We find that the frequency of the first zero-centered Lorentzian detected above the break in the power spectra of IGR~J00291+5934 also satisfies this WK correlation (Figure~\\ref{fig:BHWK}, bottom) which leads us to identify it as $L_h$. The fact that $\\nu_{QPO}$ and $\\nu_{QPO/2}$ fall much below this relation (on top of the break as discussed) suggests that they are not directly related to the low frequency QPOs ($L_{LF}$, $L_{LF/2}$) measured sometimes on top of the ``hump''. Interestingly, the same happens for XTE~J1118+480 and XTE~J1550-564: the QPOs on top of the break fall clearly below the WK correlation, whereas $L_h$ follows it.  The relations between the characteristic frequencies of several variability components of Z and atoll sources and the frequency of the upper kilohertz QPO, $\\nu_u$, were studied and synthesized by \\citet{Straaten03} (Figure~\\ref{fig:vStraaten}). Shifts in these frequency-frequency correlations were found in two AMPs by \\citet{Straaten05} and \\citet[][in SAX~J1808.4--3658 and XTE~J1807--294, respectively]{Linares05}, in a way that was best explained by the upper and lower kHz QPOs having frequencies ($\\nu_u$, $\\nu_{\\ell}$) lower than the ones measured in atoll sources by a factor $\\sim$1.5. The component $L_{\\ell ow}$ plotted in Figure~\\ref{fig:vStraaten} is the broad component present below $L_u$ (the component with the highest characteristic frequency) in extreme island states of low-luminosity atoll sources (including bursters and AMPs; \\citealt{Straaten02,Straaten03,Straaten05}). We tentatively identify the two highest frequency variability components in IGR~J00291+5934 as $L_u$ and $L_{\\ell ow}$, as these are the only broad variable-frequency components detected in extreme island states of NS-LMXBs above $L_h$, and plot the frequencies of $L_b$, $L_h$ and $L_{\\ell ow}$ against that of $L_u$ (Figure~\\ref{fig:vStraaten}a). As can clearly be seen in Figure~\\ref{fig:vStraaten}, the frequencies of IGR~J00291+5934 also vary in correlation to one another, but with slopes much less steep than those found in atoll sources and AMPs at higher frequencies. In order to compare IGR~J00291+5934 with the two shifted AMPs, we multiply $\\nu_u$ and $\\nu_{\\ell ow}$ by a trial factor of 1.5 (Figure~\\ref{fig:vStraaten}b). The tracks traced by IGR~J00291+5934 seem to connect better with the ones at higher frequencies in the shifted version of the frequency-frequency plots, but of course the difference in slopes remains. We conclude that a radical change in the slope of these correlations appears to occur at frequencies $\\nu_u \\lesssim$100~Hz. It is unclear at this point to what extent this is universal to LMXBs or specific to IGR~J00291+5934. It will be interesting to study the behavior of our source in future outbursts to see if at $\\nu_u \\gtrsim$~100Hz its components follow the usual correlations or not.  \\subsection{Mass scaling for break frequencies} \\label{ssec:scaling}  This work also sheds some light on another interesting issue regarding broad-band noise in power spectra. Break frequencies are often used to estimate masses of accreting black holes with the argument that timescales in the accretion disk scale linearly with the mass of the central object and therefore the observed frequencies should scale inversely \\citep[e.g.,][]{Uttley02}. From the theoretical point of view this argument is weakened by our poor knowledge of what produces the broad-band noise and which timescale sets its break frequency \\citep[viscous, dynamical, related with the size of the disk, with shot lifetime; see e.g.][]{Priedhorsky79,Manmoto96,Poutanen99,Merloni00,Churazov01}. From the observational point of view the application of this argument using break frequencies ($\\nu_b$) of the flat-top noise ($L_b$) observed in low-hard states of LMXBs is clearly challenged by IGR~J00291+5934, as it harbors a neutron star and shows break frequencies lower than the ones observed in most of the more massive stellar-mass black holes. Although $\\nu_b$ of Cyg~X-1 in the low-hard state was used in the past to infer masses of super-massive black holes, recent advances in active galactic nuclei (AGN) X-ray timing indicate that most AGN are in a state similar to the high-soft state of BH-LMXBs \\citep{Uttley05}. Any power spectral feature that is used in such mass-scaling arguments should i) have a rather stable frequency (like perhaps the hectohertz Lorentzian in NS systems, \\citealt{Straaten02}; see also \\citealt{Pottschmidt03}) or ii) be accurately corrected for drifts due to some parameter varying in and between systems (e.g. mass accretion rate; \\citealt{McHardy04,Uttley05}) or a combination of i) and ii). In any case the power spectra observed in IGR~J00291+5934 refute a simple mass-frequency relation for most of the variability components present in LMXBs.  \\textbf{Acknowledgments:}  We thank Phil Uttley for stimulating discussions on AGN X-ray variability. We also thank D. Altamirano, M. Klein-Wolt, M. M{\\' e}ndez and A. Patruno for useful discussions, and the anonymous referee for valuable comments that improved the clarity of this paper."
    },
    "0609/astro-ph0609366_arXiv.txt": {
        "abstract": "We present a near-infrared spectrum of the close, detached white dwarf $+$ brown dwarf binary WD\\,0137-349 \\citep{maxted06}, that directly reveals the substellar companion through an excess of flux longwards of $\\approx1.95\\um$. We best match the data with a white dwarf $+$ L8 composite model. For ages $\\sim1$~Gyr, the spectral type of the cool secondary is in agreement with the mass determined by \\citet{maxted06} from radial velocity measurements ($0.053 \\pm 0.006\\Msun$), and supports an evolutionary scenario in which the brown dwarf survived a previous phase of common envelope evolution which resulted in the formation of this close binary. The brown dwarf is the lowest mass companion to a white dwarf yet found, and the lowest mass object known to have survived a common envelope phase. At $1300 < T_{\\rm eff} < 1400$~K WD\\,0137-349B is also the coolest known companion to a white dwarf. At a separation $a = 0.65$R$_\\odot$ the hemisphere of the brown dwarf facing the $16,500$K white dwarf is being heated through irradiation. We discuss the possible effects of this additional heating, with particular relevance to those other close binaries with substellar companions, the hot Jupiters. We propose future observations to investigate the likely temperature differences between the ``day'' and ``night'' sides of the brown dwarf. ",
        "introduction": "Detached brown dwarf companions to white dwarfs are rare \\citep*{fbz05}. Proper motion surveys and searches for infra-red (IR) excesses have so far found only three confirmed examples: GD\\,165 (DA$+$L4, \\citealt{becklin88}), GD\\,1400 (DA$+$L$6/7$, \\citealt{gd1400}; \\citealt{dobbie05}), and WD\\,0137-349 \\citep{maxted06}, the subject of this paper. GD\\,165 is a widely separated system (120\\,AU); the separation of the components in GD\\,1400 is currently unknown. In contrast, WD\\,0137-349 is a close binary (orbital period $P\\approx116$~minutes).  Optical spectra of the H-rich DA white dwarf WD\\,0137-349 show a narrow H$_\\alpha$ emission line due to irradiation of the surface of the unseen companion. Radial velocities measured from this line and the white dwarf's intrinsic H$\\alpha$ absorption line allowed \\citet{maxted06} to determine the mass ratio of the system. Using the white dwarf mass ($0.39\\pm0.035\\Msun$), derived from an analysis of its optical spectrum, \\citet{maxted06} then determined the mass of the companion to be $0.053\\pm0.006\\Msun$. This is well below the limit of $0.075\\Msun$ commonly used to distinguish stars from brown dwarfs. The substellar nature of WD\\,0137-349B was reinforced by an analysis of its 2MASS near-IR $JHK$ fluxes (Figure 2, \\citealt{maxted06}). The $J$ and $H$ band photometry can be fit by a model consistent with the white dwarf alone. There is a slight excess of flux at $K_{\\rm S}$, which can be matched by a model consisting of the white dwarf plus a brown dwarf spectral type mid-L or later. This is consistent with the radial velocity determined mass measurement for WD\\,0137-349B.  Therefore, WD\\,0137-349 is the first close, detached binary to be discovered containing a confirmed substellar companion. The brown dwarf must have survived a previous phase of common envelope (CE) evolution during which it was engulfed by the red giant progenitor of the white dwarf \\citep{politano}. The drag on the brown dwarf caused it to spiral in towards the red giant core from an originally much wider orbit. Some fraction of the orbital energy was released and deposited in the envelope, which was ejected from the system, leaving a close binary. Simple physical arguments suggest that low mass companions less than some limit $m_{\\rm crit}$ will be evaporated during the CE phase. WD\\,0137-349B places an upper limit on $m_{\\rm crit}$.  Alternatively, WD\\,0137-349B may have originally been a planet which accreted a substantial fraction of its mass during the CE phase \\citep{ls83}. In this scenario, WD\\,0137-349B would be expected to have an effective temperature $T_{\\rm eff}>2000$K, and a spectral type (early~L) commensurate with the cooling age of the white dwarf ($\\sim0.25$~Gyr). \\citet{maxted06} concluded that this model was probably not applicable to WD\\,0137-349 since the 2MASS fluxes are inconsistent with a companion earlier than mid-L. However, they suggested that a definitive test would be provided by IR spectroscopy to directly determine an accurate spectral type for the brown dwarf.  In this paper we present a near-IR spectrum of WD\\,0137-349. In Section~2 we discuss the observations and data reduction, in Section~3 we present our analysis of these data, and in Section~4 we discuss the results and comment on the spectral type of WD\\,0137-349B. ",
        "conclusions": "We have obtained a near-IR spectrum of the close, detached white dwarf $+$ brown dwarf binary WD\\,0137-349 that reveals the substellar companion through an excess of flux longwards of $\\approx1.95\\um$. We best match the data with a white dwarf $+$ L8 composite model. This is the lowest mass, coolest and latest spectral type companion to a white dwarf yet discovered. For a cooling age $\\sim1.0$~Gyr, the temperature of the secondary ($1300 < T_{\\rm eff} < 1400$~K) is in agreement with the mass determined by \\citet{maxted06} from radial velocity measurements ($0.053 \\pm 0.006\\Msun$), and supports an evolutionary scenario in which the brown dwarf survived a previous phase of CE evolution $\\approx 0.25$~Gyr ago, which resulted in the formation of this close binary. The exotic alternative evolutionary scenario of \\citet{ls83}, in which a lower mass planet or brown dwarf accreted substantial mass during the CE phase, is not supported. If the brown dwarf is $\\sim1.0$~Gyr old, then the main sequence lifetime of the white dwarf progenitor was $\\approx0.75$~Myr and its mass was just over $2\\Msun$, requiring a value of the CE parameter $\\alpha_{\\rm CE}>2$ to explain the formation of the system. However, the brown dwarf may be older if it was re-heated by the very hot white dwarf immediately after the CE phase. The brown dwarf hemisphere facing the $16,500$K white dwarf is currently intercepting $\\approx1\\%$ of its light, but does not appear to be substantially heated, although our spectroscopic observations only cover $\\approx0.5$ of the orbital period."
    },
    "0609/astro-ph0609400_arXiv.txt": {
        "abstract": "{ We apply a semi-classical approach of handling waves as quasiparticle gas in a slowly varying flow -- analogous to ray tracing -- to calculate the Alfv\\'en wave transmission parameters, the resulting cross-helicity of the waves and the scattering-centre compression ratio, for cases where the shock thickness is large enough for the turbulent waves in the plasma to see the transition of the background flow parameters as smooth and slowly varying. For nonrelativistic shocks the wave transmission produces similar effects on the downstream turbulence and the scattering-centre compression ratio as does the transmission through a step shock: the downstream Alfv\\'en waves propagate predominantly towards the shock in the local plasma frame and, thus, the scattering-centre compression ratio is larger than the gas compression ratio. For thick relativistic shocks, however, we find qualitative differences with respect to the step-shock case: for low-Alfv\\'enic-Mach-number shocks the downstream waves propagate predominantly away from the shock, and the scattering-centre compression ratio is lower than that of the gas.  Thus, when taken into account, the Alfv\\'en wave transmission can decrease the efficiency of the first-order Fermi acceleration in a thick relativistic shock. ",
        "introduction": "A common approximation in shock-related particle acceleration studies involving the first-order Fermi mechanism is to assume the scattering turbulence frozen-in to the plasma. This is very well justified in cases where the speed of the turbulent waves is negligibly small compared to the flow speed. For cases where the wave speed increases and becomes a notable fraction of the flow speed, however, the deviation of the scattering centre speed from the speed of the background flow can have notable effects on the particle acceleration scenario; it was noted already by Bell (\\cite{Bell1978}) that the speed of the scattering centres is not necessarily that of the flow. Furthermore, the compression ratio (of the scattering centres) is known to have significant effect on the power-law spectral index of particles accelerated in the shock for both steplike and thick shocks (for modified shocks see Drury et al. \\cite{DruryEtAl1982} and Drury \\cite{Drury1983} for nonrelativistic analytical studies, and Virtanen \\& Vainio 2005a, hereafter \\cite{VV2005AA}% \\footnote{In Eq.~(1) of VV05 there is a typo in the exponent in the denominator of the last term: $r^2$ should be $r^3$.}% , for numerical results for all speeds, and Keshet \\& Waxman \\cite{KeshetWaxman2005} for an analytical solution for step shocks). The turbulent waves are, of course, also affected by the shock, and thus one has to start from the transmission of turbulence through the shock to arrive at describing particle acceleration in a case where wave-propagation effects are taken into account.  The effect of the transmission of Alfv\\'en waves through a step shock was solved at the nonrelativistic limit by Vainio \\& Schlickeiser (\\cite{VS1998}) and later generalised to include also relativistic shocks by Vainio et al.\\ (2003, hereafter referred to as \\cite{VVS2003}, and \\cite{VVS2005}). They found that, regardless of the shock speed, waves initially in equipartition were propagating predominantly antiparallel to the flow direction (i.e.\\ backwards toward the shock in the downstream plasma frame) after being transmitted through a step shock with Alfv\\'enic Mach number sufficiently low to yield non-negligible wave speeds compared to the speed of the plasma flow. They showed that this leads to an increased \\textit{scattering-centre compression ratio}, i.e., that the scattering-centre speed can undergo compression that is significantly higher than the compression of the flow itself. This effect thus causes the first-order Fermi acceleration to lead to accelerated particle energy spectral indices harder than $\\sim 2.2$, the well-known outcome of parallel step shocks at the ultrarelativistic limit in case of isotropic particle scattering (e.g., Kirk \\& Duffy \\cite{KD1999}, Keshet \\& Waxman \\cite{KeshetWaxman2005}). This transmission analysis is, however, valid only for shocks of small thickness and waves with wavelengths much longer than the width of the shock transition.  In this paper we continue studying the transmission of Alfv\\'en waves through parallel shocks, but take an approach opposite to the step-shock approximation of \\cite{VVS2003}. We calculate the transmission coefficients and their derivative parameters (the average cross-helicity of the waves, and the compression ratio of the scattering centres) for a continuous flow profile, approximating the transmission of waves through a shock with thickness sufficiently large for the waves to see the flow parameters varying slowly. The transmission coefficients for the nonrelativistic thick-shock case were calculated by \\cite{VV2005AA}.  In general, a single shock wave is not either a step-shock nor a thick shock: as the turbulence in the upstream medium has a broad-band spectrum, some waves will see the shock as a thick while others see it as a thin structure. In addition, some waves have wavelengths matching the shock thickness. Thus, the general theory of wave propagation can not rely on either the thick or thin shock approximation. However, the general theory is mathematically very cumbersome and we have here chosen to limit our discussion to waves with wavelengths much shorter than the shock thickness. Combining these results with our earlier ones (\\cite{VVS2003}) with the opposite assumption gives the first qualitative view on turbulence transmission through parallel shocks with finite thickness. ",
        "conclusions": "We have calculated the transmission coefficients for Alvf\\'en (test) waves with short wavelengths through a modified shock with background plasma properties changing slowly compared to the wave length scales.  While the transmission trough a step shock was found similar throughout the whole shock speed range (\\cite{VVS2003}), and, additionally, very similar behaviour was observed also for thick shocks in the nonrelativistic limit (\\cite{VV2005AA}), qualitative and signifigant differences emerge for thick shocks when the speed increases. As a consequence, from the point of view of the particle acceleration studies, the compression felt by the particles (i.e. that of the scattering centres) at a thick relativistic low-$M$ shock was shown to be weaker than the compression of the background flow. For slower (from non- to mildly relativistic) shocks the scattering centre compression ratio was shown to increase, like in step shocks."
    },
    "0609/astro-ph0609636_arXiv.txt": {
        "abstract": "We compare observations of the high redshift galaxy population to the predictions of the galaxy formation model of \\citet{Croton2006} and \\citet{Delucia2006b}. This model, implemented on the Millennium Simulation of the concordance $\\Lambda$CDM cosmogony, introduces ``radio mode'' feedback from the central galaxies of groups and clusters in order to obtain quantitative agreement with the luminosity, colour, morphology and clustering properties of the present-day galaxy population. Here we construct deep light cone surveys in order to compare model predictions to the observed counts and redshift distributions of distant galaxies, as well as to their inferred luminosity and mass functions out to redshift 5. With the exception of the mass functions, all these properties are sensitive to modelling of dust obscuration. A simple but plausible treatment agrees moderately well with most of the data. The predicted abundance of relatively massive ($\\sim M_*$) galaxies appears systematically high at high redshift, suggesting that such galaxies assemble earlier in this model than in the real Universe. An independent galaxy formation model implemented on the same simulation matches the observed mass functions slightly better, so the discrepancy probably reflects incomplete galaxy formation physics rather than problems with the underlying cosmogony. ",
        "introduction": "\\label{sec:intro} Recent work has used the very large Millennium Simulation to follow the evolution of the galaxy population throughout a large volume of the concordance $\\Lambda$CDM cosmogony \\citep{Springel2005b, Croton2006,Bower2006,Delucia2006a,Delucia2006b}. By implementing ``semi-analytic'' treatments of baryonic processes on the stored merger trees of all halos and subhalos, the formation and evolution of about $10^7$ galaxies can be simulated in some detail. The inclusion of ``radio mode'' feedback from the central galaxies in groups and clusters allowed these authors to obtain good fits to the local galaxy population and to cure several problems which had plagued earlier galaxy formation modelling of this type. In particular, they were able to produce galaxy luminosity functions with the observed exponential cutoff, dominated at bright magnitudes by passively evolving, predominantly elliptical galaxies. At the same time, this new ingredient provided an energetically plausible explanation for the failure of ``cooling flows'' to produce extremely massive galaxies in cluster cores. Most of this work compared model predictions to the systematic properties and clustering of the observed low redshift galaxy population, or studied the predicted formation paths of massive galaxies. Only \\citet{Bower2006} compared their model in detail to some of the currently available data at high redshift.  In the present paper we compare these same data and others to the galaxy formation model of \\citet{Croton2006} as updated by \\citet{Delucia2006b} and made publicly available through the Millennium Simulation data site.\\footnote{http://www.mpa-garching.mpg.de/millennium; see \\citet{Lemson2006}}  Many recent observational studies have emphasised their detection of substantial populations of massive galaxies out to at least redshift 2 and have seen this as conflicting with expectations from hierarchical formation models in the $\\Lambda$CDM cosmogony \\citep[e.g.][] {Cimatti2002c,Im2002,Pozzetti2003,Kashikawa2003,Chen2003,Somerville2004}. This notion reflects in part the fact that early hierarchical models assumed an Einstein-de-Sitter cosmogony in which recent evolution is stronger than for $\\Lambda$CDM \\citep[e.g.][]{Fontana1999}, in part an underassessment of the predictions of the contrasting toy model in which massive galaxies assemble at high redshift and thereafter evolve in luminosity alone \\citep[see][]{Kitzbichler2006}. \\citet{Bower2006} find their model to be in good agreement with current observational estimates of the abundance of massive galaxies at high redshift, while our comparisons below suggest that the model of \\citet{Croton2006} and \\citet{Delucia2006b} appears, if anything, to {\\it overpredict} this abundance. As shown by these authors and particularly by \\citet{Delucia2006a} both models predict ``anti-hierarchical'' behaviour, in that star formation completes earlier in more massive galaxies. This behaviour clearly does not conflict with the underlying hierarchical growth of structure in a $\\Lambda$CDM cosmogony.  The current paper is organised as follows. In \\secref{sec:model} we briefly describe the Millennium Simulation and the fiducial galaxy formation model we are adopting. Where we have made modifications, most significantly in the dust treatment, these are described in detail. We also give a detailed account of how we construct mock catalogues of galaxies along the backward lightcone of a particular simulated field of observation. Many of our methods resemble those which \\citet{Blaizot2005} implemented in their MOMAF facility in order to enable mock observations of simulated galaxy catalogues of the same type as (though smaller than) the Millennium Run catalogues we use here. Our results are summarised in \\secref{sec:results} where we compare number counts as a function of apparent magnitude and redshift with the currently available observational data. We also compare the predicted evolution of the luminosity and stellar mass functions to results derived from recent observational surveys, and we illustrate how the population of galaxies is predicted to shift in the colour-absolute magnitude plane. Finally in \\secref{sec:discussion} we interpret our findings and present our conclusions. ",
        "conclusions": "\\label{sec:discussion}  The model we have used in this paper is that of \\citet{Springel2005b} and \\citet{Croton2006} and as updated by \\citet{Delucia2006b} and made public through the Millennium Simulation download site \\citep[see][]{Lemson2006}. Earlier work has compared this model to a wide range of properties of low redshift galaxies: their luminosity functions, their bi-modal luminosity-colour-morphology distribution and their Tully-Fisher relation \\citep{Croton2006}; their spatial clustering as inferred from two-point correlations \\citep{Springel2005b,Li2006,Meyer2006} and from fits to halo occupation distribution models \\citep{Wang2006,Weinmann2006}; their HI gas content \\citep{Meyer2006}; and their assembly histories within clusters \\citep{Delucia2006a,Delucia2006b}. Although \\citet{Croton2006} compared the evolution of the global star formation rate and the global black hole accretion rate of the model to observation, the current paper is the first to compare its predictions in detail with observations of high redshift galaxies.  Our comparison to galaxy counts, to redshift distributions and to observational estimates of luminosity and mass functions at high redshift paints a consistent picture despite large statistical uncertainties and some significant technical issues.  Our model appears to have too many relatively massive galaxies at high redshift and these galaxies appear to be too red. Thus, while we fit optical galaxy counts well up to densities of 30 gal/mag/arcmin$^2$, we start to overpredict numbers in the $K$ band at densities above about 3 gal/mag/arcmin$^2$. This overabundance of apparently red galaxies shows up in the redshift distributions as an overprediction of the number of galaxies with $K\\sim23$ to 25 at redshift between about 1 and 3.  These correspond to moderately massive systems near the knee of the luminosity function, and indeed, while our rest-frame $B$ luminosity functions appear compatible with observation out to $z\\sim 1$, at rest-frame $K$ our luminosity functions are noticeably high beyond $z=0.5$ except possibly for the brightest objects. The problem shows up most clearly in our mass functions which overpredict observationally estimated abundances by about a factor of 2 at $z=2$. Apparently the mass function of galaxies evolved more strongly in the real Universe than in our simulation.  A galaxy formation model with similar basic ingredients to ours, but with important differences of detail has been independently implemented on the Millennium Simulation by \\citet{Bower2006}. This model is also publicly available at the download site. It fits low redshift galaxy luminosity functions as well as our model, but the comparisons which \\citet{Bower2006} show to high-redshift luminosity and mass function data (essentially the same datasets we use here) demonstrate somewhat better agreement than we find in this paper.  In the mass range $10^{10}< h^2M_*/M_\\odot< 10^{11}$ the abundances predicted by their model are lower than ours by about 20\\% at $z=1$ and by about 30\\% at both $z=2$ and $z=3.5$, despite the fact that at $z=0$ the two models agree very well. This is consistent with the fact that their model forms 20\\% of all its stars by $z=3.2$ and 50\\% by $z=1.65$ whereas the corresponding redshifts for our model are $z=3.6$ and $z=1.9$.  These differences arise from details of the star formation and feedback models adopted in the two cases.  In summary, both the \\citet{Bower2006} simulation and our own are consistent with most current faint galaxy data.  Thus there seems no difficulty in reconciling the observed properties of distant objects with hierarchical galaxy formation.  The fact that predictions from the two simulations differ at a level which can be marginally separated by the observations, means that currently accessible properties of distant galaxies can significantly constrain models of this type, and hence the detailed physics which controls the formation and the observable properties of galaxies. The fact that the model  we test here apparently {\\it overpredicts} the abundance of moderately massive galaxies at high redshift, despite the fact that late merging plays a major role in the build-up of its more massive galaxies \\citep[e.g.][]{Delucia2006a,Delucia2006b}, demonstrates that current data are still far from constraining the importance of this process.  As the data improve, the models will have to improve also to remain consistent with them.  This interplay between theory and observation should eventually lead to a more convincing and more complete picture of how galaxies came to take their present forms.   {\\bf Acknowledgements:} We thank Jeremy Blaizot, Darren Croton and Gabriella De Lucia for useful discussions of a number of technical issues which arose during this project.  MGK acknowledges a PhD fellowship from the International Max Planck Research School in Astrophysics, and support from a Marie Curie Host Fellowship for Early Stage Research Training."
    },
    "0609/astro-ph0609546_arXiv.txt": {
        "abstract": "The Primordial Anisotropy Polarization Pathfinder Array (PAPPA) is a balloon-based instrument to measure the polarization of the cosmic microwave background and search for the signal from gravity waves excited during an inflationary epoch in the early universe. PAPPA will survey a $20\\deg ~\\times ~20\\deg$ patch at the North Celestial Pole using 32 pixels in 3 passbands centered at 89, 212, and 302 GHz. Each pixel uses MEMS switches in a superconducting microstrip transmission line to combine the phase modulation techniques used in radio astronomy with the sensitivity of transition-edge superconducting bolometers. Each switched circuit modulates the incident polarization on a single detector, allowing nearly instantaneous characterization of the Stokes I, Q, and U parameters. We describe the instrument design and status. ",
        "introduction": "Linear polarization of the cosmic microwave background (CMB) carries the oldest information in the universe. Quantum fluctuations of the space-time metric (gravity waves) excited only $10^{-35}$ seconds after the Big Bang impart a distinctive signature to the CMB polarization. Detecting this signal probes physics at energies near Grand Unification and could provide insight on quantum gravity.  \\begin{figure}[b] \\centerline{ \\psfig{file=pappa_irvine_fig_1.ps,width=4.0in,angle=90}} \\caption{ Angular power spectra of the temperature and polarization anisotropy. The blue band shows the gravity-wave signal at tensor/scalar ratio $0.01 < r < 0.1$. The red line shows the PAPPA B-mode sensitivity after projecting out foregrounds. PAPPA will provide a precise measurement of the E-mode polarization and can detect B-mode polarization at levels r = 0.1. } \\label{power_spectra_fig} \\end{figure}  CMB polarization results from Thomson scattering of CMB photons by free electrons. Scattering of an isotropic radiation field produces no net polarization, but a quadrupole moment in the incident radiation yields a polarized signal. The required quadrupole can result from either intrinsic fluctuations in the radiation field itself, or the differential redshift as gravity waves propagate through an isotropic medium. Temperature or density perturbations are scalar quantities; their polarization signal must therefore be curl-free. Gravity waves, however, are tensor perturbations whose polarization includes both gradient and curl components. In analogy to electromagnetism, the scalar and curl components are often called ``E'' and ``B'' modes. Only gravity waves induce a curl component: detection of a B-mode signal in the CMB polarization field is recognized as a ``smoking gun'' signature of inflation, testing physics at energies inaccessible through any other means \\cite{kamionkowski/etal:1997,seljak/zaldarriaga:1997}.  Figure \\ref{power_spectra_fig} shows the power spectra for the temperature and polarization. E-mode polarization depends on the quadrupole anisotropy of the temperature distribution seen by each scatterer within its horizon. The temperature distribution is observed to be anisotropic and the physics of Thomson scattering is well understood; hence, the CMB {\\it must} be partially polarized with E-mode amplitude predictable from the temperature anisotropy and background cosmology. Polarization at the predicted amplitude was first observed by the DASI collaboration \\cite{kovac/etal:2002} and later confirmed by multiple experiments.  Recent results from the WMAP mission suggest that an observable gravity-wave signature should exist. Quantum-mechanical fluctuations in the inflaton field produce both density perturbations and gravity waves. The transition out of the inflationary epoch breaks pure scale invariance; the resulting signature in the scalar density perturbations can be used to estimate the accompanying signal from tensor gravity waves. Given the WMAP measurement of a deviation from scale invariance, the simplest inflation models predict gravity waves with tensor/scalar ratio in the range $0.01 < r < 0.16$ \\cite{spergel/etal:2006, linde:2005, lyth/riotto:1999}. The corresponding B-mode signal has RMS amplitude 30--100 nK at degree angular scales. Signals at this level can be detected by a dedicated polarimeter, providing a crucial test of inflationary physics.   \\begin{figure}[b] \\centerline{ \\psfig{file=pappa_irvine_fig_2.eps,width=4.0in}} \\caption{Schematic planar polarimeter. A phase switch introduces a half-wave delay in one arm before the signals are combined and detected. The detector output has a dc term proportional to the unpolarized intensity (Stokes I) plus an ac term at the switch frequency proportional to the linear polarization. } \\label{iu_schematic_fig} \\end{figure}   Detecting the gravity-wave signature in polarization will be difficult. The first challenge is sensitivity. The gravity-wave signal is faint compared to the fundamental sensitivity limit imposed by photon arrival statistics. Even noiseless detectors suffer from this photon-counting limit; the only solution is to use multiple independent detectors. Equally important is the control of systematic errors. Aliasing of power from either unpolarized sources or the dominant E-mode polarization into a spurious B-mode pattern could overwhelm the primordial signal. The Primordial Anisotropy Polarization Pathfinder Array (PAPPA) uses phase-sensitive switching techniques to combine the sensitivity of transition-edge superconducting detectors with rigorous control of potential systematic errors. With no macroscopic moving parts, PAPPA is scalable to the kilo-pixel arrays anticipated for an eventual space mission dedicated to CMB polarimetry. ",
        "conclusions": ""
    },
    "0609/astro-ph0609293_arXiv.txt": {
        "abstract": "{Gravitational lensing time delays depend upon the Hubble constant and the density distribution of the lensing galaxies.  This allows one to either model the lens and estimate the Hubble constant, or to use a prior on the Hubble constant from other studies and investigate what the preferred density distribution is. Some studies have required compact dark matter halos (constant M/L ratio) in order to reconcile gravitational lenses with the HST/WMAP value of the Hubble constant (72 $\\pm$ 8 km\\,s$^{-1}$ Mpc$^{-1}$ and 72 $\\pm$ 5 km\\,s$^{-1}$ Mpc$^{-1}$, respectively).  This is in direct contradiction with X-ray, stellar dynamical, and weak lensing studies, which all point towards extended halos and isothermal density profiles.  In this work, we examine an up-to-date sample of 13 lensing galaxies resulting in a data set consisting of 21 time delays.  We select systems in which there is a single primary lensing galaxy (e.g. excluding systems undergoing mergers). Analysis is performed using analytic models based upon a power-law density profile ($\\rho\\propto r^{-\\eta}$) of which the isothermal profile is a special case ($\\eta$ = 2). This yields a value of $\\eta$ = 2.11$\\pm0.12$ (3$\\sigma$) for the mean profile when modeling with a prior on the Hubble constant, which is only consistent with isothermality within 3 $\\sigma$.  Note that this is a formal error from our calculations, and does not include the impact of sample selection or simplifications in the lens modeling. We conclude that time delays are a useful probe of density profiles, in particular as a function of the environment in which the lens resides, when combined with a prior on the Hubble constant. ",
        "introduction": "In 1964, Refsdal first suggested that time delays between gravitationally lensed images could be used to determine the Hubble constant $H_{0}$.  The appeal of this method was the fact that it did not depend upon the local distance scale and was a relatively straightforward single-step process.  In addition, the time delay theory upon which it is based was derived from the established theory of general relativity.  Only in recent years, however, with the availability of high resolution radio and optical imaging, has the method begun to yield results.  When focusing on individual lensing galaxies, a large number of lens modeling studies (e.g. Impey et al. 1998; Williams \\& Saha 2000; Winn et al. 2002) have obtained values of $H_{0}$ that have been lower than that from the \\textit{Hubble Space Telescope (HST) Key Project} (Freedman et al. 2001) which yields a value of $H_{0}$ = 72 $\\pm$ 8 ${\\rm km}\\,{\\rm s}^{-1}{\\rm  Mpc}^{-1}$, or indeed the \\textit{Wilkinson Microwave Anisotrophy Probe (WMAP)} which yields a very similar $H_{0}$ = 72 $\\pm$ 5 ${\\rm km}\\,{\\rm s}^{-1}{\\rm Mpc}^{-1}$, although this assumes a flat universe (Spergel et al. 2003).  In particular, a few studies have inferred that if one assumes that lenses have an isothermal mass distribution (flat rotation curves) the derived value is $H_{0}\\sim 50{\\rm km}\\,{\\rm s}^{-1}{\\rm Mpc}^{-1}$, whereas if they assume a constant mass-to-light ratio (M/L), which results in a falling rotation curve, then $H_{0}\\sim 70 {\\rm km}\\,{\\rm s}^{-1} {\\rm Mpc}^{-1}$, which is in agreement with the \\textit{HST/WMAP} finding (Kochanek 2002; others e.g. Impey et al. 1998; Rusin et al. 2003).  This constant M/L ratio implies a compact dark matter halo.  However, other studies such as weak lensing (e.g. Guzik \\& Seljak 2002), stellar dynamics (Rix et al. 1997; Romanowsky \\& Kochanek 1999; Gerhard et al. 2001; Treu \\& Koopmans 2002), and X-ray galaxy measurements (e.g. Fabbiano 1999; Lowenstein \\& White 1999), are consistent with near isothermal density profiles (i.e. extended dark matter halos).  Note however that in the strong lensing regime we are probing the inner few kpc of galaxies (i.e. within the Einstein radius), whereas other studies, such as galaxy-galaxy lensing, are sensitive to more extended scales.  The aim of this paper is to study whether time delays from multiple image lens systems statistically point towards extended dark matter halos which are inferred from other measurements, when an external prior is adopted for $H_{0}$.   We perform the study without the complexity of detailed lens modeling in an effort to provide insight with a simple model.  The outline of this paper is as follows. In \\S2 we introduce two analytic models for the gravitational time delay based upon an SIS potential and general power law potential.  \\S3 will discuss the effect of ellipticity and shear on the time delays.  In \\S4 we discuss the lens sample used as the data set in the investigation, and in \\S5 we explain the analysis of this data set and present results. \\S6 draws some conclusions based upon the findings. ",
        "conclusions": "Gravitational lens time delays directly depend upon the Hubble constant and the density distribution of the lensing galaxies.  This allows us to either model the lens and infer the Hubble constant, or to marginalize over the Hubble constant derived from other studies and determine the preferred density distribution. Our goal has been to estimate the typical slope of a massive galaxy's density profile, using simple models as in Witt et al. (2000).  Although some studies have indicated that it does not seem possible to reconcile an isothermal model for galaxy halos with the HST/WMAP value for the Hubble constant, rather requiring a compact halo (Impey et al. 1998; Kochanek 2002), our analysis indicates that, on average, this is not the case.  For the sample of relatively isolated galaxies which we consider, as defined in Section 4, the mean value was found to be $\\eta$ = 2.11$\\pm0.12$ (3$\\sigma$), when incorporating a prior on the Hubble constant from the HST Key Project. The quoted error arises formally from the calculation outlined in Section 5 and does not include any bias due to sample selection, or due to the simplifications of lens modeling - for example, not fitting for ellipticity or external shear. Our finding is in agreement with a study of 22 lens galaxies using self-similar mass models, which resulted in $\\eta$ = 2.07 $\\pm$ 0.13 (1$\\sigma$) (Rusin et al. 2003). The high reduced $\\chi^{2}$ of our best fit to the data is a combination of the large amount of scatter in the density profile slopes of the sample, and over-simplistic models for the lenses. The steeper slope obtained for lenses resulting in double image configurations is perhaps surprising in that multiple imaging is more effective for steeper profiles; however, the cross-section for quads scales roughly as $\\epsilon^{2}$, so this also a competing factor. Previous studies (e.g. Witt et al. 2000; Kochanek 2002) have shown that two-image systems are relatively insensitive to quadrupole structure, and above we briefly considered how deviations from SIS impact on the time delay surface and estimates of the slope of the density profile.  We also note that the compact halo that tends to $\\eta$ = 3 is not typical. However, when we consider the four systems PG1115+080, SBS1520-530, B1600+434, and HE2149-2745, there is a tendency towards a steeper than isothermal density profile, although in this case we are dealing with small number statistics with only four systems to constrain the $\\eta$ parameter.  Kochanek et al. (2006) have recently pointed out that at least three of these four systems might have steeper density profiles as a consequence of environment - the distinction being whether at one extreme the galaxy is a satellite galaxy (and partially stripped of its dark matter halo) or at the other extreme is a central galaxy in a group. The analysis of Kochanek et al. (2006), coupled with our findings for the remaining time delay systems ($\\eta$ $\\sim$ 2), suggests that strong gravitational lensing time delays should be seen as placing valuable constraints on galaxy and group formation scenarios, reflecting the interplay between baryons and dark matter on these scales. One step to assess the impact of galaxy environment on time delays would be to investigate the lensing properties of fully simulated groups - are the profiles derived from simulated lensing time delays consistent with these? Such a study would allow us to address the impact of complex environments (e.g. for the systems we excluded from our sample) on galaxy profiles.  Ideally, more time delay measurements should be made to reduce the errors, and measurements made to increase the accuracy of the astrometry, particularly in systems where the lensing galaxy's position has notable uncertainties.  Moreover, increasing the sample size of time delay systems is crucial, as it puts further constraints on density profiles and reduces the random uncertainties due to large-scale structure. Looking towards the future, GAIA will uncover roughly fifty times more multiply imaged quasars than we have identified today (see Gaia study report); of these a few percent will have variable radio emission and hence be potential candidates for radio time delay systems."
    },
    "0609/astro-ph0609764_arXiv.txt": {
        "abstract": "We present a maximum likelihood method for fitting two-dimensional model distributions to stellar data in colour-magnitude space. This allows one to include (for example) binary stars in an isochronal population. The method also allows one to derive formal uncertainties for fitted parameters, and assess the likelihood that a good fit has been found. We use the method to derive an age of $38.5^{+3.5}_{-6.5}$Myrs and a true distance modulus of $7.79^{+0.11}_{-0.05}$ mags from the $V$ vs $V-I$ diagram of NGC2547 (the uncertainties are 67 percent confidence limits, and the parameters are insensitive to the assumed binary fraction). These values are consistent with those previously determined from low-mass isochronal fitting, and are the first measurements to have statistically meaningful uncertainties. The age is also consistent with the lithium depletion age of NGC2547, and the HIPPARCOS distance to the cluster is consistent with our value.  The method appears to be quite general and could be applied to any N-dimensional dataset, with uncertainties in each dimension. However, it is particularly useful when the data are sparse, in the sense that both the typical uncertainties for a datapoint and the size of structure in the function being fitted are small compared with the typical distance between datapoints. In this case binning the data will lose resolution, whilst the method presented here preserves it.  Software implementing the methods described in this paper is available from http://www.astro.ex.ac.uk/people/timn/tau-squared/. ",
        "introduction": "\\label{intro}  The extraction of astrophysical parameters from colour-magnitude diagrams (CMDs), has been a crucial technique for astronomy since the discovery of the CMD as a diagnostic tool \\citep[almost certainly attributable to][] {1911POPot..63.....H}. Since a coeval population of singe stars occupies a curve in a CMD, comparison with theoretical isochrones should allow a determination of global properties of the population such as age, distance and metallicity. Unfortunately such determinations have been hampered by the lack of good statistical methods for carrying out the comparison between observation and theory. For galactic astronomy, the main technique has been a visual comparison of isochrones with the data (although more sophisticated methods have been used for resolved populations in other galaxies). This not only leads to questions of objectivity, but also makes it impossible to derive statistically meaningful uncertainties for parameter estimates.  Were the problem simply fitting a set of datapoints with uncertainties in one dimension (say colour) to a curve then classical $\\chi^2$ analysis would suffice. Unfortunately a datapoint in colour-magnitude space has uncertainties in both colour and magnitude. \\citep[In addition the uncertainties are normally correlated, but as shown by][ this can be overcome by transforming the problem into a magnitude-magnitude space.]{1996ApJ...462..672T} This problem can still be solved analytically if the curve is actually a straight line \\citep[][and references therein]{Nerit}. \\cite{1982ApJ...263..166F} extended this analytical approach to the general case of a curve by a small curvature approximation. Their method has been used on a significant volume of data, including globular clusters \\citep{1997A&AS..121..499B, 1993AJ....105.1420D} single-age extra galactic populations \\citep{1999A&AS..134...21G} and young ($<$ 10Myr-old) populations \\citep{1997A&A...324..549T, 1996A&A...312..499J}. None of these studies make significant use of the uncertainty measurements, partly because of systematics, but partly there is also the comment that they produce shallow $\\chi^2$ spaces \\citep{1986ApJ...307..738H} which result in large derived uncertainties \\citep{1991MNRAS.250..314N}. This is clearly in part because the isochrones do not fit the data well, but may also be a warning that, although one can place clusters in an age sequence by eye, the absolute values of the ages, which must be derived by comparison with the model isochrones, are not as precise as we might hope.  There is a further limitation of the \\cite{1982ApJ...263..166F} technique, pointed out most explicitly by \\cite{1998A&A...337..125G}; no population of stars consists entirely of single stars. Unresolved binaries make up a significant fraction of most photometric samples, and are seen as objects which lie up to 0.75 mags brighter than the single star sequence. Indeed, in some coeval populations a distinct equal-mass binary sequence is observed 0.75 mags above the single-star sequence, with unequal-mass binaries lying between the two. Whilst one may be able to extract a single-star sequence by eye and then fit it \\citep{1992AJ....103..131H}, clearly the best way is to fit the binaries as well.  Thus one arrives at the fundamental question we address in this paper. If the model is a two dimensional distribution in the colour-magnitude plane, can we derive a statistical test to fit the data to the model? There has been some interest in using Bayesian methods to determine the age of each star in a CMD \\citep[][and references therein] {2005A&A...436..127J}, and then using the mean for the cluster age. Although \\cite{2006astro.ph..3493V} demonstrate such a technique for age determination, it is clear their work will be developed to fit other parameters as well. The problem here, though, is the absence of a goodness-of-fit parameter to choose between isochrones. Another obvious solution is to bin the data into pixels, and compare this with a model. \\cite{1997NewA....2..397D} and \\cite{1997AJ....114..680A} have developed this technique for large extra-galactic populations, with \\cite{2002MNRAS.332...91D} bringing much of the literature together into a cohesive method. The problem is, however, that our data are often sparse, by which we mean the typical separation of datapoints is large compared with their uncertainties (see Figure \\ref{first_sim}). Then binning the data simply has the effect of washing out our hard-won photometric precision.  \\cite{1996ApJ...462..672T} developed a technique which does retain the datapoints as points, and which can been seen as a relative of the method we use here. They created simulated observations with a similar number of datapoints to the observed dataset, and then used the distances in colour-magnitude space between the points of the simulated and actual observations as a fitting statistic. Our method, first presented in \\cite{2005prpl.conf.8502N}, improves on this by using a quasi-continuous 2D distribution as the model, which overcomes problems of sampling the model into a finite number of datapoints, and allows robust uncertainties to be derived.  The method we are proposing is relatively intuitive, so rather than embarking first on a formal analytical proof, we first give the intuitive interpretation (Section \\ref{intuitive}), and then discuss a numerical experiment which demonstrates the technique using a simulated observation, allowing us to conclude that it recovers the correct answer and uncertainties (Section \\ref{num_exp}). The formal proofs are in Sections \\ref{formal} and \\ref{special}, and the details of practical implementation in Section \\ref{2d_approach}. We draw all the work together in an example using real data in Section \\ref{2547}, before reaching our conclusions in Section \\ref{conclusions}. An alternative to reading the paper in this order would be to gain a working understanding from Sections \\ref{intuitive}, \\ref{num_exp}, and \\ref{2547}, and try the worked examples available with the software from\\break http://www.astro.ex.ac.uk/people/timn/tau-squared. ",
        "conclusions": "\\label{conclusions}  We have developed a maximum likelihood method for determining parameters for an isochronal population which contains binaries, from its colour-magnitude diagram. We have used numerical simulation to demonstrate it is correct, and used it on a practical example. There is clearly scope for further development. Most obviously one could search many more parameters than we have, determining, for example, binary fraction and mass ratio distribution, mass function, metallicity, or extinction. Several of these could be allowed to vary simultaneously.  One could also use this as a search statistic, looking for populations of a given age in large are surveys. Here the absolute value of $\\tau^2$ would measure how likely a given ``sequence'' is to have occurred by chance. Furthermore, one could not only search an N-dimensional colour magnitude space, but might also wish to use other parameters, such as position on the sky modelled against a clustered distribution. Finally there is also a range of other problems to which the technique might be applied such as modelling mass segregation in the mass-radius diagram \\citep[e.g.][]{2003MNRAS.345.1205L}, and one could even conceive of a replacement for the 1D Kolmogorov-Smirnov test where the datapoints had associated uncertainties."
    },
    "0609/astro-ph0609287_arXiv.txt": {
        "abstract": "We present new results on the cosmological evolution of the near-infrared galaxy luminosity function, derived from the analysis of a new sample of $\\sim 22,000$ \\kab $\\le 22.5$ galaxies selected over an area of 0.6 square degrees from the Early Data Release of the UKIDSS Ultra Deep Survey (UDS). Our study has exploited the multi-wavelength coverage of the UDS field provided by the new UKIDSS WFCAM $K$ and $J$-band imaging, the Subaru/XMM-Newton Deep Survey and the {\\it Spitzer}-SWIRE Survey. The unique combination of large area and depth provided by this new survey minimises the complicating effect of cosmic variance and has allowed us, for the first time, to trace the evolution of the brightest sources out to $z \\simeq 2$ with good statistical accuracy.  In agreement with previous studies we find that the characteristic luminosity of the near-infrared luminosity function brightens by $\\simeq 1$ magnitude between $z=0$ and $z  \\simeq 2$, while the total density decreases by a factor $\\simeq 2$. Using the rest-frame $(U-B)$ colour to split the sample into red and blue galaxies, we confirm the classic luminosity-dependent colour bimodality at $ z \\lsim 1$. However, the strength of the colour bimodality is found to be a decreasing function of redshift, and seems to disappear by $z \\simgt 1.5$. Due to the large size of our sample we are able to investigate the differing cosmological evolution of the red and blue galaxy populations. It is found that the space density of the brightest red galaxies ($M_K \\le -23$) stays approximately constant with redshift, and that these sources dominate the bright-end of the luminosity function at redshifts $z\\lsim 1$. In contrast, the brightening of the characteristic luminosity and mild decrease in space density displayed by the blue galaxy population leads them to dominate the bright-end of the luminosity function at redshifts $z\\gsim 1$. ",
        "introduction": "The potential of deep near-infrared observations to measure the cosmological evolution of the stellar mass of galaxies has long been understood (e.g. Lilly \\& Longair 1984; Dunlop et al. 1989; Glazebrook et al. 1995; Cowie et al. 1996). Now, with the advent of wide-field near-infrared detectors, this potential has begun to be realized, and several deep infrared-based studies supported by spectroscopic and/or photometric redshift information have recently been completed (e.g. Cimatti et al. 2002; Pozzetti et al. 2003; Dickinson et al. 2003; Drory et al. 2003; Fontana et al. 2004; Glazebrook et al. 2004; Caputi et al. 2006; Drory et al. 2005).  A key aim of this work is the determination of the cosmological evolution of the $K$-band galaxy luminosity function (LF), the present-day form of which is now reasonably well established (Cole et al. 2001; Kochanek et al. 2001). The most recent studies have been deep enough, and have been armed with sufficient redshift information to attempt to trace this evolution out to $z = 2-3$. The results indicate a substantial brightening in characteristic magnitude from $z =0 $ up to $z = 2-3$ along with a simultaneous decrease in galaxy number density (see e.g. Caputi et al. 2006; Saracco et al. 2006). A similar combination of luminosity and density evolution has also been observed in the optical LF (e.g. Lilly et al 1995; Poli et al. 2003; Gabasch et al. 2004; Giallongo et al. 2005).  However, even these most recent deep near-infrared surveys have still only covered relatively small areas of sky (from a few $\\rm arcmin^2$ in the Hubble Ultra Deep Field to several hundred $\\rm arcmin^2$ in the Great Observatories Origins Deep Survey (GOODS)) and are hence existing results on the evolving LF are still potentially heavily affected by cosmic variance. In this paper we exploit the order-of-magnitude improvement in areal coverage (0.6 sq. deg.) provided by the the UKIRT Infrared Deep Sky Survey (UKIDSS; Lawrence et al. 2006) Ultra Deep Survey (UDS) to constrain the evolution of the near-infrared LF and the colour bimodality up to $z \\simgt 2$.  The rest of this paper is structured as follows. In Section 2 we summarize the main properties of the datasets used in this work. Then, in Section 3 we describe in detail how our new $K$-band galaxy sample has been constructed and cleaned, and describe the method we have developed for the estimation of redshifts from the available multi-band photometry. In Section 4 we assess and discuss the reliability of our photometric redshift estimates and in Section 5 we derive the near-infrared LF. The evolution of the colour bimodality is presented in Section 6 while in Section 7 we derive separate LFs for the populations of red and blue galaxies. Our conclusions are presented in Section 8.  Throughout this paper all magnitudes are given in the AB system (Oke \\& Gunn 1983) and we have adopted a background cosmology with $\\Omega_{\\rm M}=0.3$, $\\Omega_\\Lambda=0.7$ and $H_0=70~{\\rm km~s^{-1}}$. ",
        "conclusions": "We have presented new results on the evolving near-infrared galaxy luminosity function, based on the analysis of a large sample of $K$-band selected galaxies drawn from the Early Data Release of the UKIDSS Ultra Deep Survey. The complete sample includes $\\sim 22,000$ sources down to \\kab $\\le 22.5$ over an area of 0.6 square degrees. Exploiting the complementary optical and mid-IR data provided by the Subaru/XMM-Newton Deep Survey and {\\it Spitzer}-SWIRE Survey we have obtained reliable photometric redshift estimates for the sources in the sample. This allowed us to study the evolution of the $K$-band LF in the redshift range $0.25 \\simlt z \\simlt 2.25$. In agreement with previous studies, we find that the characteristic luminosity of the near-infrared luminosity function exhibits a substantial brightening of $\\simeq 1$ magnitude between the local Universe and $z \\simeq 2$, while the total density decreases by a factor $\\simeq 2$ over the same redshift interval.  We have also used the rest frame $(U-B)$ colour to separate the sample into red and blue galaxies. We find the classical bimodal colour distribution to be clearly present at $z \\simlt 1$. However, the strength of the colour bimodality is found to decrease with increasing redshift and virtually disappears by $z \\simgt 1.5$. The progressively improvement of the UDS data will assess how the errors on the photometry affect this result.  By using the colour - magnitude relation to split the sample we have analysed the cosmological evolution of the red and blue galaxy populations. The blue objects clearly dominate the faint end of the LF over the whole redshift range explored in this work. On the other hand, up to $z \\simeq 1$ the brightest sources are mainly provided by red objects. In fact, the space density of red sources with $M_K \\le -23$ is approximately constant with redshift while the blue galaxies, with the same bright absolute magnitude, exhibit a strong increase in their space density with redshift. Therefore, this different evolutionary behaviour of the two populations determines the preponderance of blue objects in the bright-end of the luminosity function at redshifts $z \\simgt 1$.  Compared with previous studies, the unique combination of large area and depth of the UDS minimises the effect of the cosmic variance and has led to a substantial improvement in the determination of the LF. In particular, the large area coverage has allowed us to accurately trace the evolution of the brightest sources at $z \\lsim 2$ with a high level of statistical significance. As shown in Figure \\ref{sd}, the space density of the total population of bright sources ($M_K \\le -23$) is roughly constant over the redshift range $0.5 \\simlt z \\le 1.5$. This strongly suggests the vast majority of the present-day bright elliptical galaxies were already in place at $z \\simeq 1.5$.  On the other hand, our analysis has also shown a change in the fraction of red and blue galaxies with redshift. The majority of bright sources at $z \\simlt 1$ exhibit red colours consistent with having an old stellar population formed at high redshift. However, the increase of the fraction of blue galaxies with increasing redshift -- indicative of an increase of the ongoing star formation -- suggests the major formation epoch of the these bright sources to be at $z \\simgt 1.5$.  However, it is worth noting that the colour classification based on the rest-frame ($U-B$) colour is not very accurate for distinguishing between the different populations of galaxies that show the same optical colours, i.e. the old passive evolving sources and the dusty star-forming one. In fact, studies of  the population of Extremely Red Objects (EROs) at redshifts $1 \\simlt z \\simlt 2$ revealed this population to be composed by both old passive evolving galaxies and dusty star-bursts, roughly in the same proportion (Cimatti et al. 2002; Smail et al. 2002; Yan \\& Thompson 2003; Cimatti et al. 2003; Gilbank et al., 2003; Moustakas et al. 2004; Simpson et al. 2006b).  On the other hand, not all the blue objects are experiencing their major event of star formation. In fact, it only requires  less than 10\\% of the total mass of an early type galaxy to be involved in a recent episode of star formation to completely alter the colour of the objects and show it as a blue galaxy (Zepf 1997). This is because even a relatively small number of bright young stars dominates the emission at optical wavelengths. Moreover, a recent study from the VVDS in the redshift range $ 0.4 \\simlt z \\simlt 0.8$ by Ilbert et al (2006) found that $\\sim 30$\\% of bulge dominated sources exhibit blue colours, indistinguishable from the late type population (see also Im et al. 2001; Menanteau et al. 2004; Cross et al. 2004).   It is also interesting to notice that the brightening of the characteristic luminosity ($M_K^*$) of the whole $K$-band selected LF is consistent with the passive evolution of a single stellar population. Figure \\ref{evol} shows the evolution of $M_K^*$ derived from our likelihood analysis relative to $M_K^*$ at z=0 (Kochanek et al. 2001). Within the errors, our data agree remarkably well with the evolution of the $K$-band absolute magnitude displayed by a single burst model from Bruzual \\& Charlot (2003) with formation redshifts in the range $3 \\simlt z \\simlt 6$. This result, coupled with the constant space density of bright sources up to $z \\simeq 1.5$, suggests the major assembly epoch of bright sources to be at $z \\simgt 2$, even though spread over a large range in redshift. In fact, as shown from Figure \\ref{sd}, nearly 40\\% of the bright sources at $z \\simeq 1.5$ exhibit red colours. Assuming that half of them are dusty starbursts -- as suggested by studies of ERO populations (e.g. Cimatti et al. 2002) -- we end up inferring that at least 20\\% of sources are old passively evolving galaxies at these redshift. Therefore, this implies that a reasonable fraction of bright sources at $z \\simeq 1.5$ was assembled at very high redshift with very efficient star formation."
    },
    "0609/astro-ph0609078_arXiv.txt": {
        "abstract": "{ We present a systematic analysis of all the \\sax data of the soft gamma-ray repeater SGR\\,1900+14: these observations allowed us to study the long term properties of the source quiescent emission. In the observation carried out before the 1998 giant flare the spectrum in the 0.8--10 keV energy range was harder and there was evidence for a 20--150 keV emission, possibly associated with SGR\\,1900+14. This possible hard tail, if compared with the recent \\igr detection of SGR\\,1900+14, has a harder spectrum (power-law photon index $\\sim$1.6 versus $\\sim$3) and a \\mbox{20--100 keV} flux $\\sim$4 times larger. In the last \\sax observation (April 2002), while the source was entering the long quiescent period that lasted until 2006, the \\mbox{2--10 keV} flux was $\\sim$25\\% below the historical level. We also studied in detail the spectral evolution during the 2001 flare afterglow. This was characterized by a softening that can be interpreted in terms of a cooling blackbody-like component. ",
        "introduction": "Soft Gamma-ray Repeaters (SGRs) are a small group of high-energy sources, originally discovered through the emission of their characteristic short $\\gamma$-rays bursts. Only four confirmed \\sgr are known, plus two candidates. SGR bursts have typical duration of the order of 0.1 s, peak luminosity in the \\mbox{$10^{39}$\\textbf{--}$10^{42}$ erg s$^{-1}$} range, and are emitted during ``active'' periods that can last from weeks to months. Exceptionally large outbursts are also observed in SGRs. These rare events have properties clearly different from those of the usual short \\mbox{$\\gamma$-ray} bursts. Based on their intensities, they can be classified either as ``giant'' flares, with total released energy up to \\mbox{$10^{47}$ erg}, or ``intermediate'' flares, with total energy  smaller by orders of magnitude ($10^{41}$--$10^{43}$ erg). In the classical X-ray range \\mbox{($\\sim$1--10 keV)} SGRs are relatively steady sources with luminosity in the $10^{35}$--$10^{36}$ \\lum range (although fainter states have also been observed, see \\citealt{kouveliotou03} and \\citealt{mereghetti06}) and showing periodic pulsations with periods of several seconds and secular spin-down of \\mbox{$\\sim$$10^{-11}$--$10^{-10}$ s s$^{-1}$}.\\\\ \\indent It is generally thought that SGRs, as well as a group of similar pulsars known as Anomalous X-ray Pulsars (AXPs), are magnetars, i.e. isolated neutron stars with strong magnetic fields (see \\citealt{woods04} for a review of this class of objects). In the magnetar model both the persistent X-ray emission and the bursts are powered by magnetic energy \\citep{duncan92,thompson95,thompson96}. If the secular spin-down observed in \\sgr is attributed to dipole radiation losses, as in ordinary radio pulsars, magnetic fields of the order of $10^{14}$--$10^{15}$ G are inferred.\\\\ \\indent In this paper we focus on SGR\\,1900+14, reporting all the observations of this source carried out with the \\sax satellite. Although some of these data have been already published \\citep{woods99b,woods01,feroci03}, we reanalyzed all the data sets following the same procedure, in order to compare them in a consistent way. In fact these observations, spanning five years and covering different states of bursting\\,/\\,flaring activity, give the possibility to investigate the long term  spectral and flux variability of the source with a homogeneous data set.\\\\ \\indent In Section \\ref{review} we briefly review some results on SGR 1900+14, in the context of the activity history of the source. The spectral  and timing analysis are reported in Sections 3 and 4, where we focus on the long term changes in the 1--10 keV emission properties. In Section 5 we report evidence for the detection of SGR 1900+14 in the \\mbox{20-150 keV} band during one of the \\sax observations. In Section 6 we concentrate on the spectral variability on short time-scales  following the April 2001 intermediate flare. ",
        "conclusions": "Our re-analysis of the \\sax data of \\src confirms the spectral variability found in this source by \\citet{woods99b,woods01}, and \\citet{feroci03} on yearly time-scale. Since they found that in some observations an additional blackbody component was required, we were interested in a more thorough assessment of its possibly persistent presence. Such a two-components spectrum is one of the characteristics of the AXPs \\citep{mereghetti02} and has also been observed in the other well studied soft gamma-ray repeater \\mbox{SGR\\,1806--20} \\citep{mte05}. Although formally required only in two (possibly three) observations, that component might well be a permanent feature, always present in this source. In fact, except during the aftermath of the April 2001 flare, its temperature ( \\mbox{$\\sim$0.4 keV}) and emitting area (\\mbox{$\\sim$6--7 km}) are consistent with all the spectra. If we assume an underlying and nearly steady blackbody, it might be that, as proposed by \\citealt{woods99b} and \\citealt{kouveliotou01}, this spectral component is visible only in the observations that offer both a low power-law flux and good statistics.\\\\ \\indent The long term spectral variability seems to correlate with the occurrence of the giant and intermediate flares and, in a more complex way, with the ordinary bursting activity. Comparing the only \\sax pre-flare observation with the quiescent post-flare ones,  there is evidence for a softening in the spectrum. Also SGR\\,1806$-$20 after its 2004 December 27 giant flare displayed a softer spectrum with respect to the 2004 levels \\citep{rtm05}. This is qualitatively consistent with the magnetar scenario, in which the spectral hardening is linked to the increasing torque of the twisted magnetosphere, that finally drives the SGR to a giant flare \\citep{tlk02,mte05}. Then, after the flare, the source magnetosphere is foreseen to relax into a less twisted configuration, with a softer spectrum.\\\\ \\indent The most recent \\sax observation of \\src (Observation I, April 2002), shows a small but statistically significant fading compared to the preceding observations. A long term monotonic decrease of the X--ray emission has been observed in SGR\\,1627$-$41 \\citep{kouveliotou03,mereghetti06} from 1998 to 2004. During this period no bursts were recorded from SGR\\,1627$-$41, and its fading has been interpreted as due to the cooling of the neutron star surface after the heating occurred when the source was active in 1998. \\src was still moderately active during 2002 \\citep{hurley02}, but then no bursts were observed for several years. The smaller luminosity in the last \\sax observation might thus correspond to the initial part of a cooling and fading phase, at least qualitatively similar to that observed in SGR\\,1627$-$41, but now interrupted by the  recent (March 2006) reactivation  \\citep{palmer06,golenetskii06}.\\\\ \\indent During the afterglow of the 2001 April 18 flare, \\citet{feroci03} found a flux decrease and a spectral softening. Our re-analysis shows that the variable spectral component can be well modeled as an additional blackbody emitted from a smaller and hotter (but rapidly cooling) region of the neutron star surface. Successful attempts to explain observations of afterglow flux decays in magnetars by means of a cooling thermal component are described in \\citet{ibrahim01}, \\citet{lenters03}, and \\citet{woods04axp}; all these works point out evidence of cooling hot spots on the surface of the neutron star exposed to a fireball.  However we note that the occurrence of the bump in the light curve of the afterglow is an anomaly in the picture of the cooling of a thermal emission, since it requires a re-injection of energy; we refer to \\citealt{feroci03} for an extensive discussion of this issue.\\\\ \\indent Evidence for persistent emission above 20 keV for \\src has recently been obtained with \\igr observations \\citep{gotz06}. We found that a hard tail was visible also in the 1997 \\sax PDS data. If this emission is indeed due to SGR\\,1900+14, our 50 ks long observation indicates significant differences with respect to the average properties obtained with \\igr, based on the sum of many observations performed discontinuously  from March 2003 to June 2004. The PDS \\mbox{20--100 keV} flux is  $\\sim$4 times larger\\footnote{The uncertainty in the relative calibration of the two satellite in the energy band considered here is of $\\approx$10\\% \\citep{kirsch05}.} and the spectrum is harder (photon index $\\sim$1.1) than that measured with \\igr (photon index $\\sim$3). Even considering our fit based only on the PDS instrument, the difference in the hard X-ray spectral index is significant (photon index $\\sim$1.6 versus $\\sim$3). Another interesting indication from the PDS data is a decrease of the \\mbox{50--150 keV} flux of \\src after the giant flare: it is possible that the hard X-ray flux decrease and softening in \\src was  a consequence of the 1998 August 27 giant flare.\\\\ \\indent The only other SGR established as a persistent hard X-ray source to date is SGR\\,1806$-$20 \\citep{mereghetti05int,molkov05}. For this source observations carried out with \\xmm in the April 2003--October 2004 period, showed a progressive spectral hardening in the \\mbox{1--10 keV} band, as the source increased its burst rate before the giant flare \\citep{mte05}. The \\igr observations displayed some evidence of a similar behaviour above \\mbox{20 keV}. In fact its photon index varied from $\\sim$1.9 in the period March 2003--April 2004 to $\\sim$1.5 in September--October 2004 \\citep{mereghetti05int}. A comparison of the hard X-ray luminosity of the two \\sgr in the ``pre-flare'' state is subject to uncertainties in their distances. For \\src a distance of 15 kpc has been derived based on its likely association with a young star cluster \\citep{vrba00}, while for SGR\\,1806$-$20 the distance is rather debated and has been variously estimated from 6.4 kpc to 15 kpc \\citep{cameron05,mcclure05}. If we assume a distance of \\mbox{15 kpc} for both sources we obtain similar \\mbox{20--100 keV} luminosities: \\mbox{$(1.5\\pm0.3)\\times10^{36}$ \\lum} for \\src and \\mbox{$(1.2\\pm0.1)\\times10^{36}$ erg s$^{-1}$} for SGR\\,1806$-$20.\\\\ \\indent These results, together with the recent detections of several AXPs in the hard X-ray range \\citep{molkov04,kuiper04,revnivtsev04,denhatog06,kuiper06}  with \\mbox{20--100 keV} luminosities similar or larger than those below 10 keV, indicate that non thermal magnetospheric phenomena are energetically important in magnetars. Soft X-rays give only a partial view  and broad band observations are required for a better understanding of the physical processes occurring in these sources. In this respect,  \\src\\,  being probably the first magnetar showing evidence for variability in the hard X-ray range  and currently in a moderately active state, is a good target to further explore possible correlations between the persistent  emission and the bursting activity. \\indent"
    },
    "0609/astro-ph0609552_arXiv.txt": {
        "abstract": "We study the expected X-ray luminosity of stellar merger products several years after merger. The X-ray emission is assumed to result from magnetic activity. The extended envelope of the merger product possesses a large convective region and it is expected to rotate fast. The rotation and convection might give rise to an efficient dynamo operation, therefore we expect strong magnetic activity. Using well known relations connecting magnetic activity and X-ray luminosity in other types of magnetically active stars, we estimate that the strong X-ray luminosity will start several years after merger, will reach a maximum of  $L_x \\sim 3 \\times 10^{30} \\erg \\s^{-1}$, and will slowly decline on a time scale of $\\sim 100 \\yr$. We predict that X-ray emission from V838 Mon which erupted in 2002 will be detected in 2008 with 20 hours of observation. ",
        "introduction": "\\label{intro}   The eruption of V838~Mon in 2002 (Brown 2002) and subsequent studies of its observed evolution (Munari et~al. 2002; Kimeswenger et~al. 2002; Crause et~al. 2003; Kipper et~al. 2004; Tylenda 2005), as well as, of other similar objects, i.e. V4332 Sgr (Martini et~al. 1999; Tylenda et~al., 2005) and M31~RV (Mould et~al. 1990) have led to suggestions that these observed events were likely to be due to stellar mergers (Soker \\& Tylenda 2003, Tylenda \\& Soker 2006). Soker \\& Tylenda (2006), who termed these events mergebursts, discuss the different channels to produce a mergeburst.  For hours to months after merger the merger product is very luminous (e.g., Soker \\& Tylenda 2003; Bally \\& Zinnecker 2005; Tylenda 2005). For a grazing collision (namely, not a head on collision) an extended envelope is inflated by the merging stars. Still on a longer time scale, the mass loss process, both mass loss rate and geometry, is strongly influenced by the merger event (e.g. Morris \\& Podsiadlowski 2006). On a much later time of hundreds of years and longer, after the merger products reaches equilibrium, the process can alter the evolution of the star on the HR diagram (e.g., Podsiadlowski et al. 1990), like the formation of blue stragglers (e.g. De Marco et al. 2005; Sills et al. 2005).  In the present study we examine whether and when the merger product can become magnetically active, a process that might be observed in the X-ray and radio bands. ",
        "conclusions": "\\label{summary}  According to the stellar merger model of the V838 Mon outburst and similar merger products (which we term mergebursts), a large envelope is formed around the more massive of the two merging stars. The envelope then contracts on a thermal time scale. The merger remnant should become a fast rotator as it contracts. As the remnant contracts more or less along the Hayshi line, its envelope possesses a large convective region. The fast rotation and the envelope convection are the two ingredients required in the $\\alpha \\omega$ dynamo model$-$a successful model for magnetic activity of main sequence stars, pre-main sequence stars, and subgiants.  We applied the $\\alpha \\omega$ model to contracting merger products by using the Rossby number (eq. \\ref{rossby}), and the relation between the Rossby number and X-ray luminosity known for magnetically active stars, scaled according to the expression for subgiant (or G giant) stars (eq. \\ref{lxlb}). We also assumed that after the contracting product reaches some fraction $\\eta_{\\rm max}$ of its break-up (Keplerian) velocity, this ratio does not increase any more, because a stellar wind removes angular momentum from the envelope. Our final (and conservative) prediction for the X-ray luminosity of magnetically active merger products are given by equation (\\ref{lxm1}) for merger products before they reach our assumed maximum rotation rate, and by equation (\\ref{lxm2}) for merger products rotating at $\\eta_{\\rm max}$.  In section \\ref{v838} we apply  the results to our model of V838 Mon. The results are presented in Fig. \\ref{lumx} for three values of the assumed maximum rotating rate $\\eta_{\\rm max}$, as marked near the lines. For too large Rossby numbers $Ro \\ga 10$ (Pizzolato et al. 2003; we here take a stronger constraint of $Ro \\ga 5$) of  the $\\alpha \\omega$ dynamo is not efficient any more. The dotted lines are the evolutionary stages where the expected Rossby number of V838 Mon is $Ro>5$, and we expect no strong  magnetic activity.  >From Fig. \\ref{lumx} we learn the following. \\begin{enumerate} \\item  There is no magnetic activity at the first several years, and hence no X-ray emission is expected. The observation by Orio et al. (2003) was made a year after the outburst, when no magnetic activity and no X-ray emission is expected. \\item For a reasonable values of maximum rotation rate $0.4 \\la \\eta_{\\rm max} \\la 0.8$, V838 Mon will reach a maximum activity at 6-8 years after outburst. The expected X-ray luminosity then slowly declines. \\item The X-ray luminosity in the coming years will be $L_x \\sim 3 \\times 10^{30} \\erg \\s^{-1}$. At the distance of V838 Mon the expected X-ray flux is $F_X \\sim 6 \\times 10^{-16} {\\rm erg}\\,{\\rm cm}^{-2}\\,{\\rm s}^{-1}$. We estimate that with 100,000 seconds of observation this emission can be detected. \\end{enumerate}  We therefore highly encourage 100,000 seconds of X-ray observation of V838 Mon in 2008. Even if no X-ray is detected, the results is of some importance, as it can strongly constrain models for V838 Mon, e.g., rules out accreting white dwarf. Orio et al. (2003) noted that their null detection rules our a symbiotic-like event to the V838 Mon outburst.  We point out that the null detection of X-ray emission from two AGB stars (Kastner \\& Soker 2004) is not directly relevant to the case of V838 Mon. First, and most important, our prediction is based on the $\\alpha \\omega$ dynamo model, namely, the amplification of the magnetic field by the operation of both rotation and convection, which is known to be very efficient. On the other hand, predictions for AGB stars are based on the amplification of the magnetic field by convection alone (Soker \\& Zoabi 2002), which is thought to be much less efficient. Second, V838 Mon is an order of magnitude more massive than an upper AGB star. We predict the magnetic activity to take place when the radius, luminosity and temperature of V838 Mon is similar to that of upper AGB star. Due to the higher mass we expect the mass loss rate to be smaller, and the wind speed to be faster. Therefore, the column density to the expected X-ray emitting region will be much lower.  Finally, the magnetic fields might be detected also in masers spots. Deguchi (2005) and Claussen (2005) report the detection of SiO maser around v838 Mon. We predict that if maser emission, in SiO, H$_2$O, or OH, will be observed from 2007, some regions might show polarization indicating the presence of magnetic fields, similar to the case around AGB stars, e.g., Vlemmings et al. (2005)."
    },
    "0609/astro-ph0609308_arXiv.txt": {
        "abstract": "We report an apparent detection of the C\\emissiontype{VI} $4p$ to $1s$ transition line at 459~eV, during a long-term enhancement (LTE) in the {\\it Suzaku} north ecliptic pole (NEP) observation of 2005 September 2. The observed intensity of the line is comparable to that of the C\\emissiontype{VI} $2p$ to $1s$ line at 367~eV. This is strong evidence for the charge-exchange process. In addition to the C\\emissiontype{VI} lines, emission lines from O\\emissiontype{VII}, O\\emissiontype{VIII}, Ne\\emissiontype{X}, and Mg\\emissiontype{XI} lines showed clear enhancements. There are also features in the 750 to 900~eV range that could be due to some combination of Fe\\emissiontype{XVII} and \\emissiontype{XVIII} L-lines, higher order transitions of O\\emissiontype{VIII} ($3p$ to $1s$ and $6p$ to $1s$), and a Ne\\emissiontype{IX} line. From the correlation of the X-ray intensity with solar-wind flux on time scales of about half a day, and from the short-term ($\\sim 10$~minutes) variations of the X-ray intensity, these lines most likely arise from solar-wind heavy ions interacting with neutral material in the Earth's magnetosheath. A hard power-law component is also necessary to explain the LTE spectrum. The origin of this component is not yet known. Our results indicate that solar activity can significantly contaminate {\\it Suzaku} cosmic X-ray spectra below $\\sim 1$~keV. Recommendations are provided for recognizing such contamination in observations of extended sources. ",
        "introduction": "\\citet{Snowden_etal_1994} reported the existence of mysterious X-ray contamination episodes in the {\\it ROSAT} all sky survey data which they termed long-term enhancements (LTEs).  During an LTE, the X-ray counting rate in the lower energy bands as much as doubled on a time scale of 1--2~days.  However, they could not find any correlation with other observational parameters, such as spacecraft position or look direction. New insight on LTEs was obtained by the discovery of X-ray emission from comet Hyakutake (\\cite{Lisse_etal_1996}).  Following the discovery, X-rays were detected from many comets (e.g., \\cite{Dennerl_etal_1997,Cravens_2002}), and the emission mechanism is now well understood as charge exchange of solar-wind heavy ions with cometary neutrals (see \\cite{Krasnopolsky_etal_2004} for a review). Then \\citet{Cox_1998} and \\citet{Cravens_2000} suggested that solar-wind charge exchange with neutrals in the Earth's geocorona and in the heliosphere accounts for a part of soft X-ray background below 1~keV. \\citet{Robertson_etal_2001} showed that the LTEs of the {\\it ROSAT} all sky survey data were well correlated with the solar-wind proton flux, which strongly suggests the origin of the LTEs to be solar-wind charge exchange with H\\emissiontype{I} in the geocorona.  Solar-wind charge exchange results in line emission from highly ionized ions (\\cite{Krasnopolsky_etal_2004} and references therein).  {\\it ROSAT} did not have enough spectral resolution to resolve those lines.  Spectral information on geocoronal solar-wind charge exchange was first obtained during a {\\it Chandra} dark moon observation \\citep{Wargelin_etal_2004}. The X-ray photons detected in the direction of the dark moon are most likely from this source.  The emission spectrum could be described by a sum of C\\emissiontype{VI}, O\\emissiontype{VII}, and O\\emissiontype{VIII} K-lines, although the statistics and energy resolution were limited. More recently, \\citet{Snowden_etal_2004} reported time variation of soft X-ray intensity during the {\\it XMM-Newton} {\\it Hubble} deep field north observation. The enhancement was correlated with solar-wind proton flux variations. They detected C\\emissiontype{VI}, O\\emissiontype{VII}, O\\emissiontype{VIII}, Ne\\emissiontype{IX}, and Mg\\emissiontype{XI} emission lines in the enhancement.  The importance of solar-wind charge-exchange emission is three-fold. Firstly, it will enable us to remotely study low density neutrals in geocorona, in outer atmosphere of planets such as Jupiter, in interplanetary space, and in particular, around comets.  Secondly, it can be used as a highly sensitive ion probe for the solar wind.  Perhaps most importantly, it becomes a significant contaminating foreground in the study of the cosmic soft X-ray background below $\\sim$ 1~keV.  More than half of the soft X-ray background at these energies is considered to arise from hot gas in the disk and halo of our galaxy and in intergalactic space.  Although such emission was detected early in the history of X-ray astronomy (e.g. \\cite{Tanaka_1977}), its origins and the physical state of the hot gas are not yet well understood.  Geocoronal interactions with the solar wind produce at least a sporadic contamination that must be avoided in observations of any extended source.  Charge exchange with interstellar neutrals moving through the heliosphere creates a more subtle contribution where much of the diagnostic rapid time variation is washed out by the large travel time of solar-wind events through interplanetary space, and the spectral lines arise from the same ions expected in hot interstellar plasmas. \\citet{Lallement_2004} used a simplified model of interplanetary charge exchange to estimate that essentially all of the minimum flux observed near the Galactic plane in the 1/4~keV (R12) band of the {\\it ROSAT} sky survey could arise from this source (see also \\cite{Pepino_2004}).  We are far from an understanding of solar-wind charge exchange adequate to determine the true extent of this contribution.  The X-ray Imaging Spectrometer (XIS) on board {\\it Suzaku} \\citep{Suzaku} has a significantly improved spectral line response function compared to previous X-ray CCD cameras, particularly below 1~keV \\citep{XIS}. Together with the large effective area of the X-ray telescopes (XRT; \\cite{XRT}), this will open a new era for the study of soft X-ray background.  In this paper, we report on {\\it Suzaku} observations of a blank field in direction of the north ecliptic pole (NEP).  We detected a significant enhancement of the soft X-ray flux lasting for $\\sim 10$~hours.  The enhancement is mostly explained by increases in C\\emissiontype{VI} through Mg\\emissiontype{XI} emission lines.  During the enhancement, both C\\emissiontype{VI} $n=2$ to 1 and C\\emissiontype{VI} $n=4$ to 1 transition lines were clearly detected, which is a firm evidence for charge-exchange emission.  We consider that the emission is due to the charge-exchange interaction of solar-wind heavy ions with neutrals in the magnetosheath at 2--8 Earth radii ($R_{\\oplus}$).  In this paper, we will concentrate on the X-ray spectra and the emission processes, and their implications for cosmic X-ray observations.  The geophysical implications of the results will be discussed in a separate paper.  Errors quoted in the text and tables are at 90\\% confidence single parameter errors and at $1\\sigma$ confidence level in the figures, unless otherwise stated. ",
        "conclusions": "\\label{sec:discussion}  The short ($\\sim 10$~minutes) time scale variations observed during the enhancement of X-ray intensity imply that the size of the emission region is no larger than 10 light minutes. On the other hand, the apparent size of the emission region must be equal to or larger than the XIS field of view ($18'$). These require the emitter of the X-ray enhancement to be within a distance of $10~{\\rm light~minutes}/18'$, or $\\sim 10^{-3}$~pc.  Because the enhanced X-ray emission consists of emission lines from C\\emissiontype{VI} to Mg\\emissiontype{XI}, this requires an ion source within $10^{-3}$~pc. There is only one ion source in this distance range. That is the Sun.  The Sun may produce X-ray emission lines in our observations in two possible ways: scattering of solar X-rays by the Earth's atmosphere, and solar-wind charge exchange.  In the former case, the X-ray intensity is proportional to the solar X-ray intensity multiplied by the sunlit atmospheric column density.  The solar X-ray intensity is continuously monitored by the GOES (Geostationary Operational Environmental Satellites)\\footnote{The data available at http://www.ngdc.noaa.gov/stp/GOES/goes.html.}, but data show no correlation with the enhancement.  Moreover, using the MSIS atmosphere model \\citep{Hedin_1991}, we found that the column density of sunlit atmosphere varied by many orders of magnitude ($10^{9}$) during the observations, but no correlation was found with the observed X-ray intensity.  Thus scattering of solar X-rays can be excluded.  \\begin{figure} \\begin{center} \\FigureFile(80mm,60mm){figure-7.eps} \\end{center} \\caption{0.3--2 keV X-ray intensity (upper panel) and geocentric distance of the point whose geomagnetic field becomes open to the space for the first time along the line of sight from the spacecraft position in units of the earth radius, as a function of time during the observation (lower panel). See also the schematic view shown in figure~\\ref{fig:magneto_schematic}.} \\label{fig:magneto} \\end{figure}  \\begin{figure} \\begin{center} \\FigureFile(80mm,60mm){figure-8.eps} \\end{center} \\caption{Schematic view of the magnetosphere and the definition of $r_{\\rm mp}$ used in figure~\\ref{fig:magneto}.} \\label{fig:magneto_schematic} \\end{figure}  In figure \\ref{fig:light_curve}, we show the proton flux observed by the ACE (Advanced Composition Explorer)\\footnote{The data available at http://www.srl.caltech.edu/ACE/ASC/.} together with the {\\it Suzaku} X-ray counting rate.  The ACE data were shifted in time to account for propagation time from ACE to the Earth.  Clearly the proton flux was enhanced during the X-ray ``flare''. This is consistent with solar-wind charge-exchange model.  Charge-exchange X-ray emission also is strongly supported by the detection of the C\\emissiontype{VI} $n=4$ to 1 transition line (Ly$\\gamma$). In CIE (collisional ionization equilibrium) thermal emission, the Ly$\\beta$ and Ly$\\gamma$ lines of C\\emissiontype{VI} are suppressed relative to Ly$\\alpha$ by the Boltzmann factor in the distribution of exciting electrons. In charge exchange between C\\emissiontype{VII} and H\\emissiontype{I}, the electron is deposited primarily in the $n=4$ level (\\citet{Krasnopolsky_etal_2004} and references therein), and the X-ray lines are produced when it cascades to the $n=1$ level, guided only by branching ratios.  In high energy collisions, angular momentum states are tend to be populated statistically by weight of the state's degeneracy, and electrons are primarily captured into maximal $l$, where $n$ can change only by 1 unit at a time during the cascade \\citep{Beiersdorfer_2001}, again resulting in relatively weak Ly$\\beta$ and Ly$\\gamma$. Behind the Earth's bow shock, however, the solar wind velocity is reduced, which should result in recombination into low $l$-orbitals and strong Ly$\\beta$ and Ly$\\gamma$.  In the ``flare'' spectrum, both C\\emissiontype{VI} Ly$\\alpha$ and Ly$\\gamma$ lines were detected and this is a strong evidence for charge exchange.  In addition, O\\emissiontype{VIII} $n=3$ to 1 and $n=6$ to 1 transition lines are also enhanced in the comet emission model \\citep{Kharchenko_2003}.  However, we cannot separate those lines from Fe\\emissiontype{XVII}-L emission lines with the present energy resolution.  Although X-ray intensity and the proton flux are correlated on a time scale of $\\sim 10$ hours, they do not show much correlation on short time scales.  We consider that the short term X-ray intensity variation is at least partly arising from orbital motion of the spacecraft.  In figure~\\ref{fig:magneto}, we show the geocentric distance of the point whose geomagnetic field becomes open to the space for the first time along the line of sight from the spacecraft position, i.e., the point where the line of sight encounters the magnetosheath (see also the schematic view shown in figure~\\ref{fig:magneto_schematic} for the definition). We evaluated the end point of the magnetic field using the software GEOPACK-2005 and T96 magnetic field model (\\citet{Tsyganenko_2005} and references therein)\\footnote{The software package available at http://modelweb.gsfc.nasa.gov/magnetos/data-based/modeling.html.}.  We obtained the solar-wind parameters required to perform the calculation from the CDAWeb (Coordinated Data Analysis Web)\\footnote{The data available at http://cdaweb.gsfc.nasa.gov/cdaweb/sp\\_phys/.}.  We find that the line of sight during the present observation was rather special in the sense that it goes through the north magnetic pole region where charged particles of the magnetosheath can penetrate down to 2--$8R_{\\oplus}$ moving along open field lines. The short-term X-ray intensity variations during the time intervals shown by boxes in figure \\ref{fig:magneto} indicate anti-correlation with the distance to the magnetosheath.  This indicates that the charge exchange of solar-wind heavy ions is taking place at 2--$8R_{\\oplus}$ where the neutral matter density is high.  \\citet{Robertson_etal_2006} recently studied theoretically solar-wind charge-exchange emission from the magnetosheath.  Implications of the present results on the solar-wind ion composition and the Earth's magnetosheath including comparisons with the theoretical model will be reported in a separate paper.  Finally, since solar-wind charge-exchange emission can become a difficult foreground in the study of soft diffuse sources, we summarize a procedure to examine possible contamination in the {\\it Suzaku} spectra below $\\sim$ 1~keV by solar activity. \\begin{enumerate} \\item check the light curve, if time variation is not expected for the object, \\item check the solar X-ray intensity and the column density of sunlit atmosphere along the line of sight, \\item check the solar wind proton flux, \\item check the radius of the magnetosheath on the line of sight. \\end{enumerate}  \\bigskip  We would like to thank Prof.~K.~Maezawa, and Prof.~I.~Shinohara for their help in calculations of the Earth's magnetosheath and also for valuable discussions.  Thanks are also due to Prof.~T.~Mukai, Prof.~A.~Matsuoka, and Prof.~H.~Hayakawa for discussions on the relation between {\\it Suzaku} data and solar wind. We are grateful to the referee, Dr.~A.~Dalgarno, for useful comments to improve this paper."
    },
    "0609/hep-ph0609192_arXiv.txt": {
        "abstract": "\\baselineskip=22pt The process of Compton scattering $\\gamma e^{\\pm} \\to \\gamma e^{\\pm}$ in  strongly magnetized hot electron-positron plasma is considered. The analytical expressions for the partial cross sections in rarefied plasma are obtained. The numerical estimations of the absorption rates for various scattering channels are  presented  with taking into account of the photon dispersion and wave function renormalization in  strong magnetic field and plasma. The comparison of the scattering absorption rate with photon splitting probability shows the existence of the temperature range where these values are comparable with each other. The astrophysical applications of the obtaining results are discussed. ",
        "introduction": "\\unitlength 1mm Recent observations~\\cite{Kouveliotou:1998ze,Kouveliotou:1998fd,Gavriil:2002mc,Ibrahim:2002zw} give ground to believe that some astrophysical objects (SGR and AXP) are magnetars, a distinct class of isolated neutron  stars with magnetic field strength of $B \\sim 10^{14}-10^{16}$ G~\\cite{Duncan:1992,Duncan:1995,Duncan:1996}, i.e. $B \\gg B_e$, where $B_e = m^2/e \\simeq 4.41\\times 10^{13}$~G~\\footnote{We use natural units $c = \\hbar = k = 1$, $m$ is the electron mass,  $e > 0$ is the elementary charge.} is the critical magnetic field. The spectra analysis of these objects is also providing evidence for the presence of electron-positron plasma in magnetar environment. It is well-known that strong magnetic field and/or plasma could influence essentially on different quantum processes~\\cite{Duncan:2000pj,Lai:2001,KM_Book,Harding:2006qn}. Therefore the effect of  magnetized plasma on microscopic physics should be incorporated in the magnetosphere models of strongly magnetized neutron stars.   The various studies indicate that electromagnetic processes such as Compton scattering and photon splitting $\\gamma \\to \\gamma \\gamma$ (merging $\\gamma \\gamma \\to \\gamma$) could play a crucial role in these models. For example, the process of photon splitting in strong magnetic field could suppress  the production of electron-positron pair required for a detectable radio emission from magnetar~\\cite{BH:1998, BH:2000}. In turn, Compton scattering may play an important role in the models of atmosphere emission from strongly magnetized neutron stars~\\cite{Miller:1995, Bulik:1997}. The both processes are very important in the models the radiation emission from SGR burst~\\cite{Duncan:1995, HBG:1996, HBG:1997}.  Previously, Compton scattering was studied under various  conditions. The case of the magnetic field without plasma was investigated in~\\cite{Herold:1979, Melrose:1983, Harding:1986, Harding:2000}. It was found that Compton scattering becomes strongly anisotropic and essentially depends on photon polarization. In~\\cite{Bulik:1997} it was shown that the dispersion properties of photon in  strongly magnetized cold plasma could significantly influence on the process under consideration.  In the present work the process of photon scattering, $\\gamma e^{\\pm} \\to \\gamma e^{\\pm}$, is investigated in the presence of strong magnetic field and charge-symmetric electron positron plasma, when the magnetic field strength $B$ is the maximal physical parameter, namely $\\sqrt{eB} \\gg T,\\, \\omega, \\, E$. Here $T$ is the plasma temperature,  $\\omega$ and $E$ is the initial photon and electron energies. In this case almost all electrons and positrons in plasma are on the ground Landau level. We also consider the none resonance case, when electron and photon momenta satisfy the condition $eB \\gg (pk)$. Such conditions could be realized e.g. in the Thompson \\& Duncan model of SGR burst~\\cite{Duncan:1995} in which magnetically trapped high temperature ($\\sim 1$ MeV) plasma fireball could be created.   The main purposes of our work are: \\begin{itemize} \\item to investigate the influence of the strongly magnetized hot plasma on the process of Compton scattering taking account of photon dispersion and wave function renormalization; \\item to compare the Compton scattering with the process of photon splitting; \\item to estimate the possible effect of the process under consideration on spectral formation of SGR. \\end{itemize}  The plan of the paper is following. In the section 2 we consider the propagation of the electromagnetic radiation in magnetized medium and give the expression of Compton scattering amplitudes for different polarization configurations of the initial and final photons. The analytic and numerical calculations of the photon absorption rate and cross section of the process are presented in Section 3. The discussion of the possible astrophysical application of Compton scattering are given in Section 4. ",
        "conclusions": "\\label{Sec4}  In the models of soft gamma repeaters spectrum formation the dependence of photon absorption rates on energy and temperature plays an important role. It could influence on the shape of emergent spectrum and defines the temperature profile in the emission region during bursts in SGRs~\\cite{Duncan:1995, Lyubarsky:2002}. In the both figures \\ref{fig:W22} and \\ref{fig:W11} one can see that photon absorption rates corresponding to Compton scattering are the fast increasing functions of temperature. At the same time, the channels with initial photon of mode 1 and mode 2 have different character of the absorption coefficient  energy dependence. As shown in the  Fig.\\ref{fig:W22} the absorption rates for the reactions $\\gamma_2 e^{\\pm} \\to \\gamma_2 e^{\\pm}$  and $\\gamma_2 e^{\\pm} \\to \\gamma_1 e^{\\pm}$ have thresholds. They are  caused by mode 2 dispersion relation with plasma frequency and indicated the fact that the electromagnetic wave corresponding to mode 2 can't propagate with energy below $\\omega_{pl}$. In the region above plasma frequency $W_{\\gamma_2 e^{\\pm} \\to \\gamma_2 e^{\\pm}}$ is slightly depends on energy while $W_{\\gamma_2 e^{\\pm} \\to \\gamma_1 e^{\\pm}}$ tends to zero in the vicinity of $\\omega_{pl}$ and pair creation threshold, $\\omega = 2m$. We would like to note that in the vicinity of pair creation threshold taking account of wave function renormalization and photon dispersion becomes very important and defines the processes rates dependence on energy, temperature and magnetic field. In the  region $\\omega \\ge 2 m$ the photon of mode 2 is unstable due to the  process $\\gamma _2 \\to e^+ e^-$ with rate much lager than Compton scattering.  The energy dependence of $W_{\\gamma_1 e^{\\pm} \\to \\gamma_1 e^{\\pm}}$ and $W_{\\gamma_1 e^{\\pm} \\to \\gamma_2 e^{\\pm}}$ is depicted in Fig.\\ref{fig:W11}. On can see the fast increasing of absorption coefficients at low energies and rather slow dependence at $\\omega \\gtrsim 2 m$.  The previous investigation of radiation transfer problem in strongly magnetized plasma have shown that along with Compton scattering process the photon splitting $\\gamma \\to \\gamma \\gamma$ could play a significant role as mechanism of photon production~\\cite{Duncan:1995}. To compare these processes we depict the probabilities  of the photon splitting channels on the same plots (see dotted and chain lines in Fig.\\ref{fig:W22} and Fig.\\ref{fig:W11}). On can see that at temperature $T \\sim m$  and in kinematical region $\\omega \\le 2 m $ the main photon splitting process is $\\gamma_2 \\to \\gamma_1 \\gamma_1$ forbidden in pure magnetic field~\\cite{RCh05}. It is seen also that the rate of the process is much less than Compton scattering ones. Nevertheless,  it could be an effective photon production mechanism at temperatures under consideration. The probabilities of the channels $\\gamma_1 \\to \\gamma_1 \\gamma_2$ and $\\gamma_1 \\to \\gamma_2 \\gamma_2$ increase with temperature falling  and become comparable and even larger than Compton scattering rates $W_{\\gamma_1 e^{\\pm} \\to \\gamma_1 e^{\\pm}}$ and  $W_{\\gamma_1 e^{\\pm} \\to \\gamma_2 e^{\\pm}}$. As shown in Fig.~\\ref{fig:W11} the process of photon splitting strongly dominates over Compton scattering at $T=50 keV$. It was claimed previously that the effect of the cold strongly magnetized plasma on photon splitting is not pronounced and the vacuum approximation can be used in the most calculations~\\cite{Bulik:1997,Elmfors:1998}. We can see now that in the presence of hot plasma the process of photon splitting could not be only intensive source of photon production but also an effective absorption mechanism.  Let us illustrate this  fact in the framework of the magnetar model of SGR burst. It is known that the radiation transfer in the magnetically trapped plasma may be described  as s diffusion of the 1-mode photons whereas 2-mode photons are locked~\\cite{Duncan:1995,Lyubarsky:2002}. The last circumstance is connected especially with weak dependence of the 2-mode absorption coefficient on photon energy (see Fig.~\\ref{fig:W22}). In the assumption that the properties of the radiation field do not change much over absorption length $W^{-1}$ the radiation transfer may be conveniently described in terms of the Rosseland mean free path of 1-mode photons~\\cite{Duncan:1995,Lyubarsky:2002} \\begin{equation} \\lambda_R = \\frac{\\int W^{-1} \\omega^4 \\exp(\\omega/T) [\\exp(\\omega/T)-1]^{-2} d \\omega}{\\int \\omega^4 \\exp(\\omega/T) [\\exp(\\omega/T)-1]^{-2} d \\omega}, \\label{eq:lR} \\end{equation} where $W$ is the total absorption coefficient. Then the temperature distribution in the magnetically trapped plasma region could be defined from the diffusion equation: \\begin{equation} F(z)= \\frac{16 \\lambda_R \\, \\sigma}{3}\\,T^3 \\, \\frac{\\partial T}{\\partial z}, \\label{eq:Difuss} \\end{equation} where $\\sigma$ is the Stefan-Boltzmann constant, $F(z)$ is the radiation flux. If only Compton scattering process is considered then \\begin{equation} W = W_{\\gamma_1 e^{\\pm} \\to \\gamma_1 e^{\\pm}} + W_{\\gamma_1 e^{\\pm} \\to \\gamma_2 e^{\\pm}}. \\label{eq:WCompt} \\end{equation} To take into account the effect of photon splitting one should add  term $W_{\\gamma_1 \\to \\gamma_1 \\gamma_2} + W_{\\gamma_1 \\to \\gamma_2 \\gamma_2}$ to the equation (\\ref{eq:WCompt}). In the figure~\\ref{fig:RC_CSp} the ratio of the Rosseland mean free path for only Compton scattering to  Rosseland mean free path  with photon splitting is depicted. One can see that at low temperatures the additional absorption process of photon splitting leads to the significant decreasing of the Rosseland mean free path. The solution of the diffusion equation (\\ref{eq:Difuss}) is out of the scope of this article. However, we would like to note that the more detail analysis of the radiation transfer needs the consistent solution of Boltzmann equation for the photon occupation number and radiative transfer equation in the wide range of temperatures ($ 10 keV \\lesssim T \\lesssim 1 MeV$). \\begin{figure*} \\centerline{\\includegraphics{rc_csp.eps}} \\vspace*{-2mm} \\caption{The ratio of the Rosseland mean free paths without and with taking into account of photon splitting process as a function of the inverse temperature at $B = 200 B_e$} \\label{fig:RC_CSp} \\end{figure*}  In conclusion, we have investigated the influence of the strongly magnetized hot plasma on the Compton scattering process with taking into account of photon dispersion and large radiative corrections. It was found that absorption rate of the dominant scattering channel of the 2-mode photon weakly depends on energy in  wide temperature range whereas the 1-mode absorption coefficient is fast increasing function of photon energy. The comparison of the photon splitting effect and Compton scattering shows that the influences of these reaction on the 1-mode radiation transfer are competitive in rarefied plasma $(T \\ll m)$.   \\medskip  \\noindent{\\bf Acknowledgements}  We express our deep gratitude to the organizers of the Seminar ``Quarks-2006'' for warm hospitality. This work supported in part by the Council on Grants by the President of Russian Federation for the Support of Young Russian Scientists and Leading Scientific Schools of Russian Federation under the Grant No. NSh-6376.2006.2 and No. MK-1097.2005.2 (MC), and by the Russian Foundation for Basic Research under the Grant No. 04-02-16253. Work of MC is supported by the DAAD (Deutscher Akademischer Austauchdienst)."
    },
    "0609/astro-ph0609372_arXiv.txt": {
        "abstract": "We present 4.5, 8, and 16\\micron\\ photometry from the {\\it Spitzer Space Telescope} for 204 stars in the Upper Scorpius OB Association. The data are used to investigate the frequency and properties of circumstellar disks around stars with masses between \\about 0.1 and 20\\msun\\ at an age of \\about 5~Myr. We identify 35 stars that have emission at 8\\micron\\ or 16\\micron\\ in excess of the stellar photosphere. The lower mass stars (\\about 0.1-1.2\\msun) appear surrounded by primordial optically thick disks based on the excess emission characteristics. Stars more massive than \\about 1.8\\msun\\ have lower fractional excess luminosities suggesting that the inner \\about 10~AU of the disk has been largely depleted of primordial material. None of the G and F stars (\\about 1.2-1.8\\msun) in our sample have an infrared excess at wavelengths $\\le$16\\micron. These results indicate that the mechanisms for dispersing primordial optically thick disks operate less efficiently on average for low mass stars, and that longer time scales are available for the buildup of planetary systems in the terrestrial zone for stars with masses \\aboutless 1\\msun. ",
        "introduction": "Most young (\\about 1\\,Myr) stars embedded within molecular clouds are surrounded by circumstellar accretion disks \\citep{Strom89} that are potential sites of planet formation. The ubiquity of disks extends to all masses from as high as 10\\msun\\ down through brown dwarfs, and in all environments from isolated stars in Taurus to dense clusters in Orion \\citep{Lada00,Bouy06}.  By an age of 10\\,Myr, the primordial disks so ubiquitous around young stars change dramatically. The inner disk (\\aboutless 1\\,AU) dissipates in $>$90\\% of stars \\citep{Mamajek04}, accretion rates drop by an order of magnitude \\citep{Muzerolle00}, and the mass contained in small dust grains declines by at least a factor of four \\citep{Liu04,Carpenter05}. Results from the {\\it Spitzer Space Telescope} \\citep{Werner04} demonstrate an even more striking degree of evolution, as dust within a \\about 1\\,AU orbital radius is found in only a few percent of \\about 10\\,Myr stars \\citep{Silverstone06}. MIPS\\,24\\micron\\ surveys have detected dust in 7--48\\% of 10--30\\,Myr stars \\citep{Young04,Rieke05,Chen05}, but with fractional dust luminosities orders of magnitude below that found in younger sources. Together these observations have established that circumstellar disks are at an advanced evolutionary stage by an age of \\about 10\\,Myr.  Key to understanding the formation of planetary systems is examining the evolution of circumstellar disks after the main accretion phase has terminated. To measure the properties of disks during this epoch for stellar masses ranging from 0.1 to 20\\msun, we have conducted a photometric survey of 205 stars with spectral types between M5 and B0 in the 5 Myr-old Upper Scorpius OB association using the IRAC, IRS, and MIPS instruments on \\Spitzer. This Letter presents analysis of the IRAC and IRS photometry to probe for terrestrial-zone material across the stellar mass spectrum at a constant age. ",
        "conclusions": "The excess properties of the Upper Sco sources are not uniform across spectral type as demonstrated in Figure~\\ref{fig:ccd}. The 8\\micron\\ excess fraction for K+M stars (24/127) is higher than for B+A stars (5/61) at the 92\\% confidence level and higher than for F+G stars (0/30) at 99.2\\% confidence as determined from the two-tailed Fisher's Exact Test. At 16\\micron, the K+M excess fraction (23/121) is similar to that for B+A stars (10/52), but higher than for F+G stars (0/22) at 97.5\\% confidence.  More telling differences between early and late spectral types are observed in the magnitude of the excesses. Nine B+A stars have a $R_{16}$ color excess less than twice the stellar photosphere, while all K+M sources exceed this limit with excesses up to 27 times the photospheric level. Similarly, at 8\\micron, all but two B+A excess sources have a $R_8$ excess less than 10\\% of the photosphere, while 23 of the 24 K+M excess sources have larger color excesses.  In Figure~\\ref{fig:sed}, we assess the evolutionary state of the circumstellar disks in Upper Sco by comparing normalized, dereddened spectral energy distributions (SEDs) for stars in Upper Sco with a sample of well known T Tauri \\citep{Hartmann05} and Herbig Ae/Be \\citep{Hillenbrand92} stars that have tabulated photometry. The Taurus and Herbig Ae/Be objects are expected to represent young stars surrounded by primordial, optically thick, circumstellar accretion disks.  As shown in Figure~\\ref{fig:sed}, the SEDs for B and A stars in Upper Sco differ substantially from most Herbig Ae/Be stars. Herbig Ae/Be stars typically have excesses at wavelengths as short as 2\\micron\\ and have fractional excess luminosities that are 10-100 times the photosphere at 10-20\\micron. By comparison, only one B or A star in Upper Sco has a $K$-band excess (HIP~78207), and the fractional excess luminosity at 16\\micron\\ is typically less than twice the photosphere.  Surprisingly, none of the F and G stars in our Upper Sco sample exhibit a detectable excess. While \\citet{Chen05} identified a 24\\micron\\ excess around one of five F+G stars in Upper Sco, and the G6\\,V star HD~143006 is surrounded by an optically thick disk \\citep{Sylvester96}, overall infrared excesses at wavelengths $\\le$ 16\\micron\\ are relatively rare for this spectral type range at the age of Upper~Sco.  The results for the B, A, F, and G stars imply that the reservoir of small dust grains in a primordial, optically thick, inner disk have been largely depleted by an age of 5~Myr for \\about 1-20\\msun\\ stars. The inner-disk radius inferred by the weak (or lack-of) excess emission at 16\\micron\\ is \\about 4-10~AU for 6000-10,000~K photospheres assuming optically thin, blackbody dust emission \\citep{Jura03}. These results are consistent with the low fraction of accreting B and A stars found in the 4~Myr-old Trumpler~37 cluster \\citep{Sicilia06} and the 5~Myr-old Orion OBIb association \\citep{Hernandez06}. However, \\citet{Hernandez06} found that 7 of 11 F-stars in Orion~OBIb contain 24\\micron\\ excesses consistent with a debris disk, suggesting that the excess fraction for F and G-stars in Upper Sco may increase once our longer wavelength observations are obtained.  In contrast to the massive stars, the K and M stars in Upper Sco with infrared excesses have characteristics similar to optically thick primordial disks. As discussed above and demonstrated in Figure~\\ref{fig:sed}, the K+M stars have larger fractional excesses than the B+A stars, and the magnitude of the excesses overlaps with that observed toward Class~II stars in Taurus. A further connection between the disks in Upper Sco and Taurus is found by considering evidence for disk accretion as traced by H$\\alpha$ emission. Of the 100 stars in our sample with measured H$\\alpha$ line strengths (see references in \\S\\ref{sample}), nine have H$\\alpha$ equivalent widths consistent with accretion according to the criteria recommended by \\citet{White03}. Eight of these nine sources have an infrared excess and are therefore likely surrounded by accretion disks. We note however that the H$\\alpha$ equivalent widths for 14 K+M stars with excesses are consistent with active chromospheres. These similarities suggest that many of the K+M stars in Upper Sco remain surrounded by optically thick disks. Assuming optically thick blackbody emission, the 4.5\\micron\\ and 8\\micron\\ excesses may imply the presence of dust at radii as small as \\about 0.05~AU \\citep{Jura03}.  Differences in the excess characteristics between the \\about 1-2 Myr-old Taurus and 5~Myr Upper Sco populations are notable since they may reflect temporal evolution in the disk properties. While the magnitude of the excesses overlaps between the two samples, the excesses are larger on average for Taurus as seen in Figure~\\ref{fig:sed}. Furthermore, about half of the stars in Taurus exhibit a $K$-band excess \\citep{Strom89} compared to only two (1.5\\%) K+M stars in Upper Sco ($[$PBB2002$]$~USco~J161420.2$-$190648 and $[$PZ99$]$~J160421.7$-$213028). Similarly, while 68\\% of the stars in Taurus exhibit a 3.6\\micron\\ excess \\citep{Haisch01}, only $19^{+5}_{-4}$\\% of K+M stars in Upper Sco have a 8\\micron\\ excess. Therefore, not only are there fewer sources with disks in Upper Sco, but the disks that remain lack the hot dust found in younger stars. \\citet{Sicilia06} found similar tendencies for the low mass population in Trumpler~37 compared to Taurus. They attribute these differences to grain growth or dust settling in the inner disk, although further observations and modeling are needed to explore the relevance of these ideas to stars in Upper Sco.  To summarize, 19\\% of K0-M5 stars in Upper Sco possess infrared excesses similar to Class~II sources in Taurus, indicating primordial disks last around an appreciable number of 0.1-1\\msun\\ stars for at least 5~Myr. By contrast, only \\aboutless 1\\% of the more massive stars in Upper Sco contain such disks within an orbital radius of \\about 4-10~AU. Similar results have been reported for the 4~Myr-old Trumpler~37 cluster \\citep{Sicilia06}, and the 2-3~Myr-old IC~348 cluster may also contain a higher fraction of disks around later type stars \\citep{Lada06}. Our observations of Upper Sco extend these conclusions to the full range of stellar masses down to the hydrogen burning limit at an age of \\about 5~Myr. These results establish that warm dust in the terrestrial zone persist for longer times around stars with masses \\aboutless 1\\msun."
    },
    "0609/astro-ph0609691_arXiv.txt": {
        "abstract": "We explore the evolution of thermonuclear supernova explosions when the progenitor white dwarf star ignites asymmetrically off-center. Several numerical simulations are carried out in two and three dimensions to test the consequences of different initial flame configurations such as spherical bubbles displaced from the center, more complex deformed configurations, and teardrop-shaped ignitions. The burning bubbles float towards the surface while releasing energy due to the nuclear reactions. If the energy release is too small to gravitationally unbind the star, the ash sweeps around it, once the burning bubble approaches the surface. Collisions in the fuel on the opposite side increase its temperature and density and may -- in some cases -- initiate a detonation wave which will then propagate inward burning the core of the star and leading to a strong explosion. However, for initial setups in two dimensions that seem realistic from pre-ignition evolution, as well as for all three-dimensional simulations the collimation of the surface material is found to be too weak to trigger a detonation. ",
        "introduction": "Type Ia supernovae are generally associated with the thermonuclear explosion of earth-sized white dwarf (WD) stars -- the final stages of the evolution of small and intermediate-mass stars. These dense objects consist of carbon and oxygen and their pressure is dominated by degenerate electrons. In a binary system, the WD can gain mass from its companion by accretion. However, there exists a fundamental limit, the Chandrasekhar mass, beyond which the star is unstable against gravitational collapse. Before reaching this stage, the density at the center of the WD reaches values at which nuclear reactions from carbon to heavier elements ignite. This establishes a stage of convective carbon burning which is finally terminated when a thermonuclear runaway forms a flame. The thermonuclear flame propagates outward and may give rise to an explosion of the star \\cite{reinecke2002d,roepke2005b}. The process resembles premixed chemical combustion where two solutions -- supersonic detonation and subsonic deflagration -- are admissible. From observational constraints it is clear that the flame starts out as a subsonic deflagration. The ashes are buoyant and shear instabilities generate strong turbulence. The interaction of the flame with turbulence is a key feature in the model and accelerates the flame. Consequently, one of the outstanding problems is to determine, whether this turbulent flame leads to explosions energetic enough to be consistent with observations, or whether a later transition to a supersonic detonation is both possible and necessary. In the freely expanding medium of SN Ia explosions, such transitions are difficult to achieve.  The exact way of flame ignition is not well determined yet, but previous numerical studies \\cite{woosley2004a, kuhlen2006a} suggest that the ignition process may proceed on only one side of the star's center in a few patches. While it has been shown that ignitions in multiple spots distributed around the center of the WD lead to strong explosions \\cite{roepke2006a}, one-sided ignitions are expected to release less energy. The effects of the latter scenario on the explosion process are investigated in the present study. It has been speculated that collisions between gravitationally bound laterally expanding regions of fuel and ash could, in principle, give rise to a ``gravitationally confined detonation'' \\cite{plewa2004a}. This possibility is studied in detail. ",
        "conclusions": "\\begin{figure}[t] \\includegraphics[width=0.7\\textwidth]{roepke_f2.eps} \\hspace{0.01 \\textwidth} \\begin{minipage}[b]{0.299\\textwidth} \\caption{\\label{fig:et1} Maximum fuel temperature reached in the region of collision of the material sweeping around the WD as a function of the nuclear energy release prior to reaching that temperature. The 3D data point is marked, all other points correspond to 2D simulations.\\vspace*{6ex}} \\end{minipage} \\end{figure}  Initiation of a detonation wave in the collision region depends on the temperature reached, the density, and the spatial extent of the collision region. For our 2D simulations ignited in a spherical bubble, an inverse correlation exists between nuclear energy release prior to collision and the maximum temperature reached (cf.\\ Fig.~\\ref{fig:et1}). This correlation arises naturally from the expansion of the star due to the energy release. The more material is burnt in the flame on its way to the surface, the more the star expands and the weaker is the collision.  The energy release depends on the displacement of the initial flame from the center. Initiating the flame at larger radii leads to less burning and thus to stronger collisions of the surface material. The collision strength is very sensitive to perturbations of the initial flame from perfect sphericity, either directly by imposing more complex initial configurations (teardrop-like shapes, multiple initial flame kernels) or by different numerical resolution and resulting discretization errors.  \\begin{figure}[t] \\includegraphics[width=0.7\\textwidth]{roepke_f3.eps} \\hspace{0.01 \\textwidth} \\begin{minipage}[b]{0.299\\textwidth} \\caption{\\label{fig:et2} Maximum temperature reached in the collision region for fuel above the given density thresholds as a function of the nuclear energy release prior to reaching the maximum collision temperature. The shaded and dashed regions correspond to conditions where an initiation of a detonation is possible for temperatures reached in fuel of densities $\\rho > 1 \\times 10^7 \\, \\mathrm{g} \\, \\mathrm{cm}^{-3}$ and $\\rho > 3 \\times 10^6 \\, \\mathrm{g} \\, \\mathrm{cm}^{-3}$, respectively.\\vspace*{6ex}} \\end{minipage} \\end{figure}  To trigger a detonation, a collision temperature over $2 \\times 10^9 \\, \\mathrm{K}$ is necessary. However, this threshold temperature needs to be reached in fuel with density higher than $\\sim 3\\ldots10\\times 10^6 \\, \\mathrm{g} \\, \\mathrm{cm}^{-3}$. Therefore, the maximum temperature in the collision region at a density exceeding these values was measured. The results are plotted in Fig.~\\ref{fig:et2}. From this measurement, we find that a detonation initiation is only possible in those 2D models that start with a spherical initial flame placed at unrealistically large radii ($R \\ge 200\\, \\mathrm{km}$) off-center of the WD star (cf.\\ Fig.~\\ref{fig:et2}).  3D simulations are cooler in the collimation region for the same nuclear energy release. However, except for one model (included in Figs.~\\ref{fig:et1} and \\ref{fig:et2}) the 3D simulations release more energy in the nuclear burning than the 2D simulations. In the model shown in Fig.~\\ref{fig:evo}, the energy release amounted to $\\sim$$2.8 \\times 10^{50} \\, \\mathrm{erg}$ and the temperature in the collision region barely exceeded $10^9 \\, \\mathrm{K}$. For the one-sided teardrop-shaped initial flame the energy release was even greater and one singe-bubble ignition and the two-sided teardrop-ignition unbound the star. All 3D simulations clearly failed to trigger a detonation.  We conclude, that although a detonation due to the colliding surface material is in principle possible for specific (and artificial) ignition configurations, it is unlikely to happen in nature. For models that remain gravitationally bound, failures to initiate a detonation  will lead to pulsations of the WD star. This may be a second chance for triggering a detonation (\"pulsational delayed detonation scenario\", \\cite{arnett1994b}) and will be addressed in a follow-up study.    \\ack{This research used resources of the National Center for Computational Sciences at Oak Ridge National Laboratory, which is supported by the Office of Science of the U.S. Department of Energy under Contract No. DE-AC05-00OR22725.}"
    },
    "0609/astro-ph0609658_arXiv.txt": {
        "abstract": "We present a robust and fast algorithm for performing astrometry and source cross-identification on two dimensional point lists, such as between a catalogue and an astronomical image, or between two images. The method is based on minimal assumptions: the lists can be rotated, magnified and inverted with respect to each other in an arbitrary way. The algorithm is tailored to work efficiently on wide fields with large number of sources and significant non-linear distortions, as long as the distortions can be approximated with linear transformations locally, over the scale-length of the average distance between the points.  The procedure is based on symmetric point matching in a newly defined continuous triangle space that consists of triangles generated by an extended Delaunay triangulation. Our software implementation performed at the 99.995\\% success rate on $\\sim260,000$ frames taken by the HATNet project. ",
        "introduction": "\\label{sec:intro}  Cross-matching two two-dimensional points lists is a crucial step in astrometry and source identification. The tasks involves finding the appropriate geometrical transformation that transforms one list into the reference frame of the other, followed by finding the best matching point-pairs. One of the lists usually contains the pixel coordinates of sources in an astronomical image (e.g.~point-like sources, such as stars), while the other list can be either a reference catalog with celestial coordinates, or it can also consist of pixel coordinates that originate from a different source of observation (another image). Throughout this paper we denote the reference (list) as \\aref, the image (list) as \\aimg, and the function that transforms the reference to the image as \\atran.  The difficulty of the problem is that in order to find matching pairs, one needs to know the transformation, and vica versa: to derive the transformation, one needs point-pairs. Furthermore, the lists may not fully overlap in space, and may have only a small fraction of sources in common.  By making simple assumptions on the properties of \\atran, however, the problem can be tackled. A very specific case is when there is only a simple translation between the lists, and one can use cross-correlation techniques \\citep[see][]{phillips95} to find the transformation. We note, that the a method proposed by \\citet{thiebaut01} uses the whole image information to derive a transformation (translation and magnification).  A more general assumption, typical to astronomical applications, is a that $\\atran$ is a similarity transformation (rotation, magnification, inversion, without shear), i.e.~$\\atran = \\lambda \\mathbf{A}\\underline{r} + \\underline{b}$, where $\\mathbf{A}$ is a (non-zero) scalar $\\lambda$ times the orthogonal matrix, $\\underline{b}$ is an arbitrary translation, and $\\underline{r}$ is the spatial vector of points. Exploiting that geometrical patterns remain similar after the transformation, more general algorithms have been developed that are based on pattern matching \\citep{groth86,valdes95}. The idea is that the initial transformation is found by the aid of a specific set of patterns that are generated from a subset of the points on both \\aref\\ and \\aimg. For example, the subset can be that of the brightest sources, and the patterns can be triangles.  With the knowledge of this initial transformation, more points can be cross-matched, and the transformation between the lists can be iteratively refined. Some of these methods are implemented as an \\pack{IRAF} task\\footnote{IRAF is distributed by the National Optical Astronomy Observatories, which are operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation.} in \\prog{immatch} \\citep{phillips95}.   The above pattern matching methods perform well as long as the dominant term in the transformation is linear, such as for astrometry of narrow field-of-view (FOV) images, and as long as the number of sources is small (because of the large number of patterns that can be generated \\--- see later).  In the past decade of astronomy, with the development of large format CCD cameras or mosaic imagers, many wide-field surveys appeared, such as those looking for transient events (e.g.~ROTSE --- \\citealt{akerlof00}), transiting planets (e.g.~Kelt -- \\citealt{pepper04}, TrES \\--- \\citealt{alonso04}, HATNet \\--- \\citealt{bakos02,bakos04}, see \\citealt{dc06} for further references), or all-sky variability (e.g.~ASAS \\--- \\citealt{pojmanski97}). There are non-negligible, higher order distortion terms in the astrometric solution that are due to, for instance, the projection of celestial to pixel coordinates and the properties of the fast focal ratio optical systems. Furthermore, these images may contain $\\sim 10^5$ sources, and pattern matching is non-trivial.  These surveys necessitated a further generalization of the algorithm, which we present in this paper. To be more specific, we were motivated by the astrometric requirements of the Hungarian-made Automated Telescope Network (HATNet). Each HAT telescope of the Network consists of a $f=200\\mathrm{mm}$, $d=f/1.8$ telephoto lens and a \\ccdsize{2K} CCD yielding an $8\\arcdeg\\times8\\arcdeg$ FOV. In our experience, we need at least 4th order polynomial functions of the pixel coordinates in order to properly describe the distortion of the lens. With a typical exposure time of 5 minutes in I-band, in a moderately dense field ($b \\approx 15^\\circ$) there are 30000 stars brighter than I=13 for which better than 10\\% photometry can be achieved. If we consider all 3-$\\sigma$ detections, we have to deal with the identification of $\\sim$100,000 sources.  The algorithm presented in this paper is based on, and is a generalization of the above pattern matching algorithms. It is very fast, and works robustly for wide-field imaging with minimal assumptions. Namely, we assume that: i) the distortions are non-negligible, but small compared to the linear term, ii) there exists a smooth transformation between the reference and image points, iii) the point lists have a considerable number of sources in common, and iv) the transformation is locally invertible. The paper is presented as follows. First we describe symmetrical point matching in \\secr{pmatch} before we go on to the discussion of finding the transformation (\\secr{tran}). The software implementation and its performance on a large and inhomogeneous dataset is demonstrated in \\secr{exam}. Finally, we draw conclusions in \\secr{sum}. ",
        "conclusions": "\\label{sec:sum}  In this paper we present a robust algorithm for cross-matching two two-dimensional point lists. The task is twofold: finding the smooth spatial transformation between the lists, and cross-matching points. These two steps are intertwined, and are performed in an iterative way till satisfactory transformation and matching rate are reached. We make only very basic assumptions that hold for almost all astronomical applications, including wide-field surveys with distorted fields and large number of sources. Namely, the transformation between the point lists is dominantly a similarity transformation (arbitrary shift, rotation, magnification, inversion). A significant distortion term can be present, given it can be linearized on the scale-length of the average distance of neighboring points.  In \\secr{pmatch} we briefly described symmetric point-matching in one and two dimensions, because this tool is used throughout the astrometry procedure. Finding the initial transformation between the point lists is based on triangle matching. First we defined various normalized triangle spaces in \\secr{trispace}. The ``mixed'' triangle space of \\citet{valdes95} is a continuous function of the triangle parameters, but flipped triangles are not distinguished. The ``chiral'' triangle space ensures that chirality information is preserved, but this space is not continuous. We showed that it is possible to define a ``continuous'' triangle space that is both continuous and preserves chirality, and which, furthermore, spans a larger volume and diminishes confusion of triangles having similar coordinates.  Taking into account the distortion of a field and the astrometric errors, we calculated both the optimal size and number of triangles. For the typical setup of a HATNet telescope ($8\\arcdeg\\times8\\arcdeg$ FOV, distortion factor $f_d \\sim 0.005$) the optimal size is $0.2\\arcdeg$, and the optimal number of triangles is less than $2\\times 10^6$. We use Delaunay triangulation for generating the triangles of the triangle-space. This has the advantage of being fast, robust, and generating local triangles that are less prone to being distorted. In \\secr{extdel} we noted, however, that Delaunay triangulation is sensitive for removing or adding points to the list, and thus instable. We introduced an extension of this triangulation that is parametrized by an $\\ell$ level.  Having determined the transformation between the two lists, one can check how well our initial assumption of the dominant term being linear holds. In \\secr{unit} we introduced the unitarity of the transformation, a simple scalar measure of this property. We described the practical details of the algorithm in \\secr{pract}, and the actual software implementation (\\prog{grmatch}, \\prog{grtrans}) in \\secr{sw}.  Finally, we ran the above programs on some 240,000 frames taken by the wide-angle cameras of HATNet, plus 20,000 frames acquired by the TopHAT telescope. The success rate was very close to 100\\%, and the routines handled the various pointing errors, defocusing and 6th order distortions in the wide fields. Both programs will become available in binary format for a wide range of architectures upon request from the authors."
    },
    "0609/astro-ph0609144_arXiv.txt": {
        "abstract": "We study the accretion flows from the circumbinary disks onto the supermassive binary black holes in a subparsec scale of the galactic center , using a smoothed particles hydrodynamics (SPH) code. Simulation models are presented in four cases of a circular binary with equal and unequal masses, and of an eccentric binary with equal and unequal masses. We find that the circumblack-hole disks are formed around each black holes regardless of simulation parameters. There are two-step mechanisms to cause an accretion flow from the circumbinary disk onto supermassive binary black holes: First, the tidally induced elongation of the circumbinary disk triggers mass inflow towards two closest points on the circumbinary disk from the black holes. Then, the gas is increasingly accumulated on these two points owing to the gravitational attraction of black holes. Second, when the gas can pass across the maximum loci of the effective binary potential, it starts to overflow via their two points and freely infalls to each black hole. In circular binaries, the gas continues to be supplied from the circumbinary disk (i.e. the gap between the circumbinary disk and the binary black hole is always closed.) In eccentric binaries, The mass supply undergoes the periodic on/off transitions during one orbital period because of the variation of periodic potential. The gap starts to close after the apastron and to open again after the next periastron passage. Due to this gap closing/opening cycles, the mass-capture rates are eventually strongly phase dependent. This could provide observable diagnosis for the presence of supermassive binary black holes in merged galactic nuclei. ",
        "introduction": "\\label{sec:intro} {%  There is growing evidence that most galaxies have the supermassive black holes at their centers (\\cite{KR}; see also \\cite{rees} for a classical review). The remarkable evidence for supermassive black holes is provided by the existence of a gas disk with Keplerian rotation on a subparsec scale found by the radio observations \\citep{miyoshi} and by the asymmetric iron line profile discovered by the X-ray observations \\citep{tanaka}. In the vicinity of a supermassive black hole, the gas in the disk is heated up and produces very high luminosity by efficiently transforming their gravitational energy into radiation (e.g., \\cite{lynden}). Thus, such rotating gas disks with accretion flow are considered as the energy sources of active galactic nuclei (AGNs) and quasars.  Recently, it has been widely accepted that the supermassive black holes play an important role not only in the activities of AGNs and quasars but also in the formation and global evolution of galaxies (\\cite{SR}; \\cite{KH}; \\cite{matteo}; \\cite{merritt1}}, and references therein). The discovery of the tight correlation between black hole mass and the velocity dispersion of the bulge component of galaxies \\citep{ferra, geb} supports the scenario that the black holes have grown up in mass through the hierarchical galaxy mergers, just as galaxies themselves did. If this scenario is correct, the supermassive binary black holes (BBHs) will be inevitably formed during the course of galaxy mergers \\citep{milos1}.  There are actually a number of observational indications that some galaxies harbour supermassive BBHs at their centers. Main results are listed as follows (see also \\authorcite{komo1} \\yearcite{komo1,komo2}); \\begin{itemize} \\item Periodic optical and radio outbursts (e.g., OJ287) \\citep{sill,LV,valtao,valtonen}. \\item Wiggled patterns of the radio jet indicating precessional motions on a parsec scale \\citep{YI,roos2,britzen,AC,LR}. \\item X-shaped morphology of radio lobes \\citep{merritt2}. \\item Double peaked broad emission lines in AGNs \\citep{gaskell,Ho}. \\item Double  compact cores with the flat radio spectrum \\citep{maness,rodriguez}. \\item Orbital motion of the compact core \\citep{sudou}. \\end{itemize}  It is believed that the supermassive BBHs evolve mainly via three stages \\citep{begel,yu}. Firstly, each of black holes sinks independently towards the center of the common gravitational potential due to the dynamical friction with neighboring stars. When the separation between two black holes becomes as short as $1\\,\\rm{pc}$ or so, angular momentum loss by the dynamical friction slows down due to the loss-cone effect and a supermassive hard binary is formed. This is the second stage. Finally, when the semi-major axis of the binary decreases to less than $0.01\\,\\rm{pc}$, gravitational radiation dominates and then a pair of black holes, eventually, merge into a single supermassive black hole.  If there is the gas orbiting around the supermassive BBHs in the second evolutionary stage, one will be able to observe a signal arising from the interaction between the binary and its surrounding gas (i.e. a circumbinary disk). This disk-binary interaction could be also the predominant candidate to resolve the loss-cone problem (\\authorcite{armi1} \\yearcite{armi1,armi2}; see also \\cite{arty1}). An orbital angular momentum of the supermassive BBHs is transferred to the circumbinary disk, by which the gas around the supermassive BBHs will be swept away. In addition, we expect that the mass inflow will take place from the circumbinary disk to the supermassive BBHs, leading to the formation of accretion disks around each of black holes (i.e. circumblack-hole disks). A final configuration could be three-disk systems; one circumbinary disk and two circumblack-hole disks (see Fig.~\\ref{fig:system} for a schematic view of supermassive BBHs).  Such three-disk systems have been also discussed in the context of binary star formation, where the binary is composed of young stars. \\citet{arty4} found that the material can infall on to the central binary through the gap between the circumbinary disk and the central binary, and that its accretion rate modulates with the orbital phase. Since then, the theory of the young binary star formation has been extensively studied \\citep{arty5,lubow1,bate1,lubow2,GK1,GK2,ochi}. Some basic processes involved with the disk-binary interaction have been revealed through these researches. Despite its significance, however, it is poorly known how the material accretes onto supermassive BBHs from the circumbinary disk under the realistic situations. \\citet{arty4} only discussed briefly the quasi-periodic behavior of optical/infrared outbursts in a blazar OJ287 in terms of the strong phase-dependent accretion onto supermassive BBHs (see also \\cite{arty1}).  In this paper, therefore, we elucidate the theory of accretion processes in supermassive binary black-hole systems by performing smoothed particle hydrodynamics (SPH) simulations. Our ultimate goal is to give an observable diagnosis for the presence of supermassive BBHs in central region of merged galaxies. The simulation should ideally take account of all the processes at work in the three-disk systems, including the circumbinary disk evolution around the supermassive BBHs, the mass transfer from the circumbinary disk to the individual black holes, and the accretion onto each of black holes. Such a simulation, however, would require an enormous computational time. Therefore, we confine ourselves to simulate only the accretion flow from the circumbinary disk onto the supermassive BBHs. Detailed structure and evolution of the circumblack-hole disks will be reported in a subsequent paper.  The plan of this paper is as follows: We first describe our numerical model in Section~\\ref{sec:num}. Then, our numerical results will be reported in Section~\\ref{sec:acc} for the basic processes involved with mass accretion onto the circular, equal-mass binary, and in Section~\\ref{sec:uneq_ecc} for the effects of the orbital eccentricity and the unequal masses of the binary. We then discuss the observational implications and the related issues in Section~\\ref{sec:discussion}. Section~\\ref{sec:conclusions} is devoted to conclusions.   \\begin{figure} \\begin{center} \\FigureFile(80mm,80mm){kh1.eps} \\end{center} \\caption{Schematic diagram of supermassive binary black holes on a parsec/subparsec scale of merged galactic nuclei. The supermassive binary black holes are surrounded by a circumbinary disk, from which the gas overflows on to the binary, and then the circumblack-hole disks are formed around each black hole (BH). } \\label{fig:system} \\end{figure} ",
        "conclusions": "\\label{sec:conclusions}  For the purpose of providing the observable diagnosis to probe the existence of supermassive BBHs on a subparsec scale in merged galactic nuclei, we have carried out the SPH simulations of accretion flows from circumbinary disks onto the supermassive BBHs. Our main conclusions are summarized as follows:  \\begin{enumerate} \\renewcommand{\\theenumi}{(\\arabic{enumi})} \\item There are two-stage mechanisms to cause an accretion flow from the circumbinary disk onto the supermassive BBHs: First, the gas is guided to the two closest points on the circumbinary disk from the black holes by the tidal deformation of the circumbinary disk. Then, the gas is increasingly accumulated on these two points by the gravitational attraction of the black holes. Second, when the gas can pass across the maximum loci of the binary potential, the gas overflows via these two points and inspirals onto the black holes with the nearly free-fall velocity \\item In circular binaries, the gas continues to be supplied from the circumbinary disk onto the supermassive BBHs (i.e. the gap is always closed). \\item In eccentric binaries, the material supply undergoes the periodic on/off transitions with an orbital period because of the periodic variations of the binary potential. \\item While the mass-capture rates exhibit little orbital-phase dependence in circular binaries, they significantly modulate with the orbital phase due to the gap opening/closing cycles, in eccentric binaries. This could provide the observable diagnosis for the presence of supermassive BBHs in three-disk systems at the galactic center. \\item The circumblack-hole disks are formed around each black hole regardless of the orbital eccentricity and the mass ratio. \\end{enumerate}  \\bigskip  Authors are grateful to the anonymous referee for constructive comments. K.H. is grateful to Noboru Kaneko, Atsuo~T. Okazaki, James R. Murray, Yasuhiro Asano and Satoshi Tanda for their continuous encouragement. K.H. is also grateful to Atsuo~T. Okazaki, Jun Fukue and Shoji Kato for helpful discussions. Authors thank YITP in Kyoto University, where this work was also extensively discussed during the YITP-W-05-11 on September 20-21, 2005. The simulations reported here were performed using the facility of the Centre for Astrophysics \\& Supercomputing at Swinburne University of Technology, Australia and of YITP in Kyoto University. This work has been supported in part by the Grants-in-Aid of the Ministry of Education, Science, Culture, and Sport and Technology (MEXT; 14079205 K.H. \\& S.M., 16340057 S.M.), and by the Grant-in-Aid for the 21st Century COE Scientific Research Programs on \"Topological Science and Technology\" and \"Center for Diversity and Universality in Physics\" from the MEXT."
    },
    "0609/astro-ph0609002_arXiv.txt": {
        "abstract": "Deep X-ray surveys by \\chandra\\ and \\xmm\\ have resolved about 80\\% of the 2--10~\\kev\\ cosmic extragalactic X-ray background (CXRB) into point sources, the majority of which are obscured AGN. The obscuration might be connected to processes within the host galaxy, possibly the star-formation rate. Here, we use the results of CXRB synthesis calculations as input to detailed Cloudy simulations in order to predict the evolution of AGN properties at several mid-IR wavelengths. Computations were performed for three different evolutions of the AGN type 2/type 1 ratio between $z=0$ and $1$: where the ratio increased as $(1+z)^{0.9}$, as $(1+z)^{0.3}$ and one with no redshift evolution. Models were calculated with the inner radius of the absorbing gas and dust at 1~pc or at 10~pc. Comparing the results of the calculations to combined X-ray and \\spitzer\\ data of AGN shows that the predicted spectral energy distributions are a good description of average AGNs found in the deep surveys. The existing data indicates that the mid-IR emission from an average AGN is best described by models where the attenuating material is $\\sim 10$~pc from the central engine. We present the expected \\spitzer\\ cumulative number count distributions and the evolution of the total AGN (type 1 + type 2) luminosity function (LF) between $z=0$ and $1$ at rest-frame 8\\micron\\ and 30\\micron\\ for the three evolutionary scenarios. The mid-IR AGN LF will be an excellent tool to measure the evolution of the covering factor of the gas and dust from $z \\sim 0$ to $1$. ",
        "introduction": "\\label{sect:intro} The earliest optical investigations of active galactic nuclei (AGN) identified two major classes of sources: type 1 AGN exhibited both broad (FWHM $\\sim 10^4$~km ~s$^{-1}$) permitted lines and narrow (FWHM $\\lesssim 10^3$~km ~s$^{-1}$) forbidden lines in their spectra, while type 2 AGN presented only the narrow lines \\citep{kw74}.  Unification models of AGN propose that this difference can be simply explained by an obscuring medium that surrounds a basic black hole/accretion disk system, and that different lines-of-sight into and through this obscuration give rise to the variety of observational properties witnessed in AGN phenomenology \\citep[e.g.,][]{ant93}. The obscuring medium is often visualized as a geometrically thick torus comprised of dust and molecular gas and of order a parsec in size, situated between the broad-line and narrow-line regions \\citep[e.g.,][]{pk92a}. Additional obscuration is likely to occur in larger structures associated with the host galaxy \\citep[e.g.,][]{mr95}. The net unobscured view from the AGN must encompass only about 20\\% of the sky to explain the roughly 4:1 ratio of obscured type 2 to unobscured type 1 AGN in the local Universe \\citep[e.g.,][]{mr95,hfs97}. This picture is probably consistent with the recent finding from Sloan data by \\citet{hao05} that the type 2/type 1 ratio is close to unity, because about half of the AGN appear to be well hidden in their host galaxies and would not be seen by Sloan but can only be detected with very deep spectra such as those used by \\citet{mr95} and \\citet{hfs97}. The unification model explains a number of observations. For example, the near infrared (IR) emission of many AGN shows a break near 1\\micron\\ \\citep[e.g.,][]{rie78,neu79,bar87,san89}, signifying dust that is at or near its sublimation temperature and hence close to the AGN. Recent \\chandra\\ and \\xmm\\ X-ray observations have found that narrow Fe K$\\alpha$ lines are common in both type 1 and type 2 objects, which implies the presence of cold, low-velocity gas intercepting a significant fraction of the nuclear emission \\citep{yp04,nan06}.  While the existence of the absorbing gas around AGN is well established, the geometry, distribution, origin, and evolution of the material is largely unknown. Observations of the soft X-ray absorption can measure the equivalent neutral column density of the absorbing gas along the line of sight, \\nh. There is a good correspondence (at the $\\sim$ 80\\% level; \\citealt{toz06}) between the sources with \\nh\\ $\\ge$ 10$^{22}$ cm$^{-2}$ and those optically identified as being of type 2. This correspondence implies that, on average, the X-ray absorbing gas does lie outside the broad-line region. However, observations of \\nh\\ variability on week- to year-long timescales \\citep{ris02,ris05} argue against a simple geometry. Further evidence for a complicated geometry is found in the mid-IR.  The IR spectra of most AGNs (in $\\nu f_\\nu$ units) peak between 3 and 10\\micron, and roll over at $\\sim$ 100\\micron\\ \\citep{an03,kur03}. This behavior can be described by thermal emission by dust ranging in temperature from $\\sim 40$~K to the sublimation temperature of 1000--2000~K. Part of the cooler emission may be heated by star formation rather than the AGN \\citep[e.g.,][]{mr95,sch06}. Interestingly, the spectral energy distributions of type 1 and type 2 AGN are similar in the mid-IR ($\\lambda > 5$\\micron) \\citep{sm89,fad98,kur03,ah03,lutz04,rig04}, in conflict with the predictions of many compact torus models \\citep{pk92a,gd94,err95,gra97}. All of these observations argue for a more distributed or complex geometry to the absorbing material around the AGN \\citep{kur03}.  If the obscuring material in AGN is distributed in a complex manner, then large object-to-object variations in properties would be expected, complicating the interpretation of both the AGN characteristics and the relationship to the host galaxy (if any). Progress can be made, however, by determining the ratio of type 2 to type 1 objects (hereafter denoted as $R$) at differing AGN luminosities and redshifts. Under the unified models, $R$ is related to the covering factor of the absorbing gas and dust, so changes in $R$ give clues to the nature of the obscuring structures. For instance, the simplest unified model, that of an unevolving torus, predicts no evolution of $R$. However, over the last few years, surveys at multiple wavelengths have noted that the fraction of type 1 objects increases with AGN luminosity \\citep{ueda03,bar05,simp05,hao05}. This trend may simply be a selection effect \\citep{trei04}, because it is easier to find type 1 high luminosity AGN (i.e., quasars) with their broad optical lines than type 2 objects, which many models predict to be deeply buried \\citep[e.g.,][]{hopk05}. Moreover, at high redshifts dilution from galaxy light can overwhelm the narrow and relatively weak lines of a type 2 AGN, increasing the difficulty of identification \\citep{mfc02}. Alternatively, the luminosity dependence of $R$ may be real, and imply that some part of the obscuring material is close enough to the central engine to be affected by high luminosities, perhaps through dust sublimation \\citep{law91} or radiative acceleration of magnetocentrifugal winds \\citep{kk94}.  There are also hints that $R$ increases with redshift at moderate AGN luminosities. Both \\citet{bar05} and \\citet{laf05}, in their analyses of deep X-ray surveys, found indications for an increase in the type 2 fraction from $z = 0$ to $z \\sim 1$. The sources discovered in these observations account for a large fraction of the cosmic X-ray background (CXRB) in the 2--10~\\kev\\ range \\citep{mush00,bh05,wor05}, whose spectrum can be fitted only if $R \\sim 3-4$ \\citep{gill04}.  The implication of these observations is that if the obscuring medium around an AGN evolves with redshift then it must be connected in some way to the evolution of the host galaxy. The redshift distribution of the CXRB sources peaks at $z \\sim 1$ \\citep{toz01,bar02,bar05}, interestingly close to where the cosmic star-formation rate (SFR) density also reaches its maximum \\citep{hopk04}. This coincidence in redshift distribution suggests a possible connection between the absorbing gas and dust around an AGN and the SFR within the host galaxy (already seen in the Seyfert~1 NGC~3227; \\citealt{dav06}). With this motivation, \\citet{bem06} determined that a $R$ that evolves with both luminosity and $z$ (constrained by the type 1 fractions observed by \\citealt{bar05}) can fit the observed CXRB spectrum and X-ray number counts. However, the CXRB itself provides little constraint on the evolution as both rapid, slow or zero \\citep[e.g.,][]{tu05} evolution can account for the data. A measurement of the true extent of any $R$ evolution will provide vital information for models of galaxy and black hole evolution over this redshift range. It is therefore important to search for other observational signatures that can differentiate among various evolutions of the AGN obscuring material.  The natural next step is to consider the evolution of AGN, in particular obscured AGN, in the mid-IR where the absorbed radiation will be re-emitted.  With the recent launch of the \\spitzer\\ Space Telescope, deep sensitive surveys of AGN in the mid-IR are now being performed, and are finding many high-$z$ ($z \\gtrsim 0.5$) AGN \\citep{lacy04,rig04,fra05,stern05,ah06,barm06,brand06}. Therefore, this paper presents predictions for the mid-IR number counts and luminosity functions (LFs) for three different evolutions of the type 2/type 1 AGN fraction between $z=0$ and $1$. All three evolutions are chosen to fit the CXRB, and therefore the mid-IR predictions are consistent with the current best X-ray constraints on AGN properties and evolution. In contrast to most previous work \\citep[e.g.][]{silva04}, the IR predictions are not based on observed spectral energy distributions (SEDs) but are computed using the photoionization code \\cloudy. This method has the advantages of being able to directly relate the X-ray properties to other wavelengths and to allow interesting physical properties of the obscuring medium to be varied and possibly constrained.  We begin in the next section by reviewing the CXRB results of \\citet{bem06}, in particular the three evolutions of $R$ that will be investigated here. The following section (\\S~\\ref{sect:calcs}) describes the \\cloudy\\ models in detail, including a description of our assumptions, and how we average individual spectra together before estimating observables. Sect.~\\ref{sect:tests} then checks how well these model SEDs compare against the observed AGN X-ray and mid-IR properties. Sect.~\\ref{sect:agn} presents the number counts and LFs for the three AGN evolutions and compares the predictions to the available data. In Sect.~\\ref{sect:sf} we attempt to take into account the effects of star formation in the model spectra and present revised predictions. Finally, we discuss our results in Sect.~\\ref{sect:discuss} and summarize in Sect.~\\ref{sect:summ}.  This paper assumes the standard first-year \\textit{WMAP} $\\Lambda$-dominated cosmology: $H_0=70$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{\\Lambda}=0.7$, and $\\Omega_{m}=0.3$ \\citep{spe03}. ",
        "conclusions": "\\label{sect:discuss} The goal of this paper was to investigate how three possible evolutions of the AGN type 2/type 1 ratio (all of which fit the CXRB) manifest themselves in the mid-IR. The predictions were computed by employing the photoionization code \\cloudy\\ to calculate the SED from an obscured AGN over a large range of X-ray luminosity, $L_{X}$, and redshift.  In this way, a direct connection between the observed X-ray and IR fluxes could be predicted. The covering factor of the absorbing gas and dust changed with $z$ and $L_{X}$ as predicted by CXRB synthesis models and the unified model.  The results presented in the above two sections show that this photoionization-based method is very successful. The family of computed SEDs indeed capture many of the average properties in typical AGN observed in the deep surveys of \\chandra, \\spitzer\\ and \\xmm, including infrared colors and observed IR and X-ray fluxes. Typically, investigations of the IR properties of AGN involve either detailed dust radiative transfer models \\citep[e.g.,][]{gra97} that provide no connection to the X-rays, or are based on phenomenological SEDs and are unable to provide constraints on physical parameters \\citep[e.g.,][]{silva04}. Our approach is not designed to fit the SEDs of individual objects, but we can use it to constrain physically interesting quantities in the ensemble of objects. In this way, a picture of an average AGN can emerge, uncluttered from the significant source-to-source variability that affects the study of small samples.  \\subsection{The Location of the Absorbing Gas} \\label{sub:radius} Two different inner radii were used in computing the \\cloudy\\ models, 1~pc and 10~pc (see \\S~\\ref{sub:inner}), as tests for whether large samples of AGN show evidence for dust emission at these average distances. Although only these two radii were considered here, comparing the results from the two sets of models with the available data gave rise to the interesting result that almost all observational tests preferred the 10~pc models over the 1~pc ones (Figs.~\\ref{fig:nhhard} and~~\\ref{fig:brandcolor}). In particular, the GOODS number-count distributions were best described by the 10~pc models from observed-frame 5.7\\micron\\ to 24\\micron\\ (Fig.~\\ref{fig:counts}).  It is important to note here that both the $r=1$ and $10$~pc sets of models were assumed to have the same constant density of $10^4$~cm$^{-3}$. Thus, all \\cloudy\\ runs with $\\log N_{\\mathrm{H}} < 23.5$ have pathlengths less than 10~pc (ignoring ionization effects). As a result the $r=1$~pc models will be dominated by hot dust, as the column does not extend to great enough distances to produce significant warm emission. A new lower density distribution that would have an inner radius of 1~pc and be able to produce the correct 24\\micron\\ flux is certainly conceivable. On the other hand, if the obscuring medium is clumpy with each clump being roughly the same density \\citep{dv05}, then our setup may not be too unrealistic. An investigation using different density distributions will be the subject of future work.  Although the 10~pc model seems to be the best description of data observed at the mid-IR \\spitzer\\ wavelengths, the observed SEDs of many AGN show that hot dust from inside this radius is also required. The 10~pc \\cloudy\\ models do predict dust temperatures close to and at the sublimation value at the inner radius of the cloud, but only for the more luminous AGN models (see \\S~\\ref{sub:inner}). Thus, there must be dust emission over a range of radii and temperatures, in agreement with SED fitting of individual objects \\citep[e.g.,][]{kur03}. This conclusion is also consistent with observations searching for variability in the IR continua of AGN that find little to no variations at mid- to far-IR wavelengths on day to month long timescales \\citep{rl81,em87,cle88}. In contrast, near-IR (JHK band) variability of local Seyferts is observed with measured timelags from the optical/UV placing the emitting region $< 1$~pc from the continuum source \\citep{cla89,sug06,min06}.  \\citet{nm99} report long term monitoring of nearby, low-luminosity quasars at 10.6\\micron. At the typical redshifts of their sample, this wavelength is very close to 8\\micron\\ in the rest-frame. They find slow, multiple-year variability but conclude that in no case are the results in contradiction to a thermal origin for the emission in the objects with luminosities $< 10^{13}$ L$_\\odot$. Specifically, the variability time scales are consistent with our conclusion that the 8\\micron\\ emission arises from thermal re-radiation by dust in a component with a typical inner radius of $\\sim 10$~pc. They also summarize previous measurements of luminous, radio-loud quasars, such as 3C~273; in this object in particular, observations at $\\sim 8$\\micron\\ in the rest frame show variations that are too fast to interpret as thermal radiation. Since such objects are exceptional among the total quasar population, it is unlikely that such behavior would strongly affect our results for the ensemble properties of these sources.  In the future, model grids can be extended to even larger inner radii to establish a pair of bracketing values for this key parameter. It will also be possible to break observed data into redshift and luminosity slices to determine if there is an evolution of the inner radius with $z$ and/or $L_{X}$ \\citep[e.g.,][]{hatz06}.  \\subsection{Determining the Evolution of $R$} \\label{sub:evol} Sections~\\ref{sect:agn}~and~\\ref{sect:sf} showed that the mid-IR LF of AGN will be able to discriminate among different models of $R$, the type 2/type 1 ratio. Out of the wavelengths considered here, the rest frame 30\\micron\\ LF between $z=0$ and $1$ gives the largest differences among evolutionary scenarios, as well as being the least contaminated due to star-formation emission. This last point needs to be confirmed with a more sophisticated prediction of the IR emission due to star-formation.  Assuming the 10~pc models are a more accurate guide, then the right-hand panel of Fig.~\\ref{fig:30umlf} indicates that largest differences among evolutions of $R$ occurs at luminosities with $\\log \\nu L_{\\nu} > 45$. This is past the knee of the LF, and so will require wide surveys to pick up the large numbers of luminous AGN. The observations will also require significant depth in order to trace the LF out to $z \\sim 1$.  Another issue in observationally testing the LF predictions is accessible wavebands. To really understand the evolution of the obscuring region, one would like to trace the LF from low redshift (say, $z=0.2$) to $z \\sim 1$. The LF predictions are in the rest-frame, which means that for 30\\micron\\ observations at $z \\sim 0.6$, the required flux is being radiated at 48\\micron. One would then be forced to perform multi-wavelength observations and SED fitting for each object in order to predict the flux at the required wavelength. This SED fitting could be relatively simple as the largest uncertainties for determining the LF will be in the space density statistics.  Given these difficulties, it is natural to ask if the mid-IR LF is the easiest method to determine the evolution of $R$. A perhaps more straightforward strategy would be to measure the ratio of narrow-line to broad-line AGN in the optical as functions of $z$ and $L_{\\mathrm{[O\\ III]}}$, the luminosity in the [O III] $\\lambda$5007 line. Alternatively, we could measure the ratio of X-ray AGN with absorbing column densities $\\log N_{\\mathrm{H}} < 22$ to those with $\\log N_{\\mathrm{H}} > 22$ as a function of $z$ and $L_{X}$, although an accurate value of \\nh\\ requires an X-ray spectrum with sufficient counts. Since it is currently difficult to obtain enough counts for spectral fitting from faint sources, hardness ratios (which assume a spectral slope) are often used as an alternative. Thus, both of these other methods to track the evolution of $R$ require serious observational investments, but, more problematically, they also are both subject to severe selection effects. As is becoming increasingly apparent, optical selection of AGN (i.e., selection on the basis of high-ionization lines) has significant limitations \\citep[e.g.,][]{netz06} and many (mostly type 2 AGN) are missed due to obscuration or dilution \\citep{mfc02,eck06,rig06}. Likewise, X-ray selection is limited because the observable passband for the most sensitive telescopes ends at $\\sim 10$~\\kev, and the observations are thus insensitive to objects with high column densities, especially those with \\nh $> 10^{24}$~cm$^{-2}$, the Compton-thick AGN. The high energy detectors on \\textit{Swift} and \\textit{Integral} have been able to conduct surveys, but are limited to the nearby Universe due to low sensitivity \\citep{mark05,bas06}. Future observations with the Hard X-ray Detector on board \\textit{Suzaku} may help, but it is a non-imaging instrument. The cancellation of the hard X-ray imaging mission \\textit{NuStar} makes finding significant numbers of the Compton-thick AGN by their hard X-ray emission practically impossible until the launch of \\textit{Constellation-X}.  The advantage of the mid-IR method described here is that it is less affected by the above selection effects. The predicted LFs shown in Fig.~\\ref{fig:30umlf} are for both type 2 and type 1 AGN, therefore detailed optical emission line spectra are not required. Compton-thick objects should also be found in the infrared, as the absorbed X-ray radiation will be re-emitted by the dust. The difficulty, of course, is determining which IR sources are predominantly powered by AGN as opposed to star-formation, although methods that examine the spectral shape of the near- and mid-IR have had some success \\citep{ah06,brand06}, and AGN are well known to have very weak PAH emission \\citep{weed05}. However, X-ray selection is still the best way to identify AGN, as Compton-thin objects can be observed to high redshifts. Deep infrared observations are then required to find the Compton-thick sources. Comparisons between multiband observations and simulations similar to those performed here will be useful in determining the AGN contribution to these potential Compton-thick sources. This can be done by adding in contributions from the host galaxy to the photoionization models, and computing a large grid of IR properties over a range of AGN luminosities and redshifts for comparison to both near- and mid-IR data.  Another interesting effect seen in the IR and not in the other wavelengths is the dependence on the geometry of the obscuring medium. Figures~\\ref{fig:8umlf} and~\\ref{fig:30umlf} have shown different shapes for the LF depending on the inner radius of the absorbing gas and dust. For the one data point at 30\\micron, it was difficult to determine the best evolution of $R$ without knowing the geometry. Fortunately, there are significant differences in the predicted LFs and number counts for the different inner radii. Therefore, it is possible to break the degeneracy between $r$ and $R$, as was done at rest-frame 8\\micron\\ (Fig.~\\ref{fig:8umlf}).  Finally, the mid-IR LFs presented above are relatively insensitive to the actual \\nh\\ distribution assumed. We considered both a `simple' \\nh\\ distribution, where all column densities had equal weight within the type 1 or 2 designations (eqn.~\\ref{eq:simple}), and the \\citet{rms99} \\nh\\ distribution in which half of the type 2 AGN are Compton thick and 75\\% had $\\log N_{\\mathrm{H}} > 23$ (eqn.~\\ref{eq:risaliti}). These weights were used in constructing both the `unified' and \\nh-averaged SEDs. Despite this, the two distributions made only negligible differences to the predicted number counts and LFs. This is because the additional weight afforded to Compton thick AGN in the \\citet{rms99} \\nh\\ distribution has only a small effect in the mid-IR (Fig.~\\ref{fig:sedexample}). Therefore, although mid-IR LFs will not be able to quantify the relative number of Compton thick AGN, they are still very sensitive to the overall evolution of the AGN type 2/type 1 ratio.  \\subsection{The Role of Star Formation} \\label{sub:stars} Star formation is a key process in this investigation in two distinct ways. First, it acts as a contaminant of the mid-IR emission from the obscured AGN, masking the evolutionary effects we are trying to investigate. Using a straightforward estimate of the effects of star formation on the mid-IR, we found in Sect.~\\ref{sect:sf} that for a constant SFR $\\leq 10$~M$_{\\odot}$~yr$^{-1}$ the impact at 30\\micron\\ was confined to $\\log \\nu L_{\\nu} < 44.5$, which translates to X-ray luminosities of $\\log L_{X} \\lesssim 43.5$ for the 1:1/10pc models or $\\log L_{X} \\lesssim 43.7$ for the nozevol/10pc and 4:1/10pc grids. The affected luminosities are low enough that the rest-frame 30\\micron\\ LFs may still be used to discriminate possible different evolutions of the AGN type 2/type 1 ratio (Fig.~\\ref{fig:30umlf-sf}). The method may prove more problematic with galaxies with average SFRs $> 10$~M$_{\\odot}$~yr$^{-1}$.  The sources that produce the majority of the CXRB have been identified to be predominantly obscured AGN with $43 \\leq \\log L_{X} \\leq 44$ \\citep{bar05}. These objects are typically found at $z \\lesssim 1$, which may signify a relationship between the obscuring material around an AGN and the star formation in the host galaxy \\citep{fab98,fra99,bem06}. It is this possibility that makes measuring the evolution of $R$ an important goal in the study of galaxy evolution. However, there may be a potential difficulty in employing the mid-IR LFs calculated here because, as already noted, the important luminosity range to distinguish the various evolutions is $\\log L_{X} \\gtrsim 44$, larger than that of the objects that dominate the X-ray background. The reason for this is that the three evolutions considered have very different predictions for the ratio of type 2/type 1 quasars (Fig.~\\ref{fig:contours}), but quasars probably have a very different formation and evolution history than the lower luminosity Seyferts that comprise the bulk of the CXRB. Quasars are more commonly found at $z \\gtrsim 1$, are likely formed by the mergers of massive galaxies, and simulations show their host galaxies have little ongoing star formation within the galaxy at late times \\citep[e.g.,][]{hopk05}. The evolutions of $R$ described by eqns.~\\ref{eq:4to1} and~\\ref{eq:1to1} assume that the AGN covering factor increases with $z$ independent of the AGN luminosity. If the attenuating gas and dust is connected to the host galaxy SFR then this may not be correct at large $L_{X}$, and indeed very tentative evidence for a slower increase of $R$ with $z$ at larger $L_{X}$ was found by \\citet{bem06}. If this is the case, then evidence for this type of luminosity-dependent redshift evolution will appear in the 30\\micron\\ LFs, and can be measured using the techniques presented here. Furthermore, the effects of our luminosity-independent assumption are most important at high $z$, but Fig.~\\ref{fig:30umlf} show that significant variations in the rest-frame 30\\micron\\ LF due to the different evolutions of $R$ still occur at low redshifts. Thus, a determination of the LF at redshifts $z \\sim 0$--$0.2$ will result in an important constraint on the local value of $R$. Then, tracking the 30\\micron\\ LF with redshift will yield a measurement of how the AGN type 2/type 1 ratio changes with $z$ and $L_{X}$, a key parameter in understanding the formation and evolution of these galaxies.  \\subsection{Testing the Unified Model} \\label{sub:constantcovering} To this point, all the calculations presented in this paper assumed that the unified model of AGN was correct. That is, the fraction of type 2 AGN at a given $L_X$ and $z$ is equal to the covering factor of the absorbing material around the AGN. These models can also be used to test whether the unification scenario is valid at all luminosities.  As an example, we searched for \\textit{IRAS} objects satisfying $|b|>35$ degrees, $f_{\\nu}$(60\\micron)$>0.2$~Jy and with redshifts between 0.8 and 1.2. This results in 13 objects in which we identified 8 quasars using the literature. Two of them ([HB89]0235+164 and [HB89]1308+326) are blazars and are omitted from further consideration. An additional 3 (IRAS F09121+2430, PG1248+401 and B31340+407) were removed because they have weak detections at 60\\micron\\ and the association with the quasars therefore is not firmly established. The remaining three quasars (2MASSiJ0119560-201022, 2MASSiJ1543519+162422 and PG1206+459) have strong detections by \\textit{IRAS} and the optical data show them to have normal QSO SEDs.  Assuming $z \\sim 1$, the rest-frame 30\\micron\\ luminosities of these quasars are $\\nu L_{\\nu} \\sim 46.8$. Fig.~\\ref{fig:30umlf} shows that the 1:1/10pc evolutionary model, as a result of its $\\sim$3:1 type 2/type 1 ratio (Fig.~\\ref{fig:contours}), can reach a 30\\micron\\ luminosity of $\\nu L_{\\nu} = 46.5$ at $z=1$, but this occurs only when the input X-ray luminosity is $\\log L_{X} = 48$. Yet, the 2--10~\\kev\\ X-ray luminosity of PG1206+459 is $\\sim 10^{45}$~erg~s$^{-1}$ \\citep{pic05}. Indeed, using the three \\textit{IRAS} sources to estimate the 30\\micron\\ LF at $\\nu L_{\\nu} = 46.8$ results in a lower-limit of $>10^{-10}$~Mpc$^{-3}$, about 2 orders of magnitude larger than the predicted LF from Fig.~\\ref{fig:30umlf}. The best way of reducing this discrepancy would be if the models could produce these high IR luminosities at lower AGN input luminosities.  It seems that the assumed geometric configuration of a relatively compact absorbing medium situated at 1~pc or 10~pc from the AGN cannot produce the necessary 30\\micron\\ emission at high $L_X$ to account for the \\textit{IRAS} limit. Recall that because of the assumed density of 10$^4$~cm$^{-3}$, the pathlength through the absorber is only $15$~pc for $\\log N_{\\mathrm{H}} = 23$ and $\\log L_X = 48$ (taking into account the ionization effects), therefore at high AGN luminosities there is only a small column of dust cool enough to emit at 30\\micron. Material at greater distances is needed to increase the 30\\micron\\ emission. A simple way to test this idea is to calculate a new grid of \\cloudy\\ models with the inner radius remaining at 10~pc but the density now lowered to 10$^2$~cm$^{-3}$. This greatly increases the pathlength of the absorbing material (it is now over 1~kpc in the example above). Figure~\\ref{fig:iras} plots the rest-frame 30\\micron\\ LFs computed using the models with a lower density absorber as well as the previous derived ones that assume the more compact geometry. Only results for $\\log \\nu L_{\\nu} \\gtrsim 44.7$ are plotted for the low density absorber. The greater pathlength afforded by the lower density cloud now produces enough 30\\micron\\ emission to account for the \\textit{IRAS} lower-limit. The right-hand panel of Fig.~\\ref{fig:iras} shows the same results but for the rest-frame 8\\micron\\ LF. At this wavelength, the differences in the LF between the two assumed densities are small.  The implication of this simple test is that the high luminosity quasars require absorption and remission of nuclear radiation over a large range of radii, while the current data available for the lower luminosity quasars and Seyfert galaxies can be adequately described by a more compact absorber. The type 2/type 1 ratio at the \\textit{IRAS} luminosities for the three evolutions considered in this paper varies from $\\sim$1:1--4:1 (Fig.~\\ref{fig:contours}), a range that is consistent with the latest observational constraints \\citep{mart05,mart06}. Under the unification model, this translates into covering fractions of 0.5--0.8. Producing high covering fractions is significantly simpler if the obscuring material is close to the AGN, and not spread out over a large range of radii, as implied by the above results. Clearly, the change in mid-IR LF with luminosity (in addition to the possible redshift evolution concentrated on in previous sections) contains fundamental clues on the nature and geometry of the obscuring medium and therefore the host galaxy. Future observations are needed for this technique to fulfill this promise.  We conclude that the unification model based on a relatively compact absorber is consistent with the available data for moderate AGN luminosities (Seyfert galaxies) but fails when extending this to the high luminosity quasars. This adds to the evidence that the evolution and fueling of quasars are very different from those that drive the lower luminosity Seyfert galaxies."
    },
    "0609/astro-ph0609528_arXiv.txt": {
        "abstract": "We analyzed two XMM-Newton observations in the direction of the high density, high latitude, neutral hydrogen cloud MBM20 and of a nearby low density region that we called the Eridanus hole. The cloud MBM20 is at a distance evaluated between 100 and 200~pc from the Sun and its density is sufficiently high to shield about 75\\% of the foreground emission in the 3/4~keV energy band. The combination of the two observations makes possible the separation between foreground component, due to the Local Bubble and possibly charge exhange within the solar system, and the background one, due primary to the Galactic halo and unidentified point sources. The two observations are in good agreement with each other and with ROSAT observations of the same part of the sky and the O{\\tiny VII} and O{\\tiny VIII} intensities are O{\\tiny VII}=$3.89\\pm0.56\\textrm{~photons~cm}^{-2}\\textrm{~s}^{-1}\\textrm{~sr}^{-1}$, O{\\tiny VIII}=$0.68\\pm0.24\\textrm{~photons~cm}^{-2}\\textrm{~s}^{-1}\\textrm{~sr}^{-1}$ for MBM20 and O{\\tiny VII}=$7.26\\pm0.34\\textrm{~photons~cm}^{-2}\\textrm{~s}^{-1}\\textrm{~sr}^{-1}$, O{\\tiny VIII}=$1.63\\pm0.17\\textrm{photons~cm}^{-2}\\textrm{~s}^{-1}\\textrm{~sr}^{-1}$ for the Eridanus hole.  The spectra are in agreement with a simple three component model, one unabsorbed and one absorbed plasma component, and a power law, due unresolved distant point sources. Assuming that the two plasma components are in thermal equilibrium we obtain a temperature of 0.096~keV for the foreground component and 0.197~keV  for the background one.  Assuming the foreground component is due solely to Local Bubble emission we obtain a lower and upper limit for the plasma density of 0.0079 $cm^{-3}$ and 0.0095 $cm^{-3}$ and limits of 16,200 $cm^{-3}K$ and 19,500 $cm^{-3}K$ for the plasma pressure, in good agreement with theoretical predictions. Similarly, assuming that the absorbed plasma component is due to Galactic halo emission, we obtain a plasma density ranging from 0.0009 $cm^{-3}$ to 0.0016 $cm^{-3}$, and a pressure ranging from $3.8\\times10^{3}$ to $6.7\\times10^{3}$ $cm^{-3}K$. ",
        "introduction": "Optical and ultraviolet absorption line studies indicate that the Sun is located in a region of the interstellar medium of our Galaxy that is deficient in neutral matter, the Local Cavity \\cite{Cox87, Frisch83, Bochkarev87}. Observations of the X-ray background at energies below 1 keV indicate that this cavity is largely filled with million degree plasma, the Local Bubble (LB), and it extends for distances of the order of 100~pc in most directions.  Fits to data obtained with proportional counters have indicated temperatures near one million degrees for the LB, if one compares the shape of the pulse height distribution to predictions of models with normal solar abundances and collisional equilibrium ionization structure of the plasma \\cite{Snowden98}. Higher resolution spectral data over the 150-300~eV range obtained near the galactic plane with a Bragg crystal spectrometer \\cite{Sanders01} confirm the thermal nature of the LB X-ray emission, but indicate that the above model picture is a bit na\\\"ive. The data suggest that the LB plasma is not in equilibrium and the abundances are probably not solar.  High resolution observations of the diffuse background at high galactic latitude obtained by ASCA in the 500-1000~eV range using CCD detectors \\cite{Gendreau95} and by the XQC sounding rocket experiment in the 50-2000~eV range using microcalorimeter detectors \\cite{McCammon02} detect the existence of the O~{\\tiny VII} and O~{\\tiny VIII} lines, but cannot separate the LB emission from galactic halo and extragalactic components. A more recent observation of the neutral hydrogen cloud MBM12 performed with Chandra \\cite{Smith05} is also not able to clarify the picture as it seems affected by charge exchange within the solar system. A cleaner observation of MBM12 and a blank field less than 3$^\\circ$ away has recently been performed with Suzaku \\cite{Smith06}. The Suzaku observation does not seem to be affected by a strong contribution of charge exhange within the local system and the use of the second field of view helps in separating the local and distant components of the diffuse soft X-ray background. ",
        "conclusions": "In studying the Diffuse X-ray Background (DXB), a challenging task is the separation between the different components contributing to it. This is particularly difficult due to the low DXB flux compared to the relatively high detector background. The contemporary study of a high latitude, high density cloud and a low density region nearby has proven an effective tool for the separation of foreground and background components. In our investigation we found a simple, efficient method to evaluate the detector background and the comparison of the two pointing gives us good separation between the two main DXB plasma components, i.e., Local Bubble and Galactic halo. The results are consistent with a simple three components model of the DXB: an unabsorbed plasma component with temperature of 0.096~keV, an absorbed plasma component with temperature of 0.197~keV, and a power law, due to unresolved distant point sources.  Assuming the foreground component is due solely to Local Bubble emission we obtain a lower and upper limit for the plasma density of 0.0079 $cm^{-3}$ and 0.0095 $cm^{-3}$ and 16,200 $cm^{-3}K$ and 19,500 $cm^{-3}K$ for the plasma pressure, in agreement with theoretical predictions (e.g., \\cite{Smith01}) and similar observations (e.g., \\cite{smith06}). Similarly, assuming that the absorbed plasma component is due to Galactic halo emission, we obtain a plasma density ranging from 0.0009 $cm^{-3}$ to 0.0016 $cm^{-3}$, and a pressure ranging from $3.8\\times10^{3}$ to $6.7\\times10^{3}$ $cm^{-3}K$.  We also tested the possibility that the plasma is not in equilibrium using a Plane-parallel Shock Model \\cite{Borkowski00}. However, while we obtained good fit and the electron density of ranges from 0.01 $cm^{-3}$ to 0.012 $cm^{-3}$, we derived a value for the age of the LB of $\\le ~0.6~Myr$, which is quiet less in comparison to most acceptable models."
    },
    "0609/astro-ph0609573_arXiv.txt": {
        "abstract": "Development of multi-layer optics makes feasible the use of X-ray telescope at energy up to 60-80~keV: in this paper we discuss the extension of photoelectric polarimeter based on Micro Pattern Gas Chamber to high energy X-rays. We calculated the sensitivity with Neon and Argon based mixtures at high pressure with thick absorption gap: placing the MPGC at focus of a next generation multi-layer optics, galatic and extragalactic X-ray polarimetry can be done up till 30~keV. ",
        "introduction": "\\label{sect:intro}  % \\subsection{Hard X-Ray Telescopes} In the History of X-ray Astronomy a crucial turning point was the introduction of X-ray Optics, first aboard some rockets and eventually on satellite with the HEAO-2/Einstein mission. Optics play two major functions:  \\begin{itemize} \\item to image a field of X-ray Sky, to study extended sources, to resolve nearby sources, to localize sources with high angular accuracy;  \\item to increase dramatically the signal to noise ratio, also with respect to other imaging systems, such as coded masks or modulation collimators. A source is to be compared with the fluctuations of background inside the point spread function of the telescope and and not with the total background of the detector. \\end{itemize}  This tremendous power of X-ray optics has been originally applied to softer X-rays with Einsten and ROSAT, and has been extended to higher energies to include the Fe lines region, with ASCA, SAX, Newton and Chandra. The conventional X-ray optics are limited in band because of the intrinsic limits set by the laws of total reflection. The critical angle of reflection is roughly proportional to the square root of the density of the reflecting material and to the wavelength of the photons. Thence, as the margin of increase in density, with respect to gold, is very limited, the extension to higher energies is committed to inner shells with progressively smaller projected area. The extension to higher energies is therefore payed with a large number of shells and with a proportional reduction of the field of view. Historically the X-ray telescopes have been used up to 10~keV. Some extended capabilities of Newton at higher energies have not been been effective because of an area mismatched with that at lower energies and because of a high background due to the large plate-scale.  In the future the revolution of the application of optics is about to be extended to higher energies, thanks to the use of telescopes with multi-layer coatings\\cite{Pareschi03}, alternating layers of high Z and low Z materials of gradually changing thickness. This technique provides effective reflection at angles up to 3 times those of mono-layer surfaces and makes feasible telescopes effective up to 60-80~keV (in practice up to the K absorption edge of the high Z reflecting material). At lower energies they are not as good as conventional optics, due to the photoelectric absorption of the high Z component, but an improvement could be the recently proposal of over-coating the multilayer surface with a carbon layer\\cite{Pareschi04}.  New missions based on the concept of multilayer telescopes have been proposed and NUSTAR and SIMBOL-X have passed a certain level of selection. Depending on various constructive parameters these missions can gain from two to three orders of magnitude in sensitivity with respect to the best previous experiments based on collimation (PDS) or coded mask (IBIS). Also future big telescopes for X-ray Astronomy, such as NEXT, CON-X or XEUS should include such optics. The science to be performed with these telescopes is mainly based on spectroscopy (cyclotron lines, hard tails, Compton-thick objects etc.). Also imaging can play a role. E.g. the Diffuse X-ray Background peaks around 40~keV. A hard X-ray telescope could resolve the large majority of it into discrete sources. Moreover imaging of extended objects with hard non-thermal components, such as some Super Nova Remnants and some Galaxy Clusters, could help to localize the hard emission and correlate it to the dynamics of the system.  \\subsection{X-Ray Polarimetry} In this paper we want to discuss another point: the capability to perform Polarimetry of X-ray Sources in the Hard X-ray Band and the benefit to employ multi-layer optics for this purpose.  Since the birth of X-ray Astronomy, theoretical analysis has suggested that measuring the linear polarization of X-rays could significantly improve our knowledge of Physical Processes and of geometry of regions of the source emitting X-rays and its transfer % toward the observer.  Here the question rises whether the extension to higher energies is only a matter of completeness or, on the contrary, a specific science is there that can only be performed at best at higher energies. We try to lay down a few hot topics of High Energy Astrophysics, for which the polarimetry at higher energies would be more easy or more significant. \\begin{itemize} \\item X-ray pulsator, resonance frequency, inversion of the polarization plane (normal - anomalous), higher polarization;  \\item polarization by scattering on disks: LMXRB, QSO, Seyfert-1;  \\item non thermal components in extended objects: SNR, Clusters out of the thermal component; \\item effects of GR on disks: Cyg X-1, AGN (increasing effects closer to the inner orbits); \\item blazars: transition from synchrotron to inverse Compton. \\end{itemize}  From these items it is apparent the strong interest to extend the measurements of polarization to hard X-rays. Is this feasible?  The conventional approach to X-ray polarimetry includes two major techniques:  \\begin{itemize} \\item Bragg Diffraction at $45^\\circ$. It provides an excellent response to the polarization, conventionally expressed with a modulation factor $\\mu$ close to 1. Conversely the overall efficiency of the process is very poor because only the photons within a narrow band around the angle that satisfies the Bragg condition are reflected. The efficiency can be increased by the use of mosaic crystals providing a larger band-pass.  \\item Compton scattering. It provides a good analyzer of the linear polarization if the system is designed in such a way to limit the scattering angles  around $90^\\circ$. Since the scattered photons will be distributed in angle according to the Klein-Nishina cross section, this will automatically limit the efficiency of the instrument. The optimal sensitivity will derive from a trade-off on the acceptance angles. At low energies, when the energy lost in the scattering is low, the scatterer must be a passive % low Z material, usually Lithium, encased in a thin beryllium can, to give mechanical consistency and to prevent oxidation or nitridation. This device loses the imaging properties of the optics and is background dominated. If the scatterer is not geometrically perfect and if the pointing does not keep the beam exactly on the axis of the scatterer, huge spurious polarizations are generated. At higher energies an active scatterer can be used, such as a detector based on plastic scintillator or, at even more high energies, on silicon. Such a device could be in principle position sensitive. But the active scattering polarimeters built so far never arrived below 50~keV. Active scatterers based on Silicon will be likely operative at energies even higher. \\end{itemize}  Bragg diffraction allowed the first and till now unique measure of X-ray polarization from an astronomical source: the polarimeter on board OSO-8 satellite, observing Crab Nebula for 3 days, succeeded to detect a polarization of 19.22$\\pm$0.92\\% at 2.6 and 5.2~keV aligned with optical polarization\\cite{Weisskopf76}. This confirmed that X-ray emission is due to synchrotron process, the same as other wavelengths.  Bragg diffraction and Compton scattering, in a focal plane version, were employed at best in the Stellar X-Ray Polarimeter\\cite{Kaaret89}. A thin pyrolithic graphite crystal was set above the focal plane to reflect (diffract) photons and to convey them to a detector on a \\textit{secondary} focal plane. The rotation of the whole would modulate the diffracted flux according to the amount and angle of polarization of the photons. The higher energy photons would pass the graphite, impinge on a lithium stick. Four detectors all around in a well configuration would detect scattered photons. SXRP was supposed to fly aboard the SRG Mission, that was never completed. But notwithstanding the huge SODART optics it would be dominated by the background at source level of a fraction of Crab. ",
        "conclusions": "We calculated sensitivity to hard X-ray polarization of Micropattern Gas Chamber, assuming Neon and Argon based mixures at different pressure and absoprtion gap thickness. Results demonstrate that current instruments filled with Argon + DME mixture can perform astronomical sources polarimetry up till $\\sim$30~keV, allowing the study of compact galactic and extragalactic sources like X-ray pulsar or blazar. For bright sources, spectral dependence of polarization could be measured in a day of observation time, opening a new view on studies of high energy process in astrophysics.  \\subsection"
    },
    "0609/astro-ph0609090_arXiv.txt": {
        "abstract": "We present kinematic measurements of a large sample of galaxies from the Team Keck Redshift Survey in the GOODS-N field.  We measure line-of-sight velocity dispersions from linewidths of integrated emission for 1089 galaxies with median redshift 0.637, and spatially resolved kinematics for a subsample of 380 galaxies.  This is the largest sample of galaxies to $z \\sim 1$ with kinematics to date, and allows us to measure kinematic properties without morphological pre-selection. Emission linewidths provide a dynamical measurement for the bulk of the blue galaxy population.  To fit the spatially resolved kinematics, we construct models with both line-of-sight rotation amplitude and velocity dispersion as fit parameters.  Integrated linewidth correlates well with a combination of the spatially-resolved velocity gradient and dispersion, and is a robust measure of galaxy kinematics.  The spatial extents of emission and continuum are similar and there is no evidence that linewidths are affected by nuclear or clumpy emission.  The measured rotation gradient is a strong function of slit position angle alignment with galaxy major axis, but integrated linewidth is largely independent of slit alignment.  Even in a subsample of galaxies with well-aligned slits, there are galaxies whose kinematics are dominated by dispersion ($V/\\sigma<1$) rather than rotation.  These are probably objects with disordered velocity fields, not dynamically hot stellar systems.  About 35\\% of the spatially resolved sample are dispersion dominated; galaxies that are both dispersion dominated and bright exist at high redshift but appear rare at low redshift.  This kinematic morphology may yield a probe of the evolutionary state of these objects.  Kinematic morphology is linked to photometric morphology in HST/ACS images: dispersion dominated galaxies include a higher fraction of irregulars and chain galaxies, while rotation dominated galaxies are mostly disks and irregulars.  Only one-third of chain/hyphen galaxies are dominated by rotation; high redshift elongated objects cannot be assumed to be inclined disks.  In a companion paper, we use the linewidths and rotation to measure evolution in the Tully-Fisher relation. ",
        "introduction": "Galaxy redshift surveys to $z\\sim 1$ allow the measurement of evolution in properties of the galaxy population at lookback times of half the present age of the Universe. To date, a number of surveys have studied galaxy properties such as luminosities, colors, radii, morphology and environment.  The internal kinematics of galaxies are equally interesting, as a physical probe of individual objects and to allow the measurement of scaling relations of the galaxy population. However, measuring kinematics requires moderate to high-resolution spectroscopy, which has until recently been difficult to obtain for large numbers of faint galaxies.   This paper measures kinematics of galaxies from emission lines, using Keck/DEIMOS spectra from the Team Keck Redshift Survey (TKRS) in the GOODS-N (Great Observatories Origins Deep Survey) field (Wirth \\etal\\ 2004; Giavalisco \\etal\\ 2004).  The TKRS provides spectra of 1437 galaxies drawn from a magnitude limited sample to $R_{AB} < 24.4$, at resolution $R \\sim 2100$. We select 1191 galaxies with good emission line detections and measure 1089 emission line velocity dispersions from integrated emission, with median $<z> = 0.637$. A subsample of 464 are selected for modeling spatially resolved rotation and dispersion profiles, of which 380 yield good-quality measurements.  We use these spatially resolved measures to probe the kinematic properties of high redshift galaxies and to test the integrated linewidths. In Paper II (Weiner \\etal\\ 2006), we use the linewidths and rotation velocities to measure evolution in the Tully-Fisher relation.  A number of previous works have observed galaxy internal kinematics from $0.1<z<1.0$, for the purpose of measuring Tully-Fisher relations. The pioneering studies of Vogt \\etal\\ (1996, 1997) modeled rotation curves for 17 galaxies of disky morphology with median $<z> = 0.47$ by combining Keck/LRIS slitlet spectra with structural information from HST photometry.  Several subsequent studies of rotation curves with similar modeling procedures have contained 20-100 galaxies with median redshifts $\\sim 0.4 - 0.5$ (e.g. Simard \\& Pritchet 1998; Vogt 2000; Ziegler \\etal\\ 2002; Milvang-Jensen et al.\\ 2003; B{\\\"o}hm et al.\\ 2004; Bamford \\etal\\ 2005,2006; Conselice \\etal\\ 2005; Nakamura et al.\\ 2006; Metevier \\etal\\ 2006). Generally these samples have been morphologically selected to be inclined disky objects, of which a majority show measurable rotation.  Here we study a sample with emission line kinematics that is essentially selected only on magnitude and emission line strength. We measure one-dimensional integrated velocity dispersion (linewidth), and for a subsample, spatially resolved rotation profiles; we discuss the properties of these velocity measures in Sections \\ref{sec-kinprops} and \\ref{sec-spatialkin}.  We show that the integrated linewidth is a fairly robust measure of the characteristic velocity of a galaxy, expanding the scope of the study beyond galaxies that are selected to be orderly rotating disks.  Indeed, we find from the spatially resolved rotation profiles that a significant number of high-redshift galaxies show kinematics that do not appear to be orderly rotation.  We adopt a LCDM cosmology with $h=0.7$, $\\Omega_M = 0.3$, and $\\Omega_\\Lambda = 0.7$.  Magnitudes quoted in this paper are in the AB system unless explicitly indicated as Vega. Section \\ref{sec-observations} discusses the sample and measurement methods, and section \\ref{sec-complete} presents the properties of the sample and completeness. Sections \\ref{sec-kinprops} and \\ref{sec-spatialkin} study the simulated and empirical properties of the kinematic measurements: integrated linewidths and spatially resolved velocities. In Paper II we use the kinematics to measure the evolution in the Tully-Fisher relation with redshift. ",
        "conclusions": "\\label{sec-conclude}  We have measured line-of-sight kinematic linewidths \\sigoned\\ for $\\sim 1000$ galaxies in the Team Keck survey of the GOODS-N field and spatially resolved line-of-sight rotation (\\vrot) and dispersion (\\sigtwod) profiles for 380 of these.  Most galaxies with linewidths are on the blue side of the color bimodality.  Linewidths from integrated spectra are a measure of internal kinematics that is relatively robust against observational effects, based on simulations of velocity fields and on comparisons of the linewidths with the spatially resolved kinematics.  However, the unknown geometries of high-redshift velocity fields, scatter in galaxy ellipticities, and the lack of correction for projection into radial velocities mean that there is a fair amount of scatter between an individual galaxy linewidth measurement and a true circular velocity or dynamical mass estimate.  For rotating galaxies, $\\sigoned \\sim 0.6 V_c$ in the mean, due to a combination of the inclination and the fact that $\\sigoned$ is less than line-of-sight $\\vrot$; this factor should be kept in mind when estimating dynamical masses.   The rotation and dispersion profiles of spatially resolved galaxies show that not all galaxies exhibit a conventional rotation curve.  Galaxies can be roughly divided into rotation and dispersion dominated, with $\\vrot>\\sigtwod$ or vice versa.  Dispersion dominated galaxies are blue, mostly irregular, and are not ellipticals; they probably have an effective dispersion, from disordered kinematics which are integrated over by the seeing.  Dispersion dominated galaxies exist at all redshifts in our sample, but low-$z$ DDGs are quite faint.  When line-of-sight rotation and dispersion are combined to make an estimate of the overall dispersion of the potential, as in $\\whalf^2 = 0.5\\vrot^2 + \\sigtwod^2$, the result correlates well with the integrated linewidth \\sigoned, demonstrating both that \\sigoned\\ is a robust velocity indicator and that it is possible to construct scaling relations with velocity for a population of diverse kinematic properties.  In Paper II we use the linewidths to measure evolution in the Tully-Fisher relation.  At $z \\sim 1$, rotation and dispersion dominated galaxies have similar magnitudes and colors, but somewhat different size and morphology distributions.  Rotation dominated galaxies are mostly disk and irregular types; dispersion dominated galaxies are in the mean smaller in physical half-light radius, and $\\sim 60\\%$ of DDGs are irregular and chain/hyphen types.  Only a small fraction of either RDGs or DDGs are obvious mergers.  About two-thirds of chain and hyphen galaxies are dispersion dominated: elongated high-redshift objects cannot be assumed to be inclined rotating disks.  It is not clear what the local counterparts of dispersion dominated galaxies might be; integral field spectroscopy, especially with adaptive optics, may shed some light on their nature at high redshift.  There is some evidence that dispersion dominated galaxies are rare at bright magnitudes at low redshift, just as previous studies have shown that irregular photometric morphologies in bright galaxies are more common at high redshift than locally (e.g. Abraham \\etal\\ 1996; Lotz \\etal\\ 2006).  It is possible that galaxies evolve in kinematic morphology, settling into ordered rotation, or that luminosity evolution causes them to appear more ordered, or to drop out of bright emission-line samples.  These and other possibilities are avenues for further study with larger samples and quantitative morphological measurements."
    },
    "0609/astro-ph0609745_arXiv.txt": {
        "abstract": "Blue straggler stars present as secure members in the Galactic open clusters form a major challenge to the conventional picture of evolutionary population synthesis based on the stellar evolution theory of single stars, as illustrated in our previous work. Expansion of our sample in the current work to include younger age clusters provides a larger data base to expose the question raised for the simple stellar population model. The working sample now includes 97 Galactic open clusters of ages ranging from 0.1 to several Gyrs. The contributions of blue straggler stars to the integrated light of the host clusters are calculated on an individual cluster base. A data base of observational constrained simple stellar population model is made which has a larger age coverage than our previous work. It is shown in this work that the general existence of blue stragglers in star clusters of our sample dramatically altered the predictions of convectional stellar population model in terms of spectral energy distribution. The integrated spectral energy distributions of the synthetic spectra of the clusters are enhanced towards shorter wavelengths, therefore the results of the present work will cast new lights in understanding the properties of stellar populations. ",
        "introduction": "Although the evolutionary population synthesis (EPS) method has been widely applied in analyzing stellar contents of local and remote galaxies, and has been proved to be very successful in many aspects, there are still some fundamental problems indeed in the framework of EPS. One of the challenges to the EPS is that the effect of complex stellar interactions which are common in all stellar systems, such as mass exchanging and coalescence in binaries and collisions between single and/or binaries, etc., is not considered in the method. The conventional picture of EPS is mainly based on the single-star evolution theory and is performed with the evolutionary behaviors and spectra libraries of individual stars (Bressan et al. 1994).  In order to attack such a problem, getting direct observational spectra of simple stellar populations (SSPs) had been practiced by former studies, such as the early work of Bica \\& Alloin (1986a, b). They undertook a direct approach using the observational integrated spectra of star clusters, assuming that these clusters are SSPs so that the EPS can be constrained observationally. Direct observation of star clusters is a meaningful approach to discuss the integrated properties of SSPs, and it is free of any assumptions about the initial mass function (IMF) and the details of stellar evolution. However, such a treatment can not perfectly reproduce the conventional SSP model, since the dynamical evolution of star clusters has modified the content of the original population, such as gravitational evaporation of low-mass member stars. More importantly, possible foreground star contaminations to the spectra cannot be assessed in an ideal manner. A practical way out should be resolving member stars photometrically using proper motion and/or radial velocity data. By conventional SSP model through this paper, we mean that the model is built based on single star evolution theory.  We have carried out a plan to attack this problem semi-empirically. The Galactic open clusters (OCs) are taken as our working sample to explore the contributions of blue straggler stars (BSs) to the conventional SSP model. The first set of results for old OCs with ages greater than 1Gyr has been published (Xin \\& Deng 2005, hereafter XD05). All the member stars within a cluster are presumably born at the same time from the same origin, therefore should have the same age and metallicity. Star clusters have been long regarded as one of the best objects to study the population synthesis technique (Battinelli et al. 1994). As described in XD05, the bulk of cluster member stars well fitted by an isochrone represents the idea of conventional SSP model. All the other member stars straggling away from the isochrone belong to the same population, no matter how weird their positions in the color-magnitude diagram (CMD) are, and have to be included when considering the integrated light properties of the population. Among the stragglers, BSs are of special interest because they are the only luminous objects straggling away from the predictions of the theory of single star evolution. It is also important to recall that these bright objects cannot be fully understood either by current theory of binary star evolution. Baring these in mind, the semi-empirical model used in XD05 is still a working approach. Besides, the nature of BSs is a challenging subject of stellar evolution, which is under active investigations (Tian et al. 2006, Hurley et al. 2005, Chen \\& Han 2004).  Since the first identification of BSs by Sandage (1953) in the globular cluster (GC) M3, the common existence of this kind of stars has been proved in stellar systems of all scales and complexities (e.g., Ahumada \\& Lappaset 1995 for open clusters; Piotto et al. 2002 for globular clusters and Lee et al. 2003 for dwarf galaxies). The typical locations of BSs in the CMD of a cluster are at the bluer and brighter extension of the turnoff point of the host cluster, therefore they are the most luminous and bluest hydrogen burning stars of the cluster at the time of observation (Deng et al. 1999). The main formation mechanisms of BSs are generally related with dynamical interactions of close binaries (Pols \\& Marinus 1994) and stellar collisions in the high-density areas (Ferraro et al. 2003). Therefore, statistical studies of BS populations in clusters in terms of integrated light provide a clue to such stellar interaction processes, and a reference to approaches such as binary evolution and dynamical evolution within the cluster, and eventually make the EPS more realistic. A larger sample is definitely needed for this purpose.  97 Galactic OCs are studied in this paper in terms of integrated light. A short description of the working sample is given in the next section. In \\S 3, how BSs influence the integrated properties of conventional SSP model is discussed. The synthetic clusters' ISEDs involving BSs contributions show significant enhancements towards shorter wavelengths, especially in ultraviolet (UV) and blue bands, and consequently the integrated (U-B) and (B-V) colors become bluer. When measured with the modifications in (U-B) and (B-V) colors, the physical parameters of OCs, especially age and metallicity, show great uncertainties. The uncertainties are considered in more details in \\S 4. The synthetic integrated spectral energy distributions (ISEDs) of our sample OCs are fitted with those of the conventional SSPs of either younger ages or lower metallicities. Based on the results of all the sample OCs, the uncertainties in measuring the fundamental parameters of SSPs are discussed. Finally, the concluding remarks on this work are presented in \\S 5. ",
        "conclusions": "As a follow-up work of XD05, the main aim of this work is to extend the study to larger sample and to quantify the influence of BSs on the conventional SSP model. With a sample of 97 Galactic OCs of ages $\\geq$ 0.1 Gyr, we confirm the results of XD05 that the integrated spectral properties OCs are dramatically modified by the BS contents in the clusters. The ISEDs are greatly enhanced towards shorter wavelength, becoming significantly bluer, and consequently the broadband (U-B) and (B-V) colors are modified accordingly. When the conventional SSP model is adopted in understanding stellar contents in galaxies, either using spectra or broadband colors, great uncertainties in age and/or metallicity will be brought in.  Directly fitting the synthetic ISEDs of our sample OCs with the ISEDs of conventional SSP model can still result in nice fittings, if one lets the parameter of age or metallicity free. But the fitting parameters have significant underestimations when compared with the true parameters made photometrically. In the parameter regions covered by our sample OCs, the logarithmic age value underestimation rates ranges 2\\% at $\\log (\\rm{age})$=8.5, 3\\% at $\\log (\\rm{age})$=9.0 and 4\\% at $\\log (\\rm{age})$=9.5. There is a trend for such a ratio to become larger for greater ages. Also due to the presence of BSs in real stellar populations, the metallicity parameter can also be seriously underestimated by conventional SSP models. This is true at least qualitatively within the scope of current work due to the uncertainties in measuring the metallicity independently by other means.  Taking OCs as general representatives of SSPs in galaxies, and limiting our results to the age and metallicity ranges covered by the working sample of current work, the conventional SSP model derived from the single star evolution theory is seriously altered by the descendants of stellar interactions that are presumably common in normal environments. BSs discussed in our work are the most important components that are proved to be widely existing in all stellar systems (Stryker 1993). Therefore, the effect of BSs must be taken into account when EPS technique is applied to study the properties of stellar populations in galaxies.  Theoretical refinery of the conventional SSP model still needs more studies. Thorough investigations of several issues should be put forward, that including: the individual processes of stellar interactions including binary and multiple systems and the overall consequences of all these interactions to the population statistically, and effects of dynamical evolution of the stellar systems. Before such SSP model eventually become available, emprically corrected models would be a good practice. Based no current knowledge of interactive binary evolution and collisional processes, and the observational cluster data base, such models can be made. Work to this purpose is now in progress."
    },
    "0609/astro-ph0609603_arXiv.txt": {
        "abstract": "{We present results from a kinematical study of the gas in the nucleus of the active S0 galaxy NGC 3998 obtained from archival HST/STIS long-slit spectra.} {We analyzed the emission lines profiles and derived the map of the gas velocity field.  The observed velocity curves are consistent with gas in regular rotation around the galaxy's center.} {By modeling the surface brightness distribution and rotation curve of the H$\\alpha$ emission line we found that the observed kinematics of the circumnuclear gas can be accurately reproduced by adding to the stellar mass component a compact dark mass (black hole) of $M_{\\rm BH} = 2.7_{-2.0}^{+2.4}\\times 10^8 M_{\\odot}$ (uncertainties at a 2 $\\sigma$ level); the radius of its sphere of influence ($R_{\\rm sph} \\sim$ 0\\farcs16) is well resolved at the HST resolution.} {The BH mass estimate in NGC 3998 is in good agreement with both the $M_{\\rm BH}-M_{\\rm bul}$ (with an upward scatter by a factor of $\\sim$ 2) and $M_{\\rm BH}-\\sigma$ correlations (with a downward scatter by a factor of $\\sim 3-7$, depending on the form adopted for the dependence of $M_{\\rm BH}$ on $\\sigma$).} {Although NGC~3998 cannot be considered as an outlier, its location with respect to the $M_{\\rm BH}-\\sigma$ relation conforms with the trend suggesting the presence of a connection between the {\\sl residuals} \\rm from the $M_{\\rm BH}-\\sigma$ correlation and the galaxy's effective radius. In fact, NGC 3998 has one of the smallest values of $R_{\\rm e}$ among the galaxies with measured $M_{\\rm BH}$ (0.85 kpc) and it shows a negative residual. This suggests that a combination of both $\\sigma$ and $R_e$ is necessary to drive the correlations between $M_{\\rm BH}$ and other bulge properties, an indication for the presence of a black holes ``fundamental plane''.} ",
        "introduction": "\\label{intro}  It is generally believed that supermassive black holes (SMBHs; $\\sim 10^{6}-10^{10}  M_{\\odot}$) are a common, if not universal, feature in the nuclei of nearby galaxies.  Since the discovery of quasars \\citep{schmidt63}, it has been suggested that active galactic nuclei (AGNs) are powered by mass accretion onto a SMBH \\citep{salpeter64, lynden69}.  This belief, combined with the observed evolution of the space density of AGNs and the high incidence of low-luminosity AGN-like activity in the nucleus of nearby galaxies \\citep{heckman80, maoz95, ho97a, ho97b, braatz97, barth98, barth99, nagar02} implies that a significant fraction of galaxies in the Local Universe must host black holes (BHs), relics of past activity \\citep{soltan82, chokshi92, yu02, marconi04, shankar04}.  Supports to these beliefs came from studies of the centers of nearby early-type galaxies, which revealed that most contain SMBHs \\citep{Kormendy01, Merritt01b} and that the BH mass $M_{\\rm BH}$ correlates with some properties of the host galaxy, such as bulge luminosity $L_{\\rm bul}$ and mass $M_{\\rm bul}$ \\citep{kormendy95, magorrian98, marconi03}, light concentration \\citep{graham01} and bulge velocity dispersion $\\sigma_{\\rm bul}$ \\citep{ferrarese00, gebhardt00a, tremaine02}. The correlation with the bulge velocity dispersion was thought to be the tightest having the smallest scatter: rms $\\sim$ 0.3 in log $M_{\\rm BH}$. \\citet{ferrarese00} argued that, for their selected sample of 12 galaxies, thought to have the most reliable BH mass estimates, the observed scatter in the $M_{\\rm BH}-\\sigma_{\\rm bul}$ relation was fully accounted for by the assumed measurement errors, which implies that there may be no intrinsic scatter in the correlation. More recently, \\citet{marconi03} have shown that when considering only galaxies with secure BH mass and bulge parameters determinations all the above correlations have a similar observed scatter \\citep[see also][]{mclure02, erwin04, haring04}.  The existence of any intrinsic correlations of $M_{\\rm BH}$ and host galaxy bulge properties supports the idea that the growth of SMBHs and the formation of bulges are closely linked \\citep{silk98, haehnelt00}, therefore having important implications for theories of galaxy formation and evolution. Moreover, SMBH mass estimates inferred via the above correlations, when more direct methods are unplayable, enter in a variety of important studies spanning from AGNs physics and fueling to coeval formation and evolution of the host galaxy and its nuclear black hole.  These correlations need, however, to be further investigated by increasing the number of accurate BH mass determinations in nearby galactic nuclei to set these correlations on stronger statistical basis. In particular such a study has the potential to establish the precise role of the various host galaxy's parameters in setting the resulting BH mass. To date, reliable SMBH detections have been obtained for a limited number of galaxies ($\\sim$ 30, see \\citealt{ferrarese05} for a review), with a bulk of $M_{\\rm BH}$ estimates in the range of $10^{7}-10^{9}  M_{\\odot}$. Furthermore, the coverage of the mass range is particularly poor at the low-mass end, the sample being strongly biased against late Hubble types and low luminosity objects. To add reliable new points to the $M_{\\rm BH}-$ host galaxy properties planes is then a fundamental task for future developments of astronomical and physical studies, and is the aim of the work described in this paper.  Spectral information at the highest possible angular resolution is required to directly measure the mass of SMBHs: the radius of the``sphere of influence'' \\citep{bahcall76} of massive BHs is typically \\ltsima 1\\arcsec\\ even in the nearest galaxies.  Among the so far most widely used techniques to detect and estimate BHs masses is stellar dynamical modeling (e.g. \\citealt{dressler88, kormendy95, marel98a, gebhardt00a, verolme02}), but the interpretation of the data is complex involving many degrees of freedom and requiring data of very high signal-to-noise ratio \\citep{valluri04}. Radio frequency measurements of H$_{2}$O masers in disks around BHs, finally, can be applied only to the small fraction of the disks inclined such that their maser emission is directed toward us \\citep{braatz97}. A more widely applicable and relatively simple method to detect BHs is based on gas kinematics (e.g. \\citealt{harms94,ferrarese96, macchetto97, barth01}), through studies of ordinary optical emission lines from circumnuclear gas disks, provided that the gas velocity field is not dominated by non gravitational motions. A successful modeling of the gas velocity field under the sole influence of the stellar and black hole potential is needed to provide {\\sl a posteriori} support for a purely gravitational kinematics. The angular resolution of $\\sim 0\\farcs1$ of the Space Telescope Imaging Spectrograph (STIS) on-board $\\it {HST}$ is the most suitable to perform such studies. The wealth of unpublished data contained in STIS archives represents an extraordinary and still unexplored resource.  Following this chain of reasoning, we performed a systematical search for unpublished data in the $\\it {HST}$ STIS archive with the aim of finding galaxies candidates to provide a successful SMBH mass measurement.  In this paper we present the results obtained for the active galaxy NGC 3998. From the Lyon/Meudon Extragalactic Database (HyperLeda), NGC 3998 is classified as an early type (S0) galaxy with an heliocentric radial velocity of 1040 $\\pm$ 18 km s$^{-1}$. With ${\\it H}_0$ = 75 km s$^{-1}$ Mpc$^{-1}$ and after correction for Local Group infall onto Virgo, this corresponds to a distance of 17 Mpc and a scale of 83 pc arcsec$^{-1}$. NGC 3998 is spectroscopically classified as a LINER with broad H$\\alpha$ emission \\citep{heckman80, keel83, ho97a, ho97b} and no significant broad-line polarization \\citep{barth99} therefore indicating a probable direct view to a BLR. Support to this hypothesis comes from studies of $\\it {HST}$ WFPC2 optical images which revealed an unobscured nucleus and the presence of a bright circumnuclear ionized gas disk \\citep{pogge00}. The nucleus contains a variable compact flat spectrum radio source \\citep{hummel84} displaying a weak jet-like northern structure \\citep{filho02}.  A bright UV source (unresolved at $\\it {HST}$ FOC resolution) is present in the center of NGC 3998 \\citep{fabbiano94}. Its rapidly variable flux \\citep{maoz05} implies that a significant fraction of the UV output is contributed by a non-stellar AGN component. A substantial featureless continuum component is also observed in the optical \\citep{gonzalez04}. The presence of an AGN in the nucleus of NGC 3998 is further strengthened by X-ray observations, which revealed the presence of a nuclear source with a power law spectrum \\citep{roberts00, pellegrini00} and X-ray and H$\\alpha$ luminosity values consistent with the low luminosities extension of the correlation observed for Seyferts and QSOs (see \\citealt{koratkar95}).  The paper is organized as follows: in Sec. \\ref{obs} we present HST/STIS data and the reduction that lead to the results described in Sec. \\ref{results}. In Sec. \\ref{fitting} we model the observed H$\\alpha$ emission line rotation curve and  we show that the dynamics of the circumnuclear gas can be accurately reproduced by circular motions in a disk when a point-like dark mass is added to the stellar potential. Our results are discussed in Sec. \\ref{discussion}, and summarized in Sec. \\ref{summary}. ",
        "conclusions": "\\label{summary}  We have presented results from a gas kinematics study in the nucleus of the nearby S0 active galaxy NGC 3998. The analyzed data were retrieved from archival HST/STIS long-slit spectra. We performed an analysis of the H$\\alpha$, [N II]$\\lambda\\lambda$6548,6583 and [SII]$\\lambda\\lambda$6716,6731 emission lines profiles to derive the map of the gas velocity field.  The nuclear velocity curves show a general reflection symmetry and are consistent with the presence of gas in regular rotation. We used our modeling code to fit the observed H$\\alpha$ surface brightness distribution and velocity curve. The dynamics of the rotating gas can be accurately reproduced by motions in a thin disk when a compact dark mass of $M_{\\rm BH} = 2.7_{-2.0}^{+2.4}\\times 10^8  M_{\\odot}$, very likely a supermassive black hole, is added to the stellar mass component. This result is also supported by the satisfactory good match of observed and model line width distributions.  Furthermore, the black hole in NGC 3998 has a sphere of influence radius of $\\sim$ 13 pc (0$\\farcs16$). At the high HST spatial resolution, this value of $R_{\\rm sph}$ implies a resolved BH sphere of influence for our $M_{\\rm BH}$ determination (2$R_{\\rm sph}/R_{\\rm res} \\simeq$ 3.2).  For what concerns the connections of this BH mass estimate with the properties of the host galaxy, the $M_{\\rm BH}$ value for NGC 3998 is in excellent agreement (within a factor of 2) with the $M_{\\rm BH}-M_{\\rm bul}$ correlation between BH and host bulge mass. The black hole mass predicted by the $M_{\\rm BH}-\\sigma_{\\rm bul}$ correlation is a factor of 2.7 larger than our measure, adopting the relation found by \\citet{tremaine02}, or a factor of 7 using the most recent parameterization by \\citet{ferrarese05}. However, NGC 3998 cannot be considered as an outlier from the $M_{\\rm BH}$ vs. $\\sigma$ correlation considering the errors in the estimates of $M_{\\rm BH}$ and $\\sigma_{\\rm star}$, and both the uncertainties in the determination of the $M_{\\rm BH}-\\sigma$ relation as well as its scatter.  Nonetheless, the lower-than-expected value for the mass of the black hole hosted by NGC 3998 strengthens the presence of a connection between the residuals from the $M_{\\rm BH}-\\sigma$ relation and the galaxy's effective radius. In fact, NGC 3998 has one of the smallest values of $R_{\\rm e}$ among galaxies with measured $M_{\\rm BH}$ and it shows a negative residual. We also recently showed that the opposite is true for the Seyfert galaxy NGC 5252: a large effective radius corresponds in this galaxy to a large positive residual.  Apparently only with a combination of both $\\sigma$ and $R_{\\rm e}$ it is possible to account for the correlations between $M_{\\rm BH}$ and other bulge properties, an indication for the presence of a black holes ``fundamental plane''. Clearly, only by further increasing the number of direct black hole measurements it will be possible to base these conclusions on a stronger statistical foundation."
    },
    "0609/astro-ph0609329_arXiv.txt": {
        "abstract": "The ALMA North American and European prototype antennas have been evaluated by a variety of measurement systems to quantify the major performance specifications. Nearfield holography was used to set the reflector surfaces to 17~$\\mu$m RMS. Pointing and fast switching performance was determined with an optical telescope and by millimeter wavelength radiometry, yielding $2^{\\prime\\prime}$ absolute and $0.6^{\\prime\\prime}$ offset pointing accuracies. Path length stability was measured to be $\\lesssim 20~\\mu$m over 10~minute time periods using optical measurement devices. Dynamical performance was studied with a set of accelerometers, providing data on wind induced tracking errors and structural deformation. Considering all measurements made during this evaluation, both prototype antennas meet the major ALMA antenna performance specifications. ",
        "introduction": "\\label{intro}  In the early stages of the collaboration between the National Radio Astronomy Observatory (NRAO) and the European Southern Observatory (ESO) toward the joint design, construction and operation of a large millimeter array (now called the Atacama Large Millimeter Array, or ``ALMA'', project), it was agreed that each of the partners would acquire from industry a prototype of the planned 12 m diameter antennas. The specifications of these antennas are beyond the current state of the art in accurate reflector antenna technology. Considering that 64 of these antennas would eventually be needed, this step appeared warranted to mitigate performance risk and increase competition among prospective manufacturers. It was also agreed to place the antennas next to each other at a suitable site and subject them to an extensive evaluation program to be carried out by a joint team of scientists and engineers drawn from the institutes of the ALMA collaboration.  The two ALMA prototype antennas are located at the site of the Very Large Array (VLA) in New Mexico. This site provides a reasonable compromise between the quality of the atmosphere, necessary for millimeter wavelength operation, and ease of construction and operation. The latter is provided by the existing infrastructure of the VLA site. The atmospheric characteristics however allow only operation at relatively long millimeter wavelengths during the dry winter months. Thus, the antennas have been equipped with evaluation receivers for the wavelength bands near 3 and 1.2 mm.  The evaluation of the antennas has been carried out by the Antenna Evaluation Group (AEG), consisting of experienced ``antenna integrators and commissioners'' from both organizations, Associated Universities Incorporated (AUI)/National Radio Astronomy Observatory (NRAO) and the European Southern Observatory (ESO). The charge of the AEG was to subject both antennas to a series of identical tests that would indicate the compliance (or not) of the antennas with the specifications. The core of the AEG is composed of the authors of this paper. ",
        "conclusions": "The overall design and performance of the ALMA prototype antennas makes either of them attractive options for production antennas that will satisfy the stringent ALMA requirements.   \\textit{Acknowledgements: The contributions of the following individuals were necessary for the success of the ALMA prototype antenna evaluation process: Marc Rafal (Commonwealth Technical Associates LLC); Fritz Stauffer, Nicholas Emerson, Jinquan Cheng, Jack Meadows (NRAO); Angel Ot\\'arola (ESO); Jos\\'e Lopez-Perez (OAN); David Smith (MERLAB); Michael Bremer (IRAM); and Henry Matthews (HIA).}"
    },
    "0609/astro-ph0609479.txt": {
        "abstract": "In this work, we study the formation and evolution of dark matter halos by means of the spherical infall model with shell-crossing. We present a framework to tackle this effect properly based on the numerical follow-up, with time, of that individual shell of matter that contains always the same fraction of mass with respect to the total mass. In this first step, we do not include angular momentum, velocity dispersion or triaxiality. Within this framework - named as the Spherical Shell Tracker (SST) - we investigate the dependence of the evolution of the halo with virial mass, with the adopted mass fraction of the shell, and for different cosmologies. We find that our results are very sensitive to a variation of the halo virial mass or the mass fraction of the shell that we consider. However, we obtain a negligible dependence on cosmology. Furthermore, we show that the effect of shell-crossing plays a crucial role in the way that the halo reaches the stabilization in radius and the virial equilibrium. We find that the values currently adopted in the literature for the actual density contrast at the moment of virialization, $\\delta_{vir}$, may not be accurate enough. In this context, we stress the problems related to the definition of a virial mass and a virial radius for the halo. The question of whether the results found here may be obtained by tracking the shells with an analytic approximation remains to be explored. ",
        "introduction": "In the hierarchical scenario for structure formation in the Universe, the small primordial density fluctuations grow %%@ due to non-linear gravitational evolution and finally become the first virialized structures (halos). In this %%@ picture, larger Cold Dark Matter (CDM) halos will be formed by the accretion and merger of those first smaller halos, %%@ forming in this way massive structures, and so on. This scenario, that constitutes the actual paradigm of %%@ hierarchical structure formation, is able to explain in general terms the universe that we see today. Yet, we do not %%@ have a framework or theory capable of reproducing this picture accurately. In this context, N-body cosmological %%@ simulations are a powerful tool to try to understand the formation and subsequent evolution of CDM halos. They %%@ constitute a very important help to build any theoretical model and their predictions explain many of different %%@ observations.  Basically, there are two analytical approaches that make the problem tractable, although some simplifications have to %%@ be made and, as it was said, comparison between these analytical studies and simulations are crucial to make %%@ progress: the Press-Schechter formalism (Press \\& Schechter 1974), based on the role of mergers (Nusser \\& Sheth %%@ 1999; Manrique et al. 2003), and the spherical infall model (SIM) focused on the understanding of the collapse %%@ of individual objects. We must note that in the Press \\& Schechter formalism, the \\emph{SIM} has also been widely %%@ used, but from a statistical point of view, to treat problems related to mass accretion histories, mass function, %%@ etc. The \\simm, first developed by Gunn \\& Gott (1972) and Gunn (1977), describes the collision-less collapse of a %%@ spherical perturbation in an expanding background. In those two articles, they introduced for the first time the %%@ cosmological expansion and the role of adiabatic invariance in the formation of individual objects. Fillmore \\& %%@ Goldreich (1984) and Bertschinger (1985) found analytical predictions for the density profiles of collapsed objects %%@ seeded by scale-free primordial perturbations in a flat universe. Hoffman \\& Shaham (1985) generalized these %%@ solutions to realistic initial profiles in flat and open Friedmann models, and Baarden et al. (1986) (hereafter BBKS) %%@ improved this work introducing the peak formalism. Later, some studies have been done to include more realistic %%@ dynamics of the growth process of dark matter halos (e.g. Padmanabhan 1996; Avila-Reese, Firmani \\& Hern\\'andez 1998; %%@ Lokas 2000; Subramanian, Cen \\& Ostriker 2000).  In parallel, a large amount of numerical work have been done. Quinn, Salmon \\& Zurek (1986) and Frenk et al.~(1988) obtained isothermal density profiles ($\\rho \\propto r^{-2}$) of CDM haloes, and Dubinski \\& Carlberg (1991) and Crone, Evrard \\& Richstone (1994) basically reproduced the predictions of Hoffman \\& Shaham (1985) and found some evidence for no pure power-law density profiles. Later, it was established that the density profiles of CDM halos have an universal form (Navarro, Frenk \\& White (1996, 1997, hereafter NFW), with $\\rho \\propto r^{-1}$ in the inner regions and $\\rho \\propto r^{-3}$ in the outskirts, although there is still controversy about the shape of the profile near the center and recently it has been found that the profiles flatten out close to $\\rho \\propto r^{-2}$ beyond $\\sim 2$ virial radius. Moore et al.~(1998,1999), amongst others, find $\\rho \\propto r^{-1.5}$ in the very center (although in a more recent work by Graham et al.~(2006) they find $\\rho \\propto r^{-0.7}$ at 0.1 kpc from the center) and other authors (Jing \\& Suto 2000; Klypin et al.~2001; Ricotti 2003) find an inner slope ranging from $-1$ to $-1.5$ depending on halo mass, merger history and substructure. Concerning the outskirts of dark matter halos, Prada et al.~(2006) carried out a detailed study and concluded that a 3D S\\'ersic three parameter approximation provides excellent density fits up to $\\sim 2$ times the virial radius, although these profiles differ considerably from the NFW ones beyond $2$ virial radius.  There are also plenty of works in the literature using the \\emph{SIM} to predict the density profiles of dark matter %%@ halos mainly focused on explaining their central regions. Moreover, the \\emph{SIM} has been widely used to obtain %%@ some quantities specially relevant and directly related to crucial stages in the formation and evolution of CDM %%@ haloes for different cosmologies, redshifts, etc. A particularly relevant quantity is the value of the overdensity at %%@ the moment of virialization $\\delta_{vir}$ ($\\Delta_{vir}$ usually in the literature), where overdensity is defined %%@ here as a number of times the background density, and its linear counterpart $\\delta_{l,vir}$. The values of %%@ $\\delta_{l,vir}$ and $\\delta_{vir}$ were obtained introducing the virial theorem into the \\emph{SIM} formalism. This %%@ has important implications in the way we define the virial radius (the radius that attains an overdensity %%@ $\\delta_{vir}$ inside) of dark matter halos in N-body cosmological simulations. $\\delta_{vir}$ is conventionally %%@ chosen to be near $180$ for an Einstein-deSitter cosmology (e.g. Peebles 1980), or $340$ for the %%@ $\\Omega_{\\Lambda}=0.7$ cosmology (e.g. Bryan \\& Norman 1998; Lokas \\& Hoffman 2000).  In the standard derivation of $\\delta_{l,vir}$ and $\\delta_{vir}$, the typical way to proceed is to assume that a %%@ shell of matter stabilizes at an epoch twice the time of turn-around (i.e. the time predicted by the standard %%@ \\emph{SIM} to collapse into a point), and in average with a radius that is $1/2$ the turn-around radius (e.g. Peebles %%@ 1980; Lacey \\& Cole 1994; Eke, Cole \\& Frenk 1996). This $1/2$ factor (in the Einstein-deSitter cosmology; for other %%@ cosmologies we need to use the Lahav equation, Lahav et al.~1991) is called the \\emph{collapse factor}. However, the %%@ justification to introduce this collapse factor and to suppose the time of virialization as twice the time of %%@ turn-around, is poor and lack a solid theoretical background. In contrast, in this work we will study the spherical %%@ collapse without supposing any collapse factor, only taking into account the shell-crossing as the dominant effect. %%@ The angular momentum and velocity dispersion may also play an important role. The question is: if these effects were %%@ included in the model, would we obtain the same values for $\\delta_{vir}$ and $\\delta_{l,vir}$ that those found in %%@ the most simplistic scenario described by the standard \\simm? This issue is one of the aims of the present work.  The main goal of this line of work is to develop a theoretical framework that help us to understand the dynamical %%@ elements that determine the process of formation of structures (collapsed objects) using spherical symmetry to %%@ explain main properties of dark matter halos. In this first work we will tackle these questions by means of a \"cold\" %%@ collapse, that is, without including the effects of the velocity dispersion and angular momentum. The point is to %%@ ascertain if the non-uniformity of the density profiles generated via shell-crossing is able to provide the radial %%@ motions necessary to produce the virialization and stabilization in an appropriate time scale. In a future work, we %%@ will include the angular momentum and velocity dispersion to go a step further.  There are some issues that it is worth mentioning and that make this work different from previous works that also %%@ included the shell-crossing in their formalism (e.g. Lokas \\& Hoffman 2000; Nusser 2001; Hiotelis 2002; Ascasibar et %%@ al.~2004). The way to proceed in these works is to handle the effect of shell-crossing by means of an adiabatic %%@ invariant, once the standard \\emph{SIM} becomes incorrect for late stages of the evolution. This adiabatic invariant, %%@ also known as \\emph{radial action}, makes the problem analytically tractable, and is based on the fact that the %%@ potential evolves in a time larger than the orbital period of the most inner shells. In contrast, we will study the %%@ shell-crossing effect doing a follow-up of the radius that contains inside always the same fraction of the virial %%@ mass. This way to tackle the problem is only one of the possible options, but is essential, for example, in order to %%@ build and study the relationship between the actual enclosed density contrast $\\delta$, defined as $\\delta = %%@ \\frac{\\rho(<r)-<\\rho_m>}{<\\rho_m>}$, with $<\\rho_m>$ the mean matter density of the Universe, and the linear density %%@ contrast, $\\delta_l$, obtained from the linear theory. Only Gehard Lemson did something similar, although using %%@ N-body simulations and mainly focused on showing how accurate are the predictions of the standard \\emph{SIM} compared %%@ to his simulations (Lemson 1995). Despite the fact that he showed that the \\emph{SIM} is a powerful tool to %%@ understand the evolution of halos, he never provided detailed quantities and relations for the actual and linear %%@ density contrasts. The function $\\delta_l(\\delta)$  is very important to obtain the density profiles of dark matter %%@ halos, as we discussed in previous works (Prada et al.~2006; Betancort-Rijo et al.~2006). Sheth \\& Tormen (2002) %%@ parametrized this function for the standard \\emph{SIM}. The framework presented here will allow us in the near future %%@ to provide also a simple parametrization for $\\delta_l(\\delta)$, but taking into account the important effect of %%@ shell-crossing, together with others relevant effects such as the angular momentum and velocity dispersion. This will %%@ lead us, for example, to obtain $\\delta_{l,vir}$ and its corresponding $\\delta_{vir}$, to explain the shape of the %%@ dark matter density profiles or to shed light on the mass functions. All of this without supposing any collapse %%@ factor, as pointed before, or other vague assumptions. Nevertheless, it will not be possible to obtain useful %%@ applications for the moment, since in this first work we will include in our study only the shell-crossing, which is %%@ the dominant effect. The full treatment will be done and presented in an upcoming work. Here we will provide the %%@ first results of our theoretical framework related to the role played by the shell-crossing.  The paper is organized as follows. In section \\ref{Sec2} we briefly describe the \\emph{SIM} and explain the formalism %%@ and unities that we will use in the rest of the work. In section \\ref{Sec3} we will study in detail the dependence of %%@ the way that the evolution occurs varying some parameters, in particular, the virial mass of the halo, the fraction %%@ of mass for a given virial mass, and the cosmology. Section \\ref{Sec4} will be specially dedicated to the moments of %%@ virialization and stabilization according to a given criterion, and their dependence with the same parameters %%@ described above. We will also emphasize here the difference between these two concepts. Finally, in section %%@ \\ref{Sec5}, we address the main results and ideas of the work, and point the lines for a future work. ",
        "conclusions": "\\label{Sec5}  In this work we have studied the effect of shell-crossing in the formation and subsequent evolution of dark matter %%@ halos. To do that, we have used the spherical collapse model, which has been widely used in the literature for more %%@ than thirty years to manage and solve questions related to these processes. Despite of the large amount of works that %%@ have used this model or many others that have improved it by introducing in the formalism more and more complex %%@ considerations, only a few of them have included the effect of shell-crossing. Moreover, most of these works have %%@ managed this effect analytically using the adiabatic invariant as a good approximation.  Here we handle the effect of shell-crossing numerically. This allows us to study individually any shell of matter %%@ involved in the process of formation of the halo. Doing so, we can extract multiple conclusions about the way in %%@ which this process occurs, like the relation between the linear and actual density contrasts, the process of %%@ stabilization of a shell of matter, the virialization, etc. Most of these issues have been treated in the present %%@ work to a greater or lesser extent, although the main goal have been always the developing of an adequate framework %%@ - named as Spherical Shell Tracker (SST) - in which we can study in depth the shell crossing and also other secondary effects.  It is possible to summarize the main conclusions of this work as follows:  \\begin{enumerate} \\item The \\emph{SST} framework is adequate to tackle the effect of shell-crossing in a way that allow us to extract %%@ exact results for different issues related to the evolution of the halo: the way that the radius of a given shell %%@ evolves with time, the relation between the linear and actual density contrasts, the stabilization, the %%@ virialization, etc. \\item The shell-crossing by itself is able to produce stabilization and virialization. Nevertheless, for the moment, %%@ it is not possible to obtain the exact values of the linear and actual density contrasts related to both moments of %%@ evolution. It is necessary to take into account also other important effects, such as angular momentum, velocity dispersions or triaxiality. \\item Concerning the relation between the linear and actual density contrasts, the dependence of this relation %%@ with the cosmology is very small and practically negligible. This conclusion is contrary to most of previous works, %%@ which find in general a large dependence with cosmology. However, the dependence with the virial mass or the fraction %%@ of virial mass that we consider, is large. \\item Neither stabilization nor virialization are reached in a time according to that given by the common assumptions %%@ related to the collapse factor and the time of virialization. In all the cases, we find that both stabilization and %%@ virialization occur at later times. \\item The values typically used in the literature for $\\delta_{l,vir}$ and $\\delta_{vir}$ seem to be clearly %%@ inadequate and incorrect, and are based on not very solid assumptions. In this work, new values of $\\delta_{l,vir}$ %%@ and $\\delta_{vir}$ are presented, but only taking into account the effect of shell-crossing. It will be necessary to %%@ include in our framework other effects also relevant to be able to provide useful and final values for %%@ $\\delta_{l,vir}$ and $\\delta_{vir}$. \\end{enumerate}  It is worth to emphasize that this work constitutes only a first step in our attempt to obtain exact and precise predictions related to the formation and evolution of dark matter halos. In a future work we plan to include in the \\emph{SST} framework other important effects that it will be absolutely necessary. In particular, including the angular momentum and velocity dispersion will be the next step. Furthermore, in parallel, we will implement a more sophisticated initial density profile than that used in this work, which fits better that found in the simulations and could change the results presented here only slightly. It is also in our mind to use cosmological N-body simulations, since comparison between both analytical and simulation studies will be, by sure, crucial to reach a better and deeper understanding of the processes involved in the formation and evolution of dark matter halos."
    },
    "0609/astro-ph0609165_arXiv.txt": {
        "abstract": "{} {The so-called Limber equation is widely used in the literature to relate the projected angular clustering of galaxies to the spatial clustering of galaxies in an approximate way. This note gives estimates of where the regime of applicability of Limber's equation stops.} {This paper revisits Limber's equation, summarises its underlying assumptions, and compares its predictions to the accurate integral of which the Limber equation is an approximation for some realistic situations. Thereby, the accuracy of the approximation is quantified. Different cases of spatial correlation functions are considered.} {Limber's equation is accurate for small galaxy separations but breaks down beyond a certain separation that depends mainly on the ratio \\mbox{$\\sigma/r_{\\rm m}$} and to some degree on the power-law index of spatial clustering; $\\sigma$ is the $1\\sigma$-width of the galaxy distribution in \\emph{comoving distance}, and $r_{\\rm m}$ the mean comoving distance. Limber's equation introduces a $10\\%$ systematic error that may be important at a scale of about a few degrees or in some cases, narrow galaxy distributions for example, even at arcmin scales. As rule-of-thumb, a $10\\%$ relative error is reached at \\mbox{$\\sim260\\,\\sigma/r_{\\rm m}$ arcmin} for $\\gamma\\sim1.6$, if the spatial clustering is a power-law \\mbox{$\\xi\\propto r^{-\\gamma}$}. More realistic $\\xi$ are discussed in the paper.} {Limber's equation becomes increasingly inaccurate for larger angular separations. Ignoring this effect and blindly applying Limber's equation can possibly bias results for the inferred spatial correlation.  It is suggested to use in cases of doubt, or maybe even in general, the exact equation that can easily be integrated numerically in the form given in the paper. This would actually eliminate all problems discussed in this paper.} ",
        "introduction": "If the distance of individual galaxies is not known, but only statistically for a sample of galaxies on the whole, one has to infer the real-space correlation from the observed angular correlation on the sky. This task has been performed many times since the early days of galaxy surveys as, for example, in the pioneering work of \\citet{tot69}, \\citet{groth77} and \\citet{davis77} based on the Lick galaxy survey. To relate the projected correlation function, $\\omega(\\theta)$, to the spatial, three-dimensional correlation function, $\\xi$, many authors use an approximation that was introduced by \\citet{limber53}. This approximation is nowadays referred to as Limber's equation.  It will be shown in Sect. \\ref{limberapprox} that Limber's approximation diverges when galaxies are closely distributed. That Limber's equation breaks down for narrow distributions is no surprise, as it is assumed that the distributions do no vary considerably over the coherence length of the spatial correlation function. But how wide has a distribution to be?  How much do we bias our estimate for the angular correlation by using Limber's equation, and is there an alternative?  These questions are addressed by this paper.  \\citet[ p.199]{peebles80} argued that this break-down occurs at separations of roughly $30^\\circ$, but made assumptions about galaxy surveys that do not meet the standards of many contemporary wide-field surveys. Due to the availability of spectroscopic or photometric redshifts quite narrowly binned distributions are often used in current studies.  This paper delineates the circumstances in which Limber's approximation is valid and gives an alternative for situations where it is not valid. Throughout this paper a flat universe will be assumed.  The exact equation for $\\omega$ is given in Sect. \\ref{exactsolution}. The thereafter following Sect. \\ref{limberapprox} summarises Limber's approximation and investigates why it breaks down for peaked galaxy distributions or beyond a certain angular separation. Alternative approximations for cases of narrow distributions of galaxies or for larger $\\theta$ are given inside the Sect.  \\ref{thinlayer}.  By looking at different distributions with varying widths and means, the accuracy of the Limber approximation is assessed in Sect.  \\ref{accuracy}. The paper finishes with conclusions in Sect.  \\ref{conclusions}. ",
        "conclusions": "The Limber equation is an approximate solution for the angular correlation function of objects, mainly applied in the context of galaxies, for a given spatial correlation function.  In this paper, it was shown that the assumptions made for Limber's approximation become inaccurate \\changed{beyond a certain angular separation, which essentially depends on \\mbox{$\\sigma/r_{\\rm m}$} and the shape of the spatial correlation function}. It should be emphasised that this inaccuracy has no relation to the flat-sky approximation (small separations) that is commonly adopted when working with Limber's equation.  Beyond what separation $\\theta$ does Limber's approximation become wrong by $10\\%$?  For power-law $\\xi$'s this separation depends entirely on $\\gamma$ and $\\sigma/r_{\\rm m}$ of the filter (selection function) used. Spatial correlations with larger $\\gamma$ are treated more accurately by Limber's equation than $\\xi$ with shallower $\\gamma$. As a rule-of-thumb, a systematic error of $10\\%$ is reached at about \\mbox{$260\\,\\frac{\\sigma}{r_{\\rm m}}~\\rm arcmin$} for $\\gamma\\sim1.6$.  For more realistic $\\xi$'s with a running slope $\\gamma$, the separation is also dependent on $r_{\\rm m}$, the centre of the filter. Even for relatively wide filters with \\mbox{$\\sigma/r_{\\rm m}\\sim0.2$}, Limber's equation becomes inaccurate on the $10\\%$ level at a scale of a few degrees separation.  On the other hand, for narrow filters of about \\mbox{$\\sigma/r_{\\rm m}\\lesssim0.02$}, a systematic error of $10\\%$ is already reached after $\\sim10^\\prime$ (or even a few arcmin depending on $r_{\\rm m}$).  \\begin{figure} \\begin{center} \\psfig{file=fig4.eps,width=85mm,angle=0} \\end{center} \\caption{\\label{fig2} \\changed{Filter (galaxy distributions) $p(r)$ appear ``wider'' in redshift, \\mbox{$p_{\\rm z}(z)=p(r(z))\\frac{\\d r}{\\d z}$}, than in comoving distance, i.e. usually \\mbox{$(\\theta_{\\rm z}\\equiv\\sigma_{\\rm z}/\\bar{z})\\ge(\\theta_{\\rm r}\\equiv\\sigma_{\\rm r}/r_{\\rm m})$} where $\\sigma_{\\rm z}$ and $\\sigma_{\\rm r}$ are the $1\\sigma$-widths in redshift, and comoving distance, respectively; $r_{\\rm m}$ is the mean of the weight in comoving distance, while $\\bar{z}$ is the mean in redshift. To give an estimate of the correction $\\theta_{\\rm r}/\\theta_{\\rm z}$ ($y$-axis) as function of $\\theta_{\\rm z}$ ($x$-axis) $p_{\\rm z}$'s are assumed that are close to Gaussians. Each solid line corresponds to different means $\\bar{z}\\in[0.05,7]$ (upper to lower line).}} \\end{figure}  \\begin{table} \\center \\begin{tabular}{lll} galaxy sample&$\\bar{z}\\pm\\sigma_z$&$\\sigma/r_{\\rm m}$\\\\ \\hline\\\\ \\textbf{\\cite{2007ApJ...654..138Q}}\\\\ ~$K$-dropouts&$z=2\\ldots3.5$&$\\sim0.16$\\\\ \\textbf{\\cite{lee06}}\\\\ ~$U$-dropouts&$3.2\\pm0.3$&$\\sim0.04$\\\\ ~$B_{435}$-dropouts&$3.8\\pm0.3$&$\\sim0.03$\\\\ ~$V_{606}$-dropouts&$4.0\\pm0.3$&$\\sim0.02$\\\\ \\textbf{\\cite{2006ApJ...648L...5O}}\\\\ ~$i_{775}$-dropout&$6.0\\pm0.5$&$\\sim0.03$\\\\ \\textbf{\\cite{adel05}}\\\\ ~BM&$1.69\\pm0.30$&$\\sim0.12$\\\\ ~BX&$2.24\\pm0.37$&$\\sim0.09$\\\\ ~LBGs&$2.94\\pm0.30$&$\\sim0.05$\\\\ \\textbf{\\cite{hildebrandt05}}\\\\ ~LBGs&$2.96\\pm0.24$&$\\sim0.04$\\\\ \\textbf{\\cite{ouchi04}}\\\\ ~$BRi\\rm-LBGs$&$4.0\\pm0.5$&$\\sim0.05$\\\\ ~$Viz\\rm-LBGs$&$4.7\\pm0.5$&$\\sim0.04$\\\\ ~$Riz\\rm-LBGs$&$4.9\\pm0.3$&$\\sim0.02$\\\\ ~LAE&$4.86\\pm0.03$&$\\sim0.003$ \\end{tabular} \\caption{\\label{examples}\\changed{List of typical values for the relative width of the survey selection function, \\mbox{$\\sigma/r_{\\rm m}$} (comoving distance), in recent works using the dropout technique (LBGs: Lyman-break galaxies, LAE: Ly$\\alpha$-emitter). Note that typically only separations less than about one arcmin are considered so that the accuracy of Limber's equation is, apart from the only LAE study in the list, no real issue.}} \\end{table}  Note that a filter can appear quite wide, i.e. \\mbox{$\\sigma_{\\rm z}/\\bar{z}$} is large, as a function of redshift $z$ while it is actually quite narrow as a function of comoving distance, the latter being the relevant parameter for the accuracy of Limber's equation. In particular, this means that applying Limber's equation to high redshift galaxies may be a problem. The reader might find Figure \\ref{fig2} useful to quickly look up the approximate conversion between $\\sigma_{\\rm z}/\\bar{z}$ and $\\sigma_{\\rm r}/r_{\\rm m}$.  Table \\ref{examples} shows a small list of some recent studies of the high-redshift universe using the dropout-technique. According to the table the relative width of the selection filter frequently reaches values near $0.03$. Therefore, Limber's equation should in these and similar studies not be used when separations larger than $\\sim20^\\prime$ are considered.  In cases where Limber's approximation is in doubt, that one uses the exact Eq.  \\Ref{omega0} that can easily be integrated numerically.  Limber's approximation is also used in gravitational lensing to quantify correlations between the gravity-induced distortions of faint galaxy images \\citep[see e.g. ][]{bs01, kaiser98}. In this case, the filter function relevant for lensing is the lensing efficiency that, quite naturally, has a wide distribution; one can roughly estimate \\mbox{$\\sigma/r_{\\rm m}\\approx0.22$} if the source galaxies have a typical value of \\mbox{$z_{\\rm s}=1$}. The centre of this filter is at about \\mbox{$\\bar{z}\\approx0.5$}. These values will depend on the distribution of source galaxies in redshift and the fiducial cosmological model. Taking these values as a rough estimate of the relative lensing filter width and centre, and assuming that the dark matter clustering is somewhere between the clustering of red and blue galaxies, one can infer from Figure \\ref{fig4} that a two-point \\emph{auto-correlation} function of the convergence, based on Limber's equation, is accurate to about $10\\%$ for separations less than several degrees; beyond that, an alternative description should be used."
    },
    "0609/astro-ph0609679_arXiv.txt": {
        "abstract": "We suggest that the Orion A cloud is gravitationally collapsing on large scales, and is producing the Orion Nebula Cluster due to the focusing effects of gravity acting within a finite cloud geometry. In support of this suggestion, we show how an elliptical rotating sheet of gas with a modest density gradient along the major axis can collapse to produce a structure qualitatively resembling Orion A, with a fan-shaped structure at one end, ridges or filaments along the fan, and a narrow curved filament at the other end reminiscent of the famous integral-shaped filament. The model produces a local concentration of mass within the narrow filament which in principle could form a dense cluster of stars like that of the Orion Nebula.  We suggest that global gravitational contraction might be a more common feature of molecular clouds than previously recognized, and that the formation of star clusters is a dynamic process resulting from the focusing effects of gravity acting upon the geometry of finite clouds. ",
        "introduction": "The dynamical state of molecular clouds clearly plays an important role in star formation.  Older concepts of long-lived clouds, maintained in rough equilibrium by balancing turbulence and magnetic forces against gravity, are being challenged by theoretical results suggesting rapid dissipation of magnetohydrodynamic turbulence which could lead to collapse (e.g. Mac Low \\etal 1998; Stone, Ostriker, \\& Gammie 1998; Ostriker, Gammie, \\& Stone 1999) as well as observational constraints from stellar populations which indicate short cloud lifetimes (Ballesteros-Paredes, Hartmann, \\& Vazquez-Semadeni 1999; Elmegreen 2000; Hartmann, Ballesteros-Paredes, \\& Bergin 2001; see Ballesteros-Paredes \\etal 2006 and Burkert 2006 for a recent reviews).  In addition, we recently pointed out the difficulties in setting up a velocity field which would globally stabilize a finite cloud containing many Jeans masses against its self-gravity (Burkert \\& Hartmann 2004; BH04); our results suggested that molecular clouds of appreciable mass are likely to be collapsing, at least in some regions.  These developments suggest that star-forming molecular clouds are relatively transient, rapidly evolving objects.  The dynamics of molecular clouds are also likely to be highly relevant to the formation of clusters, in which most newly-formed stars reside (e.g., Lada \\& Lada 2003). A number of dynamical simulations of star formation have been carried out (e.g., Klessen \\& Burkert 2000, 2001; Bate, Bonnell, \\& Bromm 2002, 2003), with some semi-analytic investigations (e.g., Krumholz, McKee, \\& Klein 2005), but the processes responsible for assembling the protocluster gas have been given much less attention; for example, in the numerical work of Bonnell, Bate, \\& Vine (2003), the protocluster aggregation of gas is simply assumed as an initial condition.  In BH04 we suggested that the assemblage of protocluster gas might be due to the focusing effects of gravity in finite, collapsing clouds.  The Orion A molecular cloud is a favorable site for studying the relationship between cloud dynamics and cluster formation, as it contains the Orion Nebula Cluster (ONC), the closest, relatively populous ($\\gtrsim 2000$ members) young cluster containing relatively massive (O, early B) stars.  In BH04 we presented an extremely simplified model of the gravitational collapse of an elliptical sheet which exhibited features in rough qualitative agreement with the structure of the Orion A cloud, forming a mass concentration near one end of the cloud which might represent the accumulation of gas for a large cluster like the ONC. In this paper we present a modified model of sheet collapse which exhibits a remarkable similarity to the structural features of the Orion A cloud. This picture makes predictions for the expected proper motions of the stellar members and suggests new perspectives both on the nature of \"turbulence\" in the region and on the conditions under which star clusters form. ",
        "conclusions": "We have developed a simple model of a gravitationally collapsing molecular cloud which accounts the overall morphology of the Orion A cloud and suggests the formation of a mass concentration in a location similar to that of the ONC.  Further kinematic studies, especially stellar proper motions, can be used to test this picture.  Our model implies that star cluster formation at least in some cases is a dynamic, complex process, with gravitational collapse playing an essential role in assembling the protocluster gas.  Future simulations in three dimensions including stars as well as gas will be needed to test new and improved kinematic measurements of the stellar population, providing additional insight into the formation of the ONC."
    },
    "0609/astro-ph0609353_arXiv.txt": {
        "abstract": "We have mosaiced 2MASS images  to derive surface brightness profiles  in \\irfilt\\ for 104 Galactic globular clusters.  We fit these with King profiles, and show that the core radii are identical to within the errors for each of these IR colors, and are identical to the core radii at $V$ in essentially all cases.  We derive integrated light colors $V-J, V-H, V-K_s, J-H$ and $J-K_s$ for these globular clusters.  Each color shows a reasonably tight relation between the dereddened colors and metallicity.  Fits to these are given for each color. The IR--IR colors have very small errors due largely to the all-sky photometric calibration of the 2MASS survey, while the $V-$IR colors have substantially larger uncertainties. We find fairly good agreement with measurements of integrated light colors for a smaller sample of Galactic globular clusters by Aaronson, Malkan \\& Kleinmann from 1977. Our results provide a calibration for the integrated light of distant single burst old stellar populations from very low to Solar metallicities.  A comparison of our dereddened measured colors with predictions from several models of the integrated light of single burst old populations shows good agreement in the low metallicity domain for $V-K_s$ colors, but an offset at a fixed [Fe/H] of $\\sim$0.1 mag in $J-K_s$, which we ascribe to photometric system transformation issues. Some of the models fail to reproduce the behavior of the integrated light colors of the Galactic globular clusters near Solar metallicity. ",
        "introduction": "Galactic globular clusters (GCs) are of great interest as nearby representatives of simple stellar systems all of whose stars share the same age and initial chemical composition\\footnote{We ignore the anomalous GC $\\omega$ Cen here.}. They are close enough that individual stars can be studied in detail with spectroscopy and photometry, while far enough away that, with some difficulty, their integrated light can be measured as well. Their ages and initial mass functions can be determined through analysis of deep high spatial resolution imaging, primarily from HST; mass segregation can be studied for these objects as well. Integrated light measurements of Galactic GCs are of key importance as they provide calibration data for the study of more distant early type galaxies with predominantly old populations as well as the GCs of distant galaxies for which only the integrated light can be observed.  The Two Micron All Sky Survey (2MASS) affords us a wonderful opportunity to study the surface brightness profiles of nearby GCs  in the infrared.  As described in detail by \\cite{2mass1}, two 1.3m diameter telescopes were used, one in the northern and one in the southern hemisphere. The whole sky was observed in three colors, $J$, $H$, and $K_s$ (a variant of the $K$ filter described in detail therein).  The advantages for our purposes of the 2MASS data over any previously existing are many.  The photometry is all-sky with careful attention to calibration issues, ensuring uniformity over all the frames.  The database is digital, and hence background subtraction and sophisticated image analyses are feasible. There are, of course, disadvantages as well. The effective exposure time for each point in the sky in this survey was short, only 7.8 sec in each of the three colors, so these images reach a relatively shallow limiting magnitude. Furthermore the spatial resolution is limited by the adopted detector pixel size of 2.0'' on a side. ",
        "conclusions": "We have mosaiced 2MASS images  to derive surface brightness profiles  in \\irfilt\\ for 104 Galactic globular clusters, incorporating algorithms to reduce the impact of bright field stars.  We fit these with empirical \\cite{king62} profiles, adopting tidal radii and cluster center positions from the literature.  This leaves only the central surface brightness and the core radius as parameters to be determined.  We show that the resulting core radii for each of these three IR colors are identical to within the errors.  We therefore set $r_c$ for each of \\irfilt\\ to be $r_c(J)$.  We then show that the $r_c(J)$ for each GC are identical to the core radii at $V$ in essentially all cases if the uncertainty for $r_c(V)$ is taken to be 2''.  The only discrepant cases are core collapsed GCs, where the lower spatial resolution of 2MASS combined with the small optical core radii produce smaller measured core radii at $V$ than in the IR from 2MASS.  We derive integrated light colors $V-J, V-H, V-K_s, J-H$ and $J-K_s$ for these globular clusters.  We do this by directly integrating the surface brightness profiles in most cases, equivalet to slightly censored aperture photometry, as this leads to smaller statistical measurement uncertainties than does integrating the fit King profiles. Each color shows a reasonably tight relation between the dereddened colors and metallicity.  Fits of these are given for each color. Linear fits suffice when the most metal-rich GCs are not considered.  Once the four GCs in our sample with [Fe/H]$ > -0.2$ dex are included, a quadratic fit is necessary.  We use our derived $V-K_s$ colors, combined with total $M_V$ from the database of H96, to find the luminosity function at $K_s$ of the Galactic globular cluster system.  The IR--IR colors have very small errors and very low dispersions about the fits due largely to the all-sky photometric calibration of the 2MASS survey.  These errors are consistent with the expected random fluctuations of the measurements based solely on the measured signal levels, indicating that other sources of error do not contribute much. The $V-$IR colors have substantially larger uncertainties due in part to the lack of an all-sky photometric calibration for surface brightness measurements at optical wavelengths. Incorrect choices for reddening, discrepancies in the adopted position of the center of a cluster, and, for the least populous GCs, stochastic errors due to the small number of luminous stars near the tip of the red giant branch, also contribute to the uncertainties in the $V-$IR colors.  We find good agreement with measurements of integrated light colors for a much smaller sample of Galactic globular clusters by Aaronson, Malkan \\& Kleinmann from 1977.  Small constant offsets between the two datasets of $\\sim$0.1 mag in IR--IR colors are required; we ascribe them to the difficulties of transforming between the filter and detector system used in 1977 and the 2MASS system. We find excellent agreement with the IR--IR colors of \\cite{nantais06}, not surprising since they too use 2MASS images from which to derive their colors.  But a comparison at $V-K_s$ of our colors with theirs shows very poor agreement in the mean, and with a very large dispersion; we suspect that they did not correctly match the optical and IR magnitudes in many cases.  Our results provide a calibration for the integrated light of distant single burst old stellar populations from very low to Solar metallicities.  We compare our dereddened measured colors with predictions from several models of the integrated light of single burst old populations, bearing in mind that the models have almost certainly been set up to reproduce the  data of \\cite{frogel83} and \\cite{cohen83} for colors of individual red giant branch stars in Galactic GCs.  While there is reasonable agreement for $V-K_s$ colors, a $\\sim$0.1 mag offset is required in $J-K_s$, with the model predictions being redder than our colors. Until the origin of this problem is understood, any determination of [Fe/H] (or age) in old populations based on IR--IR colors cannot be considered valid. In addition, some of the models fail to reproduce the behavior of the integrated light  $V-K_s$ colors of the Galactic globular clusters near Solar metallicity."
    },
    "0609/astro-ph0609812_arXiv.txt": {
        "abstract": "We consider the dynamical evolution of a disk of stars orbiting a central black hole.  In particular, we focus on the effect of the stellar mass function on the evolution of the disk, using both analytic arguments and numerical simulations.  We apply our model to the ring of massive stars at $\\simeq0.1$pc from the Galactic Center, assuming that the stars formed in a cold, circular disk, and find that our model requires the presence of a significant population of massive ($>100$M$_{\\odot}$) stars in order to explain the the observed eccentricities of 0.2--0.3.  Moreover, in order to limit the damping of the heavier stars' eccentricities, we also require fewer low-mass stars than expected from a Salpeter mass function, giving strong evidence for a significantly ``top-heavy'' mass function in the rings of stars seen near to the Galactic Center.  We also note that the maximum possible eccentricities attainable from circular initial conditions at ages of $<10$Myr are around 0.4--0.5, and suggest that any rings of stars found with higher eccentricities were probably not formed from circular disks. ",
        "introduction": "The relative proximity of the Galactic Center (henceforth GC) provides a unique opportunity for ``close-up'' study of processes that are expected to be crucial in the formation of galaxies and black holes, yet remain essentially unobservable in more distant galaxies.  Recent advances in telescope technology have enabled us to resolve individual stars in the crowded GC environment, and the development of adaptive optics has allowed determination of both velocities and positions of such stars to be made with ever-increasing accuracy \\citep[e.g.][]{ghez98,genzel03,ghez05,paumard06}.  These new data have presented several puzzles, notably the presence of a number of B-type stars very close ($\\lesssim 0.01$pc) to the GC, and also the detection of one, and possibly two, coherent rings of massive O- \\& B-type stars at somewhat larger radii of $\\simeq 0.1$pc \\citep{genzel03,paumard06}.  These stars are known to be young, and are therefore presumed to be the result of recent star formation at or close to the GC.  The environment at the GC is vastly different from the typical environment of ongoing star formation in the solar neighborhood, with much larger pressures, densities and temperatures, as well as strong tidal forces, so study of the formation of these stars raises a number of interesting issues \\citep[see also the recent review by][]{alexander05}.  A popular theory for the origin of these rings of stars is that they formed via fragmentation of accretion disks around the central black hole \\citep[e.g.][]{lb03,goodman03,gt04,nayak06}.  The stars are assumed to form on nearly circular orbits, as a result of gravitational instabilities in the disk.  \\citet{nayak06} suggests that this will lead to a stellar mass function that is significantly more top-heavy (i.e.~with significantly more massive stars) than that seen elsewhere in the Galaxy, and recent observational studies suggest that this is indeed the case \\citep{ns05,paumard06}.  Other suggestions for the formation of these stellar rings exist, however, most notably the infalling star cluster scenario \\citep[e.g.][]{gerhard01,mpz03,bh06}.  This scenario can also result in a significantly top-heavy mass function, due to mass segregation and stellar collisions within the cluster \\citep*[e.g.][]{bd98,gr05,fgr06}.  %  In this paper we consider the dynamical evolution of a ring of stars such as those observed around the GC.  Previous models of the stellar dynamics of such a ring have considered a single stellar mass population \\citep{nc05}, or the effect of processes such as resonant relaxation \\citep{ha06}.  Here we investigate the effects of the stellar mass function on the evolution of the system, and use observational determinations of the stellar orbital parameters to constrain the mass function of the stellar rings.  We find that the distribution of eccentricities can tell us about both the upper and lower ends of the stellar mass function, and discuss the consequences of this result for theories of star formation in the GC environment.  The structure of the paper is as follows.  In Section \\ref{sec:analytic} we present a simple analytic model for the dynamical evolution of a mass-segregated ring of stars orbiting a massive central black hole.  In Section \\ref{sec:numerical} we describe numerical simulations of such a system, and compare the results to the predictions of our analytic model.  We then apply our model to the GC system (Section \\ref{sec:gc_app}), and derive constraints on the initial conditions by comparing our model to recent observations.  We discuss the consequences of our results, and the limitations of our analysis, in Section \\ref{sec:dis}, and summarize our conclusions in Section \\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc} In this paper we have considered the dynamical evolution of rings of stars around a massive black hole.  Through analytic arguments and numerical simulations we have constructed a model for the evolution of a disk of stars of different masses, and shown that the stellar mass function is the dominant factor in determining the evolution of such a system.  We have then applied our analysis to rings of stars observed to orbit the  Galactic Center system.  We find, in agreement with previous studies, that the total mass in low-mass ($\\lesssim 5$M$_{\\odot}$) stars must be significantly lower than expected from a Salpeter mass function, and also find that a significant population of massive ($>100$M$_{\\odot}$) stars must have been present in order to produce eccentricities in the range 0.2--0.3, as observed by \\citet{paumard06}.  However, we find that dynamical relaxation alone is unlikely to produce rms eccentricities larger than $\\simeq 0.4$ in the GC system.  Consequently we conclude that rings with larger eccentricities, such as the counter-clockwise system observed by \\citet{paumard06}, are unlikely to have originated in a circular disk, and suggest that some other dynamical process must be responsible for such systems.  Alternatively, we suggest that star formation by fragmentation of an eccentric accretion disk could produce the observed eccentricities."
    },
    "0609/astro-ph0609023_arXiv.txt": {
        "abstract": "Recent simulations of solar active regions have shown that it is possible to reproduce both the total intensity and the general morphology of the high temperature emission observed at soft X-ray wavelengths using static heating models. There is ample observational evidence, however, that the solar corona is highly variable, indicating a significant role for dynamical processes in coronal heating. Because they are computationally demanding, full hydrodynamic simulations of solar active regions have not been considered previously. In this paper we make first application of an impulsive heating model to the simulation of an entire active region, AR8156 observed on 1998 February 16. We model this region by coupling potential field extrapolations to full solutions of the time-dependent hydrodynamic loop equations. To make the problem more tractable we begin with a static heating model that reproduces the emission observed in 4 different \\textit{Yohkoh} Soft X-Ray Telescope (SXT) filters and consider dynamical heating scenarios that yield time-averaged SXT intensities that are consistent with the static case. We find that it is possible to reproduce the total observed soft X-ray emission in all of the SXT filters with a dynamical heating model, indicating that nanoflare heating is consistent with the observational properties of the high temperature solar corona. ",
        "introduction": "Understanding how the Sun's corona is heated to high temperatures remains one of the most significant challenges in solar physics. Unfortunately, the complexity of the solar atmosphere, with its many disparate spatial and temporal scales, makes it impossible to represent with a single, all encompassing model. Instead we need to break the problem up into smaller, more manageable pieces (e.g., see the recent review by \\citealt{klimchuk2006}). For example, kinetic theory or generalized MHD is used to describe the microphysics of the energy release process. Ideal and resistive MHD are used to study the evolution of coronal magnetic fields and the conditions that give rise to energy release. Finally, one dimensional hydrodynamical modeling is employed to calculate the response of the solar atmosphere to the release of energy.  This last step is a critical one in the process of understanding coronal emission because it links theoretical models with solar observations. Even here, however, most previous work has focused on modeling small pieces of the Sun, such as individual loops \\citep[e.g.,][]{aschwanden2001b,reale2000}. Though understanding the heating in individual structures is an important first step, it has been difficult to apply this information to constrain the properties of the global coronal heating mechanism.  Recent advances in high performance computing have made it possible to simulate large regions of the corona, at least with static heating models. \\cite{schrijver2004}, for example, have coupled potential field source-surface models of the coronal magnetic field with parametric fits to solutions of the hydrostatic loop equations to calculate visualizations of the full Sun. Comparisons between the simulation results and full-disk solar images indicate that the energy flux ($F_H$) into a corona loop scales as $B_F/L$, where $B_F$ is the foot point field strength and $L$ is the loop length. \\cite{schrijver2005b} also find that this form for the heating flux is consistent with the flux-luminosity relationship derived from X-ray observations of other cool dwarf stars \\cite{schrijver2005b}.  \\cite{warren2006b} have performed similar simulations for 26 solar active regions using potential field extrapolations and full solutions to the hydrostatic loop equations. These simulation results indicate that the observed emission is consistent with a volumetric heating rate ($\\epsilon_S$) that scales as $\\bar{B}/L$, where $\\bar{B}$ is the field strength averaged along the field line. In the sample of active regions used in that study $\\bar{B}\\sim B_F/L$ so that $F_H\\sim \\epsilon_S L \\sim \\bar{B}\\sim B_F/L$, and this form for the volumetric heating rate is consistent with the energy flux determined by \\cite{schrijver2004}.  In these previous studies it was possible to use static heating models to reproduce the high temperature emission observed at soft X-ray wavelengths, but not the lower temperature emission typically observed in the EUV. The static models are not able to account for the EUV loops evident in the solar images. Recent work has shown that the active region loops observed at these lower temperatures are often evolving \\citep{urra2006,winebarger2003b}. Simulation results suggest that these loops can be understood using dynamical models where the loops are heated impulsively and are cooling \\citep[e.g.,][]{spadaro2003,warren2003}. Furthermore, spectrally resolved observations have indicated pervasive red shifts in active regions at upper transition region temperatures \\citep[e.g.,][]{winebarger2002}, suggesting that much of the plasma in solar active regions near 1\\,MK has been heated to higher temperatures and is cooling.  Finally, \\cite{warren2006b} found that static heating in loops with constant cross section yields footpoint emission that is much brighter than what is observed. This suggests that static heating models may not be consistent with the observations, even in the central cores of active regions.  The need for exploring dynamical heating models of the solar corona is clear, but there are a number of problems that make this difficult in practice. One problem is the many free parameters possible in parameterizations of impulsive heating models. In addition to the magnitude and spatial location of the heating, it is possible to vary the temporal envelope and repetition rate of the heating \\citep[e.g.,][]{patsourakos2006,testa2005}. Furthermore, dynamical solutions to the hydrodynamic loop equations are much more computationally intensive to calculate than static solutions, limiting our ability to explore parameter space.  In this paper we explore the application of impulsive heating models to the high temperature emission observed in active region 8156 on 1998 February 16. To make the problem more tractable we begin with a static heating model that reproduces the emission observed in 4 different \\textit{Yohkoh} Soft X-Ray Telescope (SXT) filters and look for dynamical models that yield time-averaged SXT intensities that are in agreement with those computed from the static solutions. Relating the time-averaged intensities derived from the full dynamical solutions to the observed intensities is based on the idea that the emission from a single feature results from the superposition of even finer, dynamical structures that are in various stages of heating and cooling. This idea is similar to the nanoflare model of coronal heating \\citep[e.g.,][]{parker1983,cargill1994}. Other nanoflare heating scenarios are possible, such as heating events on larger scale threads that are distributed randomly in space and time, but are not considered here. We find that it is possible to construct a dynamical heating model that reproduces the total soft X-ray emission in each SXT filter. This indicates that nanoflare heating is consistent with the observational properties of the high temperature corona. ",
        "conclusions": "We have investigated the use of static and dynamic heating models in the simulation of AR8156. Recent work has shown that static models can capture many of the observed properties of the high temperature soft X-ray emission from solar active regions and our results confirm this. We are able to reproduce both the total intensity and the general morphology of this active region with a static heating model. Furthermore, our results show that this agreement extends to the the hotter SXT filters which have not been considered before.  The application of dynamic heating models to active region emission on this scale has not been considered previously. Only the properties of individual loops have been explored \\citep[e.g.][]{warren2002b}. The computational complexity of the dynamical simulations precludes the calculation of individual solutions for each field line and we have utilized interpolation grids for estimating the expected SXT intensities for each field line. We find that it is possible to reproduce the observed SXT intensities in 4 filters, including the high temperature Al12 and Be119 filters, using the dynamical model.  Conceptually, the simple dynamical heating model investigated here, where we assume that the emission from a solar feature results from the superposition of many, very fine structures that are in various stages of heating and cooling, is closely related to the nanoflare model of coronal heating \\cite[e.g.,][]{parker1983,cargill1994}. The use of time-averaged intensities computed from the dynamical simulations implicitly assumes that the heating is very coherent, with each infinitesimal thread being heated once and then allowed to cool and drain before being heated again. Other scenarios are possible, such as heating events on larger scale threads that are distributed randomly in time. The spatial and temporal characteristics of coronal heating is likely to fall somewhere in between these extremes. The analysis of high spatial resolution (0.5\\arcsec) EUV images suggests that current solar instrumentation may be close to resolving individual threads in the corona \\citep{aschwanden2005a,aschwanden2005b}, but considerable work remains to be done to determine the fundamental spatial scale for coronal heating.  The geometrical properties of coronal threads is also unclear at present. We have assumed constant cross sections in our modeling, consistent with the observational results \\citep{klimchuk1992,watko2000}.  In the static modeling of solar active regions there has been some evidence that the loops with expanding cross sections better reproduce the observations \\citep{schrijver2004,warren2006b}. Detailed comparisons between simulated and observed solar images are needed to resolve this issue.  Despite the many limitations to our modeling the results that we have presented are encouraging and provide a framework for further exploration. The highest priority for future work is the full dynamical simulation of solar active regions without the use of interpolation grids so that synthetic soft X-ray and EUV images can be computed and compared with observations. Another priority is the comparison of active region simulations with spatially and spectrally resolved observations from the upcoming Solar-B mission. Spectral diagnostics, such as Doppler velocities and nonthermal widths, are another dimension that have not been explored in the context of this modeling."
    },
    "0609/astro-ph0609215_arXiv.txt": {
        "abstract": "We have observed the diffuse X-ray emission from the Galactic center (GC) using the X-ray Imaging Spectrometer (XIS) on Suzaku. The high-energy resolution and the low-background orbit provide excellent spectra of the GC diffuse X-rays (GCDX). The XIS found many emission lines in the GCDX near the energy of K-shell transitions of iron and nickel. The most pronounced features are Fe\\emissiontype{I}~K$\\alpha$ at 6.4~keV  and K-shell absorption edge at 7.1~keV, which are from neutral and/or low ionization states of iron, and the K-shell lines at 6.7~keV and 6.9~keV from He-like (Fe\\emissiontype{XXV}~K$\\alpha$) and hydrogenic (Fe\\emissiontype{XXVI}~Ly$\\alpha$) ions of iron. In addition, ~K$\\alpha$ lines from neutral or low ionization nickel (Ni\\emissiontype{I}~K$\\alpha$) and He-like nickel (Ni\\emissiontype{XXVII}~K$\\alpha$), and Fe\\emissiontype{I}~K$\\beta$, Fe\\emissiontype{XXV}~K$\\beta$, Fe\\emissiontype{XXVI}~Ly$\\beta$,  Fe\\emissiontype{XXV}~K$\\gamma$ and Fe\\emissiontype{XXVI}~Ly$\\gamma$ are detected for the first time. The line center energies  and widths of Fe\\emissiontype{XXV}~K$\\alpha$ and Fe\\emissiontype{XXVI}~Ly$\\alpha$ favor a collisional excitation (CE) plasma for the origin of the GCDX. The electron temperature determined from the line flux ratio of Fe\\emissiontype{XXV}-K$\\alpha$/ Fe\\emissiontype{XXV}-K$\\beta$  is similar to the ionization temperature determined from that of Fe\\emissiontype{XXV}-K$\\alpha$/Fe\\emissiontype{XXVI}-Ly$\\alpha$. Thus it would appear that the GCDX plasma is close to ionization equilibrium. The 6.7~keV flux and temperature distribution to the galactic longitude is smooth and monotonic, in contrast to the integrated point source flux distribution. These facts support the hypothesis that the GCDX is truly diffuse emission rather than the integration of the outputs of a large number of unresolved point sources. In addition, our results demonstrate that the chemical composition of Fe in the interstellar gas near the GC is constrained to be about 3.5 times solar. ",
        "introduction": "The iron K-shell lines near the Galactic center (GCDX) was first detected with the Ginga satellite \\citep{Ko89, Ya90}. The line energy is about 6.7~keV and the emission extends about 1--2$^\\circ$ over the Galactic center (GC) region. The ASCA satellite observed more details of the GC region, and confirmed the nature of the GCDX \\citep{Ko96,Ta00}. Moreover, the ASCA CCD resolved  the iron K-shell line detected with Ginga into three distinct lines at 6.4, 6.7 and 6.9~keV \\citep{Ko96}. The 6.7 and 6.9~keV lines are likely due to He-like (Fe\\emissiontype{XXV}) and hydrogenic (Fe\\emissiontype{XXVI}) ions of iron. There are three plausible scenarios that have been proposed for the origin of the iron K-shell lines. One is collisional excitation  (CE)  by low energy electrons (LEE), similar to the Galactic ridge emission. This process, however, contributes mainly to the 6.4 keV line emission, and not to the 6.7 and 6.9~keV lines (Valinia et al. 2000). With respect to the 6.7 keV and 6.9 keV lines, CE in a thin, high temperature plasma, and charge exchange (CX) recombination are both considered. For the CE case, the best-fit thin thermal plasma temperature of the GCDX spectrum was $\\sim$10~keV. If the GCDX is really due to a thin hot and uniformly distributed plasma, the total thermal energy is estimated to be 10$^{53-54}$~erg, or equivalent to 10$^2$--10$^3$ supernovae explosions. The temperature of $\\sim$10~keV is higher than that bounded by the Galactic gravity, hence the plasma should escape from the GC region. The time scale estimated by the plasma size (1$^\\circ$ or $\\sim$ 150~pc at 8.5 kpc) divided by the sound velocity of the $\\sim$10~keV plasma is $\\sim$10$^5$ years. Therefore the energy of $\\sim$10$^{53-54}$~erg should be supplied in the past $\\sim$10$^5$ years, e.g. 1 SN every 100--1000 year is required in the GC region. This huge energy budget in the GC region suggests that CX recombination may be a more reasonable explanation for the GCDX. In this scenario, during a collision between neutral hydrogen and Fe\\emissiontype{XXVI} or Fe\\emissiontype{XXVII},  the bound hydrogen electron is captured into an excited state of the iron, and then the excited state decays producing an X-ray. The X-ray spectral signature is distinct from the signature produced by CE  in a  collisional ionization equilibrium (CIE) plasma, such as the high-temperature plasma mentioned above. In the CX case, the forbidden line at $E_{\\rm f}=6636$~eV, is stronger than the resonance line at $E_{\\rm r}=6700$~eV \\citep{Ot06,Be05,Wa05}, whereas the resonance line is stronger in the case of a plasma in CIE. Thus, even if the resonance and forbidden lines are not completely resolved, the energy centroid of the K$\\alpha$ line of Fe\\emissiontype{XXV} produced by CX will be lower than produced by CE. We also note that an energy  shift towards lower energy as  the CX  would be observed in the case of X-ray emission following recombination with low energy electrons, as in a photo-ionized plasma \\citep{Be03}.  The 6.4~keV line is likely to be K$\\alpha$ from the neutral and/or a blend of  low charge states of iron (here, Fe\\emissiontype{I}). In fact, the 6.4~keV line is localized near molecular clouds such as the radio complex Sgr B2 and the Radio Arc regions. Sgr B2, the most massive giant molecular cloud with ultra compact H\\emissiontype{II} (UCH\\emissiontype{II}) and many maser sources, has been extensively studied with ASCA and Chandra (\\cite{Ko96};  \\cite{Mu00}, 2001). As a result of those studies, the authors proposed that the 6.4~keV line arises in an X-ray reflection nebula (XRN) irradiated by Sgr A$^*$ during an exceptionally bright period that occurred about 300 years ago, i.e., a time span equivalent to  the light traveling time between Sgr B2 and Sgr A$^*$. This putative high activity of Sgr A$^*$ could provide the energy supply to the GCDX.  Accordingly, a high quality observation of the iron and nickel K-shell lines can provide key information for understanding the origin of the GCDX.  Using the X-ray Imaging Spectrometer (XIS) (Koyama et al.~2006a), we have measured  X-ray emission for the spectral band containing these lines. The XIS is the X-ray CCD camera system on board Suzaku (\\cite{Mi06}). The excellent energy resolution and well-calibrated performance for diffuse sources, coupled with the low background orbit of Suzaku, make the XIS well-suited to provide crucial information about the GCDX. Here we present the results of the first measurement of the GCDX with the XIS focusing on the K-shell emission from iron and nickel. ",
        "conclusions": "%  \\subsection{Over Abundance of Fe in the Interstellar Gas}%  The continuum shape of the GCDX spectrum is well reproduced with the addition of an Fe K-edge at 7.1~keV. The best-fit column density, $N_{\\rm Fe}$, is determined to be 9.7$^{+0.7}_{-0.4}\\times 10^{18}$~cm$^{-2}$. This corresponds to an iron abundance of 3.5 times solar assuming a typical hydrogen column density ($N_{\\rm H}$) toward the GC of 6$\\times 10^{22}$~cm$^{-2}$ \\citep{Sa02} and the solar abundances of  Anders and  Grevesse (1993). We expect that a large fraction of this absorption takes place in the dense gas clouds prevailing near the GC region.  \\subsection{K$\\alpha$ Line from Fe\\emissiontype{I}}%  The energy of Fe\\emissiontype{I}~K$\\alpha$ is 6400~eV (Kaastra and Mewe 1993), which is lower than that found in the GCDX of 6409$\\pm$1~eV by 9~eV.  However, because the energy difference of 9 eV is larger than the systematic error of $^{+3}_{-6}$ eV, we suspect that this difference is a result of blending of Fe\\emissiontype{I}~K$\\alpha$ with line emission from low charge states of iron. The line width of 33$^{+2}_{-4}$~eV is comparable to, but slightly larger than,  the systematic broadening of $\\sim$30~eV, which may also be the result of the line blending.\\\\  \\subsection {K$\\alpha$ Line from Fe\\emissiontype{XXVI}}%  The line width of Fe\\emissiontype{XXVI}~K$\\alpha$ is determined to be $\\le$23~eV, which is smaller than the systematic broadening of$\\sim$30 eV . This unnaturally narrow line could be due to the coupling with the nearby line Fe\\emissiontype{I}~K$\\beta$. We therefore examine the coupling effect by fixing to the theoretically predicted flux of 0.125$\\times$  K$\\alpha$  (Kaastra and Mewe 1993). The best-fit  result of 25$^{+5}_{-8}$ eV is consistent with the  systematic broadening. Thus we can safely conclude that the line width broadening of Fe\\emissiontype{XXVI}~K$\\alpha$ is instrumental.  \\subsection{The Center Energy and Width of Fe\\emissiontype{XXV}~K$\\alpha$}  The Fe\\emissiontype{XXV}~K$\\alpha$ line is a blend of the resonance, inter-combination, and forbidden lines from Fe\\emissiontype{XXV} (see table \\ref{tbl:tripletline}), and depending on the conditions of the source plasma, may also contain satellite lines produced by dielectronic recombination and innershell excitation \\citep{Be92,Be93}. The fact that the observed line centroid and  width of Fe\\emissiontype{XXV}~K$\\alpha$ depends on the flux ratio of these lines, which, in turn, are a function of the  plasma conditions, make the K$\\alpha$ line feature an excellent diagnostic. As discussed above, a centroid shifted towards lower energy relative to the CIE case is an indication of a recombining  plasma. For the GCDX, the observed center energy of Fe\\emissiontype{XXV}~K$\\alpha$ is 6680$\\pm$1~eV. We compare our result to the laboratory values produced by the CX recombination, i.e., a recombination dominated plasma, and also to calculations using APEC and MEKAL of the line emission produced by a CIE plasma. The laboratory measured line centroid of Fe\\emissiontype{XXV}~K$\\alpha$ produced by the CX recombination alone, is $6666\\pm 5$ eV \\citep{Wa05}, significantly lower than the value measure in the GCDX. The APEC and MEKAL predicted centroid for the CIE case are 6685~eV and 6680~eV, respectively, in agreement with the GCDX value of 6680~eV.  We also note that the line broadening of 39$\\pm$2~eV for Fe\\emissiontype{XXV}~K$\\alpha$  is larger than the systematic broadening of $\\sim$30 eV (see section 3.2). This must be due to the fact that several lines contribute to this line feature.  \\begin{table}[!ht] \\begin{center} \\caption{Table of He-like Iron triplets} \\label{tbl:tripletline} \\begin{tabular}{lll} \\hline\\hline Name & Transition & Energy (eV)$^a$ \\\\ \\hline Resonance        &$1s^2$ $^1S_0$--$1s2p$ $^1P_1$&6701\\\\ Intercombination &$1s^2$ $^1S_0$--$1s2p$ $^3P_1$&6682\\\\ Intercombination &$1s^2$ $^1S_0$--$1s2p$ $^3P_2$&6668\\\\ Forbidden        &$1s^2$ $^1S_0$--$1s2s$ $^3S_1$&6637\\\\ \\hline \\\\ \\end{tabular}\\\\ $^a$ from Wargelin et~al. (2005).\\\\ \\end{center} \\end{table}  \\subsection{Constraints on the Cosmic Ray Iron Velocity in Charge Exchange Scenario} % The relatively high resolution of the XIS also affords us the ability to extract crucial information from the line feature of Fe\\emissiontype{XXVI}~Ly$\\alpha$ provides further evidence against the CX process for X-ray production.  As discussed in section 4.3,  the intrinsic line broadening  of Ly$\\alpha$ is nearly zero, and it follows that if CX is the source of the X-ray emission, then recombination would occur  when the bare iron  velocity is nearly zero, i.e., the collision energy between the iron ion and neutral hydrogen is small. In such a low-collision energy case, electrons are non-statistically captured into levels with large principal quantum numbers ($n\\geq10$) \\citep{Wa05,Pe01,Ot06}.  We, however, see no enhancement at the Fe\\emissiontype{XXVI}~Rydberg series limit at 9.2~keV. Our measurements gives an upper limit of $9\\times 10^{-6}$~photons~s$^{-1}$~cm$^{-2}$ for the high-$n$ Rydberg series line emission (we assumed a Gaussian line with width of 30~eV), only about 6\\% of the intensity of Ly$\\alpha$. This upper limit corresponds to a lower limit of collision energy of 100~eV amu$^{-1}$ (\\cite{Pe01};  \\cite{Wa05}) between bare Fe nuclei and neutral H. For that collision energy, the velocity of bare iron nuclei must be larger than $\\sim$150~km~s$^{-1}$. Therefore the CX process, if present in the GCDX, gives a lower limit on the velocity of $\\sim$150~km~s$^{-1}$. \\citet{Ta02} suggested that the CX process may have the maximum cross section at a higher iron velocity of $\\sim 5000$km s$^{-1}$, hence the subsequent decay lines have widths of about 100~eV. No such large broadening is observed in Fe\\emissiontype{XXVI}~Ly$\\alpha$.  Together with the discussion of section 4.4, we therefore propose that the 6.7~keV and 6.9~keV lines in the GCDX are likely due to a CIE plasma and would not arise from a CX process.  \\subsection{The Flux Ratio of K-shell Transition Lines from Fe\\emissiontype{XXV, XXVI} and Ni\\emissiontype{XXVII}}  The K-shell lines of Fe\\emissiontype{XXV}~K$\\alpha$ and Fe\\emissiontype{XXVI}~Ly$\\alpha$ were discovered with the data from ASCA~\\citep{Ko96}, Chandra~\\citep{Mu04}  and XMM-Newton~\\citep{Pr03}. With Suzaku, we have, for the first time, detected Ni\\emissiontype{XXVII}~K$\\alpha$, and Fe\\emissiontype{XXV}~K$\\beta$ and Fe\\emissiontype{XXVI}~Ly$\\beta$. We have also tentatively identified  Fe\\emissiontype{XXV}~K$\\gamma$  and  Fe\\emissiontype{XXVI}~Ly$\\gamma$ (see table 2). As discussed in sections 4.4 and 4.5, the emission lines from these highly ionized atoms are likely due to collisional excitation  in a thin thermal plasma. The best-fit flux ratio of Fe\\emissiontype{XXVI}-Ly$\\alpha$/Fe\\emissiontype{XXV}-K$\\alpha$ is 0.33$\\pm$0.02, which provides an ionization temperature of  6.4$\\pm$0.2~keV. The best-fit flux ratio of Fe\\emissiontype{XXV}-K$\\beta$/Fe\\emissiontype{XXV}-K$\\alpha$ is 0.092$^{+0.016}_{-0.012}$, which gives the electron temperature of 6.6$^{+4.2}_{-1.7}$~keV. Although the error is still large, this electron temperature is similar to the ionization temperature. The line flux ratio of Ni\\emissiontype{XXVII}-K$\\alpha$/Fe\\emissiontype{XXV}-K$\\alpha$ is 0.078. With the further assumption that the relative abundances of iron and nickel are proportional to the solar value, this ratio is converted  to the plasma temperature of $\\sim$5.4~keV. We therefore conclude that the GCDX plasma has a temperature of $\\sim$5--7 keV in collisional ionization equilibrium. The observed flux ratios of Fe\\emissiontype{XXVI}-Ly$\\beta$/Fe\\emissiontype{XXVI}-Ly$\\alpha$, Fe\\emissiontype{XXV}-K$\\gamma$/Fe\\emissiontype{XXV}-K$\\alpha$ fr and Fe\\emissiontype{XXVI}-Ly$\\gamma$/Fe\\emissiontype{XXVI}-Ly$\\alpha$ from are consistent with this plasma condition.  \\subsection{Temperature and Flux Variations of the GCDX Plasma}  Although we show that the GCDX  is attributable to a high temperature plasma of  $\\sim$ 5--7 keV in collisional ionization equilibrium, whether the origin is diffuse or the integrated  emission of a large number of unresolved point sources, is not yet clear.  We therefore investigated the spatial variation of the GCDX in temperature and flux. We made a series of X-ray spectra in $3\\times6$ arc-min rectangles along the Galactic plane of $b=-\\timeform{0D.05}$ (figure \\ref{fig:boxreg}), then extracted  the fluxes of the 6.7 and 6.9~keV lines, applying the same model as is given in table \\ref{tabl:lines}. The plots of the flux ratio between the 6.7 and 6.9~keV lines, and the fluxes of the 6.7 keV line are given in figure \\ref{fig:ratio} and figure \\ref{fig:ps_diff}, respectively, in which the flux and flux ratio of the bright SNR Sgr A East are excluded. We see that the 6.7 and 6.9~keV line ratio is constant at $\\sim0.38$ in the negative longitude region from $l= -\\timeform{0D.1}$ to $l=-\\timeform{0D.4}$, but decreases to $\\sim0.30$ at the positive longitude near at $l = \\timeform{0D.25}$. The ratios of 0.38 and 0.30 suggest ionization temperatures of 6.8~keV and 6.2~keV, respectively. The smooth and monotonic variation of the ionization temperature along the galactic longitude argues that the GCDX is not due to the superposition of the outputs of a large number of unresolved point sources.  We simulated the point-source flux in the 4.7--8.0~keV band observed with XIS, putting the positions and fluxes (2.0--8.0~keV band) of the  Chandra cataloged sources \\citep{Mu03} in the same rectangles of figure 4. To preserve completeness limit, we selected the sources with flux larger than 5$\\times 10^{-7}$ photons~cm$^{-2}$~s$^{-1}$ in the 2.0--8.0~keV band.  After convolving the response of the XRT + XIS, we extracted the integrated point-source fluxes from each rectangle. The results are plotted in the same figure of the 6.7 keV line (figure \\ref{fig:ps_diff}). The integrated point-source flux is symmetrically distributed between positive and negative galactic longitude with a sharp peak at Sgr A$^*$  ($l = -\\timeform{0D.06}$), which is significantly different from the 6.7 keV line flux distribution. The 6.7 keV line flux  shows an asymmetric distribution (larger flux at the positive galactic longitude than the negative side) and a gradual increase towards Sgr A$^*$. These facts suggest that the major fraction of the 6.7 keV line flux is not due to the integrated emission of unresolved point sources, but is truly diffuse.  \\begin{figure}[!ht] \\begin{center} \\FigureFile(70mm,50mm){figure4.eps} % \\caption{Image of the 2 fields of view (FOV) in the 6.62--6.74~keV band, overlaid with the rectangular regions from which the spectra were extracted.} \\label{fig:boxreg} \\end{center} \\end{figure}  \\begin{figure}[!ht] \\begin{center} \\FigureFile(70mm,50mm){figure5.eps} % \\caption{The line flux ratios of Fe\\emissiontype{XXVI}-Ly$\\alpha$/Fe\\emissiontype{XXV}-K$\\alpha$  of measured in the series of rectangular regions along the Galactic plane  given in figure 4 (the constant  $b=-\\timeform{0D.05}$ line).} \\label{fig:ratio} \\end{center} \\end{figure}  \\begin{figure}[!ht] \\begin{center} \\FigureFile(70mm,50mm){figure6.eps} % \\caption{The line fluxes of Fe{XXV} K$\\alpha$ are given by the squares, while the integrated point source fluxes in the 4.7--8~keV band are plotted by the crosses. The horizontal axis is the same as figure \\ref{fig:ratio}, but the vertical axis is a logarithmic scale.} \\label{fig:ps_diff} \\end{center} \\end{figure}  \\subsection{Hard X-ray Tail}  From the line diagnostics, the GCDX is most likely produced by a thin hot plasma with a temperature of $\\sim$5--7~keV in CIE, while the bremsstrahlung temperature determined with the continuum shape in the 5.5--11.5~keV band is $\\sim$15~keV(see table 2), significantly higher than $\\sim$5--7~keV. This suggests that the continuum flux may contain an additional hard component. We therefore fit the spectrum with a model of a CIE plasma and  power-law component with 3 Gaussians for the K-shell emissions from Fe\\emissiontype{I} and Ni\\emissiontype{I}. In this model, the K-edge absorption with $N_{\\rm Fe}$ is also implemented. Although its contribution is minor, we added the model spectrum of the cosmic X-ray background (CXB). As discussed in section 3.2, there are systematic errors in the line width ($\\sim$30 eV) and center energy ($^{+3}_{-6}$~eV). We therefore introduced an artificial random motion and small red-shift to cancel these systematic errors. The fit is reasonably good and the best-fit results and parameters are given in figure \\ref{fig:apecfit} and table \\ref{tabl:vapec}. We note that the forbidden line of Ni\\emissiontype{XVII} and some of the satellite lines are missing in the APEC model, hence the Ni abundance may be overestimated.  \\begin{figure}[!ht] \\begin{center} \\FigureFile(70mm,50mm){figure7.eps} % \\caption{Same as figure \\ref{fig:hardspec}, but the model is a CIE plus  a power-law with 3 Gaussian lines and an iron absorption edge.} \\label{fig:apecfit} \\end{center} \\end{figure}  \\begin{table*} \\begin{center} \\caption{Best-fit parameter with the VAPEC model} \\label{tabl:vapec} \\begin{tabular}{lcccc} \\\\ \\hline\\hline Parameter & Unit   &&             &  Value\\\\ \\hline $N_{\\rm H}$       & cm$^{-2}$ &&& 6$\\times 10^{22}$ (fixed)\\\\ $N_{\\rm Fe}$      & cm$^{-2}$ &&& 8.5$^{+0.8}_{-0.5}\\times 10^{18}$ \\\\ $kT_{\\rm e}$    & keV            &&     & 6.5$^{+0.1}_{-0.1}$\\\\ $Z_{\\rm Fe}$ & solar             &&  &1.2$^{+0.5}_{-0.3}$ \\\\ $Z_{\\rm Ni}$ & solar             &&  &2.1$^{+0.7}_{-0.5}$ \\\\ Velocity     & km~s$^{-1}$ &&&1192\\\\ Redshift    &                    && &6.0$\\times 10^{-4}$\\\\ \\hline \\multicolumn{5}{c}{Hard Tail}\\\\ \\hline $\\Gamma$ & & & & 1.4$^{+0.5}_{-0.7}$ \\\\ Flux at 8 ~keV & photons~s$^{-1}$~cm$^{-2}~$keV$^{-1}$ & & &4.5$^{+1.9}_{-2.1}\\times 10^{-4}$ \\\\ \\hline \\multicolumn{5}{c}{Emission lines}\\\\ \\hline Center Energy  &   \\multicolumn{2}{c}{Identification}& Width & Intensity \\\\ (eV)             &         Line    ~        &  Energy (eV) ~ & (eV)  &(photons~s$^{-1}$~cm$^{-2}$)\\\\ \\hline 6409$^{+1}_{-1}$  &  Fe\\emissiontype{I}$^\\dag$~K$\\alpha$& & 36$^{+1}_{-3}$ & 4.41$^{+0.05}_{-0.08}\\times 10^{-4}$\\\\ 7069.1$^*$  &  Fe\\emissiontype{I}$^\\dag$~K$\\beta$ & &39$^*$&7.8$^{+0.4}_{-0.6}\\times 10^{-5}$\\\\ 7492$^{+14}_{-18}$&  Ni\\emissiontype{I}$^\\dag$~K$\\alpha$&   & 0($<$30) &2.3$^{+0.7}_{-0.4}\\times 10^{-5}$\\\\ \\hline \\end{tabular} \\\\ $^\\dag$ or low ionization state\\\\ $^*$ fixed to 1.103~$\\times E$ (Fe\\emissiontype{XXV}~K$\\alpha$) .\\\\  \\end{center} \\vspace{-0.5cm} \\end{table*}  One notable finding is the presence of power-law hard tail. The power-law index $\\Gamma$ is sensitive to the slope of the high energy continuum. Because the effective area of the XRT-XIS combination drops rapidly at high energy, our determination of the power-law index $\\Gamma$ has a relatively large statistical error  of  $\\Gamma\\sim$1.4$^{+0.5}_{-0.7}$. The low numbers of counts make it challenging to determine the error contribution from systematic effects. To provide a more reliable estimate of the systematic error, we checked it using the spectrum of the Crab nebula. We examined the variation of the best-fit $\\Gamma$ by fitting the Crab spectrum in the 4--10, 5--10, 6--10 and 7--10~keV bands, and found no systematic variations with the increasing energy band. The mean variations of $\\Gamma$ are within 5\\% for all the XISs, which is far smaller than the statistical error of $\\sim$40\\%.  Since the flux level of the NXBG becomes comparable to the hard tail flux at 8--10 keV, it may be argued that the error due to the NXBG-subtraction should not be ignored. The possible flux error of the NXBG is estimated to be $\\sim\\pm$3\\% (Koyama et al. 2006a). Including  this error in the spectrum, the best-fit $\\Gamma$ are 1.4 (adding 3\\% of the NXBG) and 1.2 (3\\% reduction of the NXBG). Further systematic analysis of the NXBG found a hint that the flux  of the night Earth data is smaller than that from a blank sky.  The details of the flux difference is still in study, but is at a level of $\\leq$ 10\\% at the 8--11 keV band. If we use a NXBG of 1.1$\\times$ of the night Earth, then the  best-fit $\\Gamma$ becomes 1.5. Thus any possible systematic error of NXBG does not change the conclusion that a flat hard tail is present in the GCDX. In addition, the flux of the CXB is more than  one order of magnitude smaller than the hard tail in the hard X-ray band, so that the fluctuation and/or error of the CXB can be ignored.  A possible contributor of the power-law hard tail is non-thermal filament structures found with Chandra  (\\cite{Mo02}; \\cite{Se02}; \\cite{Pa04}). In fact, the power-law flux in the positive galactic longitude (GC\\_SR1) is significantly larger than that of the negative galactic longitude (GC\\_SR2), which is a trend similar to the population of the non-thermal filaments (\\cite{Mo02};  \\cite{Se02}; \\cite{Pa04}). The  origin of the non-thermal emission is a key for understanding the presence of high energy particles, which are suggested by the detection of very high energy Gamma-rays (\\cite{Ah06}; \\cite{Ts04}; \\cite{Ko04}). Further analysis of the XIS data in combination with those of the Hard X-ray Detector(HXD) (\\cite{Ta06}) could yield an important contribution to the study of this hard component."
    },
    "0609/astro-ph0609037_arXiv.txt": {
        "abstract": "We present the results of a {\\it Chandra} observation of the cooling core cluster Abell 2063. Spectral analysis shows that there is cool gas (2 keV) associated with the cluster core, which is more than a factor of 2 cooler than the outer cluster gas (4.1 keV). There also is spectral evidence for a weak cooling flow, $\\dot M \\approx 20 \\, \\msun$ yr$^{-1}$. The cluster exhibits a complex structure in the center that consists of several bright knots of emission, a depression in the emission to the north of the center of the cluster, and a shell of emission surrounding it. The depression in the X-ray emission is coincident with the position of the north-eastern radio lobe of the radio source associated with the cluster-central galaxy. The shell surrounding this region appears to be hotter, which may be the result of a shock that has been driven into the gas by the radio source. The power output of the radio source appears to be sufficient to offset the cooling flow, and heating of the gas through shocks is a possible explanation of how the energy transfer is established. ",
        "introduction": "\\label{sec:intro}  Clusters of galaxies contain very hot gas, the intracluster medium (ICM). Due to the extremely strong gravitational fields of clusters, plasma that falls towards the center will gain significant gravitational potential energy, which results in temperatures in the ICM of $\\sim$$10^{8}$ K. With time, energy is lost due to X-ray radiation (thermal bremsstrahlung and line emission from metals) from the ICM, which acts to cool the X-ray emitting gas unless these energy losses are balanced by some heating process. One would expect this cooling gas to flow toward the dense core of the cluster. The cooling process is known as a ``cooling flow.'' This process can potentially trigger star formation in the core. Evidence for such cooling flows have been observed by the {\\it Einstein}, {\\it ROSAT}, and {\\it ASCA} observatories \\citep[e.g.,][]{fabian94}. The predicted cooling rates from these observations are hundreds of solar masses per year. More recent observations made with the {\\it XMM-Newton} and {\\it Chandra} X-ray observatories suggest cooling rates that are at least an order of magnitude lower \\citep[e.g.,][]{peterson03}, with most of the gas only cooling down to about one third of the ambient cluster temperature.  One of the proposed explanations of this phenomenon and a solution to the discrepancies in measured cooling rates is heating of the ICM by a central radio source. These radio sources are thought to be the result of accretion events by active galactic nuclei (AGN) within the central cD galaxies that are generally present in the centers of cooling core clusters. X-ray observations show that the lobes of the central radio sources can displace the X-ray emitting gas and leave ``holes'' in the ICM, which are filled by the synchrotron plasma \\citep[e.g.,][]{fabian00}. Some of the theoretical models for these ``radio bubbles'' predicted that the ICM would be heated by strong shocks from the radio emitting plasma \\citep[e.g.,][]{heinz98}. However, in most cases the observed shells of higher surface brightness surrounding the holes have been found to have lower temperatures than the surrounding gas \\citep[e.g.,][]{blanton01}. More recently, a few cases of radio lobes bounded by strong shocks have been observed \\citep{mcnamara05,nulsen05a,nulsen05b}. In addition, weak shocks or sound waves have been detected around the radio lobes in Perseus \\citep{fabian03}. The energy output contained in the radio outbursts is found to be sufficient to counterbalance the cooling flow in many cases \\citep[e.g.,][]{birzan04}.  In this paper, we present an analysis of a {\\it Chandra} observation of the cooling core cluster Abell~2063. Abell~2063 is a rich cluster of galaxies at a redshift of $z = 0.0349$. The central cD galaxy in the cluster is host to the radio source [OL97]1520+087 \\citep{owen}. Abell~2063 has been previously observed in the X-ray with {\\it Einstein} \\citep{david93,jf84}, {\\it ROSAT} \\citep{david95,peres98}, and {\\it ASCA} \\citep{white00}. The derived temperature from the {\\it Einstein} and {\\it ROSAT} data is $kT = 4.1$ keV, whereas the {\\it ASCA} data gives a temperature of $kT = 3.9$ keV and an average chemical abundance of 0.24 times the solar value. The bolometric luminosity as computed from the {\\it Einstein} data is $L_{\\rm bol} = 1.48 \\times 10^{44}$ ergs s$^{-1}$. The cooling rate derived from these observations is $\\dot M=50^{+39}_{-47} \\, M_{\\odot}\\, {\\rm yr}^{-1}$ \\citep{white00} for the {\\it ASCA} data and $\\dot M = 19^{+4}_{-6} \\, M_{\\odot}\\, {\\rm yr}^{-1}$ \\citep{peres98} within a radius of $r_{\\rm cool} = 68^{+9}_{-21}$ kpc for the {\\it ROSAT} data. Throughout this paper we adopt WMAP cosmological parameters \\citep{wmap} of $H_0$ = 71 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\Lambda}=0.73$, and $\\Omega_M=0.27$. At the redshift of Abell~2063, this corresponds to a scale of 0.686 kpc/arcsec. The uncertainties quoted in this paper are 90\\% confidence intervals unless otherwise stated. ",
        "conclusions": "\\label{sec:concl}  The {\\it Chandra} observation and analysis of the core of the cooling flow cluster Abell 2063 has been presented above. The X-ray data reveals complex structure near the center of the cluster, which consists of a central bar of high surface brightness, a blob of emission to the south of the center of the cluster, and a depression in the emission to the north of the center of the cluster which is surrounded by a shell of higher surface brightness. The position of the hole in the X-ray emission is coincident with the northeastern radio lobe of the radio source associated with the central cD galaxy. The temperature map of the central region of the cluster and detailed spectral analysis of these features suggest that the northern shell of emission is hotter than the surrounding gas. This may be evidence that a shock has been driven into the gas by the radio source. The knot of emission to the south appeared hotter as well, but it is surrounded by a cooler region.  From the surface brightness profile of the cluster, we determined the total mass and the gas mass out to radius $r = 66$ kpc. We found $M_{\\rm tot} = ( 1.31 \\pm 0.06 ) \\times 10^{13} \\, \\msun$ and $M_{\\rm gas} = ( 1.69 \\pm 0.03 ) \\times 10^{11} \\, \\msun$.  A cooling flow model fitted to the spectrum of the cluster revealed that gas near the center of the cluster is cooling from 4.06 keV down to 1.17 keV at a mass deposition rate of $\\dot M = 20^{+29}_{-16} \\, \\msun$ yr$^{-1}$. The abundance was found to be 0.51 times the solar value. There was no strong indication of gas cooling to very low temperatures.  The estimated power output of the radio source in the center of Abell 2063 was found to be sufficient to offset the cooling flow as is the case with sources in other clusters. The minimum energy nonthermal pressure provided by the radio source was found to be orders of magnitude lower and insufficient to balance the pressure exerted by the gas from the bright shell surrounding the hole. There is some evidence that the radio source is expanding rapidly and driving a shock into the X-ray gas. This suggests that the radio source is relatively young, and has just started its current period of activity. Low frequency radio observations would be useful to search for more extended radio emission within the northern radio bubble, and for ghost bubbles due to previous episodes of AGN activity."
    },
    "0609/astro-ph0609806_arXiv.txt": {
        "abstract": "% The key project HERLOGA to be proposed for the Herschel mission is aimed to provide an homogeneous data set on local active galaxies that could be the reference for further studies on the interplay between star formation processes and accretion onto massive black holes in galactic nuclei. The selected sample will be the overlap between the 12$\\mu$m sample and an all-sky hard-X ray selected sample of AGNs. The Herschel data will follow extensive observational work that is already available from radio to X-rays frequencies and will provide the baseline of zero-redshift objects for comparison with the high redshift Universe which will be revealed by the Herschel cosmological surveys. ",
        "introduction": "ISO opened a new window on the spectra of IR-bright, ultraluminous infrared galaxies and active galactic nuclei, followed in the last years by \\textit{Spitzer}. The first mid-infrared spectra of active and starburst galaxies were obtained with ISO-SWS in Seyfert (Sturm et al. 2002) and starburst galaxies (Verma et al. 2003). \\textit{Spitzer} spectroscopy followed up with more systematic studies of samples of Seyfert (e.g. Buchanan et al. 2006) and starburst galaxies. The MIR spectral range includes many fine-structure lines excited by the hard radiation produced by black hole accretion as well as those excited by stellar ionization (e.g. Spinoglio \\& Malkan 1992) and MIR spectra can distinguish between the two processes, especially in obscured nuclei. However, while MIR spectroscopy can study the Narrow Line Regions (NLR) in obscured AGN and the effects of HII/starburst emission in actively star forming regions, FIR spectroscopy is the best way to study the influence of both the accretion process and stellar evolution on the geometry and physics of the densest, inner regions of galaxies. The geometry, and in particular the presence of the hypothetical torus around the nuclear engine, could be studied spectroscopically in the FIR, in a variety of molecular (high-J CO, OH, H$_2$O) and low excitation ionic transitions that trace the expected conditions of X-UV illuminated dusty tori (see e.g. Krolik \\& Lepp 1989; for predictions of high-J CO lines).  \\begin{figure}[h] \\centerline{\\includegraphics[angle=0,width=6cm]{spinoglio_l_fig1.eps}} \\caption{\\footnotesize The full ISO-LWS spectra of six IR-bright galaxies.  The spectra have been shifted and ordered vertically according to apparent excitation (Fischer et~al. 1999) and are not in order of relative luminosity or brightness.}\\label{fig1} \\end{figure}   As can be seen in Fig.1, the ISO-LWS observations showed a dramatic progression in ionic fine-structure emission line and molecular/atomic absorption line characteristics from strong [OIII]52, 88$\\mu$m and [NIII]57$\\mu$m emission to detection of only faint [CII]157$\\mu$m emission (Fischer et al.1999). The [CII]157$\\mu$m line in 15 ULIRGs revealed an order of magnitude deficit compared to normal and starburst galaxies relative to the FIR continuum (Luhman et al 2003). LWS observations of Arp 220 show absorption in molecular lines of OH, H$_2$O, CH, NH, and NH$_3$, as well as in the [OI]63$\\mu$m line and faint emission in the [CII]158$\\mu$m line (Gonzalez-Alfonso et al 2004).  The complete 2-200$\\mu$m spectrum of the prototype Seyfert 2 galaxy NGC1068 has been decomposed in the AGN and starburst components. Moreover, in the FIR it shows the 79, 119 and 163$\\mu$m OH rotational lines in emission, not in absorption as in every other galaxy yet observed. Modeling the three FIR lines of OH suggests the gas lies in small (0.1pc) and dense clouds ($\\sim 10^4 {\\rm cm^{-3}}$)  most probably in the nuclear region (potentially a signature of the torus) (Spinoglio et al. 2005).  ISO FIR spectroscopic observations have been collected only for a few objects and a systematic work on a representative sample of galaxies has yet to be done. High sensitivity and high resolution FIR spectroscopy is needed to understand the effects of AGN and starbursts on the geometry and physics of galactic nuclear and circumnuclear regions. The interrelationship between star formation and accretion onto massive black holes and their connection with galaxy formation and evolution can be understood through detailed observations of active and starburst galaxies in the local universe. Because dust and active nuclei (AGNs) coexist even at very high redshifts (e.g., Priddey et al. 2003; Bertoldi et al. 2003), it is imperative to study nuclear activity in a way that minimizes the obscuring effects of dust. This implies that both the sample and the observations must be defined carefully to overcome the biases associated with dust, extreme star-formation rates, and nuclear activity. ",
        "conclusions": ""
    },
    "0609/astro-ph0609347_arXiv.txt": {
        "abstract": "The shock model of gamma-ray bursts (GRBs) contains two equipartition parameters: the magnetic energy density and the kinetic energy density of the electrons relative to the total energy density of the shock, $\\epsilon_B$ and $\\epsilon_e$, respectively. These are free parameters within the model. Whereas the Weibel shock theory and numerical simulations fix $\\epsilon_B$ at the level of $\\sim$few$\\times(10^{-3}...10^{-4})$, no understanding of $\\epsilon_e$ exists so far. Here we demonstrate that it inevitably follows from the theory that $\\epsilon_e\\simeq\\sqrt{\\epsilon_B}$. The GRB afteglow data fully agree with this theoretical prediction. Our result explains why the electrons are close to equipartition in GRBs. The $\\epsilon_e-\\epsilon_B$ relation can potentially be used to reduce the number of free parameters in afterglow models. ",
        "introduction": "\\label{s:intro}  Shocks in gamma-ray bursters (GRBs) as well as in many other astrophysical sources are collisionless. Their physics is much more complex than that of standard hydrodynamic shocks. The Weibel shock theory developed by \\citet{ML99} has  been confirmed by a large number of numerical simulations both in the ultra-relativistic regime \\cite{Silva+03,Nish+03,Fred+04,Spit05} and in the non-relativistic one \\cite{MSK06}. In particular, it has been predicted and then numerically confirmed that the magnetic equipartition parameter, $\\epsilon_B$, defined as the ratio of the magnetic energy density to the total kinetic energy density of a shock, \\beq \\epsilon_B=\\frac{B^2/4\\pi}{m_pc^2n\\Gamma}, \\eeq falls in the range of $\\sim$few$\\times(10^{-3}...10^{-4})$, where $n$ is the co-moving particle density and $\\Gamma$ is the shock Lorentz factor; we assumed $\\Gamma\\gg1$.  Multiwavelength spectral fits to afterglows allowed reliable determination of micro-physics parameters for  a number of GRBs \\cite{PK01a,PK01b,PK02,CL00,LC01,LC03,Yost+03,Frail+03,Berger+04,% McMahon+04,P05}. The analyses found that $\\epsilon_B$ is indeed falls in the range $\\sim10^{-2}...10^{-5}$, however with relatively large errorbars --- typically of about a decade, depending on quality of data (afterglow temporal and spectral coverage) and the method of analysis. These results confirm (within uncertainty) the predictions of the Weibel shock theory. We note that this is the only theory, which explains the origin of magnetic fields in relativistic shocks.  The afterglow analyses also yield the values of $\\epsilon_e$, defined as the ratio of the kinetic energy density of electrons to the total energy density of the shock:\\footnote{ Often, $\\epsilon_e$ is defined for a power-law distributed electrons, so the power-law index enters the definition: $\\gamma_{e,min}=\\left((p-2)/(p-1)\\right)(m_p/m_e)\\Gamma\\epsilon_e$, where $\\gamma_{e,min}$ is the low-energy cut-off of the electron distribution and $p$ is the power-law index. In this case, to measure  the total energetics one uses another (related) parameter $\\epsilon_i$. Here we make no assumption about the electron energy distribution, hence our $\\epsilon_e$ measures the overall energetics of the electrons. Thus our $\\epsilon_e$, defined in Eq. (\\ref{ee}), is identical to $\\epsilon_i$ used by some authors. } \\beq \\epsilon_e=\\frac{U_e}{m_pc^2n\\Gamma}. \\label{ee} \\eeq The values turn out to be clustered at $\\sim$few$\\times(10^{-1}...10^{-2})$. Given the fact that before being shocked the electrons carry only about $\\sim m_e/m_p\\sim10^{-3}$ of the total energy of the ejecta, they must be accelerated somehow in the downstream region. Traditionally, one invokes diffusive Fermi acceleration for this purpose. However, the Fermi theory cannot accurately predict the value of $\\epsilon_e$ in GRB shocks. Moreover, recent studies \\cite{NO06,LW06} put the efficiency, or even the very presence, of Fermi acceleration at GRB shocks into question. Thus, until now, we had little understanding of the electron acceleration in GRBs and we did not know why the electrons are close to equipartition.  In this paper, we demonstrate that $\\epsilon_e\\simeq\\lambda\\sqrt{\\epsilon_B}$ (with the constant $\\lambda\\sim1$) in relativistic, baryon-dominated shocks (i.e., shocks in electron-proton plasma). The result is very robust and is based solely on well-known properties of collisionless shocks. ",
        "conclusions": "\\label{s:concl}  In this paper, we have shown that in relativistic, baryon-dominated shocks $\\epsilon_e\\simeq\\lambda\\sqrt{\\epsilon_B}$ with the constant $\\lambda\\sim1$. The result inevitably follows from the micro-structure of collisionless shocks. No {\\it ad hoc} assumptions were made. Interestingly, the values of $\\epsilon_e/\\sqrt{\\epsilon_B}$ derived from afterglow data of ten GRBs are clustered around unity, thus supporting our theory. Since the typical value of $\\epsilon_B$ is about $10^{-3}$, the corresponding value of $\\epsilon_e$ should be $\\sim0.03$. Indeed, the afterglow fits show that 0.03 is the typical value of the $\\epsilon_e$ parameter. Thus, our theory explains why the electrons in GRB shocks are close to equipartition. Interestingly, the typical uncertainly in $\\epsilon_B$ determined from afterglow fits is often rather large, -- of about an order of magnitude or even more. In contrast, the theoretical uncertainly in the value of $\\lambda$ is only a factor of few at most, and is very likely even less. Therefore, one can use the obtained relation $\\epsilon_e\\sim\\sqrt{\\epsilon_B}$ to reduce the number of free parameters in afterglow models. It will be very interesting to investigate how the goodness of afterglow model fits changes when the above relation is used.  A number of important questions are left outside of the present study. It would be interesting to calculate the energy distribution of electrons. Some simple estimates can, however, be given as follows. Far away from the filaments, the electrons form the isotropic distribution with a typical Lorentz factor $\\gamma_{e,min}\\sim\\Gamma$, because they have been pitch-angle scattered in small-scale magnetic fields (these scatterings do not change particle energy). Inside the filaments the $\\gamma$-factor of the electrons shall be maximum, $\\gamma_{e,max}\\sim(m_p/m_e)\\Gamma\\epsilon_e$. Since these electrons contribute the most to the observed emission, the emission-weighted (that is, ``observed'') distribution will be peaked at $\\gamma_{e,max}$, which shall correspond to $E_{peak}$ in the spectral distribution of GRB emission. Other mechanisms are likely needed in order to explain the power-law spectra above $E_{peak}$. Numerical PIC simulations \\citep{Silva+03} indicate that reconnection events occuring during current filament coalescences lead to local acceleration of electrons. This can possibly produce a power-law distribution of electrons. Alternatively, a power-law radiation spectrum can be produced even by a monoenergetic electron distribution, provided the spatial spectrum of small-scale magnetic fields is a power-law \\citep{M06}. PIC simulations indicate that power-law spectra of magnetic felds indeed form at relativistic shocks \\citep{Fred+04}. A related question is how large the fraction of the energetic electrons is. A simple estimate indicates that is should be of order the filling factor of current filaments, which can be determined from PIC simulations: it varies with the downstream distance from $\\sim1/few$ to $\\sim1/100$  in typical electron-proton runs. However, more accurate answers to the above questions require detailed modeling of the electron dynamics, taking into account the residence time of electrons inside the filaments (the electrons are deflected in the direction of the ion current as they move toward a filament). We are developing a detailed model and the results will be presented in forthcoming publications."
    },
    "0609/astro-ph0609171_arXiv.txt": {
        "abstract": "\\noindent We present results from a 2300 arcmin$^2$ survey of the Orion A molecular cloud at 450 and 850\\,$\\mu$m using the Submillimetre Common-User Bolometer Array (SCUBA) on the James Clerk Maxwell Telescope.  The region mapped lies directly south of the OMC1 cloud core and includes OMC4, OMC5, HH1/2, HH34, and L1641N.  We identify 71 independent clumps in the 850\\,$\\mu$m map and compute size, flux, and degree of central concentration in each. Comparison with isothermal, pressure-confined, self-gravitating Bonnor-Ebert spheres implies that the clumps have internal temperatures $T_d \\sim 22 \\pm 5\\,$K and surface pressures ${\\rm log}(k^{-1}\\,P\\,{\\rm cm}^{-3}\\,{\\rm K}) = 6.0 \\pm 0.2$.  The clump masses span the range $0.3 - 22\\,M_\\odot$ assuming a dust temperature $T_d \\sim 20\\,$K and a dust emissivity $\\kappa_{850} = 0.02\\,$cm$^{2}\\,$g$^{-1}$.  The distribution of clump masses is well characterized by a power-law $N(M) \\propto M^{-\\alpha}$ with $\\alpha = 2.0 \\pm 0.5$ for $M > 3.0\\,M_\\odot$, indicating a clump mass function steeper than the stellar Initial Mass Function. Significant incompleteness makes determination of the slope at lower masses difficult. A comparison of the submillimeter emission map with an H$_2$ 2.122\\,$\\mu$m survey of the same region is performed. Several new Class 0 sources are revealed and a correlation is found between both the column density and degree of concentration of the submillimeter sources and the likelihood of coincident H$_2$ shock emission. ",
        "introduction": "Star formation takes place within molecular clouds, in regions where the local pull of gravity overwhelms the random and ordered motions of the gas. Thus, in order to understand the physical process of star formation it is necessary to observe the formation and evolution of structure inside these clouds. The advent of sensitive millimeter and submillimeter continuum array detectors has provided an opportunity for large area mapping of the optically thin emission from cold dust within molecular clouds, yielding observations of the column density distribution within these regions. The most common features observed in these maps are dense condensations of dust, ``clumps\"\\footnote{Unfortunately, the words ``cores\" and ``clumps\" have ambiguous meaning and are used to denote various types of objects by different authors. In this paper ``clumps\" refer to stellar-mass sized condensations of dust and gas.}. These clumps reside in areas where stars are either forming or are expected to form. When combined with observations being taken at shorter wavelengths by the Spitzer Space Telescope, it should be possible to unambiguously classify these objects as either starless (possibly prestellar) or protostellar (Class 0 or Class I) (Jorgensen et al. 2006). Alternatively, comparison with other evolutionary tracers, such as H$_2$ shock emission (Walawender et al.  2005, 2006), provides evidence for which clumps already contain forming stars.  Continuing an effort to quantify the necessary pre-conditions for star formation, and in anticipation of the Spitzer observations, in this paper we present submillimeter dust continuum maps of the Orion A South region, including OMC4, OMC5, HH1/2, HH34, and L1641N. Previous papers in the series investigated the northern portion of the Orion A cloud containing the Orion Nebula (Johnstone \\& Bally 1999), the  Ophiuchus cloud (Johnstone et al. 2000b, hereafter Paper II), the Orion B North region (Johnstone et al. 2001, hereafter Paper III), and the Orion B South region (Johnstone, Matthews, \\& Mitchell 2006, hereafter Paper IV).  The Orion A molecular cloud south of the Orion Nebula contains several active sites of ongoing star and star cluster formation. The region dubbed OMC4 is located about 10\\arcmin\\ south of the Orion Nebula and consists of a V-shaped cluster of cores located near the top portion of the SCUBA survey. The double star $\\iota$ Orionis (spectral types O9III, B1III), about 30\\arcmin\\ south of the Orion Nebula, is the brightest member of the $>$3 Myr old NGC 1980 cluster located in the southern part of the Integral Shaped Filament (ISF) in Orion A. The $\\iota$ Orionis system has been implicated in a four-body dynamical interaction that launched the famous high-velocity run-away stars AE Auriga and $\\mu$ Columbae 2.6 Myr ago (Hoogerwerf, de Bruijne,  \\& de Zeeuw 2001).  This portion of the Orion OB1c association must lie at least 10 pc in front of the Orion A cloud and must represent a considerably older site of star birth than either the Orion Nebula or the other currently active star forming regions in Orion A (Brown et al. 1994).  The embedded L1641N cluster contains at least 30 mid-infrared sources detected by ISO (Ali \\& Noriega-Crespo 2004; Chen \\& Tokunaga 1994). Young stars in this cluster power dozens of Herbig-Haro (HH) flows, including two giant flows that can be traced for several parsecs from the L1641 cluster core (Reipurth, Devine, \\& Bally 1998).  The L1641N cluster is the most active site of on-going star formation in the fields covered by this SCUBA survey.  The western side of the Orion A cloud near the L1641N cluster contains a number of Herbig-Haro objects and young stars not associated with rich clusters.    The HH 34 jet, located about 16\\arcmin\\ west of the L1641 region, emerges from a star embedded within a small cloud which hosts a small group of young stars (e.g. Devine et. al. 1997).    The source and true nature of the bright but enigmatic bow-shaped HH~222,  located about 5\\arcmin\\ north-northwest of HH 34, remains unknown (Castets, Reipurth, \\& Loinard, 2004).  However, another giant bow, HH~401, located about 5\\arcmin\\ southwest of HH 34 appears to trace a parsec-scale outflow possibly powered by the HH~1/2 jet (Ogura 1995).  Near the southern periphery, our survey field contains the bright NGC1999 reflection nebula and its exciting star V380 Ori (spectral type A0e).  Many molecular outflows and Herbig-Haro objects are observed here, suggesting that this region also contains a small embedded group of young stars (Chen \\& Tokunaga 1994). NGC1999 is associated with a dense cloud core and molecular ridge, seen in millimeter CO and CS transitions, which extends about 10\\arcmin\\ towards the south (Morgan et al. 1991). The southern end of this ridge contains the famous Herbig-Haro objects HH~1, 2, and 3, as well as a number of fainter shocks (e.g. HH 144/145 and HH~147).  This region contains another compact group of young stars. ",
        "conclusions": "\\label{s_sum}  Using the techniques developed in Papers I-IV we have identified 71 submillimeter clumps within a  2300 arcmin$^2$ region of Orion A South. We have been able to derive estimated masses and temperatures for the majority of the clumps by noting that they can be modeled as Bonnor-Ebert spheres; that is, they can be approximated by almost constant density models with low internal velocities. On this basis the typical clump temperature is found to be 22$\\pm$5~K; this value is in reasonable agreement with that (18$\\pm$9~K) derived from the spectral indices obtained from the 850 and 450\\,$\\mu$m data.  Thus for most of the clumps a temperature of 20~K is a good working approximation, similar to the results in Orion B (Papers III and IV).  The mass function of submillimeter clumps found in Orion A South is steeper than the molecular cloud mass spectrum, suggesting that most of the  mass of submillimeter clumps lies in small mass objects. This is in agreement with the results in other nearby regions (Motte et al. 1998; Paper II; Paper III; Motte et al. 2001; Paper IV) although the mass function found here is somewhat steeper still than the mass functions derived for these other regions.  A comparison of the appearance of molecular hydrogen shock excited emission and the location of submillimeter clumps reveals that concentration and central column density of the clumps are correlated with evidence of internal protostellar activity. Several new Class O sources have been identified in Orion A South using this technique."
    },
    "0609/astro-ph0609617_arXiv.txt": {
        "abstract": "The program CORSIKA, usually used to simulate extensive cosmic ray air showers, has been adapted to a water medium in order to study the acoustic detection of ultra high energy neutrinos. Showers in water from incident protons and from neutrinos have been generated and their properties are described. The results obtained from CORSIKA are compared to those from other available simulation programs such as Geant4. ",
        "introduction": "The air shower program, CORSIKA (version 6204) \\cite{CORSIKA}, has been adapted to run in sea water i.e. a medium of constant density of 1.025 g per cm$^3$ rather than the variable density needed for an air atmosphere. Sea water was assumed to consist of a medium in which $66.2\\%$ of the atoms are hydrogen, $33.1\\%$ of the atoms are oxygen and $0.7\\%$ of the atoms are made of common salt, NaCl. The salt was assumed to be a material with atomic weight and atomic number A=29.2 and Z=14, the mean of sodium and chlorine. The purpose of this is to maintain the structure of the program as closely as possible to the air shower version which had two principal atmospheric components (oxygen and nitrogen) with a trace of argon. The presence of the salt component had an almost undetectable effect on the behaviour of the showers.  Other changes made to the program to accomodate the water medium include a modification of the stopping power formula to allow for the density effect in water. This only affects the energy loss in hadrons since the stopping powers for electrons are part of the EGS \\cite{EGS} package which is used by CORSIKA to simulate the propagation of the electromagnetic component of the shower. Smaller radial binning of the shower was also required since shower radii in water are much smaller than those in air. In addition the threshold for the LPM effect \\cite{LPM}, which suppresses pair production from photons and bremsstrahlung from electrons at high energy, was reduced to the much lower value necessary for water. Similarly, the interactions of $\\pi^0$s had to be simulated at lower energy than in air because of the higher density water medium. In all about 100 detailed changes needed to be made to the CORSIKA program to accomodate the water medium.  To test the implementation of the LPM effect \\cite{LPM} in the program 100 showers from incident gamma ray photons were generated and the mean depth of the first interaction (the mean free path) calculated. The observed mean free path was found to be in agreement with the expected behaviour when both the suppression of pair production and photonuclear interactions were taken into account (see figure \\ref{LPMfig}). This showed that the LPM effect had been properly implemented in CORSIKA.  Considerable fluctuations between showers occurred giving the following observed values of the ratios of the root mean square deviations to the mean value in proton showers: rms peak energy deposit to the peak energy deposit was observed to be $14\\%$ at $10^5$ GeV reducing to $4\\%$ $10^{11}$ GeV, that for the depth of the peak position varied from $19\\%$ to $7.4\\%$ and for the shower width from $63\\%$ to $18\\%$. To smooth out such fluctuations averages of 100 generated showers will be taken in the following. The statistical error on the averages is then given by these RMS values divided by $\\sqrt 100$. The hadronic energy contributes only about 10$\\%$ to the shower energy at the shower peak, the remainder being carried by the electromagnetic part of the shower. This is a well known effect in calorimeters. ",
        "conclusions": "The program CORSIKA has been modified to work in a water medium. Hadron and electromagnetic showers can be generated routinely. Neutrino interactions at low energy can be generated by the HERWIG package interfaced to CORSIKA  and it may be possible to extend this to higher energies.    \\subsection"
    },
    "0609/astro-ph0609421_arXiv.txt": {
        "abstract": "We present Berkeley Illinois Maryland Association (BIMA) millimeter interferometer observations of giant molecular clouds (GMCs) along a spiral arm in M31.  The observations consist of a survey using the compact configuration of the interferometer and follow-up, higher-resolution observations on a subset of the detections in the survey.  The data are processed using an analysis algorithm designed to extract GMCs and correct their derived properties for observational biases thereby facilitating comparison with Milky Way data.  The algorithm identifies 67 GMCs of which 19 have sufficient signal-to-noise to accurately measure their properties.  The GMCs in this portion of M31 are indistinguishable from those found in the Milky Way, having a similar size-line width relationship and distribution of virial parameters, confirming the results of previous, smaller studies.  The velocity gradients and angular momenta of the GMCs are comparable to the values measured in M33 and the Milky Way; and, in all cases, are below expected values based on the local galactic shear.  The studied region of M31 has a similar interstellar radiation field, metallicity, Toomre $Q$ parameter, and midplane volume density as the inner Milky Way, so the similarity of GMC populations between the two systems is not surprising. ",
        "introduction": "As the instrumentation for millimeter-wave telescopes improves, it becomes progressively more straightforward to study individual molecular clouds in other galaxies.  Recent studies of Local Group galaxies have surveyed large numbers of molecular clouds in the Large Magellanic Cloud \\citep{nanten}, the Small Magellanic Cloud \\citep{nanten-smc}, M33 \\citep{eprb03}, and a bevy of Local Group dwarf galaxies \\citep[e.g.~][]{wilson_6822,taylor_1569}.  These recent studies explore the nature of star formation on galactic scales by studying the properties of giant molecular clouds (GMCs, $M>10^5~M_{\\odot}$) throughout their host galaxies.  Such GMCs contain the majority of the molecular mass in the Milky Way's ISM and are responsible for most of the star formation in the Galaxy \\citep{psp3}.  The Andromeda Galaxy (M31) is the second largest disk galaxy in the Local Group, after the Milky Way, and it subtends over 2 deg$^{2}$ on the sky.  Its proximity \\citep[770 kpc,][]{m31-dist} makes it an excellent target for studying extragalactic molecular clouds. Numerous surveys of CO emission have been conducted over a portion of M31 and a comprehensive list of the 24 CO studies published up to 1999 is given in \\citet{fcrao-m31}. This extensive list of surveys can be supplemented with a few major studies that have occurred since then.  \\citet{sheth,sheth06} used the BIMA millimeter interferometer to study a $3'$ field in the outer region of the galaxy ($R_{gal}=12$~kpc) and find 6 molecular complexes similar to those found in the Milky Way. An extensive survey covering the entirety of the star-forming disk of M31 has been completed using the IRAM 30-m by \\citet[][see also references therein]{iram-m31-aa}. Finally, \\citet{muller-thesis} used the Plateau de Burre interferometer to examine the properties of molecular clouds in 9 fields.  Using the GAUSSCLUMPS \\citep{gaussclumps,gaussclumps2} algorithm, they decompose the emission into 30 individual molecular clouds.   Previous high-resolution observations of CO in M31 indicate that a large fraction of the molecular gas is found in GMCs.  Identifying individual GMCs requires a telescope beam with a projected size $\\lesssim 50$ pc, the typical size of a GMC in the Milky Way \\citep{psp3}, which requires an angular resolution of $14''$ at the distance of M31.  There have been seven observational campaigns that observed CO $(1\\to 0)$ emission from M31 at sufficient resolution to distinguish molecular clouds: \\citet{ich85, vbb87, lada_m31, wilson_m31, loinard-ovro, sheth, muller-thesis}.  With the exception of \\citet{loinard-ovro}, all of these studies have found GMCs with properties similar to those found in the inner Milky Way and \\citet{iram-bulge} have argued that the differences observed by \\citet{loinard-ovro} can be attributed to observational errors. Indeed, \\citet{vbb87} presented the first direct observations of GMCs in any external galaxy using interferometric observations.  Subsequent studies with interferometers and single-dish telescopes confirmed that most CO emission in M31 comes from GMCs and that the GMCs properties were similar to those found in the Milky Way \\citep{lada_m31, wilson_m31, sheth, muller-thesis}.  Although the molecular gas in M31 has been extensively studied, there remains a gap connecting the large-scale, single-dish observations and the small-scale, interferometer observations.  To address this gap, we completed CO($J=1\\to 0$) observations of a large (20~kpc${^2}$ region) along a spiral arm of M31 with high resolution ($\\sim 50$ pc).  We then followed up on these observations using a more extended configuration of the interferometer yielding data with a resolution of $\\sim 25$ pc.  This paper presents the observational data of the both the survey and the follow-up observations (\\S\\ref{obs}).  Using only the follow-up data, we present the first results, namely a confirmation of previous studies that find GMCs in M31 are similar to those in the Milky Way (\\S\\S \\ref{analysis},\\ref{larson-sec}).  Notably, this paper utilizes the techniques described in \\citep{props} to correct the observational biases that plague extragalactic CO observations, thereby placing derived cloud properties on a common scale that can be rigorously compared with GMC data from other galaxies.  The follow-up observations are also used to examine the velocity gradients and angular momentum of the GMCs, which are then compared to the remainder of gas in the galaxy for insight into the GMC formation problem (\\S\\ref{spang}).  We conclude the paper by examining the larger galactic environment of M31 to explore connections between the GMCs and the larger ISM (\\S\\ref{environment}). Subsequent work will explore the star formation properties of these GMCs and the formation of such clouds along the spiral arm using the data from the spiral arm survey. ",
        "conclusions": "  1. The GMC population is very similar to that found in the inner Milky Way.  Both the size-line width relationship and distribution of virial parameters are statistically indistinguishable from those derived from the population of clouds in the inner Milky Way.  This confirms the work of several authors \\citep{vbb87,wilson_m31, iram-m31-orig, iram-m31,sheth} while using techniques that correct for the biases imprinted by using interferometric observations.  2. We measured the velocity gradients and angular momenta across the GMCs.  Like the Milky Way and M33, there is significantly more angular momentum in possible progenitor material than in the resulting GMCs, though the discrepancy is smallest in M31 owing to the relatively low galactic shear at large galactocentric radii.  All three systems suggest braking mechanisms act during the process of cloud formation to dissipate angular momentum.  3. The galactic environment in M31 where the GMCs are found is similar in most respects that of the solar neighborhood and inner Milky Way. That GMCs in both environments are similar is consistent with the idea that the galactic environment regulates GMC properties.  In systems where the galactic environment differs significantly from those studied here, the GMC properties are found to be different \\citep[e.g.,][]{psp5}."
    },
    "0609/astro-ph0609751_arXiv.txt": {
        "abstract": "{We present a new cluster-finding algorithm based on a combination of the Voronoi Tessellation and Friends-Of-Friends methods. The algorithm utilises probability distribution functions derived from a photometric redshift analysis. We test our algorithm on a set of simulated cluster-catalogues and have published elsewhere its employment on UKIDSS Ultra Deep Survey infrared $J$ and $K$ data combined with 3.6 $\\mu$m and 4.5 $\\mu$m Spitzer bands and optical $BVRi'z'$ imaging from the Subaru Telescope. This pilot study has detected clusters over 0.5 square degrees in the Subaru XMM-Newton Deep Field. The resulting cluster catalogue contains 13 clusters at redshifts 0.61 $\\leq z \\leq$ 1.39 with luminosities $10 {\\rm L^*} \\lesssim L_{\\rm tot} \\lesssim 50 \\rm L^*$.}  \\authorrunning{Caroline van Breukelen et al.} ",
        "introduction": "Clusters of galaxies are the largest virialized structures in the Universe and play an important role in our understanding of how dark-matter haloes collapse and large-scale structure evolves. Their number density can place constraints on the mass density of the universe and the amplitude of the mass fluctuations (e.g. Eke et al. 1998). Clusters also act as astrophysical laboratories for understanding the formation and evolution of galaxies and their environments. For instance the study of high-redshift clusters can help us gain an understanding of the feedback processes caused by star-formation and Active Galactic Nuclei (e.g. Silk \\& Rees 1998). It is therefore desirable to have a large, homogeneous catalogue of clusters at a range of redshifts in the universe.  At present, however, there are only few clusters known at $z>1$ and the vast majority of these are from X-ray surveys (e.g. using XMM-Newton; Stanford et al. 2006). The paucity of known, distant clusters stems from various selection effects or deficiencies. For instance, optical searches for clusters, that work so efficiently at $z<1$, are ineffective once the 4000 \\AA\\ break falls outside the I-filter pass band ($z>1$) given the predominance of early-type, red galaxies in clusters. One solution to this problem is to select clusters in the near-infrared. Recent developments have seen the advent of large near-infrared cameras like the Wide Field Infrared Camera (WFCAM) on the United Kingdom Infrared Telescope (UKIRT). WFCAM is now undertaking the UKIRT Infrared Deep Sky Survey (UKIDSS, Lawrence et al. 2006), which is a suite of large area and deep near-infrared sky surveys. UKIDSS provides the ideal opportunity to search for high-redshift clusters in the infrared wavelength regime. The survey has a high efficiency as it provides over an order of magnitude increase in survey speed over existing near-infrared imagers.  There are numerous methods for detecting clusters in optical/infrared imaging surveys. The problem is easier with spectroscopic redshifts, but currently spectroscopic information for infrared-selected galaxies is in short supply, time consuming, and impractical over large areas. However, approximate redshifts can be calculated via photometric redshift estimation.  Despite being a popular method, remarkably little attention is paid to using photometric redshifts to isolate clusters. Further, with the exceptions of Kim et al. (2002), Goto et al. (2002), Bahcall et al. (2003) and Lopes at al. (2004), very little work has been done to compare the various cluster detection method that exist (for a review see Gal 2005). ",
        "conclusions": ""
    },
    "0609/astro-ph0609798_arXiv.txt": {
        "abstract": "We simulate the early stages of the evolution of turbulent, virialized, high-mass protostellar cores, with primary attention to how cores fragment, and whether they form a small or large number of protostars. Our simulations use the Orion adaptive mesh refinement code to follow the collapse from $\\sim 0.1$ pc scales to $\\sim 10$ AU scales, for durations that cover the main fragmentation phase, using three-dimensional gravito-radiation hydrodynamics. We find that for a wide range of initial conditions radiation feedback from accreting protostars inhibits the formation of fragments, so that the vast majority of the collapsed mass accretes onto one or a few objects. Most of the fragmentation that does occur takes place in massive, self-shielding disks. These are driven to gravitational instability by rapid accretion, producing rapid mass and angular momentum transport that allows most of the gas to accrete onto the central star rather than forming fragments. In contrast, a control run using the same initial conditions but an isothermal equation of state produces much more fragmentation, both in and out of the disk. We conclude that massive cores with observed properties are not likely to fragment into many stars, so that, at least at high masses, the core mass function probably determines the stellar initial mass function. Our results also demonstrate that simulations of massive star forming regions that do not include radiative transfer, and instead rely on a barotropic equation of state or optically thin heating and cooling curves, are likely to produce misleading results. ",
        "introduction": "The previous generation of telescopes revealed a great deal about the gas from which massive stellar clusters form. With them, observers were able to survey the dense clumps of thousands of $\\msun$ that are likely the progenitors of clusters, using molecular line emission \\citep[e.g.][]{plume97,shirley03}, thermal dust emission \\citep[e.g.][]{carey00,mueller02}, or infrared absorption \\citep[e.g.][]{menten05,rathborne05,rathborne06a,simon06}. However, the large distances to these regions, their high extinctions, and the confusion produced by their density prevented these observations from directly probing structures with masses comparable to individual stars, data that have been available since the 1980s for nearby, low-mass star-forming regions. In the last few years, millimeter interferometers and the Spitzer Space Telescope have started to change that situation by making available information on the structure of massive star forming regions comparable to that previously available only for low-mass regions. These observations have identified a population of high-mass cores that are comparable in mass to individual massive stars. They are dense (mean densities $\\sim 10^6$ H nuclei cm$^{-3}$, rising strongly towards their centers), cold (temperatures $\\approx 10-40$ K), turbulent (linewidths $\\sim 1$ km s$^{-1}$), compact (radii $\\sim 0.1$ pc), and round (aspect ratios of $2:1$ or less) \\citep{reid05,sridharan05,beuther05b,beuther06a,garay05, pillai06a}. In some cases they show no mid-infrared emission or even mid-infrared absorption, indicating that they have not yet converted a significant fraction of their mass into stars, and are therefore near the onset of star formation.  The idea that these cores might be the progenitors of individual massive stars is bolstered by two pieces of circumstantial evidence. First, the mass function of these cores appears to match the \\citet{salpeter55} slope of roughly $-1.3$ in the logarithmic distribution observed for high-mass end of the stellar initial mass function \\citep[IMF,][]{beuther04b, reid05, reid06a, reid06b}. This extends earlier observations showing that in nearby low-mass star-forming regions the core mass function matches the IMF as well \\citep{motte98,testi98,johnstone01, onishi02}. Second, cores are mass segregated in such a manner that the mass function is the same throughout a protocluster gas clump, with the exception that the most massive cores, those greater than several $\\msun$ in mass, are found only near the center \\citep{elmegreen01,stanke06}. Star clusters exhibit a very similar pattern of mass segregation \\citep{hillenbrand98,huff06}, and while some of this may be dynamically produced, much of it is likely a result of the locations where the stars formed \\citep{bonnell98b,tan06a}.  However, a direct mapping from core mass to star mass is only possible if massive cores collapse to form individual stars or small-multiple systems, as proposed by \\citet[hereafter MT02 and MT03]{mckee02, mckee03}, rather than fragmenting into many objects and producing a cluster of low-mass stars. Whether this happens or not is quite uncertain. \\citet{bate05} argue that dense cores are likely to produce small objects because the Jeans mass decreases with density at fixed temperature \\citep[although see][]{martel06}. \\citet{dobbs05} simulate the collapse of massive, turbulent cores and find that they generally fragment into as many as 20 objects, depending on initial conditions and on the assumed gas equation of state. However, \\citet{krumholz06b} uses one dimensional analytic calculations to show that radiation feedback from accreting protostars can substantially inhibit fragmentation even at early times, because at the high accretion rates and opacities expected in massive cores, accretion luminosity can heat gas to hundreds of Kelvin out to distances of $\\gtsim 1000$ AU from an accreting protostar. \\citeauthor{krumholz06b} also finds that using an isothermal or barotropic equation of state, as \\citet{bate05} and \\citet{dobbs05} do, is likely to produce misleading results on fragmentation because it misses this effect.  While these calculations are suggestive, because they are analytic they are necessarily limited in how they deal with real, turbulent cores. The best means of settling the question of how massive cores fragment is direct numerical simulation, including a treatment of radiative feedback from embedded protostars. However, such simulations have to this point not been reported in the literature. Some simulations of massive star formation with radiation use quiescent initial conditions \\citep{yorke99, yorke02} in two dimensions, and are therefore incapable of answering questions about the fragmentation of turbulent structures. (Those calculations focus on the effects of radiation pressure, an important effect in the later evolution of massive protostars which we do not consider in detail in this paper.) In three dimensions, some simulations of massive cores use local cooling functions rather than solving the radiative transfer problem \\citep{banerjee06}, and are therefore unable to study the effects of feedback. Moreover, the modifications to the cooling function the authors use to approximate the behavior of optically thick gas are of unknown accuracy. Simulations of star formation with feedback and a treatment of radiative transfer have been limited to low-mass, non-turbulent initial conditions \\citep{whitehouse06}. Furthermore, both \\citeauthor{whitehouse06} and \\citeauthor{banerjee06}\\ only advance their simulations to the point where the first collapsed object forms, and for this reason they are incapable of studying accretion and fragmentation, or the effects of radiative feedback on either of these processes.  Here we report the first three-dimensional gravito-radiation hydrodynamic simulations of the collapse and fragmentation of turbulent high-mass protostellar cores. Our simulations include radiative transfer and the effects of feedback from both accretion onto and nuclear burning within embedded protostars. We run through the main fragmentation phase and follow the accretion process to the point where deuterium burning begins in the most massive stars, thereby greatly heating the gas and strongly suppressing further fragmentation. This enables us to address the question of how fragmentation of massive cores proceeds, and how radiative feedback influences it. In \\S~\\ref{method} we discuss the methodology for our simulations, and in \\S~\\ref{results} we present our results. We discuss the implications of these results for the mechanism of massive star formation and the origin in of the IMF in \\S~\\ref{discussion}, and summarize in \\S~\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  We report the results of simulations of the collapse and fragmentation of massive protostellar cores using gravito-radiation hydrodynamics with adaptive mesh refinement. We find that including radiative transfer in our simulations produces dramatic effects on the evolution of these objects. When radiation is included, massive protostellar cores with the properties of observed cores do not fragment strongly. They collapse to a handful of objects, with the majority of the mass accreting onto one primary object. The object gains mass by accretion of gas that is prevented from collapsing by radiative heating. Some low-mass stars do form in addition to the primary massive star, through a combination of fragmentation at sites sufficiently far from the protostar to be fairly cool, and fragmentation of unstable protostellar disks around the massive star. However, these do not contain significant mass. The disks that form are able to transport mass onto the central star very rapidly due to a large-scale gravitational instability, which appears to form due to the SLING mechanism. Overall, our results are consistent with the turbulent core model of MT02 and MT03, with the analytic treatment of radiative suppression of fragmentation by \\citet{krumholz06b}, and with the analytic massive protostellar disk models of \\citet{kratter06}.  Our results suggest that massive cores are the direct progenitors of massive stars, without an intermediate phase of competitive accretion or stellar collisions. Both of these mechanisms require that a large number of protostars form out of dense gas in clusters $\\sim 1000$ AU in size or smaller. However, such strong fragmentation seems not to occur in the environments where massive stars form, because rapid accretion rapidly raises the gas temperature and prevents nucleation of small protostars. Instead, most gas in a massive cores accretes onto the first object to form, and this object prevents other, comparable mass stars from forming after it.  One additional conclusion from our work is that one must include radiative transfer in simulations in order to obtain the correct behavior on small scales. Prescriptions for the equation of state based on the barotropic approximation or on optically thin heating and cooling models severely underestimate gas temperatures because they do not, and cannot, include heating by accretion luminosity onto a central protostar. The high densities found in massive protostellar cores mean that this luminosity can reach thousands of $\\lsun$ even for $\\ltsim 1$ $\\msun$ protostars, an effect that cannot be neglected."
    },
    "0609/astro-ph0609084_arXiv.txt": {
        "abstract": "% The winds of stars with very specific temperatures and luminosities are ideal for determining the magnitude and nature of mass loss in OB stars. I identify these stars and analyze their wind lines.  The results are discussed within the context of recent findings which appear to indicate that the mass-loss rates of OB stars may as much as an order of magnitude less than commonly accepted values. ",
        "introduction": "Recent measurements of O stars winds indicate that their mass loss rates, $\\dot{M}$, may be significantly less than previous estimates and expectations (e.g., Massa et al.\\ 2003, Repolust et al.\\ 2004, Bouret et al.\\ 2005, Fullerton et al.\\ 2006).  The primary mass loss diagnostic used in these analyses is P~{\\sc v}~$\\lambda \\lambda 1118, 1128$.  Although wind models predict that the ion fraction of P$^{+4}$ should approach 100\\% in the winds of mid- to late O stars, this is a Na-like ion, and can be quite fragile.  Consequently, it is important to verify the results based on P~{\\sc v}. ",
        "conclusions": "Preliminary Sulfur results give {\\em total} sulfur mass loss rates, $\\dot{M}q(S^{+3})+\\dot{M}q(S^{+4})+\\dot{M}q(S^{+5})$, between 0.04 -- 0.18 times smaller than those expected from either theoretical (Vink et al.\\ 2000) or radio or H$_\\alpha$ determinations (see, Fullerton et al.\\ 2006). These results are similar those determined by the P~{\\sc v} analyses. Clumping and porosity can influence the results, and must be incorporated into the analysis to determine exact factors.  Nevertheless, it seems that the mass loss rates of O stars will have to be revised downward by some amount. We intend to extend the current analysis to a much larger sample and to investigate the effects of clumping more closely.  \\begin{center} \\centering\\epsfig{figure=fits.ps,angle=90, width=4.9in}% \\end{center} \\vspace{-0.2in}\\noindent {\\bf Figure 1:} Fits to the program stars.  The data are faint, solid curves, and the fits are heavy solid curves.  Models with $\\tau_{rad}$ equal to twice and half of the best fit values are shown as dashed curves.  The values of $\\dot{M}q(S^{+4})$ implied by the fits are between 0.04 and 0.18 of the expected values."
    },
    "0609/astro-ph0609567_arXiv.txt": {
        "abstract": "We present single-dish and interferometric millimeter line observations of the HH~111 outflow and its driving source. The physical conditions of the protostellar core have been determined from the emission of the millimeter line emission of CO and its isotopomers and CS with the IRAM 30m telescope, and the CO \\jseventosix\\ line with the Caltech Submm Observatory. The molecular gas emission reveals a small condensation of cold ($\\rm T= 20-25\\K$) and dense gas ($n(\\htwo)= 3\\times 10^5\\cmmt$). The low-velocity outflowing gas has been mapped with the IRAM Plateau de Bure interferometer. The cold gas is distributed in a hollow cylinder surrounding the optical jet.  The formation of this cavity and its kinematics are well accounted for in the frame of outflow gas entrainment by jet bow shocks. Evidence of gas acceleration is found along the cavity walls, correlated with the presence of optical bow shocks. The cavity has been expanding with a mean velocity of $4\\kms$ on a timescale of $8700\\yr$, similar to the dynamical age of the optical jet. The separation of the inner walls reaches $8\\arcsec-10\\arcsec$, which matches the transverse size of the wings in the bow shock. CSO observations of the \\jseventosix\\ line show evidence of a high-velocity and hot gas component ($\\rm T=300-1000\\K$) with a low filling factor. This emission probably arises from shocked molecular gas in the jet. Observations of the $\\rm ^3P_2-^3P_1$ [CI] line are consistent with C-type non-dissociative shocks. Mapping of the high-velocity molecular bullets B1-B3, located beyond the optical jet, with the IRAM PdBI reveals small structures of $3\\arcsec \\times 7\\arcsec$ flattened perpendicular to the flow direction. They are made of cold ($\\rm T\\sim 30\\K$), moderate density gas ($\\rm n(\\htwo)= (0.5-1.0)\\times 10^4\\cmmt$). We find evidence that the bullets are expanding into the low-density surrounding medium. Their properties are consistent with their being shocked gas knots resulting from past time-variable ejections in the jet. ",
        "introduction": "Herbig-Haro jets are manifestations of the accretion processes that are involved in the birth of stars, and as such represent a fossil record of the recent activity of their driving sources. The study of jets thus enables us to understand the past history of young stars as they build up their mass. As these jets penetrate their ambient medium, they transfer momentum and accelerate their surroundings and it appears that it is this interaction which causes the ubiquitous molecular outflows. For a recent review on these issues, see Reipurth \\& Bally (2001).  One of the finest HH jets known is the HH 111 jet, located in the L1617 cloud in Orion (Reipurth 1989). It has a very highly collimated body, with a large number of individual knots, and has large proper motions (300 to $600\\kms$). The jet flows at an angle of only 10$^o$ to the plane of the sky (Reipurth, Raga, \\& Heathcote 1992). Beyond the collimated jet the flow continues, and it is in fact only a minor part of a truly gigantic bipolar HH complex stretching over 7.7 pc (Reipurth, Bally, \\& Devine 1997a). Optical spectra are discussed by Morse et al. (1993a,b) and Noriega-Crespo et al. (1993). Near-infrared observations have revealed a remarkable symmetry between the optical jet and a near-infrared counterjet, and that a second bipolar flow, HH~121, emerges from the source, suggesting that it is a binary (Gredel \\& Reipurth 1993, 1994; Davis, Mundt, \\& Eisl\\\"offel 1994, Coppin et al. 1998, Reipurth et al. 2000). The optical jet is co-axial with a major well collimated molecular outflow (Reipurth \\& Olberg 1991, Cernicharo \\& Reipurth 1996, Nagar et al. 1997, Lee et al. 2000).  Beyond the visible jet, Cernicharo \\& Reipurth (1996) (henceforth CR) found a set of three high velocity CO bullets, which apparently represent working surfaces that no longer produce optical shock emission.  The jet is driven by IRAS 05491+0247, a Class~I source with a luminosity of about 25 L$_\\odot$, surrounded by cold dust and gas (Reipurth et al. 1993, Stapelfeldt \\& Scoville 1993, Yang et al. 1997, Dent, Matthews, Ward-Thompson 1998). A small 3.6 cm radio continuum jet was detected around the source by Rodr\\'\\i guez \\& Reipurth (1994). Recently, infrared images obtained with {\\em HST} have revealed that another source is located only 3 arcsec from the main one, which itself is a binary judging from the quadrupolar outflows; so altogether it appears that the driving source is part of a hierarchical triple system (Reipurth et al. 1999). In fact, the entire giant jet flow may be a result of the dynamical interactions between the three members of this triple system (Reipurth 2000).  The short blue lobe found by Reipurth \\& Olberg (1991) appears to be a highly collimated molecular jet about $90\\arcsec$ long with a collimation factor of 9 (see their Figs.~2-3-4) in the $12\\arcsec$ data of CR. The blue lobe ends at the bow shock P and the optical knots Q,R,S. Nagar et al. (1997) showed an interferometric map at $6\\arcsec$ resolution of the $\\jonetozero$ line, which allowed them to resolve the structure of the low-velocity gas around the optical jet.  The main features of the ionized and molecular gas in HH~111 are summarized in Fig.~1.  In this paper we present single dish and high resolution interferometric observations of the body of HH~111, the protostellar core and the high velocity CO bullets.  We characterize the gas kinematics and the physical conditions in the gas and report the detection of a hot temperature component in the outflow. We discuss these results in relation to current theories on the coupling of HH jets and molecular outflows.  We find that the molecular gas entrainment in the outflow is well accounted for in the frame of jet bow shock driven outflows, as modelled by Raga \\& Cabrit (1993) (henceforth RC93). Preliminary results were reported by Reipurth \\& Cernicharo (1995) and Cernicharo, Neri \\& Reipurth (1997). ",
        "conclusions": "We have presented single-dish and interferometric line observations of the HH~111 outflow and its driving source. The physical conditions of the protostellar core were determined from the emission of the millimeter lines of CO and its isotopomers, the millimeter lines of CS, all of them observed with the IRAM 30m telescope, and the CO \\jseventosix\\ line observed with the CSO. The molecular gas emission reveals a small condensation of cold ($\\rm T= 20-25\\K$) and dense gas ($n(\\htwo)= 3\\times 10^5\\cmmt$).  The low-velocity outflowing gas has been mapped with the IRAM Plateau de Bure interferometer. It is distributed in a hollow cylinder that surrounds the optical jet. The formation of this cavity and its kinematics are well accounted for in the frame of the model of RC93. The jet entrains ambient molecular material swept-up in the propagation of bow shocks, which disturb the ambient gas at large distances from the jet axis. Our observations indicate that bow shock V played a major role in shaping the global cavity. The gas velocity is found to vary almost linearly with distance to the source over the whole flow length, except close to the source. The cavity has been expanding with a mean velocity of $4\\kms$ on a timescale of $8700\\yr$, similar to the dynamical age of the optical jet. The separation of the inner walls reaches $8-10\\arcsec$ near knots P and S, which corresponds to the transverse size of the wings in these bow shocks.  The more recent ejections detected close to the source are still efficient at transferring momentum to the ambient material. Regions of strong gas acceleration are detected along the jet, which are correlated with the optical bow shocks at knots F and L.  Observations of the high-excitation line CO \\jseventosix\\ reveal high-velocity outflowing gas ($\\rm v= 150\\kms$ after deprojection) near the central source. This high-velocity component is hot ($T= 300-1000\\K$) and fills only a small fraction of the beam. We propose that this emission arises from shocked gas in the optical jet.  Emission of the molecular bullets B1-2-3 was mapped with the IRAM PdBI in the CO \\jtwotoone\\ transition. The emission is distributed in structures of $3\\arcsec \\times 7\\arcsec$ flattened perpendicularly to the jet main axis. Their kinematics is complex but an overall side-ways expanding motion is detected inside all of them. Their physical properties have been reassessed; we find a low temperature for the bullets ($\\sim 30\\K$), in agreement with the cooling timescale for shocked gas initially heated at temperatures of $100\\K$ or beyond, and cooling down while flying freely.  Their density lies now in the range $(0.5-1.0)\\times 10^4\\cmmt$. These molecular bullets probably form in the shocks resulting from time-variable ejections in the powering source, in the manner proposed by Raga et al. (2002) to account for the presence of mini bow shocks and knots in the optical jet. In such a scenario, the weak high-velocity molecular jet detected between the molecular bullets represents intershocked gas ejected from the central source."
    },
    "0609/hep-th0609150_arXiv.txt": {
        "abstract": "We introduce a model of scalar field dark energy, {\\it Cuscuton}, which can be realized as the incompressible (or infinite speed of sound) limit of a scalar field theory with a non-canonical kinetic term (or k-essence). Even though perturbations of {\\it Cuscuton} propagate superluminally, we show that they have a locally degenerate phase space volume (or zero entropy), implying that they cannot carry any microscopic information, and thus the theory is causal. Even coupling to ordinary scalar fields cannot lead to superluminal signal propagation. Furthermore, we show that the family of constant field hypersurfaces are the family of Constant Mean Curvature (CMC) hypersurfaces, which are the analogs of soap films (or soap bubbles) in a Euclidian space. This enables us to find the most general solution in 1+1 dimensions, whose properties motivate conjectures for global degeneracy of the phase space in higher dimensions. Finally, we show that the {\\it Cuscuton} action can model the continuum limit of the evolution of a field with discrete degrees of freedom and argue why it is protected against quantum corrections at low energies. While this paper mainly focuses on interesting features of {\\it Cuscuton} in a Minkowski spacetime, a companion paper \\cite{companion} examines cosmology with {\\it Cuscuton} dark energy. ",
        "introduction": "There has been much recent interest in field theories with non-canonical kinetic terms. Many of these theories are inspired by purely phenomenological motivations. K-essence \\cite{Armendariz-Picon:2000dh,Chiba:1999ka} and k-inflation \\cite{Armendariz-Picon:1999rj} have been designed in order to solve the cosmological coincidence and inflationary fine tuning problems, while Bekenstein's theory of gravity \\cite{Bekenstein:2004ne} attempts to accommodate apparent deviations from Newtonian gravity on galactic scales, within a relativistically covariant framework.  More theoretical motivations have led to theories such as ghost condensate \\cite{Arkani-Hamed:2003uy}, which is considered as an analog of Higgs boson for general relativity, or variations of Born-Infeld action that describe non-perturbative objects in string theory \\cite{Silverstein:2003hf,Sen:2002nu}.  Here we present a new class of actions with a non-canonical kinetic term which is characterized by the Hamiltonian symplectic structure of the theory degenerating in a cosmologically interesting homogeneous limit. In other words, in the limit that the field degree of freedom becomes homogeneous in the locally freely falling frame (where the metric is locally a Minkowski metric), the equation of motion does not have any second order time derivatives and the field becomes a non-dynamical auxiliary field, which merely follows the dynamics of the fields that it couples to. Thus we call this field {\\it Cuscuton} (pronounced {\\tt k\\\"{a}s-k\\\"{u}-t\\\"{a}n}), after the Latin name for the parasitic plant of dodder, Cuscuta \\footnote{The word Cuscuta is most likely of Aramaic or Hebrew origin, which has entered Latin through Arabic or Greek \\cite{cusc}.}.  Nonetheless, the {\\it Cuscuton} action may a priori appear to have dynamical degrees of freedom because there is a non-degenerate conjugate momentum degree of freedom if the field configuration has a non-vanishing spatial gradient.  It is this feature which justifies exploration of field theoretic features of this class of theories presented here.  For concreteness, we will focus on scalar field actions of the k-essence \\cite{Armendariz-Picon:2000dh,Armendariz-Picon:2000ah} form\\begin{equation} S_{\\varphi}=\\int d^{4}x\\sqrt{-g}\\left[\\frac{1}{2}F(X,\\varphi)-V(\\varphi)\\right],\\label{eq:simpformstart}\\end{equation} where $X\\equiv\\partial_{\\mu}\\varphi\\partial^{\\mu}\\varphi$ with a particular class of choice for $F$ such as to satisfy our degeneration of symplectic structure feature that will be defined in the next section.  There has been some controversy in the literature regarding the causality of dynamical k-essence models with $c_s > 1$. While the original literature on k-essence (e.g. \\cite{Armendariz-Picon:2000dh,Armendariz-Picon:2000ah}), as well as some follow-up studies \\cite{2002astro.ph.12302D,Bruneton:2006gf}, argue that superluminal modes of k-essence cannot carry information on closed loops, and thus do not break causality, others \\cite{Adams:2006sv,Bonvin:2006vc} take $c_s>1$ at its face value, arguing that it cannot be realized as the IR limit of an inherently causal field theory. Superluminal scalar field models are also claimed to lead to signal propagation out of the horizon of a black hole \\cite{Babichev:2006vx}, as well as large tensor-to-scalar ratios for inflationary perturbations \\cite{Mukhanov:2005bu}.  In Section\\ref{causality}, we show that although at face value {\\it Cuscuton} seems to possess dynamics, which allows superluminal signal propagation, it actually contains no local dynamical degree of freedom for $X>0$. What that means is that the theory even with a kinetic term acts like a pure constraint system, modifying the dynamics of whatever it couples to. In that sense, this is a k-essence theory which behaves like modified gravity, since if it couples to gravity and other fields, it merely provides a means of changing those dynamical sectors through constraints without adding any new local degrees of freedom of its own. Moreover, the superluminal modes cannot carry any information due to this degeneration of dynamical degrees of freedom \\footnote{However, note that although the corner that smoothly connects to the phase space of degenerating symplectic structure contains no local dynamics, this is not the case in the region where the variation of $S_{\\varphi}$ with respect to $X$ becomes singular. We leave this for a future investigation since that part of the theory is not fixed by the definition of the class of theories that we are proposing.}.  Section \\ref{CMC discussion} demonstrates that constant field hypersurfaces have constant mean curvature, making them Minkowski space analogs of soap bubbles/films in a Euclidian space. This leads to exact solutions of theory in 1+1 dimensions in Section \\ref{exact}, which enables us to make general conjectures about uniqueness and generic singularity of the solutions. We go on to briefly consider the coupling of {\\it Cuscuton} to gravity (in the context of homogeneous cosmology), as well as ordinary scalar fields, in Section \\ref{coupled}. The latter, in particular, shows that propagating degrees of freedom remain subluminal, even after coupling to {\\it Cuscuton}.  In Section \\ref{instanton} we discuss a possible physical model for {\\it Cuscuton action}, and argue why it may be protected against quantum corrections at low energies. Finally, Section~\\ref{conclude} summarizes our results and concludes the paper. A companion paper \\cite{companion} examines cosmology with a {\\it Cuscuton} dark energy fluid. ",
        "conclusions": "\\label{conclude}  In this paper, we have analyzed the flat space field theoretic aspects of {\\it Cuscuton} action, which we define as a scalar field theory whose kinetic term reduces to a total derivative in the homogenous limit.  This theory was defined in this way and may be of interest to cosmology because it is equivalent to a k-essence fluid with an infinite speed of sound.  The surprise is that even though the dynamical equations seem to admit superluminal signal propagation, there is \\emph{no} physical violation of causality, as perturbations have degenerate phase space and thus transport no information.  That means this can be used to modify gravity and other field theories to which it couples in a novel manner since the fields have a kinetic term, unlike usual Lagrange multipliers or auxiliary fields. We have also verified that this modification does not lead to superluminal propagation in ordinary scalar fields.  We have found some general interesting features to classical solutions to this theory.  Hypersurfaces of constant field have constant mean curvatures, which makes them the Minkowski space analogs of soap films (or soap bubbles) in the Euclidian space.  We can also solve the theory completely in 1+1 dimensions, and thereby prove non-local degeneracy of phase space (which is equivalent to the overconstraining behavior of Cauchy initial conditions), as well as generic presence of singularities in the solutions. Extension of these results to higher dimensions is a plausible conjecture.  As far as the qualitative behavior of the physics is concerned, {\\it Cuscuton} action can be viewed as a continuum limit of an action of the form Eq.~(\\ref{seff}) governing a large number of discrete degrees of freedom. It is the most general (local) action that only depends on the area and volume of constant field hypersurfaces, which is why it may probably be protected against quantum corrections at low energies. This is consistent with the absence of quantum corrections to the {\\it Cuscuton } action as a result of its degenerate phase space.  Cosmology literature does not lack in abundance of models of dark energy (or its alternatives).  The features that still make {\\it Cuscuton} interesting are the following:  \\begin{itemize} \\item  Even though it has a kinetic term and modifies the cosmic dynamics, it does not introduce any additional (perturbative) degree of freedom. Therefore, it can be considered a minimal modification of a cosmological constant, or a minimal model for evolving dark energy.  \\item It is protected against quantum corrections at low energies. \\end{itemize} A companion paper \\cite{companion} examines cosmology with a {\\it Cuscuton} dark energy fluid."
    },
    "0609/astro-ph0609641_arXiv.txt": {
        "abstract": "As a neutron star spins down, the nuclear matter continuously is converted into quark matter due to the core density increase and then latent heat is released. We have investigated the thermal evolution of neutron stars undergoing such deconfinement phase transition. We have taken into account the conversion in the frame of the general theory of relativity. The released energy has been estimated as a function of change rate of deconfinement baryon number. Numerical solutions to cooling equation are obtained to be very different from the without heating effect. The results show that neutron stars may be heated to higher temperature which is well-matched with pulsar's data despite onset of fast cooling in neutron stars with quark matter core. It is also found that heating effect has magnetic field strength dependence. This feature could be particularly interesting for high temperature of low-field millisecond pulsar at late stage. The high temperature could fit the observed temperature for PSR J0437-4715. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609688_arXiv.txt": {
        "abstract": "We have detected low-amplitude radial-velocity variations in two stars, USNO-B1.0 1219-0005465 (GSC 02265-00107 = WASP-1) and USNO-B1.0 0964-0543604 (GSC 00522-01199 = WASP-2). Both stars were identified as being likely host stars of transiting exoplanets in the 2004 SuperWASP wide-field transit survey. Using the newly-commissioned radial-velocity spectrograph SOPHIE at the Observatoire de Haute-Provence, we found that both objects exhibit reflex orbital radial-velocity variations with amplitudes characteristic of planetary-mass companions and in-phase with the photometric orbits. Line-bisector studies rule out faint blended binaries as the cause of either the radial-velocity variations or the transits. We perform preliminary spectral analyses of the host stars, which together with their radial-velocity variations and fits to the transit light curves, yield estimates of the planetary masses and radii. WASP-1b and WASP-2b have orbital periods of 2.52 and 2.15 days respectively. Given mass estimates for their F7V and K1V primaries we derive planet masses 0.80 to 0.98  and 0.81 to 0.95 times that of Jupiter respectively. WASP-1b appears to have an inflated radius of at least 1.33~$\\rmsub{R}{Jup}$, whereas WASP-2b has a radius in the range 0.65 to 1.26~$\\rmsub{R}{Jup}$. ",
        "introduction": "Extra-solar planets that transit their parent stars are of key interest because their masses and radii can be determined directly, providing clues to their internal compositions \\citep{guillot2006}. They define the mass-radius-separation relation for irradiated giant planets \\citep{mazeh2005}. They provide unique insights into their thermal properties (e.g. \\citealt{charbonneau2005}; \\citealt{deming2005}; \\citealt{deming2006}) and the chemical compositions of their atmospheres (\\citealt{charbonneau2002}; \\citealt{vidalmadjar2003}; \\citealt{vidalmadjar2004}).  The first exoplanet found to exhibit transits, HD 209458b (\\citealt{charbonneau2000}; \\citealt{henry2000}), was initially discovered using the radial-velocity method that has to date yielded the vast majority of the 210 known exoplanets. Its inflated radius \\citep{brown2001} presents a challenge to theories of the structure and evolution of irradiated exoplanets. New radial-velocity discoveries are routinely subjected to careful photometric followup at around the predicted times of transit. Indeed, recent radial-velocity surveys targeting bright stars of high metallicity with the specific goal of discovering new hot Jupiters  have revealed two new transiting planets in the last two years (\\citealt{sato2005}; \\citealt{bouchy2005}). Transiting planets thus comprise roughly 10\\% of all the hot Jupiters with orbital periods under 5 days or so, which is consistent with expectations for randomly-oriented orbits.  The complementary approach is to look for transits first, then to seek evidence of their planetary nature in radial-velocity followup observations. The OGLE project pioneered this approach, which has to date yielded five transiting exoplanet candidates for which radial-velocity variations have been detected with amplitudes indicative of planetary-mass companions (\\citealt{konacki2003}; \\citealt{bouchy2004}, \\citealt{konacki2004}; \\citealt{pont2004}; \\citealt{konacki2005}).  Several teams have embarked on ultra-wide field searches for transiting exoplanets. In this complementary approach, small-aperture CCD imaging systems coupled to commercial camera optics secure light curves of millions of stars. The TrES survey has so far yielded two transiting planets (\\citealt{alonso2004}; \\citealt{odonovan2006b}). The XO survey and HAT surveys have so far found one each (\\citealt{mccullough2006}; \\citealt{bakos2006}). Here we present the first two transiting exoplanets from the SuperWASP survey. The hardware, data analysis pipeline and archive methodology for this project are described in detail by  \\citet{pollacco2006}.  In a series of recent papers, \\citet{christian2006}, \\citet{street2006} and \\citet{lister2006} presented transit candidates from the inaugural 2004 May--September observing season of the five SuperWASP wide-field survey cameras that were operating at that time. The light curves of some $1.1\\times 10^6$ stars in the magnitude range $8<V<13$ were searched for periodic shallow transits, and several dozen plausible candidates were identified for detailed followup. The candidate selection methodology is described in detail in these papers and by \\citet{cameron2006}. ",
        "conclusions": "We  have detected the presence of radial-velocity variations in two exoplanetary transit candidates. Our preliminary analysis yields masses between 0.8 and 1.0 $\\rmsub{M}{Jup}$ for both planets. WASP-1b and WASP-2b lie between the ``hot Jupiters\" and the ``very hot Jupiters\" in the mass-orbital separation diagram for transiting hot Jupiters,   (Fig.~\\ref{fig:sepmass}).Their intermediate masses appear consistent with the general trend toward high masses at the smallest orbital separations noted by \\citet{mazeh2005}. WASP-2b in particular lies close to the minimum separation at which planets in this mass range survive, making it  a good candidate for future mass-loss studies. The radius of WASP-1b, which orbits an F7V star, is poorly constrained by the SuperWASP data alone, but appears to be at least 1.33~$\\rmsub{R}{Jup}$. WASP-1b seems thus to be an expanded, low-density planet, similar to HD 209458b and HAT-P-1b \\citep{bakos2006}. Additional photometry of WASP-2b, which orbits a K1 dwarf at a slightly smaller orbital separation yields a radius close to that of Jupiter, suggesting a substantially higher density. Additional high-precision photometry is needed to refine the radii and densities of both planets; indeed, precise radius estimates have recently been derived from high-precision photometry by \\citet{shporer2006} and \\citet{charbonneau2006}, confirming the oversized nature of WASP-1b."
    },
    "0609/astro-ph0609477_arXiv.txt": {
        "abstract": "We examine the luminosity and dynamical mass estimates for young massive stellar clusters.  For many young ($<50$~Myr) clusters, the luminosity and dynamical mass estimates differ by a significant amount.  We explain this as being due to many young clusters being out of virial equilibrium (which is assumed in dynamical mass estimates) because the clusters are undergoing violent relaxation after expelling gas not used in star formation.  We show that, if we assume that luminous mass estimates are correct (for a standard IMF), at least 50~per cent of young clusters for which dynamical masses are known are likely to  be destroyed within a few 10s~Myr of their formation.  Even clusters which will retain a bound core may lose a large fraction of their stellar mass.  We also show that  the core radius and other structural parameters change significantly during the violent relaxation that follows gas expulsion and that they should be considered instantaneous values only, not necessarily reflecting the final state of the cluster. In particular we note that the increasing core radii observed in young LMC/SMC clusters can be well explained as an effect of rapid gas loss. ",
        "introduction": "Almost all stars form in star clusters which have a mass function of the form $N(M)  \\propto M^{-2}$ (Lada \\& Lada 2003).  This means that most clusters are fairly small ($10^2$ -- $10^3 M_\\odot$), however some of the most interesting young clusters are the rare, massive `super star clusters' (SSCs) with masses comparable to, or greater than, the Galactic globular clusters.  SSCs are found in merging and interacting galaxies such as NGC~7252 (Miller et al.~1997) and NGC~6745  (de Grijs et al.~2003)\\footnote{See Larsen~(2004) for an extensive list of starburst, merger, spiral, dwarf, and irregular galaxies which contain populations of young massive clusters.}, as well as the LMC (Elson, Fall \\& Freeman 1989; Elson 1991; Mackay \\& Gilmore 2003), and even within our own Galaxy (esp. Wd~1, see Clark et al. 2005).  Star clusters form within giant molecular clouds and remain embedded in their natal clouds for $\\sim 2$ -- $5$~Myr before the combined effect of massive stellar winds and supernovae drive-out the gas not used in star formation.  This gas represents $\\sim 50$~\\% of the initial mass of the cluster and its rapid loss (on less than a crossing time, Goodwin 1997a; Melioli \\& de Gouveia dal Pino 2006) causes a huge change to the potential of the cluster.  The stars left behind after gas expulsion violently relax to the new potential and attempt to return to virial equilibrium.  The effect of gas expulsion on the (stellar) dynamics of star clusters is well understood theoretically (e.g. Tutkov 1978; Hills 1980; Goodwin 1997a,b; Boily \\& Kroupa 2003a,b).  The stars have too large a velocity dispersion for their new potential and the cluster expands, and stars with a velocity greater than the new escape velocity are lost.  The loss of stars is not instantaneous, however, and escaping stars remain physically associated with the cluster for 10s~Myr.  Recently Bastian \\& Goodwin (2006, hereafter BG06) have shown that a number of massive young clusters show the signature of gas expulsion: an excess of light at large radii over and above a relaxed (King or Elson-Fall-Freeman (EFF)) surface brightness profile.  They note that such signatures are also seen in many young LMC clusters (see Elson et al. 1989; Elson 1991; Mackay \\& Gilmore 2003).  BG06 show that a consequence of gas expulsion is that many young clusters {\\em are not in virial equilibrium}: thus dynamical mass estimates based on virial equilibrium may be highly inaccurate.  In particular, claims of unusual IMFs based on dynamical masses being inconsistent with luminous mass estimates are probably incorrect.  In this paper we take the approach that luminosity mass estimates are correct and the discrepancy between luminosity and dynamical masses can be used as an indicator of the extent to which young clusters are out of virial equilibrium.  This may then be used to estimate the star formation efficiency (SFE; see below for details)  of young clusters and their survivability.  Additionally, we will show that rapid gas removal can explain the trend of increasing core radii with age observed in young LMC clusters.  This paper is organised as follows; in \\S~\\ref{sec:sims} we introduce the simulations of rapid gas loss in cluster and address the caveats in the current study. In \\S~\\ref{sec:results} we present the results of the simulations placing particular emphasis on the dynamical state of young clusters and on the effects of ``infant-weight loss'' and ``infant mortality''.  We summarise our results in \\S~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have simulated the evolution of massive young star clusters that are far-out of virial equilibrium due to the expulsion of the residual gas left-over after star formation.  Our results are in agreement with previous simulations of this phase (e.g. Lada et al. 1984; Goodwin 1997a,b; Geyer \\& Burkert 2001; Kroupa et al. 2001; Kroupa \\& Boily 2002; Boily \\& Kroupa 2003a,b).  We parameterise the initial state of a cluster with an effective star formation efficiency (eSFE), $\\epsilon$, where the initial virial ratio of a cluster is $Q = 1/2\\epsilon$ (we define virial equilibrium to be $Q=0.5$).  We explore clusters with eSFEs between 10 and 60 per cent.  Clusters expand after gas expulsion in an attempt to reach a new equilibrium.  If eSFE $\\leq 30$ per cent then the clusters are destroyed within a few 10s~of Myr, for higher eSFEs a bound core remains, but there may still be a significant degree of (stellar) mass loss and the final cluster may only be a small fraction of its initial mass.  Those stars that escape cause an excess of light at large radii that is observed in many young clusters (BG06).  During its expansion a cluster is not in virial equilibrium.  Its velocity dispersion retains a memory of the initial (gas plus stars) mass of the cluster and so dynamical mass estimates -- which assume virial equilibrium -- may be wrong by a very significant factor.  We also note that during the expansion phase cluster parameters, such as the core radius, may change significantly on a short timescale.  In particular they may not change in a linear fashion.  The core radii of intermediate eSFE clusters measured when escaping stars are still physically associated with the cluster may be significantly greater than the `final' (relaxed) core radii (see also Goodwin 1997b). Many properties are {\\em instantaneous} and it may be difficult to extrapolate to future values without knowing the initial conditions.  This includes the mass of clusters which can change significantly as stars are lost (see also Kroupa \\& Boily 2002).  We compare the dynamical (assuming virial equilibrium) and luminous (assuming a standard Kroupa IMF) masses for a large sample of young clusters (limited by the number of clusters for which reliable values for both parameters are available).  As noted by Bastian et al. (2006), young ($<50$~Myr) clusters tend to have too much dynamical mass for their luminosity, whilst older clusters have dynamical masses that match their expected luminosity for a Kroupa IMF.  We suggest that the discrepancy between the dynamical and luminous masses for young clusters is due to the incorrect assumption of virial equilibrium.  We assume that young clusters -- the same as old clusters -- have a standard IMF and use the discrepancy between the mass estimates to calculate their eSFEs which appear to fall between $20$ and $\\sim 100$ per cent (an upper limit is very difficult to establish as for eSFEs$>60$ per cent the effect of gas expulsion is difficult to observe).  This wide range of eSFEs may be due to a real difference in the star formation efficiencies of different clusters, or to differences in the initial equilibrium between the stars and gas (or, most likely, a combination of the two).  At least $50$ per cent of the young clusters have eSFEs $<30$ -- $40$ per cent and so we expect the infant mortality rate in this sample to be {\\em at least} $50$ per cent (as we ignore further destructive effects such as stellar evolutionary mass loss and tidal fields). Further, there is an observational bias that means that low surface brightness -- hence more extended, or more expanded -- clusters which are likely to be destroyed are not observed.  In summary, young clusters are most likely to be re-virialising after gas expulsion.  This means that (a) dynamical mass estimates may be significantly wrong, and (b) measured parameters (including mass) are instantaneous values and may not reflect the `final' values.  Assuming that luminosity-derived masses for young clusters are correct allows us to determine the effective star formation efficiencies of young clusters with dynamical mass estimates and we find that at least $50$ per cent of young clusters in our sample are unlikely to survive for more than a few 10s~Myr.  The true infant mortality rate is probably much higher than this."
    },
    "0609/astro-ph0609194_arXiv.txt": {
        "abstract": "Blank field sources (BFS)  are good candidates for hosting dim isolated neutron stars (DINS). The results of a search of BFS in the ROSAT HRI images are revised. We then focus on transient BFS, arguing that they belong to a rather large population. The perspectives of future research on DINS are then discussed. ",
        "introduction": "\\label{intro} Isolated neutron stars, which have overcome the pulsar phase are elusive sources. In principle they can shine from some residual internal energy (coolers), or because of their interaction with the interstellar medium (e.g. accretors). Their number in the Galaxy should be very high, about one percent of the total number of stars. Their significance as a population is of the utmost interest: they are the end-point of the evolution of a vast class of stars.  It is just because of these considerations that the discovery of dim isolated neutron stars by the ROSAT satellite has been a major achievement (see  e.g. Treves et al. 2000, Haberl 2005, Zane et al. 2005).  DINS are one of the main attractions of this meeting, see in particular the presentations by Cropper and Popov. Their properties can be summarized as follows: softness $T\\sim100\\,$eV; closeness $d\\sim100\\,$pc; extremely dim optical counterparts ($V>25$), periodicities of $5-10\\,$s; absorption features below 1 keV. The seven DINS discovered thus far are probably a mixed bag, in the sense that the above properties may not appear all together. For instance the prototype of the class 1856--37, has a perfect black body spectrum in the $X$-ray band with neither absorption lines  nor indications of pulsations. Most likely they are all coolers (Neuh\\\"auser \\& Tr\\\"umper 1999, Popov et al.~2002).  In order to further improve our knowledge of DINS it is mandatory to enlarge their sample. The procedure followed up to now to discover new DINS has been to search the ROSAT images for the so called ``Blank Field Sources'' (BFS, Cagnoni et al. 2002), i.e. $X$-ray sources without counterparts in other spectral bands, and use the properties listed above to argue that the candidate belongs to the class. This line was pursued by a number of authors, we mention in particular the recent paper by Agueros et al. (2006) where the ROSAT PSPC images (RASS) are compared with the Sloan Digital Sky Survey.  Here we focus on progresses of our search of BFS in the ROSAT HRI images (Chieregato et al. 2005), concentrating on the possible detection of transient BFS. ",
        "conclusions": "\\label{sec:4}  In the fifteen years of research about DINS the progress was remarkable, yet there are two basic points that have not yet been achieved:  \\begin{itemize} \\item There is still no example of a DINS which is convincingly powered by the accretion of the interstellar medium. While there are a number of arguments indicating that these objects should be much rarer than originally estimated (Treves \\& Colpi 1991, Blaes \\& Madau 1993, Perna et  al.~2003 and references therein), these objects should finally show up, and their emission should be largely independent of the neutron star age. Neutron stars as old as the Galaxy could be a part of the lot. \\item We have not yet any information on the cousins of DINS, i.e. isolated black holes (see e.g. Agol \\& \\linebreak Kamionkowski 2002). \\end{itemize}  The challenge for the future is obviously to explore the sky for DINS and their cousins, at a flux threshold which is an order of magnitude lower than that of ROSAT, and which is easily accessible to present generation $X$-ray telescopes. There is no doubt that the activity up to now has been rather slow , since rich $X$-ray and optical archives are already available.  Let us summarize the hopes for the future, which can derive from a thorough study of the existing data:  \\begin{itemize} \\item discovery of accretion fed DINS; \\item establishing or excluding the existence of transient BFS; \\item discovery of isolated black holes; \\item discovery of intermediate mass black holes possibly related to ultraluminous X-ray sources (e.g. Mapelli, Ferrara \\& Rea 2006). \\end{itemize}"
    },
    "0609/astro-ph0609707_arXiv.txt": {
        "abstract": "In this paper the N$-$S asymmetry of the soft X-ray flare index ($FI_{SXR}$) during the solar cycles 21, 22 and 23 has been analyzed. The results show the existence of a real N$-$S asymmetry which is strengthened during solar minimum. The slope of the regression lines fitted to the daily values of asymmetry time series has been found to be negative in all the three cycles. The yearly asymmetry curve can be fitted by a sinusoidal function with a period of eleven years. The power spectral analysis of daily asymmetry time series reveals the significant periods of around 28.26 days, 550.73 days and 3.72 years. ",
        "introduction": "It is well known that many types of solar phenomena exhibit some form of north-south (N$-$S) asymmetry. Bell (1962) found a long term asymmetry in the sunspot area data for cycle 8 through 18. Roy (1977) studied the N$-$S distribution of major solar flares from 1955 to 1974, covering cycles 19 and 20, and found an asymmetry in favour of northern hemisphere which is more pronounced during the minimum of solar cycle. Kno\\v{s}ka (1985) studied the asymmetry of flare activity during the years 1937$-$1978 using the H$\\alpha$ flare index ($FI_{H\\alpha}$) introduced by Kleczek (1952, 1953) and found no unique relationship between the asymmetry of flare activity and eleven year solar cycle. Vizoso and Ballester (1987) have studied the sudden disappearance (SD) of solar prominence during cycles 18$-$21 and found that the N$-$S asymmetry curve is not in phase with the solar cycle and peaks at or around the solar minimum and that the asymmetry changes its sign during the maximum, at the time of reversal of the Sun's general magnetic field. Verma (1993) reported a study of various solar phenomena occurring in both northern and southern hemisphere of the Sun for solar cycles 8$-$22. Joshi (1995); Li, Schimider, and Li (1998) and Ata\\c{c} and \\\"{O}zg\\\"{u}\\c{c} (1998) studied the N$-$S asymmetry during the solar cycle 22 with different manifestations of solar activity. In all these studies the asymmetry during this cycle was found to be in favour of southern hemisphere. Ata\\c{c} and \\\"{O}zg\\\"{u}\\c{c} (2001) studied the N$-$S asymmetry of $FI_{H\\alpha}$ during the rising phase of solar cycle 23 (1996$-$1999) and found that the activity prefers the northern hemisphere in general during the present cycle. Recently Li et al. (2002) have presented the dominant hemisphere of solar activity in each of solar cycle 12 to 22 by calculating the actual probability of the hemisphere distribution of sunspot groups, sunspot areas, the relative number of sunspots, flare index and number of SDs.  Studies of periodicity in time series of N$-$S asymmetry data have also been made by several authors. Vizoso and Ballester (1989) found a significant peak of around 3.1$-$3.2 years in the power spectrum performed with the values of flare number and flare index N$-$S asymmetry. Carbonell, Oliver, and Ballester (1993) and Oliver and Ballester (1994) studied the periodicity in the N$-$S asymmetry of sunspot area from 1874$-$1989 and 1874$-$1993 respectively.  They obtained statistically significant period of 163.93 rotations (12.1 years). \\\"{O}zg\\\"{u}\\c{c} and Ata\\c{c} (1996) performed the power spectral analysis of N$-$S asymmetry time series of flare index values of solar cycle 22 and confirmed the 25.5 day fundamental period of the sun which was discovered by Bai and Sturrok (1991). Taking all this into account, here we have made an attempt to investigate the N$-$S asymmetry of soft X-ray flare index ($FI_{SXR}$), which is based on the continuous record of SXR flares observed by GOES during solar cycles 21, 22 and 23. The daily values of $FI_{SXR}$ include the contribution of SXR flares of different classes (B to X) and are suitable for a short and intermediate term solar activity analysis. In the paper we have also searched for the periodicity in the daily asymmetry time series of $FI_{SXR}$ for the period 1975$-$2003.  \\begin{figure} \\epsfig{file=fig1.ps,width=8.0cm, height=12.0cm,angle=270} \\vspace{0.5cm} \\caption{Plot of the monthly values of $FI_{SXR}$ for the period 1975$-$2003. Dashed line shows the smooth spline curve.} \\end{figure} ",
        "conclusions": "In this paper we have calculated the daily $FI_{SXR}$ for solar cycles 21 to 23. The plot (Figure 1) of monthly values of $FI_{SXR}$ indicates a lower level of activity during solar cycle 23 compared to previous two cycles. This has been pointed out earlier by \\\"{O}zg\\\"{u}\\c{c} et al. (2002) in his study during the rising phase of solar cycle 23. The scatter diagram (Figure 2) between $FI_{SXR}$ and $FI_{H\\alpha}$ shows that there exists a good correlation between these two solar activity indices. Ata\\c{c} and \\\"{O}zg\\\"{u}\\c{c} (1989, 1998, 2003) in their various papers have made extensive statistical analysis of $FI_{H\\alpha}$, finding its good correlation with a number of solar indices that arise under different physical conditions. Therefore, a good correlation of $FI_{SXR}$ with $FI_{H\\alpha}$ also makes the former an interesting parameter which is good enough to describe the solar activity.  We have found that there exists a real N$-$S asymmetry in $FI_{SXR}$ which is not due to random fluctuations. The fitted straight line to asymmetry time series (Figure 3) for the cycles 21 and 23 shows that the activity in the northern hemisphere is more important during the ascending branch of cycle whereas during the descending branch the activity becomes more important in the southern hemisphere. In the cycle 22 there was a southern dominance throughout. The slope of regression lines for all the three cycles is found to be negative. Vizoso and Ballester (1990), in his study of N$-$S asymmetry of sunspot area during solar cycles 12$-$20 and sunspot number for cycle 21, found that the slope of regression lines fitted to the yearly values of N$-$S asymmetry change its sign each four cycles suggesting a long term periodic behaviour in the N$-$S asymmetry of around eight cycles. They also found that the slope of regression lines has changed from positive to negative values after cycle 19. Ata\\c{c} and \\\"{O}zg\\\"{u}\\c{c} (1996) made similar conclusions with the N$-$S asymmetry of $FI_{H\\alpha}$ during cycles 17$-$22. Our study confirms and extends these earlier results. Moreover, Ata\\c{c} and \\\"{O}zg\\\"{u}\\c{c} (1996) suggested that, if the period of eight cycles is real, the behaviour of asymmetry during cycles 20$-$22 should be maintained in cycle 23, which has also been confirmed by the present study.  On the other hand, if we consider the plot of yearly N$-$S asymmetry index (Figure 4) the change in the predominance of activity in the two hemisphere can be studied for every year starting from 1976. Initially during the solar cycle 21, the northern hemisphere was more active but after 1980 it was shifted slightly towards south. In the year 1983 the activity moved towards south strongly and prevailed there during most of the solar cycle 22. In the solar cycle 23, we observe a strong peak in southern hemisphere during the minimum phase in the year 1996 and then the activity is moved towards the north. In the years 1998, 1999 and 2000 the northern hemisphere was dominated. These results are in agreement with the work done by Temmer et al. (2001), who analyzed the N$-$S asymmetry of $H\\alpha$ flares from 1975 to 1999. The preference for northern hemisphere during the rising phase of cycle 23 is also reported by Ata\\c{c} and \\\"{O}zg\\\"{u}\\c{c} (2001) with the data of $FI_{H\\alpha}$. Here we also notice that, during cycle 23, the behaviour of asymmetry is changed in the year 2000 and it has dominated towards southern hemisphere in the successive years. The Figure also reveals that the fitted sinusoidal curve with a period of eleven years reflects the behaviour of asymmetry in a good manner. This suggests that N$-$S asymmetry of $FI_{SXR}$ has a periodic behaviour shifted in phase with respect to solar cycle. The shape of the curve also indicates that the asymmetry has peaked at or around the minimum of solar activity. This result is complementary to the study of N$-$S asymmetry made by several authors (Swinson, Koyama and Saito, 1986; Vizoso and Ballester, 1990; Joshi, 1995; Ata\\c{c} and \\\"{O}zg\\\"{u}\\c{c}, 1996).  Finally, we have studied periodicities in the daily asymmetry time series for three frequency ranges. In the power spectrum for the frequency range of 116.0$-$463.0 nHz (Figure 5), we get a significant peak at 409.5 nHz (28.26 days). It is close to the 27 days solar rotational period which is well known in different solar indices. Antonucci et al. (1990) investigated the rotation of photospheric magnetic field during solar cycle 21, and obtained a dominant period of 26.9 days for the northern and 28.1 days for the southern hemisphere. A similar outcome is obtained by Temmer et al. (2002) from the power spectral analysis of daily sunspot numbers during 1975$-$2000 in which they derive a rigid rotation with 27.0 days for the northern hemisphere, while the southern hemisphere reveals a dominant period of 28.2 days. Pap, Tobiska, and Bouwer (1990) and Bouwer (1992) reported periods between 26$-$28 days from the power spectra of various solar indices. Joshi (1999) reported peaks at 26.5 days and 28.3 days in the power spectrum of solar radio flux at 10.7 cm. Lou et al. (2003) found a prominent peak of around 28 days in the daily averaged Ap index for geomagnetic disturbances, which is physically identified with solar rotation. The power spectrum for the 16.5$-$116.0 nHz range (Figure 6) shows a significant peak at 21.02 nHz (550.73 days). Ichimoto et al. (1985) reported that a peak in the power spectra appears near 510$-$540 days when northern and southern hemisphere flares are analyzed separately. Bai (1987) found an important peak at 552 days in the power spectra of major flares belonging to northern hemisphere. A significant peak of around 20 rotations (about 540 days) is found by Akioka et al. (1987) in the power spectra for TA (Sum of maximum areas of sunspot groups per solar rotation) and MA (mean area of sunspot groups per solar rotations) during solar cycle 21. \\\"{O}zg\\\"{u}\\c{c} and Ata\\c{c} (1989) reported a peak at 564 days in the $FI_{H\\alpha}$ during cycle 20 and 21 considered together. Oliver, Carbonell, and Ballester (1992) found the periodicities around 540 days in the cycles 12, 14 and 17 in sunspot areas, while during cycles 18 and 19 it is present, with a very high significance, in sunspot and Zurich sunspot number. They also found a peak at 528 days in the periodogram for sunspot areas during cycles 12 to 21. In the power spectra for the frequency range of 7.7$-$16.5 nHz (Figure 7), we have a significant peak at 8.5 nHz (3.72 years). Rao (1973) studied periodicities in several indices of solar activity and found that the index P (average area of all spot groups), $T_{0}$ (average spot group life time) and $f_{0}$ (number of all spot groups formed all over the Sun within a unit time) presented periodicities around 3.9, 3.5 and 3.48 years, respectively. Vizoso and Ballester (1989) found a periodicity of around 3.1$-$3.2 years in the values of $FI_{H\\alpha}$ and flare numbers N$-$S asymmetry. A periodicity of 3.27 years is reported by Vizoso and Ballester (1990) in the N$-$S asymmetry of sunspot areas. In all the three power spectra presented by us, the subtraction of sine curve indicates that the peaks are probably real and not spurious or due to aliasing (Figure 8). We, therefore, see that the periodicities in the asymmetry time series of $FI_{SXR}$ is consistent with the earlier results."
    },
    "0609/astro-ph0609531_arXiv.txt": {
        "abstract": "An exact model is proposed for a gray, isotropically scattering planetary atmosphere in radiative equilibrium. The slab is illuminated on one side by a collimated beam and is bounded on the other side by an emitting and partially reflecting ground. We provide expressions for the incident and reflected fluxes on both boundary surfaces, as well as the temperature of the ground and the temperature distribution in the atmosphere, assuming the latter to be in local thermodynamic equilibrium. Tables and curves of the temperature distribution are included for various values of the optical thickness. Finally, semi-infinite atmospheres illuminated from the outside or by sources at infinity will be dealt with. ",
        "introduction": "The modern phase in modeling radiation transfer through planetary atmospheres started with the seminal work of Humphreys \\cite{humphreys1909}, and the contemporary elaborate attempt of Gold \\cite{gold1909}, followed by the critical reappraisal of both attempts by Emden \\cite{emden1913}. Of these works, Emden's is directly based on previous work by Schuster \\cite{schuster1905} and Schwarzschild \\cite{schwarzschild1906}, which, with Schwarzschild's second paper \\cite{schwarzschild1914}, laid the foundation for the theory of radiation transfer. While those three papers may doubtlessly be regarded as ground-breaking, they necessarily remained rudimentary from our present vantage point in radiative transfer theory. It was not until the next decade that Milne \\cite{milne1922} would deal with radiative energy transfer through planetary atmospheres in a mathematically rigorous manner. His important developments would not be superseded before the second half of the twentieth century, with the contributions by King \\cite{king1963}, Ostriker \\cite{ostriker1963}, Wildt \\cite{wildt1966}, Schultis and Kaper \\cite{schultisandkaper1969}, Stibbs \\cite{stibbs1971} and Barkstrom \\cite{barkstrom1974}, to name but those authors who attempted to model exactly the transfer through gray atmospheres.\\\\ By gray atmosphere we understand, as is customary, an atmosphere whose opacity does not depend on the frequency of the electromagnetic radiation involved. A semi-gray atmosphere is, on the other hand, one in which the opacity assumes just two different values in adjacent frequency intervals, conveniently distinguished as the ``visible'' and the ``infrared'' regions of the spectrum. It was Emden \\cite{emden1913} who established this distinction, which, besides being more realistic, is susceptible to analytical treatment. The present paper is concerned with gray atmospheres in radiative equilibrium, while two further papers will be devoted to the case of semi-gray atmospheres.\\\\ In defining the model features, we shall assume the usual ingredients of the classical theory of planetary atmospheres to hold true: a plane-parallel atmosphere, subject to both illumination by a parallel beam of light impinging on one of its boundaries and to radiative heating through the other boundary, which is assumed contiguous to a partially reflecting and gray-emitting ground. Our purpose is to take the analytical developments as far as possible, under the assumption of a gray or semi-gray atmosphere in radiative equilibrium. We thereby wish to push the line of investigations recalled above farther, drawing from the theory of radiation transfer as it evolved in the last three decades. Our goal is to develop elementary but rigorous models which may be used not only to validate more realistic models based on numerical resources, but also to foster the understanding of radiative transfer in more comprehensive models of, say, the climates of planets, where it plays an important part in long-term evolutions.\\\\ The paper may be outlined as follows. Section 2 introduces the standard problem of planetary atmospheres. In Section 3, the fundamental integral equation underlying this problem is solved. Then follows the calculation of radiative flux densities at both boundary planes (Section 4), which enables the evaluation of the ground temperature (Section 5) and the atmospheric temperature (Section 6) of the planet. Next, in Section 7, we calculate the radiance of radiation leaving the atmosphere, as it represents the fundamental quantity in solving the inverse problem of the theory of planetary atmospheres. In the subsequent section, we deal with semi-infinite atmospheres heated either from the outside or from within by sources at infinity. Finally, we conclude the work with a summary, addressing at the same time some issues to be taken into account in forthcoming work. ",
        "conclusions": "The generalized problem we have outlined originated with two prototype problems of the classical theory of planetary atmospheres. The first one is of relevance to situations where shortwave radiation interacts with the planet's atmosphere, modeled as a slab under parallel-beam illumination from the sun and overlying a partly reflecting ground. In the second problem, there is no incident radiation on the upper surface of the atmosphere but only on the lower one, issuing from an isotropically emitting ground, a model appropriate to radiation transfer through a planetary atmosphere in the thermal (or infrared) regime. In a strictly gray model, in which both the sun's radiation and the planet's ground emission are scattered and absorbed by the planet's atmosphere, the combined boundary conditions prevail simultaneously, providing us with the model that we offer in mathematical elaboration in the present article.\\\\ Of course, this physically rather unsophisticated model is not realistic, but in exchange it is susceptible to analytical tractability. It yields the radiance field in a semi-infinite atmosphere illuminated from the outside or heated by sources at infinity, for which the analytical developement could be pushed as far as it goes. In the case of an optically finite atmosphere, we were able to calculate exactly the source function for each depth, as well as the upward and downward fluxes across the boundaries. As a result, it was possible to determine the temperature of the atmosphere as well as that of the ground.\\\\ Models more realistic than the one here presented may inspire further developments and extensions. In particular, we think of the possibility of applying each set of boundary conditions to two separate frequency domains in such a manner that the radiation interactions remain coupled by the radiative equilibrium constraint. Models of that kind were proposed early, notably by Emden \\cite{emden1913}, and Milne \\cite{milne1922}, who realized that the angular variable $\\mu_0$ of the gray models could be extended to the range $]1,+\\infty[$ in what is sometimes termed semi-gray models. This opens up the possibility of applying the results for gray media to those of semi-gray media, which represents a first step towards more realistic models.\\\\ Semi-gray models are thus seen to constitute the natural next step in a hierarchy of analytical models of radiative transfer within planetary atmospheres. We intend to propose such a semi-gray model in two companion papers. The present study can therefore be judged as an intermediate step in the application of the exact solution to the thermal radiation problem, as given recently in \\cite{rutilyetal2004}, and the forthcoming semi-gray model.   \\appendix \\renewcommand{\\theequation}{\\Alph{section}\\arabic{equation}} \\setcounter{equation}{0} \\setcounter{section}{1}"
    },
    "0609/astro-ph0609257_arXiv.txt": {
        "abstract": "We give a brief overview of the current experimental and theoretical status of cosmic rays above $\\sim10^{17}\\,$eV. We focus on the role of large scale magnetic fields and on multi-messenger aspects linking charged cosmic ray with secondary $\\gamma-$ray and neutrino fluxes. ",
        "introduction": "High energy cosmic ray (CR) particles are shielded by Earth's atmosphere and reveal their existence on the ground only by indirect effects such as ionization and showers of secondary charged particles covering areas up to many km$^2$ for the highest energy particles. In fact, in 1912 Victor Hess discovered CRs by measuring ionization from a balloon~\\cite{hess}, and in 1938 Pierre Auger proved the existence of extensive air showers (EAS) caused by primary particles with energies above $10^{15}\\,$eV by simultaneously observing the arrival of secondary particles in Geiger counters many meters apart~\\cite{auger_disc}.  After almost 90 years of research, the origin of cosmic rays is still an open question, with a degree of uncertainty increasing with energy~\\cite{crbook}: Only below 100 MeV kinetic energy, where the solar wind shields protons coming from outside the solar system, the sun must give rise to the observed proton flux. Above that energy the CR spectrum exhibits little structure and is approximated by broken power laws $\\propto E^{-\\gamma}$: At the energy $E\\simeq4\\times 10^{15}\\,$eV called the ``knee'', the flux of particles per area, time, solid angle, and energy steepens from a power law index $\\gamma\\simeq2.7$ to one of index $\\simeq3.0$. The bulk of the CRs up to at least that energy is believed to originate within the Milky Way Galaxy, typically by shock acceleration in supernova remnants. The spectrum continues with a further steepening to $\\gamma\\simeq3.3$ at $E\\simeq4\\times 10^{17}\\,$eV, sometimes called the ``second knee''. There are experimental indications that the chemical composition changes from light, mostly protons, at the knee to domination by iron and even heavier nuclei at the second knee~\\cite{Hoerandel:2004gv}. This is in fact expected in any scenario where acceleration and propagation is due to magnetic fields whose effects only depend on rigidity, the ratio of charge to rest mass, $Z/A$. This is true as long as energy losses and interaction effects, which in general depend on $Z$ and $A$ separately, are small, as is the case in the Galaxy, in contrast to extra-galactic cosmic ray propagation at ultra-high energy. Above the so called ``ankle'' or ``dip'' at $E\\simeq5\\times10^{18}\\,$eV, the spectrum flattens again to a power law of index $\\gamma\\simeq2.8$. This latter feature is often interpreted as a cross over from a Galactic component, which steepens because cosmic rays are not confined by the Galactic magnetic field any more or because Galactic sources do not accelerate beyond the ankle, to a harder component of extragalactic origin. However, the dip at $E\\simeq5\\times10^{18}\\,$eV could also be explained by pair production by extra-galactic protons, if the extra-galactic component already starts to dominate below the ankle, for example, around the second-knee~\\cite{Aloisio:2006wv} at a few times $10^{17}\\,$eV. This requires a relatively steep injection spectrum $\\propto E^{-2.6-2.7}$. Below a few times $10^{17}\\,$eV this extra-galactic component would become unobservable at Earth due to diffusion in extra-galactic magnetic fields (EGMF)~\\cite{Lemoine:2004uw}. In addition, the effective volume-averaged injection spectrum has to become flatter somewhere below $\\sim10^{18}\\,$eV in order to avoid excessive power going into cosmic rays and to avoid overproduction of GeV--TeV $\\gamma-$rays from $pp$ interactions with the ambient gas.  The low cross-over scenario also requires the dominance of protons around the dip. Theoretically, this can be achieved either because preferentially protons are accelerated or because extended EGMF lead to strong photo-spallation during propagation~\\cite{Sigl:2005md}. Experimentally, above $\\simeq10^{17}\\,$eV the chemical composition is basically unknown~\\cite{Watson:2004ew}. Around $10^{18}\\,$eV the situation is particularly inconclusive as HiRes~\\cite{Abbasi:2004nz} and HiRes-MIA~\\cite{Abu-Zayyad:2000ay} data suggest a light (proton dominated) composition, whereas other experiments indicate a heavy composition~\\cite{Hoerandel:2004gv}. In any case, the cosmic ray flux should be extra-galactic at least above the ankle, where a galactic origin would predict an anisotropy toward the galactic plane because galactic magnetic fields can no longer isotropize the cosmic rays. No such anisotropy is seen. There are also experimental indications for a chemical composition becoming again lighter above the ankle, although a significant heavy component is not excluded and the inferred chemical composition above $\\sim10^{18}\\,$eV is sensitive to the model of air shower interactions and consequently uncertain presently~\\cite{Watson:2004ew}. In addition, should a substantial heavy composition be experimentally observed up to the highest energies, some sources would have to be surprisingly nearby, within a few Mpc, otherwise only low mass spallation products would survive propagation~\\cite{er}. In the following we will restrict our discussion on extra-galactic ultra-high energy cosmic rays (UHECRs).  Although statistically meaningful information about the UHECR energy spectrum and arrival direction distribution has been accumulated~\\cite{Cronin:2004ye}, no conclusive picture for the nature and distribution of the sources emerges naturally from the data. There is on the one hand the approximate isotropic arrival direction distribution~\\cite{bm} which indicates that we are observing a large number of weak or distant sources. On the other hand, there are also indications which point more towards a small number of local and therefore bright sources, especially at the highest energies: First, the AGASA ground array claimed statistically significant multi-plets of events from the same directions within a few degrees~\\cite{Shinozaki:2006kk,bm}, although this is controversial~\\cite{Finley:2003ur} and has not been seen so far by other experiments such as the fluorescence experiment HiRes~\\cite{finley}. The spectrum of this clustered component is $\\propto E^{-1.8}$ and thus much harder than the total spectrum~\\cite{Shinozaki:2006kk}. Second, nucleons above $\\simeq70\\,$EeV suffer heavy energy losses due to photo-pion production on the cosmic microwave background (CMB) --- the Greisen-Zatsepin-Kuzmin (GZK) effect~\\cite{gzk} --- which limits the distance to possible sources to less than $\\simeq100\\,$Mpc~\\cite{stecker}. This predicts a ``GZK cutoff'', a drop in the spectrum, whose strength depends on the source distribution and may even depend on the part of the sky one is looking at: The ``cutoff'' could be mitigated in the northern hemisphere where more nearby accelerators related to the local supercluster can be expected. Apart from the SUGAR array which was active from 1968 until 1979 in Australia, all UHECR detectors completed up to the present were situated in the northern hemisphere. Nevertheless the situation is unclear even there: Whereas a ``cut-off'' is consistent with the few events above $10^{20}\\,$eV recorded by the fluorescence detector HiRes~\\cite{hires}, there is a tension with the 11 events above $10^{20}\\,$eV detected by the AGASA ground array~\\cite{agasa}. Still, this could be a combination of statistical and systematic effects~\\cite{DeMarco:2003ig}, especially given the recent downward revision of the energy normalization in AGASA~\\cite{teshima}. The solution of this problem will have to await more analysis and more statistics and, in particular, the completion of the Pierre Auger project~\\cite{Kampert:2006ec} which will combine the two complementary detection techniques adopted by the aforementioned experiments and whose southern site is currently in construction in Argentina. ",
        "conclusions": ""
    },
    "0609/astro-ph0609241.txt": {
        "abstract": "We perform a study of the ultra high energy neutrino detection performances of a km$^3$ Neutrino Telescope sitting at the three proposed sites for \\verb\"ANTARES\", \\verb\"NEMO\" and \\verb\"NESTOR\" in the Mediterranean sea. We focus on the effect of the underwater surface profile on the total amount of yearly expected $\\tau$ and $\\mu$ crossing the fiducial volume in the limit of full detection efficiency and energy resolution. We also emphasize the possible enhancement of matter effect by a suitable choice of the geometry of the Telescope. ",
        "introduction": "Neutrinos are one of the main components of the cosmic radiation in the ultra-high energy (UHE) regime.  Although their fluxes are uncertain and depend on the production mechanism, their detection can provide information on the sources and origin of the UHE cosmic rays. For example, UHE neutrinos can be produced via $\\pi$-photo\\-production by strongly accelerated hadrons in presence of a background electromagnetic field.  This scenario is expected to occur in extreme astrophysical environments like the jets of active galactic nuclei, radio galaxies and gamma ray burst sources as well as in the propagation of UHE nucleons scattering off the cosmic background radiation (known as {\\it cosmogenic neutrinos} \\cite{[1],[2]}).  {}From the experimental point of view, after the first pioneering and successful achievements, neutrino astronomy in the high energy regime \\cite{[3],[4],[5],[6],[7]} is a rapidly developing field, with a new generation of neutrino telescopes on the way. A benchmark result was obtained by the \\verb\"DUMAND\" \\cite{[8]} collaboration, followed by the successful deployments of \\verb\"NT-200\" at Lake Baikal \\cite{[9]} and \\verb\"AMANDA\" \\cite{[10]} at the South Pole, which have shown the feasibility of large optical Cherenkov neutrino telescopes (NT) in open media like sea- or lake-water and glacial ice. These experiments observed atmospheric neutrinos \\cite{[11]} and set bounds on their extraterrestrial flux \\cite{[12],[13],[14]} which are much more constraining than the corresponding bounds obtained by underground neutrino detectors \\cite{[15]}. These interesting results and the perspective to perform astronomical studies using UHE neutrinos stimulated several proposals and R\\&D projects for neutrino telescopes in the deep water of the Mediterranean sea, namely \\verb\"ANTARES\" \\cite{[16]}, \\verb\"NESTOR\" \\cite{Anassontzis:1999gp,Aggouras:2004mh,Aggouras:2005bg} and \\verb\"NEMO\" \\cite{[18]}, which in the future could lead to the construction of a km$^3$ telescope as pursued by the \\verb\"KM3NeT\" project \\cite{[19],[20]}. Actually, the \\verb\"ANTARES\" collaboration is in a more advanced phase, with a telescope with an area of $\\sim$\\,0.1 km$^2$  already under construction \\cite{Aguilar:2006rm}. A further project is \\verb\"IceCube\", a cubic-kilometer under-ice neutrino detector \\cite{[21],[22],[23]} currently being deployed in a location near the geographic South Pole in Antarctica. \\verb\"IceCube\" applies and improves the successful technique of \\verb\"AMANDA\" to a larger volume.  Until now, the possibility to perform astronomy with neutrinos has been seriously limited by the presence of the heavy atmospheric background in the energy range presently explored. The start of $\\nu$-astronomy is eagerly waiting the completion of the new km$^3$ projects. In fact, according to theoretical expectations, the km$^3$ is the minimum detector size required to detect with reasonable chances of success point-like sources at the TeV scale and, more relevant for this paper, to explore energies above about 100 TeV, where extraterrestrial diffuse fluxes should start to dominate over the steeper atmospheric spectrum. In the following we shall focus the attention mainly on this energy range, although many of our results are valid also at lower energies.  Although NT's were originally thought as $\\nu_{\\mu}$ detectors, their capability as $\\nu_{\\tau}$ detectors has become a hot topic \\cite{Gandhi:1995tf,[24],[25],Anchordoqui:2005is,Yoshida:2003js,Beacom:2003nh,Athar:2000rx,Bugaev:2003sw,Ishihara}, in view of the fact that neutrino oscillations lead to nearly equal astrophysical fluxes for the three neutrino flavors\\footnote{This statement may not hold for exotic neutrino models \\cite{Beacom:2002vi,Beacom:2003eu} or for peculiar astrophysical sources \\cite{Serpico:2005sz,Kashti:2005qa,Serpico:2005bs}.}. Despite the different behavior of the produced tau leptons with respect to muons in terms of energy loss and decay length, both $\\nu_\\mu$ and $\\nu_{\\tau}$ detection are sensitive to the matter distribution near the NT site. Thus, a computation of the event detection rate of a km$^3$ telescope requires a careful analysis of the surroundings of the proposed site. The importance of the elevation profile of the Earth surface around the detector was already found of some relevance in Ref.\\ \\cite{Miele:2005bt}, where some of the present authors calculated the aperture of the Pierre Auger Observatory \\cite{Auger,Abraham:2004dt} for Earth-skimming UHE $\\nu_{\\tau}$'s. Indeed, air shower experiments can be used as NT's at energies$\\agt 10^{18}\\,$eV, a topic recently reviewed in \\cite{Zas:2005zz}.  In particular, the possibility of detection of the $\\tau$ leptons produced by Earth-skimming UHE $\\nu_\\tau$'s has been analyzed in a series of papers \\cite{Miele:2005bt}, \\cite{Capelle:1998zz}-\\cite{Miele:2004ze}. In Ref.\\ \\cite{Miele:2005bt} the use of a Digital Elevation Map (DEM) of the geographical area of the experiment proved useful to characterize peculiar matter effects in Earth-skimming events.  The aim of this paper is to estimate the effective aperture for $\\nu_\\tau$ and $\\nu_\\mu$ detection of a km$^3$ NT in the Mediterranean sea placed at any of the three locations proposed by the \\verb\"ANTARES\", \\verb\"NEMO\" and \\verb\"NESTOR\" collaborations. We do not consider any detail related to the experimental setup nor the detector response. In particular, we assume full detection efficiency via Cherenkov radiation for muons and taus crossing the NT fiducial volume. This means that we consider a lepton as detected if it crosses in any point the surface delimiting our fiducial volume, without e.g. taking into account further requirements or cuts needed for a good directional or energy reconstruction. These would depend on several parameters, like the spacing of the strings, the distribution of photomultipliers along the string, etc. which are characteristics of the apparatus and thus beyond the aim of the present analysis. We rather compare the site characteristics by using the DEM of the different areas. We shall therefore characterize and quantify the importance of ``matter effects'' for the three sites, and focus on the role played by the geometry of the experiment in enhancing the effect. These considerations may provide an important ingredient in shaping the final design of a km$^3$ Mediterranean NT.  A detailed DEM of the under-water Earth surface is available from the Global Relief Data survey (ETOPO2) \\cite{ETOPO2}, a grid of altimetry measurements with a vertical resolution of 1 m averaged over cells of 2 minutes of latitude and longitude. In Figures \\ref{Antares}, \\ref{Nemo} and \\ref{Nestor} we show the 3D maps of the areas around the three NT sites. The black curve represents the coast line, whereas the red spot stands for the location of the apparatus. By following the same approach developed in \\cite{Miele:2005bt}, we use this DEM to produce a realistic and statistically significant sample of $\\nu_\\tau$/$\\tau$ and $\\nu_\\mu$/$\\mu$ tracks crossing the fiducial volume of the NT that are then used to evaluate the effective aperture of each detector. % \\begin{figure}[t] \\begin{center} \\hspace{2cm} \\includegraphics[width=.70\\textwidth]{AntaresSiteLowRes.eps} \\vspace{-0.0cm} \\caption{The surface profile of the area near the \\texttt{ANTARES} site (red spot) at 42$^\\circ$ 30\u0092 N, 07$^\\circ$ 00\u0092 E. The black curve represents the coast line. The sea plateau depth in the simulation is assumed to be 2685 m. The effective volume starts at an height of 100 m from the seabed, to account for the spacing of the first photomultipliers as foresee by the current designs. The km$^3$ detector is oriented along the E-W/S-N directions.} \\label{Antares} \\end{center} \\end{figure} %  We note that when the events are reconstructed only in terms of the energy loss along the track, the UHE taus can not be distinguished from less energetic muons. This implies that the reconstruction analyses of UHE $\\nu_\\mu$ and $\\nu_\\tau$ events are highly entangled issues. We shall consider both of them, although for the sake of clarity we shall first focus on $\\nu_\\tau$ detection. Of course, when considering shower events or more in general contained events---where the charged lepton production happens inside the instrumented volume---including events with peculiar topologies like ``lollipop\" or ``double bang\", there are realistic chances of flavor-tagging in the detector. However, in the UHE range above $\\sim 10^7$ GeV these kind of events are subdominant with respect to the bulk of tau track-events. Further details on the flavor discrimination possibilities are discussed in Section \\ref{nu_mu}.  % \\begin{figure}[t] \\begin{center} \\hspace{2cm} \\includegraphics[width=.70\\textwidth]{NemoSiteLowRes.eps} \\vspace{-0.8cm} \\caption{The surface profile of the area near the \\texttt{NEMO} site (red spot) at 36$^\\circ$ 21\u0092 N, 16$^\\circ$ 10\u0092 E. The black curve represents the coast line. The sea plateau depth used in the simulation is 3424 m. The effective volume starts at an height of 100 m from the seabed, to account for the spacing of the first photomultipliers as foresee by the current designs. The km$^3$ detector is oriented along the E-W/S-N directions.} \\label{Nemo} \\end{center} \\end{figure}  The structure of the paper is as follows. In Section~\\ref{ev_rate} we introduce the formalism and definitions used in the analysis and define the aperture for a NT. Our results for $\\nu_\\tau$ induced events are reported and discussed in Section~\\ref{nu_tau} for various incoming neutrino fluxes, while $\\nu_\\mu$/$\\mu$ events are described in Section~\\ref{nu_mu}. Finally, we report our conclusions in Section~\\ref{concl}.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "Ultra-high energy neutrinos represent one of the main targets for several experiments which adopt a variety of detection techniques. Among these, the optical Cherenkov neutrino telescopes deployed under water or ice look for the tracks of charged leptons produced by the high energy neutrinos that reach the Earth. In this paper, we have presented a new study of the performance of a km$^3$ neutrino telescope to be located in any of the three sites in the Mediterranean sea proposed by the \\texttt{ANTARES}, \\texttt{NEMO}, and \\texttt{NESTOR} collaborations. Our main goal is to compare the performances of the different sites, keeping apart detector-specific features, like partial detection efficiency or the architecture adopted for the towers of strings. We concentrated instead on the details of the under-water surface profile of each of the three sites, using the data from a Digital Elevation Map. By generating a realistic and statistically significant sample of $\\nu_\\tau$/$\\tau$ and $\\nu_\\mu$/$\\mu$ tracks crossing the fiducial volume of the km$^3$ neutrino telescope, we have calculated its effective aperture to UHE $\\nu_\\tau$ and $\\nu_\\mu$ neutrinos and the expected number of events for different UHE neutrino fluxes, for both cases where the neutrino/charged lepton track is crossing the rock (denoted as {\\it rock events}) or the water only (denoted as {\\it water events}). Our results can be summarized as follows: \\begin{itemize} \\item The impact of the site geography (or matter effects) on observables such as the ``rock fraction\" of the total event rate  or asymmetries in the event direction can be important, particularly at high energies.  \\item Even for a fixed instrumented volume, these matter effects can be enhanced by a suitable choice of the geometry of the telescope, maximizing the lateral surface of the fiducial volume.  \\item The continental crust provides an absolute orientation, hence the matter effects may provide a mean for calibrating the pointing capabilities of the detector, even when no point-source identification is possible, like for diffuse cosmic fluxes at UHE.  \\item  We analyzed briefly the dependence of the rock and water events from the neutrino fluxes and the neutrino-nucleon cross section. We found that the ratio of rock to water events may provide an additional way to disentangle the two unknowns, in addition e.g. to the well-known zenith angle dependence of the rates due to the screening effect of the average spherical Earth. Although a detailed analysis would be needed, we stress that this may be important for constraining neutrino-nucleon cross section at UHE energies, otherwise unaccessible at the Lab.  \\item While below the PeV scale the aperture for muon tracks is one order of magnitude larger than the aperture for contained events, above $\\sim 10^8$ GeV the numbers of muon and tau track events are comparable. We have briefly addressed the problem of whether it is possible to distinguish $\\mu$'s and $\\tau$'s in the detector. Apart for the sub-dominant fraction of contained events with specific signatures, we stressed that a realistic prescription at UHE is to sum the bulk of $\\mu$ and $\\tau$ events, the natural variables for describing the events being the arrival direction and the energy loss rate in the fiducial volume. \\end{itemize}  The main conclusion one can draw from our analysis is that the optimization for a telescope aiming at the $E>$PeV region is significantly different from one whose target is the $E\\sim$TeV range: in the first case, the search is basically background free and even a relatively poor angular and energy resolution may be acceptable. The crucial goal is to maximize the event rates, and the discrimination among models may be based on ``counts\" and a very rough directional and energy binning of the events. In this respect, one should maximize the instrumented volume---compatibly with the experimental requirements for a meaningful reconstruction of the event---and also carefully design the geometry to maximize the lateral surface, in order to exploit the matter effects provided by the underwater profile. On the other hand, at TeV energies angular and energy resolutions are crucial to improve the very signal to noise ratio, and may help identifying point-like sources. At the same time, in the TeV range the matter effects we have stressed on in this paper are less relevant, and should not influence significantly the choice of the site.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %% Acknowledgments %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0609/gr-qc0609026_arXiv.txt": {
        "abstract": "#1{\\vskip 7mm \\begin{center}{\\large Abstract}\\par \\smallskip \\begin{minipage}[c]{12cm} \\small #1 \\end{minipage} \\end{center} } \\def\\title#1{\\begin{center}{\\Large\\bf #1}\\end{center}} \\def\\author#1{\\vskip 5mm \\begin{center}{#1}\\end{center}} \\def\\address#1{\\begin{center}{\\it #1}\\end{center}} \\def\\lsim{\\mathrel{\\rlap{\\lower4pt\\hbox{\\hskip1pt$\\sim$}} \\raise1pt\\hbox{$<$}}}                % \\def\\gsim{\\mathrel{\\rlap{\\lower4pt\\hbox{\\hskip1pt$\\sim$}} \\raise1pt\\hbox{$>$}}}                %  \\makeatletter \\@ifundefined{lesssim}{\\def\\lesssim{\\mathrel{\\mathpalette\\vereq<}}}{} \\@ifundefined{gtrsim}{\\def\\gtrsim{\\mathrel{\\mathpalette\\vereq>}}}{} \\def\\vereq#1#2{\\lower3pt\\vbox{\\baselineskip1.5pt \\lineskip1.5pt \\ialign{$\\m@th#1\\hfill##\\hfil$\\crcr#2\\crcr\\sim\\crcr}}} \\makeatother   \\begin{document}  \\title{% Spectroscopy of Cosmic topology } \\author{% {\\bf Tarun Souradeep} } \\address{% Inter-University Centre for   Astronomy and Astrophysics,\\\\ Post Bag 4, Ganeshkhind, Pune 411007,  India\\\\ {E-mail:tarun@iucaa.ernet.in}}  \\abstract{ Einstein's theory of gravitation that governs the geometry of space-time, coupled with spectacular advance in cosmological observations, promises to deliver a `standard model' of cosmology in the near future.  However, local geometry of space constrains, but does not dictate the topology of the cosmos. hence, Cosmic topology has remained an enigmatic aspect of the `standard model' of cosmology. Recent advance in the quantity and quality of observations has brought this issue within the realm of observational query.  The breakdown of statistical homogeneity and isotropy of cosmic perturbations is a generic consequence of non trivial cosmic topology arising from to the imposed `crystallographic' periodicity on the eigenstates of the Laplacian. The sky maps of Cosmic Microwave Background (CMB) anisotropy and polarization most promising observations that would carry signatures of a violation of statistical isotropy and homogeneity.  Hence, a {\\em measurable} spectroscopy of cosmic topology is made possible using the Bipolar power spectrum (BiPS) of the temperature and polarization that quantifies violation of statistical isotropy.  }   \\footnotetext[1]{Invited contribution to a special issue of IJP in memory of Prof. A.K. Raychaudhuri. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609061_arXiv.txt": {
        "abstract": "We study the evolution of bound pairs of star clusters by means of direct N-body simulations. Our simulations include mass loss by stellar evolution.  The initial conditions are selected to mimic the observed binary star cluster NGC\\,2136 and NGC\\,2137 in the Large Magellanic Cloud.  Based on the rather old ages ($\\sim 100$\\,Myr), masses, sizes of the two clusters and their projected separation, we conclude that the cluster pair must have been born with an initial separation of 15--20\\,pc.  Clusters with a smaller initial separation tend to merge in $\\aplt 60$\\,Myr due to loss of angular momentum from escaping stars.  Clusters with a larger initial separation tend to become even more widely separated due to mass loss from the evolving stellar populations. The early orbital evolution of a binary cluster is governed by mass loss from the evolving stellar population and by loss of angular momentum from escaping stars. Mass loss by stellar winds and supernovae explosions in the first $\\sim 30$\\,Myr causes the binary to expand and the orbit to become eccentric. The initially less massive cluster expands more quickly than the binary separation increases, and is therefore bound to initiate mass transfer to the more massive cluster.  This process is quite contrary to stellar binaries in which the more massive star tends to initiate mass transfer.  Since mass transfer proceeds on a thermal timescale from the less massive to the more massive cluster, this semi-detached phase is quite stable, even in an eccentric orbit until the orbital separation reaches the gyration radius of the two clusters, at which point both clusters merge to one. ",
        "introduction": "\\label{sec:intro}  The large Magellanic cloud (LMC) is a rich environment with many relatively young star clusters \\citep{2003MNRAS.338...85M,2003MNRAS.338..120M}. Several of these clusters appear closer together on the sky than expected from statistical arguments, i.e, they look like binary clusters \\citep{1991A&AS...87..335B,2002A&A...391..547D}.  The LMC contains a total of 69 cluster pairs with a separation $R \\aplt 1'.3$ \\citep{1988MNRAS.230..215B} which corresponds to $\\aplt 19$ pc by adopting distance modulus of 18.5 mag to the LMC \\citep{2003A&A...410..887G}, many of which have similar colors \\citep{1988AJ.....96.1625K}.  The small Magellanic cloud also contains a relatively large number of paired clusters \\citep{1990A&A...230...11H}.  Binary star clusters are expected to merge on a time scale of a few times 10\\,Myr \\citep{1988MNRAS.230..215B}, and NGC\\,2214 may be an LMC cluster in the process of merging \\citep{1988A&A...203L...5B} (but see \\cite{1995MNRAS.274.1225B} for counter arguments).  Well known examples of such binary clusters are SL\\,538/NGC\\,2006 (SL 537) \\citep{1998A&A...339..773D}, NGC\\,1850 \\citep{1993AJ....105..938F} and NGC\\,2136 (SL\\,762, at $\\alpha=$ 5$^{\\rm h}$\\, 53$^{\\rm m}$\\, 17$^{\\rm s}$; $\\delta= -69^\\circ$\\, 32'')/NGC\\,2137 \\citep{1995A&A...299L..37H}.  Confirmation of the binarity of such apparent pairs of star clusters should come from detailed measurements of their orbital parameters, rather than statistical arguments. However, with an orbital period of tens of million years such measurements are not trivial.  In the Antennae galaxies \\citep{2005AJ....130.2104W,2005ApJ...631L.133F} and in the young starburst galaxy M51 \\citep{2000MNRAS.319..893L,2005A&A...443...79B}, star clusters appear to be formed in groups often consisting more than two.  Apparently, star clusters in general are born in conglomerates in which several clumps tend to form within a relatively small time interval presumably initiated by galaxy-galaxy interactions. The individual star clusters in such conglomerates may subsequently merge in due time \\citep{2002Ap&SS.281..355F,2002CeMDA..82..113F}.  Cluster binaries may form in a comparable environments, in which case their ages have to be at least comparable.  Alternatively, a more massive cluster could possibly captures a lower mass cluster in a tidal event, resulting in two cluster with different ages.  However, this process is likely to result in a single cluster rather than a binary \\citep{1999IAUS..190..443D}.  In this study we mainly concentrate on the two star clusters NGC\\,2136 and NGC\\,2137, which attract attention because the projected distance between them is only about 20\\,pc (1'.34) \\citep{1994AGAb...10..198S}. In addition, both clusters have similar ages of about $100\\pm20$\\,Myr \\citep{2000A&A...360..133D} and identical metalicity \\citep{1995A&A...299L..37H}. With a primary mass of 26300--28200\\,\\msun\\, \\citep{2003MNRAS.338..120M} and a mass ratio $q \\simeq 0.17$, their orbital period is about 46\\,Myr. If bound, both clusters have orbited each-other about twice within their lifetime. Regretfully it is currently not known whether or not the clusters really form a bound pair, but at the moment this seems to be the most logical conclusion.  In this study we explore the initial conditions for which a pair of star clusters could survive until an age of about 100\\,Myr. We focus in particular on the observed pair NGC\\,2136 and NGC\\,2137 \\citep{1995A&A...299L..37H}, as they form an observational motivation for this exercise.  Assuming that the two clusters form a bound pair we are able to limit the initial cluster parameters. We identify two main regimes in binary cluster evolution. Small initial orbital separation ($\\aplt 12$\\,pc) tends to lead to a merger on a time scale of less than about 60\\,Myr, whereas large initial separation ($\\apgt 17$\\,pc) causes the two star clusters to recede. Clusters with initially a larger separation become very vulnerable to disruption by passing other star clusters, giant molecular clounds or the background potential of the host galaxy.  Between about 12\\,pc and 17\\,pc is a range of orbital separations between which the period changes little and in which the cluster binary is able to survive for more than 100\\,Myr. In these cases the change in orbital angular momentum by stellar mass loss and escapers is compensated by the redistribution of internal angular momentum in the rotation of the clusters. ",
        "conclusions": "\\label{sec:Discussion}  We studied the evolution of bound pairs of star clusters by means of direct N-body simulations while including the effects of stellar mass loss.  The initial parameters selected for our study are loosely based on the binary star cluster NGC\\,2136 and NGC\\,2137 in the large Magellanic cloud.  When starting with an initial cluster mass of about 40800\\,\\msun\\, and 6800\\,\\msun\\, in a circular orbit with a separation of 15--20\\,pc our simulations at about 100\\,Myr are consistent with the masses and orbital separation of the observed cluster pair.  The observations are reproduced most satisfactorily when we adopt a Salpeter mass function with a lower mass limit of 0.5--0.6\\,\\msun, but lower mass cut-offs may also produce a satisfactory comparison with the observations.  A word of caution, however, is well placed here, as we did not cover the entire parameters space in excruciating detail. The initial density profile, not varied in our study, may have a profound effect on the evolution of the binary cluster. The main aim of this paper is therefore not to constrain the initial conditions for NGC\\,2136/NGC\\,2137, but rather to obtain a better understanding of the general evolution of binary star clusters. Earlier numerical studies of binary star clusters ignored the stellar mass spectrum, stellar evolution and used approximate N-body techniques \\citep{1989PASJ...41.1117S,1991Ap&SS.185...63M,1998MNRAS.295..921D}. In our simulations all these effects are included in a self consistent fashion.  We are pleasantly surprised by the rich dynamics embedded in the evolution of binary star clusters. In particular the finding that stellar mass loss in the early evolution of the cluster is rapid enough to be considered a shock for the orbital elements of the binary cluster came somewhat unexpected. The consequence is that the orbits of star clusters with an age $\\apgt 10$\\,Myr are not circular, but small eccentricities $e\\apgt 0.2$ are induced upon the orbit, even if the two clusters are in Roche-lobe contact. The evolution of the distance between the two clusters is than a consequence of the orbital evolution (adiabatic/impulsive expansion) and the fact that cluster pair aproaches apoclustron.  Note, however, that all our models started with circular orbits. This choise was to limit parameters space as introducing an initial eccentricity also requires the choise of an eccentric anomaly.  Our initial clusters are not rotating. Since the redistribution of angular momentum turns out to be a mayor effect in the evolution of the cluster binary, future simulation may study the evolution of star cluster binaries by taking this extra parameter into account.  Another interesting aspect of our study that the less massive cluster is more likely to overfill its Roche lobe. The secondary cluster experiences core collapse in an earlier stage than the primary clusters. The resulting expansion of the post collapsed secondary cluster subsequently drives the transfer of mass through the first Lagrangian point to the primary cluster. Mass transfer proceeds from the less massive to the more massive component but the total number of stars transferred is relatively small (1--3\\%). The orbital evolution is dominated by stellar mass loss and by redistribution of angular momentum and by escaping stars  During the exchange of stars from one cluster to the other hardly any stars are lost through the second and third Lagrangian points. Stars do escape the clusters through the first Lagrangian point and isotropically. The isotropic escapers are mainly neutron stars which often receive high kick velocities in a supernova explosion. We therefore see no indication that tidal tails form in binary clusters. The origin of the tidal tails in several observed multiple clusters in the small Magellanic cloud by \\cite{1999A&A...344..450L} and \\cite{2000A&AS..146...57D} is therefore unlikely to be caused by the binarity of these clusters.  Binary star clusters in the LMC with a separation exceeding about 20\\,pc are vulnerable to ionization by giant molecular clouds. Our initial conditions approach this boundary. Such widely separated clusters can survive because the encounter rate between clusters in the LMC is rather small, making it unlikely that a cluster binary is ionized within a few times 10\\,Myr. However, as we demonstrated with our simulations, the orbital separation may increase considerably, making such clusters much more vulnerable to close (ionizing) encounters with other clusters or giant molecular clouds. Ionization by the background tidal field of the LMC is also possible. We did not take these effects into account in our simulations, but the consequences may be profound. One consequence is that there are (at least) two reasons why binary clusters in older populations are rare: close cluster pairs tend to merge and wide cluster pairs are likely to be ionized by the background tidal field or by passing Molecular clouds.  Of course, we cannot exclude that binary clusters are simply rarely formed.  If the distance between the two clusters is initially sufficiently large that the redistribution of angular momentum and escapers has little effect on the orbital evolution, semi-analytic calculations about the orbital evolution of the binary cluster produce satisfactory results. As soon as dynamical effects start to become important (at separations $\\aplt 15$\\,pc) these approximations break down."
    },
    "0609/astro-ph0609582_arXiv.txt": {
        "abstract": "The bright short-hard GRB~051103 was triangulated by the inter-planetary network and found to occur in the direction of the nearby M81/M82 galaxy group. Given its possible local-Universe nature, we searched for an afterglow associated with this burst. We observed the entire 3-$\\sigma$ error quadrilateral using the Palomar 60-inch robotic telescope and the Very Large Array (VLA) about three days after the burst. We used the optical and radio observations to constrain the flux of any afterglow related to this burst, and to show that this burst is not associated with a typical supernova out to $z\\approx 0.15$. Our optical and radio observations, along with the Konus/Wind gamma-ray energy and light curve are consistent with this burst being a giant flare of a Soft Gamma-ray Repeater (SGR) within the M81 galaxy group. Furthermore, we find a star forming region associated with M81 within the error quadrilateral of this burst which supports the SGR hypothesis. If confirmed, this will be the first case of a soft gamma-ray repeater outside the local group. ",
        "introduction": "\\label{Introduction}   The recent discovery of short-hard Gamma-Ray Bursts (GRB) afterglows in the X-ray (Gehrels et al. 2005; Fox et al. 2005), optical (Hjorth et al. 2005), and radio (Berger et al. 2005) demonstrates that short-hard GRBs reside in both early- and late-type galaxies with at least some preference for old systems (e.g. Gal-Yam et al. 2005; Tanvir et al. 2005) and release energy of the order of $10^{50}$~erg. With the detection of the giant flare of the 2004 December 27, from the Soft Gamma-ray Repeater (SGR)~1806$-$20 (e.g. Hurley et al. 2005), it has been suggested that a fraction (or all) of the short-hard bursts population originates from giant flares of SGR in nearby galaxies (e.g. Dar 2005; Hurley et al. 2005). This idea was based on the observed Galactic rate of SGR giant flares and the star formation rate in the local Universe. However, the fraction of extragalactic SGR flares among short-hard GRBs was shown to be small (Nakar et al. 2006; Gal-Yam et al. 2005; Palmer et al. 2005; Popov \\& Stern 2005; Lazzati et al. 2005; Ofek 2006). In contrast with the evidence associating short-hard GRBs with old stellar population (e.g. Nakar et al. 2005a), SGRs seem to emerge from young population objects (e.g. Gaensler et al. 2001; see however Levan et al. 2006; for a recent review on SGRs see Woods \\& Thompson 2006).   On UTC 2005 November 3 09:25:43.785, a short-hard GRB with $0.17\\,$s duration was detected (Golenetskii et al. 2005) by five satellites, of the Inter-Planetary Network (IPN; e.g. Hurley et al. 1999), carrying gamma-ray detectors: Konus/Wind; HETE/Fregate; Mars Odyssey/GRS; Mars Odyssey/HEND; RHESSI; and Swift/BAT. The GRB fluence in the Konus/Wind $20$~keV-$10$~MeV band was $2.34_{-0.28}^{+0.31}\\times10^{-5}$~erg~cm$^{-2}$, and its peak flux on two-milliseconds time scale was $1.89_{-0.35}^{+0.25}\\times10^{-3}$~erg~cm$^{-2}$~s$^{-1}$ ($90\\%$~confidence). This peak flux was one of the largest ever observed for Konus-Wind short GRBs, second only to GRB~031214 (Hurley et al. 2003). Moreover, comparing the burst fluence with the distribution of the integrated four-channel fluence of all BATSE short-hard bursts, we find it to be in the $99$th percentile. The light curve, shown in Figure~\\ref{GRB051103_LC_comparison} (solid line), has a very steep rise on $\\sim4$~milliseconds time scale and a weak decaying tail on time scale of $0.1\\,$s. \\begin{figure} \\centerline{\\includegraphics[width=8.5cm]{f1.eps}} \\caption{The Konus/Wind gamma-ray light curve of GRB~051103 (solid line), compared with the light curve of the 2004 December 27 SGR giant flare (dashed line). The light curve of the 2004 December 27 SGR giant flare is based on a digitization of Figure~1 in Terasawa et al. (2005), while the light curve of GRB~051103 is based on a digitization of the 18-1160 keV-band light curve from the Konus/Wind website. \\label{GRB051103_LC_comparison}} \\end{figure} The IPN 3-$\\sigma$ error box, totaling 260 square arcminutes, includes the outskirts of the nearby galaxies M81 and M82. At first glance, the error region seems to exclude the main body of M81 or M82 (see however \\S\\ref{Obs}).   In this paper, we present new optical and radio observations of M81 and M82 region which cover the entire IPN error quadrilateral, taken several days after the burst, and use them to put limits on any afterglow emission related to this GRB. We combine these with Galaxy Evolution Explorer (GALEX) archival data, and discuss the implications for the nature of this burst. ",
        "conclusions": "\\label{Disc}  The IPN error quadrilateral of the bright GRB~051103 includes the outskirts of the nearby ($3.63\\pm0.34$~Mpc; Freedman et al. 1994) galaxies M81 and M82 which are among the ten optically-brightest galaxies in the sky.   Assuming that GRB~051103 is related to the M81/M82 group, and using a distance modulus of $27.8$ (Freedman et al. 1994), and $A_R=0.2$ Galactic extinction (Schlegel et al. 1998), our null detection of any optical transients implies an upper limit of $-7.5$ on the R-band absolute magnitude of an optical transient. This limit is comparable to the absolute magnitude of novae, that peaks at $M_{V}=-6$ to $-10$ (e.g. della Valle \\& Livio 1995) and rule out the possibility that this event is associated with a typical supernova (unless $z\\gtorder0.15$, which rules out a source in the local Universe). The optical limit also rules out some of the macronovae models (assuming the distance of M81) recently explored by Kulkarni (2005).   Given the observational properties of GRB~051103, two questions arises: is it associated with the M81/M82 group? And is it an SGR flare or ``genuine'' short-hard GRB? Below, we explore the different possibilities and confront them with the observational facts.   \\subsection{An SGR in M81}  Assuming GRB~051103 was originated in M81, the energy from this burst, $(3.7\\pm0.8)\\times10^{46}$~erg (including the uncertainty in distance to M81; Freedman et al. 1994), is comparable to the energy of the most luminous SGR giant flare yet detected. For comparison, the isotropic energy release of the 1979 March 5 SGR flare was $>6\\times10^{44}$~erg, that of 1998 August 27 flare was $2\\times10^{44}$~erg, while the energy release from the 2004 December 27 giant flare was as much as $(3.7\\pm0.9)\\times10^{46}d_{15}^{2}$~erg (Hurley et al. 2005; Palmer et al. 2005), where $d_{15}$ is the distance to SGR~1806$-$20 in $15$~kpc units. As shown in Fig.~\\ref{GRB051103_LC_comparison}, the gamma-ray light curve of this burst has a rise-time of about a millisecond, which is characteristic of SGR giant flares (as well as less energetic flares). Moreover, it is composed of a single major peak -- similar to other SGR flares. Note that many, but not all, short-hard GRB light curves show a more complex temporal structure as observed in BATSE bursts (Nakar \\& Piran 2002), as well as in recent Swift bursts (e.g. GRB~051221; Parsons et al. 2005).  Known SGRs are associated with star-forming regions (Gaensler et al. 2001). Therefore, if GRB~051103 was an SGR giant flare in M81 or M82, then we expect it to occur within a young star forming region. Indeed, we have found UV sources indicating a young stellar population, in an extension of the M81 northern spiral arm, within the IPN error quadrilateral (see Fig.~\\ref{GALEX_M81} and Table~\\ref{Tab-UV}).  The SGR-origin idea is consistent with our radio observational limits. An extrapolation of the $1.4$~GHz specific flux of SGR~1806$-$20 (Cameron et al. 2005) to three days after the burst, suggests that the specific flux of a 2004 December 27-like event in M81, three days after the GRB, would be $\\sim0.02$~d$_{15}^{-2}$~mJy, significantly below our VLA limit.   The gamma-ray spectrum is an important clue for the nature of such bursts (e.g. Lazzati, Ghirlanda, \\& Ghisellini 2005). Unfortunately, the gamma-ray spectrum of GRB~051103 was not published. Therefore, it is not clear if the spectrum can be described by a black body spectrum, as observed in the 2004 December 27 giant flare. Assuming the spectrum is consistent with that of a black-body, the peak energy of GRB~051103, $1.92\\pm0.40$~MeV (Golenetskii et al. 2005), is translated to a black body peak energy ($3kT$) of $0.64\\pm0.13$~MeV. Given the peak flux (on a two-milliseconds time scale), and the distance to M81, the peak luminosity of GRB~051103 is about $2\\times10^{48}$~erg~s$^{-1}$. The black-body radius inferred from this luminosity is $12\\pm4$~km. This radius is comparable to the radius inferred for the 2004 December 27 giant flare, obtained in a similar way (Hurley et al. 2005; Nakar, Piran, \\& Sari 2005b). Intriguingly, these radii are comparable to the typical radius of neutron stars.    \\subsection{Short-hard GRB in the background}  Another possibility, already suggested by Lipunov et al. (2005), is that GRB~051103 is actually a genuine short-hard GRB. In this case, given the observed energy release from other short-hard GRBs (e.g. $\\sim10^{50}$~erg; Berger et al. 2005), this burst is more likely a background event unrelated to M81/M82. If however, it was a genuine short-hard GRB in M81 then we need to invoke a fainter population of short-hard GRBs, for which the implied rate will be higher than current estimates (e.g. Nakar et al. 2006).  Scaling the radio and optical properties of the afterglow of the relatively well observed short-hard GRB~050724 (Berger et al. 2005), by the ratio of gamma-ray fluence of GRB~051103 to GRB~050724, we estimate that at the time of our VLA observation, GRB~051103 should have had a $\\sim40$~mJy radio afterglow (at 8.46~GHz) -- an order of magnitude above our detection limit. Moreover, the non-detection of an optical afterglow, assuming a power-law decay, implies that the optical decay power-law was $\\ltorder-2$. However, the afterglow properties of short-hard GRBs are non-homogeneous (e.g. Berger et al. 2005; Hjorth et al. 2005; Bloom et al. 2006). Therefore, our observations can not rule out the possibility that this event was a genuine short-hard GRB.   If GRB~051103 was a genuine short-hard GRB in the background of M81, then we expect the spectrum to be non-thermal. In this case, the fact that this GRB is optically thin to gamma-ray photons allows to put a lower limit on its Lorentz factor (i.e. the ``compactness problem''). We used the recipe of Lithwick \\& Sari (2001) to estimate a lower limit on the Lorentz factor, $\\gamma$. We assumed that the gamma-ray spectrum of GRB~051103 is described by a Band spectrum (Band et al. 1993), and used the duration of the event ($\\delta{T}=0.17$~s), the peak energy, and the fluence within the Konus/Wind spectral-band. Assuming the gamma-ray spectrum of GRB~051103 can indeed be described by a Band-spectrum, and the distance to the GRB is $3.6$~Mpc (500~Mpc) we can set a lower limit of $\\gamma\\gtorder12$ ($\\gtorder18$), on the Lorentz factor of GRB~051103.    To conclude, although we can not rule out the possibility that GRB~051103 was a genuine short-hard GRB, the SGR giant flare appears to be the simplest interpretation of this event. The decay time, the peak luminosity and isotropic energy release, the presence of UV sources in the error quadrilateral, and the lack of optical and radio afterglow are all consistent with this burst being a giant flare from an SGR in M81. The spectrum of GRB~051103 has not been published. If a future analysis of the gamma-ray spectrum of GRB~051103 will show it is consistent with a black body spectrum, this will be another important evidence in favor of the SGR origin of this burst."
    },
    "0609/astro-ph0609127_arXiv.txt": {
        "abstract": "{ It is still a matter of debate if the group of \\boo stars is homogeneously defined. A widely discussed working hypothesis formulates that two apparent solar abundant stars of an undetected spectroscopic binary system mimic a single metal-weak spectrum preventing any reliable analysis of the group characteristics. }{ Is the proposed spectroscopic binary model able to explain the observed abundance pattern and photometric metallicity indices for the group members? What is the percentage of undetected spectroscopic binary systems? }{ We have used the newest available stellar atmospheres to synthesize 105 hypothetical binary systems in the relevant astrophysical parameter range. These models were used to derive photometric indices. As a test, values for single stellar atmospheres, Vega and two typical \\boo stars, HD\\,107233 and HD\\,204041, were generated. }{ The synthesized indices fit the standard lines and the observations of the three stars excellently. For about 90\\% of the group members, the spectroscopic binary hypothesis can not explain the observations. }{ A carefully preselection of \\boo stars results in a homogeneous group of objects which can be used to investigate the group characteristics. } ",
        "introduction": "More than 60 years after the first notification of the peculiar nature of \\boo (HR\\,5351) itself was addressed in the spectral classification survey by Morgan et al. (\\cite{morg43}) the origin and even the existence of a homogeneous group of $\\lambda$\\,Bootis stars is still a matter of debate.  Several theories were developed to explain the main characteristic of this group which comprise of main sequence late B to early F type stars: the lighter elements (C, N, O and S) are solar abundant whereas the heavier elements are significantly underabundant (Paunzen \\cite{pa04}).  In this letter we quantify the hypothesis formulated by Faraggiana \\& Bonifacio (\\cite{farag99}) and Gerbaldi et al. (\\cite{ger03}) that some, if not all, $\\lambda$\\,Bootis stars are in fact undetected spectroscopic binary systems with two solar abundant components simulating a combined single lined metal weak spectrum. ",
        "conclusions": "Figure\\,\\ref{fig1} shows the results of the relevant photometric diagrams. The observational data for the \\boo stars were taken from Paunzen et al. \\cite{pa02}. The classical metallicity sensitive diagrams m$_1$ versus (b\\,-\\,y) and m$_2$ versus (B2\\,-\\,V1) show that for cooler objects, the \\boo stars are nicely separated from the models on the main sequence (see c$_1$ versus (b\\,-\\,y)). These observational facts were also described in Paunzen et al. (\\cite{pa97}).  \\begin{table}[h] \\begin{center} \\caption{The results for $\\lambda$ Bootis stars taken from Paunzen et al. \\cite{pa02} according to Figure\\,\\ref{fig1} and Table \\ref{legend}. We would only consider the stars of group (\\,5\\,) as being good candidates for undetected spectroscopic binary systems. No conclusion can be drawn for group (\\,6\\,), the abbreviations are {\\bf{P}}assed or {\\bf{F}}ailed. The last row contains the number of objects as part of the total sample.} \\begin{tabular}[h]{rrrrrr} \\hline \\hline \\multicolumn{1}{c}{(\\,1\\,)} & \\multicolumn{1}{c}{(\\,2\\,)} & \\multicolumn{1}{c}{(\\,3\\,)} & \\multicolumn{1}{c}{(\\,4\\,)} & \\multicolumn{1}{c}{(\\,5\\,)} & \\multicolumn{1}{c}{(\\,6\\,)} \\\\ 6870 &     319 &   31295 &  111604 &   23392 &   13755\\,{\\bf{P}} \\\\ 11413 &    7908 &   35242 &  156954 &  110377 &   15165\\,{\\bf{P}} \\\\ 24472 &   30422 &  101108 &  193256 &  170680 &   54272\\,{\\bf{P}} \\\\ 75654 &   74873 &         &  193281 &  198160 &   87271\\,{\\bf{P}} \\\\ 81290 &   91130 &         &         &         &   90821\\,{\\bf{F}} \\\\ 83041 &  110411 &         &         &         &  105759\\,{\\bf{P}} \\\\ 83277 &  125162 &         &         &         &  111005\\,{\\bf{P}} \\\\ 84123 &  130767 &         &         &         &  120500\\,{\\bf{F}} \\\\ 102541 &  183324 &         &         &         &  120896\\,{\\bf{P}} \\\\ 105058 &  204041 &         &         &         &  175445\\,{\\bf{F}} \\\\ 106223 &  221756 &         &         &         &         \\\\ 107233 &         &         &         &         &         \\\\ 109738 &         &         &         &         &         \\\\ 125889 &         &         &         &         &         \\\\ 142703 &         &         &         &         &         \\\\ 142994 &         &         &         &         &         \\\\ 149130 &         &         &         &         &         \\\\ 153747 &         &         &         &         &         \\\\ 154153 &         &         &         &         &         \\\\ 168740 &         &         &         &         &         \\\\ 168947 &         &         &         &         &         \\\\ 184779 &         &         &         &         &         \\\\ 192640 &         &         &         &         &         \\\\ 210111 &         &         &         &         &         \\\\ 216847 &         &         &         &         &         \\\\ \\hline 44\\,\\% & 19\\,\\% & 5\\,\\% & 7\\,\\% & 7\\,\\% & 18\\,\\% \\\\ \\hline \\hline \\end{tabular} \\label{result} \\end{center} \\end{table}  Very interesting are the other three diagrams which include directly metallicity depend indices (Paunzen et al. \\cite{pa05}). The most important index is $\\Delta$a, for which, unfortunately, the least measurements are available whereas $\\Delta$(V1\\,-\\,G) seems less sensitive with a larger scatter. We notice that Vega as metal-deficient object (Hill \\& Landstreet \\cite{hill93}, Garc\\'ia-Gil et al. \\cite{gar05}) shows, as expected, a reduced $\\Delta$a value. Otherwise, the standard lines and observations of Vega fit the synthetic values very good. The differences of the synthetic and observed values for the two \\boo stars are only 1 to 2\\,mmag for $\\Delta$a, Z, and $\\Delta$(V1\\,-\\,G).  According to Figure\\,\\ref{fig1} we have divided our sample in six different groups numbered from (\\,1\\,) to (\\,6\\,) depending if the placement is compatible with the spectroscopic binary models taking an error of $\\pm$5\\,mmag into account. The final division is listed in Table \\ref{legend} whereas Table\\,\\ref{result} shows the results for the individual stars. The first two groups include objects which are not located in the spectroscopic binary area (SBA hereafter) in any diagram. The third and fourth group include stars which are situated in the SBA either in the $\\Delta$(V1\\,-\\,G) or Z diagram, but not in the classical metallicity and the $\\Delta$a diagram or do not have an available $\\Delta$a measurement (fourth group). The fifth group comprise good candidates for undetected spectroscopic binaries among the \\boo group because these stars compare well to the synthetic photometry of the binary models. HD\\,198160 is already known as close visual binary star whereas inconsistent $v \\sin i$ measurements were reported by Heiter et al. (\\cite{he02}) for HD\\,170680. The four objects of this group certainly deserve further attention in the future. The last group includes those stars which could only be tested within the m$_1$ versus (b\\,-\\,y) diagram which prevents any clear conclusion.  If we compare the list of Table\\,\\ref{result} with the results of Faraggiana et al. (\\cite{farag04}), only two stars, HD\\,11413 and HD\\,210111, were reported as spectroscopic binary systems on the basis of radial velocity shifts of two and three spectra, respectively. Both are well investigated $\\delta$ Scuti pulsators (Koen et al. \\cite{ko03}, Breger et al. \\cite{br06}) which makes them especially interesting for further investigations. We conclude that if these stars are indeed binary systems, at least one component is of \\boo type similar to the systems investigated by Iliev et al. (\\cite{il02}).  The percentage of undetected spectroscopic binary systems mimicking a single, metal-weak object seems very low. From 47 well investigated stars, groups (\\,1\\,) to (\\,5\\,), only four objects seem good candidates for a further investigation which is below 10\\% of the complete sample.  A carefully preselection of \\boo stars results in a homogeneous group of intrinsic \\boo stars which can be used to investigate the group properties in more detail."
    },
    "0609/astro-ph0609419_arXiv.txt": {
        "abstract": "We report the discovery of a bright ($H=12.77$) brown dwarf designated SIMP~J013656.5+093347. The discovery was made as part of a near-infrared proper motion survey, SIMP (Sondage Infrarouge de Mouvement Propre), which uses proper motion and near-infrared/optical photometry to identify brown dwarf candidates. A low resolution ($\\lambda/\\Delta\\lambda\\sim40$) spectrum of this brown dwarf covering the 0.88-2.35 $\\mu$m wavelength interval is presented. Analysis of the spectrum indicates a spectral type of T$2.5\\pm0.5$. A photometric distance of $6.4\\pm0.3$ pc is estimated assuming it is a single object. Current observations rule out a binary of mass ratio $\\sim$1 and separation $\\gtrsim5$ AU. SIMP 0136 is the brightest T dwarf in the northern hemisphere and is surpassed only by $\\varepsilon$~Indi~Bab over the whole sky. It is thus an excellent candidate for detailed studies and should become a benchmark object for the early-T spectral class. ",
        "introduction": "Numerous L and T dwarfs have been found during the last decade through the large scale Two Micron All Sky Survey (2MASS, \\citealp{Skrutskie2006}), the Deep Near-Infrared Survey of the Southern Sky (DENIS, \\citealp{Epchtein1997}), and the Sloan Digital Sky Survey (SDSS, \\citealp{York2000}). These objects have effective temperatures ranging from $\\sim$2300 to $\\sim$750 K and, with the exception of some early-L dwarfs, have substellar masses and are thus brown dwarfs (BDs). The L dwarfs are characterized by the presence of dust condensates suspended in their atmosphere and by increasingly red near-infrared (NIR) colors while the T dwarfs are distinguished by the appearance of broad methane absorption features in their NIR spectrum and by bluer NIR colors. The L/T transition is marked by a rapid decrease of $J-K_{\\rm s}$, from $\\sim$2 to $\\sim$0 between late-L and mid-T, and by an increase of the $J$-band flux, in contrast with the overall decreasing trend from early-L to late-T \\citep[see][and references therein]{Kirkpatrick2005}. Several scenarios have been proposed to explain these rapid variations of the NIR properties of BDs. \\citet{Burgasser2002a} propose that they are due to the clearing of patches in the cloud decks, with increasingly clearer atmospheres for later types. \\citet{Tsuji2003} suggest that, at the L/T transition, the $M_J$ versus $J-K_{\\rm s}$ diagram does not constitute a single evolutionary sequence. Rather, it is populated by objects of different surface gravities, for which the sinking of dust clouds below the photosphere occurs at different values of $M_J$ and temperature. \\citet{Knapp2004} offer an explanation based on a rapid rain-out of condensates from the photosphere occurring from late-L to mid-T. These scenarios make predictions that can be tested given a sufficiently large sample of early-T dwarfs.  Almost all known BDs were initially selected based on their NIR colors, and/or optical colors when available. Although this is an effective mean of identifying L or late-T dwarfs, early-T dwarfs are difficult to differentiate from M dwarfs when only NIR data are available because both types share similar $J-H$ and $H-K_s$ colors. Deep $i$- or $z$-band imaging, as in SDSS, removes the degeneracy since $i-z$ and $i-J$ differ significantly between M and T dwarfs. So far SDSS has covered 25\\% of the sky and found eight\\footnote{http://dwarfarchives.org/ \\label{bdarchive}} T dwarfs of spectral type T4.5 or earlier and brighter than $J=16$ (a conservative completeness limit for 2MASS). On the other hand, BD searches based on the 2MASS point source catalog (2MASS PSC, \\citealp{Cutri2003}) have unveiled only nine\\footnotemark[\\ref{bdarchive}] such objects over an area three times larger. This suggests that many early-T dwarfs remain to be indentified in the 2MASS PSC.  Proper motion (PM) is an efficient mean of discriminating BDs from M dwarfs. For a given magnitude, BDs are much closer and therefore have, on average, larger PMs. We have undertaken a NIR PM survey aimed at identifying new BDs and, in particular, increasing the sample of early-T dwarfs. This Letter reports the discovery, as part of this survey, of the brightest T dwarf in the northern hemisphere. ",
        "conclusions": "\\label{discussion}  SIMP 0136 was not identified as a BD by previous 2MASS searches because its NIR colors fall outside of the selection criteria used, e.g. $J-K_{\\rm s} \\le 0.3$ or $\\ge 1.30$ \\citep{Kirkpatrick1999}, $J-H\\le0.3$ and $H-K_{\\rm s}\\le0$ \\citep{Burgasser2003a}, or $J-K_{\\rm s} \\ge 1.0$ \\citep{Cruz2003}. These selection criteria were defined to minimize the number of M dwarfs in the candidate sample. SIMP 0136 would have been easily detected by SDSS but this region of the sky was not surveyed. SIMP 0136 is the first T dwarf discovered through PM in the NIR; the only other T dwarf discovered through PM is $\\varepsilon$~Indi~B \\citep{Scholz2003}, a discovery made with optical data. The discovery of SIMP 0136 within a dataset covering only $\\sim10\\%$ of the sky lends credency to the claim that many relatively bright ($J<16$) early-T dwarfs have yet to be discovered.  \\begin{table} \\begin{center} \\caption{\\label{tbl-2} Spectral indices} \\begin{tabular}{lcccc} \\hline\\hline Index & Numerator & Denominator & Value & Spectral type \\\\ \\hline CH$_4$-$H$ & $f_{1.635-1.675}$ & $f_{1.560-1.600}$ & 0.87 & T2\\\\ CH$_4$-$K$ & $f_{2.215-2.255}$ & $f_{2.080-2.120}$ & 0.42 & T3\\\\ \\hline \\end{tabular} \\end{center} \\end{table}  The photometric distance of SIMP 0136 can be estimated using the distances and magnitudes of known T dwarfs of similar spectral types. Table~\\ref{tbl-3} gives the NIR absolute magnitudes of the four T dwarfs in the T1-T4 range with distances known with an accuracy better than $\\sim10\\%$. Using their mean absolute magnitudes in $J$, $H$ and $K_s$, and assuming that SIMP 0136 has similar absolute magnitudes, we find photometric distances of $6.12\\pm0.97$, $6.26\\pm0.34$ and $6.61\\pm0.41$~pc respectively. The weighted average photometric distance is $6.4\\pm0.3$~pc.  \\begin{table*} \\begin{center} \\caption{\\label{tbl-3} Near-IR absolute magnitudes of nearby T dwarfs} \\begin{tabular}{lcr@{$\\:\\pm\\:$}lccc} \\hline\\hline Name & Sp & \\multicolumn{2}{c}{Distance} & \\multicolumn{3}{c}{Absolute magnitude}\\\\ &    & \\multicolumn{2}{c}{(pc)}     & $J$ & $H$ & $K_s$ \\\\ \\hline SDSS0151+12          &T1   &  21.4&1.5$^{\\ a}$   &$14.91\\pm 0.20$&$13.95\\pm0.19$&$13.53\\pm 0.25$\\\\ SDSS1254-01          &T2   & 11.78&0.26$^{\\ b}$  &$14.54\\pm0.06$&$13.73\\pm0.05$&$13.48\\pm 0.07$\\\\ SDSS1750+17          &T3.5 &  27.6&3.5$^{\\ a}$   &$14.14\\pm0.29$&$13.75\\pm0.30$&$13.27\\pm 0.33$\\\\ $\\varepsilon$~Indi~Ba &T1   & 3.626&0.009$^{\\ c}$ &$14.49\\pm0.02$&$13.71\\pm0.02$&$13.55\\pm0.02$\\\\ \\hline Mean&\\multicolumn{3}{c}{}&$14.52\\pm0.32$ & $13.79\\pm0.11$&$13.46\\pm0.13$\\\\ \\hline \\end{tabular} \\tablenotetext{a}{Parallax from \\citet{Vrba2004}.} \\tablenotetext{b}{Parallax from \\citet{Dahn2002}.} \\tablenotetext{c}{Parallax from the Hipparcos catalog \\citep{Perryman1997}.} \\end{center} \\end{table*}  Considering that $\\sim$40\\% of T dwarfs are resolved binaries \\citep{Burgasser2006b}, it is crucial to test SIMP 0136 for binarity to ensure a proper interpretation of the observations and a legitimate comparison with models. Astrometric observations are underway to determine the parallax of SIMP 0136 using the NIR camera WIRCAM \\citep{Puget2004} at the Canada-France-Hawaii Telescope. This camera features a field of view of $20^\\prime\\times20^\\prime$ sampled with 0\\farcs30 pixels. These observations have provided the best angular resolution observations for this object (0\\farcs6 seeing) and no binarity or PSF elongation has been detected. Thus, SIMP 0136 does not appear to be a binary of mass ratio $\\sim$1 and separation $\\gtrsim$5 AU. However, observations at a higher resolution are needed to further constrain its multiplicity since most binary T dwarfs have separations smaller than this \\citep{Burgasser2006b}. If SIMP 0136 is a binary, its photometric distance could be as large as $\\sim$9 pc. The two other brightest early-T dwarfs, $\\varepsilon$~Indi~Bab and SDSS0423-04 are both binaries \\citep{Volk2003, McCaughrean2004, Burgasser2005}; thus, if SIMP 0136 is a single object, it would be the brightest isolated T dwarf, and as such it would become a benchmark object for the early-T spectral class, enabling very high-resolution spectroscopy, time-resolved spectroscopy and polarimetry observations. Its location near the celestial equator makes it all the more appealing as a benchmark object since it can be observed from all major observatories around the world."
    },
    "0609/astro-ph0609133_arXiv.txt": {
        "abstract": "We observed the oxygen-rich Large Magellanic Cloud (LMC) supernova remnant N132D (SNR\\,0525--69.6), using all instruments onboard the {\\it Spitzer Space Telescope}, IRS, IRAC, and MIPS (Infrared Spectrograph, Infrared Array Camera, Multiband Imaging Photometer for Spitzer). The 5--40\\mic\\ IRS spectra toward the southeastern shell of the remnant show a steeply rising continuum with \\neiiif\\ and \\oivf\\ as well as PAH emission. We also present the spectrum of a fast moving ejecta knot, previously detected at optical wavelengths, which is dominated by strong \\neiiif\\ and \\oivf\\ emission lines. We interpret the continuum as thermal emission from swept-up, shock-heated dust grains in the expanding shell of N132D, which is clearly visible in the MIPS 24\\mic\\ image. A \\mbox{15--20\\mic} emission hump appears superposed on the dust continuum, and we attribute this to PAH \\mbox{C-C-C} bending modes. We also detect the well-known 11.3\\mic\\ PAH C-H bending feature, and find the integrated strength of the 15--20\\mic\\ hump about a factor of seven stronger than the 11.3\\mic\\ band in the shell of the remnant. IRAC 3--9\\mic\\ images do not show clear evidence of large-scale, shell-like emission from the remnant, partly due to confusion with the ambient ISM material. However, we identified several knots of shocked interstellar gas based on their distinct infrared colors. We discuss the bright infrared continuum and the polycyclic aromatic hydrocarbon features with respect to dust processing in young supernova remnants. ",
        "introduction": "\\label{sec;introduction} Supernovae influence the chemistry and physics in the interstellar medium (ISM) from galactic scales down to the atomic level. Primarily, it is the energy input of supernova explosions that drives important processes in the ISM: strong shock waves accelerate, heat, and compress interstellar gas, destroy dust grains, accelerate cosmic rays, and affect interstellar chemistry \\citep[cf.][]{draine1993}.  The role that supernovae play in the formation, destruction, and processing of polycyclic aromatic hydrocarbons (PAHs) is largely unexplored. PAHs are believed to be the carrier of the unidentified infrared (UIR) bands, well-studied emission features near 3.3, 6.2, 7.7, 8.6, and 11.3\\mic\\ attributed to the C-C and C-H stretching and bending modes of these molecules \\citep{allam1989,puget1989}. It is established that interstellar grains are readily destroyed by strong shock waves \\citep[e.g.][]{tiele1994}, and it was suggested that grain destruction through shocks may be a source of PAH molecules \\citep{jones1996}. Although PAH features have been observed in a wide variety of sources \\citep[cf.][]{vandi2004}, no PAH emission has yet been detected in SNRs as far as we are aware, however, \\citet{reach2006} identified four SNRs whose infrared colors suggest emission from PAHs.  More recently, spectral features at 15--20\\mic\\ including  broad plateaus and distinct, narrower features, most notably near 16.4 and 17.4\\mic, were detected in conjunction with the well-known PAH features \\citep[][and see \\citealp{peete2004} for a summary]{mouto2000,vanke2000,armus2004,morri2004,smith2004,werne2004}. The 15--20\\mic\\ features are generally attributed to \\mbox{C-C-C} bending modes of large PAH molecules \\citep[cf.][]{allam1989,vanke2000,peete2004}. The relatively narrow observed widths of the discrete 16.4 and 17.4\\mic\\ bands suggest that they arise from individual PAH molecules in the gas phase \\citep{peete2004}. There is also evidence that the appearance, relative strengths, and profiles of the PAH features are coupled to the local physical conditions \\citep{hony2001,peete2002,vandi2004,werne2004}.   In this paper, we report the detection of weak PAH features near 6.2 and 7.7--8.6\\mic, at 11.3\\mic, and a prominent, broad hump at 15--20\\mic\\ from the young supernova remnant N132D (SNR\\,0525$-$69.6) in the Large Magellanic Cloud (LMC). The age of N132D determined from the ejecta kinematics is $\\sim$\\,2500\\,yr \\citep{morse1995}, and its spatial extent is about 100\\arcsec, or 25\\,pc assuming a LMC distance of 50\\,kpc. N132D belongs to the small class of young, oxygen-rich SNRs that are presumably the product of core-collapse supernovae, with the most famous Galactic example being Cas\\,A. For the progenitor of N132D, \\citet{blair2000} suggested a W/O star of 30--35\\msun\\ based on relative elemental abundances and comparisons to stellar nucleosynthesis models, which may have exploded as a Type Ib supernova. N132D shows a radio spectral index of $-0.7$ \\citep{dicke1995}, consistent with its overall shell-like morphology. It is one of the brightest X-ray sources in the LMC \\citep{hughe1987,hwang1993,favat1997,hughe1998}, and has also been investigated in great detail in the UV/optical \\citep{blair1994,morse1995,suthe1995,morse1996,blair2000}. The large-scale morphology of the gas in N132D shows three types of line-emitting regions in the optical \\citep{morse1996}: (1) red- and blueshifted oxygen-rich ejecta filaments in the center, with a total velocity range $\\sim$\\,4400\\km\\ \\citep{morse1995}, (2) shocked interstellar clouds within the remnant perimeter and in the outer rim, and (3) a photoionized precursor around the outer edge of the remnant.  We observed N132D in the infrared with the {\\it Spitzer Space Telescope} \\citep{werne2004i}. In Section~\\ref{sec;observations}, we describe our observations and the data reduction procedure. We took particular care of the background subtraction in order to avoid confusion with foreground and background emission (Section \\ref{sec;background}). We present our results from the imaging and the spectroscopy in Section \\ref{sec;results}, and give a detailed discussion of the IR emission features and a comparison to PAH features observed in other astronomical objects in Section~\\ref{sec;discussion}. ",
        "conclusions": "Figure~\\ref{fig:n132dspec-bbfit} shows the IRS spectrum of the southern rim. The foreground/background and the fitted dust continuum have been subtracted (cf.~Sect.~\\ref{sec;background} and \\ref{sec;IRS}). The quoted total error consists of the pipeline-delivered statistical ramp uncertainties, combined with an estimate of the error due to the local background variation (cf.~Sect.~\\ref{sec;background}). \\begin{figure}[!tp] \\epsscale{1.15} \\centering \\plotone{f6_color.eps} \\caption{N132D southeastern rim spectrum (color online version: magenta=SL2, red=SL1, green=LL2, blue=LL1): the local background and the modified-blackbody continuum from Fig.~\\ref{fig:n132dspec+bbfit} have been subtracted. The total error includes the reduction pipeline delivered errors and an estimate of the variation in the local background (see text for details). } \\label{fig:n132dspec-bbfit} \\end{figure} We detect several emission lines from atomic ions, tentative weak PAH features near 6.2 and 7.7--8.6\\mic, a weak PAH 11.3\\mic\\ feature, and a prominent, broad hump at 15--20\\mic. The ionic emission lines are from Ne$^{2+}$, O$^{3+}$, and S$^{2+}$ ions. We can not distinguish between \\oivf\\,25.91\\mic\\ and \\feiif\\,25.99\\mic\\ in our low-resolution spectra, but argue that most of the emission is from O$^{3+}$ (cf.~Sect.~\\ref{sec;atomiclines}). The PAH features at 6.2, between 7.7 and 8.6\\mic, and at 11.3\\mic\\ are the well-studied PAH C-H in-plane and out-of-plane bending modes \\citep[e.g.][]{hony2001,peete2002,vandi2004}. The presence of discrete 6.2 and 7.7--8.6\\mic\\ bands is uncertain due to the low quality of our spectra in that region (cf.~total error in Fig.~\\ref{fig:n132dspec-bbfit}). There is probably an underlying emission plateau at 6--9\\mic, which has been observed in other sources and is probably associated with broad continua due to large PAHs, clusters of small PAHs, or very small grains \\citep[e.g.][]{allam1989,peete2002,rapac2005}.  The broad hump at 15--20\\mic\\ is presumably due to PAH \\mbox{C-C-C} in- and out-of-plane bending modes \\citep{allam1989,vanke2000,peete2004}. Figure~\\ref{fig:n132dspec_gaussfit} shows the prominent emission hump between 15 and 20\\mic\\ in detail. The spectrum without background subtraction clearly shows the 16.4\\mic\\ PAH feature \\citep[cf.][]{mouto2000,vanke2000}, and additional features near 17.1 and 17.9\\mic. These distinct bands with widths $\\sim$\\,0.5\\mic\\ may arise from smaller PAHs in the gas phase. However, the individual bands are hardly distinguishable after background subtraction, indicating that these features are primarily due to the background (cf.~background spectrum in Fig.~\\ref{fig:n132dspec+backgr}).  \\begin{figure}[!htbp] \\epsscale{1.15} \\centering \\plotone{f7.eps} \\caption{N132D southeastern rim spectrum (LL2): detailed view of 15--20\\mic\\ region from Fig.~\\ref{fig:n132dspec-bbfit}. For comparison, we show the spectrum without the local background subtraction (dot-dashed line), and the total error from Fig.~\\ref{fig:n132dspec-bbfit} as a gray-shaded error band. } \\label{fig:n132dspec_gaussfit} \\end{figure}  Table~\\ref{tab;gaussfit} lists the fit parameters for all detected ionic lines and PAH features in the rim, the local background, and the molecular cloud/\\hii-region near N132D. For the local background spectrum, we assumed a flat continuum of $1.2\\times10^{-8}\\,\\rm W\\,m^{-2}\\mic^{-1}sr^{-1}$. \\begin{table}[!tbp] \\centering \\caption{N132D fit data with errors given in parentheses: southeastern rim (Fig.~\\ref{fig:n132dspec-bbfit}), local background (Fig.~\\ref{fig:n132dspec+backgr}), and the molecular cloud/{\\rm H}\\,{\\sc ii}-region (Fig.~\\ref{fig:n132dspec_rim+knot+cloud+ngc} and \\ref{fig:n132dspec_sources}). } \\begin{tabular}{lcc@{\\hspace{0.3cm}}c@{\\hspace{0.1cm}}c@{\\hspace{0.3cm}}cc}  % \\hline\\hline  ID      &   \\multicolumn{2}{c}{$\\lambda$}        &  \\multicolumn{2}{c}{\\hspace{-0.25cm}Flux$\\times 10^{-8}$}          &    \\multicolumn{2}{c}{FWHM}       \\\\ &    \\multicolumn{2}{c}{[$\\mu$m]}   &  \\multicolumn{2}{c}{\\hspace{-0.25cm}[$\\rm W\\,m^{-2}\\,sr^{-1}$]}    &   \\multicolumn{2}{c}{[$\\mu$m]}    \\\\ \\hline\\hline  \\multicolumn{7}{c}{\\it N132D southeastern rim}     \\\\ PAH\\,6.2  &         &      &   \\multicolumn{2}{l}{\\raisebox{0.5mm}{\\ \\ \\ \\tiny ${\\lesssim}$}1.8} &  &  \\\\  PAH\\,11.3  &      11.28 &       (0.01) &        1.06 &        (0.12) &       0.20 &       (0.05) \\\\  [Ne\\,{\\sc iii}] & 15.60 &       (0.01) &        0.81 &        (0.01) &       0.27 &       (0.03) \\\\  PAH\\,15--20  &       &         &        7.47 &        (0.67) &       &        \\\\  [S\\,{\\sc iii}] 2--1&    18.74 &       (0.01) &        0.49 &        (0.04) &       0.40 &       (0.03) \\\\  [O\\,{\\sc iv}] &    25.96 &       (0.01) &        0.30 &        (0.02) &       0.36 &       (0.02) \\\\  \\multicolumn{7}{c}{\\it N132D local background}     \\\\ PAH\\,6.2  &         &       &   \\multicolumn{2}{l}{\\raisebox{0.5mm}{\\ \\ \\ \\tiny ${\\lesssim}$}5.8} &  &  \\\\  PAH\\,11.3  &      11.26 &       (0.01) &        4.28 &        (0.15) &       0.28 &       (0.02) \\\\  [Ne\\,{\\sc iii}] & 15.58 &       (0.01) &        0.30 &        (0.02) &       0.23 &       (0.01) \\\\  PAH\\,15--20  &       &         &        3.20 &        (0.68) &       &        \\\\  [S\\,{\\sc iii}] 2--1&    18.74 &       (0.01) &        0.24 &        (0.02) &       0.30 &       (0.03) \\\\  [S\\,{\\sc iii}] 1--0&    33.52 &       (0.01) &        0.34 &        (0.01) &       0.46 &       (0.02) \\\\  [Si\\,{\\sc ii}]&    34.86 &       (0.01) &        0.77 &        (0.01) &       0.43 &       (0.01) \\\\  \\multicolumn{7}{c}{\\it Molecular cloud/{\\it H\\,{\\it\\scriptsize II}}-region near N132D (LL only)}    \\\\  [Ne\\,{\\sc iii}]? &15.63 & (0.02) & 0.16 & (0.04) & 0.24 & (0.06) \\\\ PAH? &    15.92 &          (0.01) &  0.25     &       (0.05) &   0.23   & (0.05) \\\\ PAH\\,16.4 &  16.46 &       (0.01) &   1.67     &       (0.08) &    0.29  & (0.01) \\\\ PAH?    &  16.75 &        (0.01) &   0.83     &       (0.13) &    0.35  & (0.04)\\\\ \\htwo\\ S(1)& 17.06       & (0.01) &   1.02    &       (0.01) &  0.27    & (0.01)\\\\ PAH\\,17.4  & 17.38 &       (0.01) &    1.57    &       (0.11) &    0.41    & (0.02)\\\\ PAH\\,17.8  &   17.76 &     (0.01) &   0.48     &       (0.07) &   0.32    &  (0.04)\\\\ PAH? &    18.27 &          (0.03) &  0.14     &       (0.04) &   0.30   & (0.09) \\\\  [S\\,{\\sc iii}] 2--1 & 18.77 & (0.01) &  0.68     &       (0.01) &  0.26    & (0.01) \\\\  \\htwo\\ S(0)&   28.25 &       (0.01) &  0.26     &       (0.01) &  0.38    & (0.01) \\\\  [S\\,{\\sc iii}] 1--0 &33.52 &(0.01) &  0.71      &       (0.01) &  0.41    &  (0.01)\\\\  [Si\\,{\\sc ii}] &   34.86 &  (0.01) &  0.97      &       (0.01) &  0.38    &  (0.01)\\\\ \\hline\\\\ \\end{tabular} \\label{tab;gaussfit} \\end{table} The detected PAH features in the southeastern rim of N132D are the first evidence of PAH molecules surviving a supernova blast wave to our knowledge. PAH features also appear in lines-of-sight outside the remnant perimeter, and it is necessary to make certain that the PAH emission originates from the blast wave, and not from foreground, background, or ambient ISM material. There are three points of evidence that the PAH emission is indeed originating from material associated with the blast wave. Firstly, we removed the foreground and background emission by subtracting the flux from lines-of-sight just outside the blast wave perimeter (Sect.~\\ref{sec;background}). Secondly, we analyzed the spatial variation of emission line fluxes along the LL spectrograph slit (cf.~Sect.~\\ref{sec:discussion_spatialvariation}, Figure~\\ref{fig:n132dspec_spatialvariation}), which shows that the PAH features between 15 and 20\\mic\\ have a well-defined local flux maximum at the position of the rim. This is expected for material that has been swept up by the blast wave and is not yet completely destroyed by the shock. If the majority of the flux would be from unrelated foreground or background emission, there would be no local maximum at the position of the rim. Some fraction of the detected PAHs may reside in the radiative precursor (cf.~Sect.~\\ref{sec;IRAC}), where they would be excited/processed by energetic photons emanating from the supernova remnant. Finally, the 15--20\\mic\\ PAH emission changes notably when comparing features from the supernova remnant to the adjacent molecular cloud/\\hii-region and the local background (cf.~Fig.~\\ref{fig:n132dspec_rim+knot+cloud+ngc} and \\ref{fig:n132dspec_sources}), which we interpret as PAH/grain processing in N132D (cf.~Sect.~\\ref{sec;IRSgrainprocessing_discussion} and \\ref{sec:previousobservation}).    \\subsection{Spatial variation of the emission features} \\label{sec:discussion_spatialvariation} \\label{sec;atomiclines} \\label{sec:PAHlines} In Figure~\\ref{fig:n132dspec_rim+knot+cloud}, we show the spectrum of the southeastern rim of N132D in comparison to the ejecta knot detected in the LL spectra, the molecular cloud/\\hii-region south of N132D, and the average spectrum of the galaxy NGC\\,7331. The characteristic features of the rim spectrum have been described earlier in this paper. The spectrum of the ejecta knot lacks the steeply rising dust continuum and the 15--20\\mic\\ PAH features, but shows strong ionic emission lines from Ne$^{2+}$ and O$^{3+}$, which is consistent with the strong neon and oxygen lines seen in the optical for the same knot \\citep[N132D-P3,][]{blair2000}. The line fluxes and FWHM are $\\rm 1.6\\times10^{-8}\\wms$ (0.15\\mic) for the neon and $\\rm 8.4\\times10^{-9}\\wms$ (0.34\\mic) for the oxygen line. The spectrum of the molecular cloud/\\hii-region shows strong PAH bands as well as ionic and \\htwo\\ lines. The association of the CO emission south of N132D with an \\hii-region seen in H$\\alpha$ was noticed by \\citet{morse1995}. \\citet{banas1997} estimated the mass of the CO cloud to be about $3\\times10^5\\msun$, and argued that its spatial proximity and matching LSR velocity provide strong evidence that this giant cloud is associated with N132D. This provides an opportunity to relate our observations of dust and PAHs in the supernova remnant N132D to the nearby environment its progenitor star probably originated from. Such a comparison will also be useful to interpret spectroscopic observations of star-forming galaxies \\citep[e.g.~NGC\\,7331,][see Sect.~\\ref{sec:previousobservation}]{smith2004}.  \\begin{figure}[!tbp] \\epsscale{1.15} \\centering \\plotone{f8_color.eps} \\caption{Comparison of different {\\it Spitzer }IRS spectra for N132D, its nearby molecular cloud/H\\,{\\sc ii}-region, and an archival spectrum of the galaxy NGC\\,7331 (dashed line). We extracted and processed the average spectrum of the nucleus and inner ring of NGC\\,7331 \\citep[][]{smith2004} in the same way as our spectra, and applied a scaling factor of 0.5 for illustration purposes (see Sect.~\\ref{sec:discussion_spatialvariation} and \\ref{sec:previousobservation} for further discussion). } \\label{fig:n132dspec_rim+knot+cloud} \\label{fig:n132dspec_rim+knot+cloud+ngc} \\end{figure}   We present a systematic assessment of the PAH 15--20\\mic, ionic, and molecular line fluxes from positions across N132D and the nearby molecular cloud in Figure~\\ref{fig:n132dspec_spatialvariation}. We extracted spectra along the LL slit with $\\approx25\\arcsec$ separation starting at the northern edge of N132D and ending at the southern edge of the molecular cloud about 3.3\\arcmin\\ south of N132D, resulting in a total number of 15 spectra. We converted the spatial distance along the slit into a parsec scale assuming a LMC distance of 50\\,kpc, and adopted the location of the southeastern rim of N132D as the zero position. Each node in Figure~\\ref{fig:n132dspec_spatialvariation} represents a line flux derived from continuum and line fitting, with errors given for the positions covering the supernova remnant (cf.~Sect.~\\ref{sec:analysis_emissionlines}).  \\begin{figure}[!htbp] \\epsscale{1.15} \\centering \\plotone{f9.eps} \\caption{Spatial variation of line vs.~PAH 15--20\\mic\\ fluxes, and total integrated flux from continuum and lines in the 14--21\\mic\\ range as a function of projected distance from the supernova shock. [Si\\,{\\sc ii}] and [O\\,{\\sc iv}] appear only in LL1 (19.5--38\\mic), limiting the spatial coverage to about 20\\,pc. The \\htwo\\ flux across the supernova remnant is uncertain due to strong blending with PAH features but most likely very small (see text for further discussion). } \\label{fig:n132dspec_spatialvariation} \\end{figure}   The ionic and molecular lines detected and examined in our spectra are \\neiiif\\,15.6,  \\siiif\\,18.7 and 33.5, \\oivf\\,25.9, \\siliif\\,34.8, and \\htwo\\,0--0\\,S(0)\\,28.2 and S(1)\\,17.0\\mic. The lines from atomic ions, particularly the strong \\neiiif\\  and \\siiif\\ transitions, are mostly unblended and straightforward to fit, which makes them excellent references in terms of flux and central wavelengths.  \\neiiif\\ and \\oivf\\ lines only appear across the supernova remnant and in the precursor, and are particulary strong at the position of the ejecta knot. Generally, \\neiiif\\ and \\oivf\\ trace the hot ionized gas of N132D. Note that we can not distinguish between \\oivf\\,25.91\\mic\\ and \\feiif\\,25.99\\mic\\ in our low-resolution spectra, but we argue that most of the flux is likely to be from \\oivf\\ because the flux follows \\neiiif\\ closely (cf.~Fig.~\\ref{fig:n132dspec_spatialvariation}); under the conditions where we observe Ne$^{2+}$, iron would be mostly present as Fe$^{3+}$ and Fe$^{4+}$. There is a weak line near 17.9\\mic, which may be \\feiif\\,17.9, but possible confusion with the PAH 15--20\\mic\\ features renders this line inconclusive. We do not see \\feiif\\,24.5/35.3\\mic, \\feiif/\\feiiif\\,22.9, and \\fevif\\,19.6\\mic, which are probably too weak to be detected in our spectra. However, X-ray emission from Fe$^{16+}$ and more highly ionized states was detected in the southern rim \\citep{behar2001}.  \\siiif\\ is a common line seen in PDRs and \\hii-regions. Here, it not only appears across the remnant, but also across the molecular cloud tracing the \\hii-region. The \\siiif\\,1-0 transition at 33.5\\mic\\ follows the same general trend, with the \\mbox{2-1/1-0} ratio increasing from 0.08 to about 0.2 from the northern edge of the molecular cloud into the remnant.  \\htwo\\ traces the  molecular cloud seen in CO (Fig.~\\ref{fig:n132d}). We checked for \\htwo\\,0-0\\,S(0)\\,28.2, S(2)\\,12.3, and S(3)\\,9.7\\mic, and did not detect these lines across the remnant; weak S(1)\\,17.0\\mic\\ emission is not detectable because of strong blending with PAH emission. We conclude that \\htwo\\ is mostly absent across the remnant, where it was probably destroyed by the fast shock and the strong radiation field, but apparently present in the precursor. The \\htwo\\,S(1)/S(0) ratio increases from about 1.4 at the northern edge of the molecular cloud/\\hii-region to \\raisebox{0.5mm}{\\scriptsize ${\\gtrsim}$}\\,11 at the shock precursor of the rim, which implies an increase in the UV radiation field as expected for the presence of a radiative precursor (cf.~Sect.~\\ref{sec;IRAC}).   \\siliif\\ shows a nearly constant flux, except a noticeable minimum at the position of the rim. This is most likely an effect of the increased radiation field and the elevated temperature. Silicon will be much more highly ionized under those conditions, which is confirmed by the detection of X-ray emission from Si$^{12+}$ in the southern rim by \\citet{behar2001}. Unfortunately, this conceals a possibly increased amount of silicon in the gas phase due to the destruction of dust grains.   There are three well-defined PAH components at 16.5, 17.4, and 17.8\\mic\\ in the region of the molecular cloud (Fig.~\\ref{fig:n132dspec_rim+knot+cloud+ngc} and \\ref{fig:n132dspec_sources}, and Table~\\ref{tab;gaussfit}), which have all been previously observed in other sources (see Sect.~\\ref{sec:previousobservation}). The fluxes of these features follow the \\htwo\\ flux very well as expected if the PAHs are well-mixed with the gas (Fig.~\\ref{fig:n132dspec_spatialvariation}). The feature at 16.8\\mic\\ is blended with the stronger PAH feature at 16.5\\mic\\ and the strong \\htwo\\,S(1) line at 17.1\\mic, and therefore more uncertain. The situation changes for the supernova remnant, where the discrete 15--20\\mic\\ PAH features all but disappear, and a nearly featureless 15--20\\mic\\ emission hump remains (Fig.~\\ref{fig:n132dspec_gaussfit} and Table~\\ref{tab;gaussfit}).         \\subsection{IRAC emission knots} \\label{sec;IRACdiscussion} We detect shocked interstellar clouds with IRAC due to their distinct emission in the IRAC channels (cf.~Fig.~\\ref{fig:n132dknots}). The IRAC emission knots seen in the northern central part and the northwestern rim of the remnant coincide with previously detected, bright optical \\oiiif\\ and \\siif\\ emission complexes named ``Lasker's Bowl'' and ``West Complex''. They were identified as shocked interstellar clouds/circumstellar material based on optical line diagnostics, elemental abundances, and their radial velocities similar to the local ISM surrounding N132D \\citep[][]{morse1995,morse1996,blair2000}. The small knot visible close to the center of the remnant is probably a smaller analogue of these two complexes. We do not detect the red- and blueshifted, oxygen-rich ejecta of N132D with IRAC (cf.~the dotted contours in Fig.~\\ref{fig:n132dknots}).  The measured IRAC band ratios for the shocked clouds are roughly consistent with thermal emission from hot silicate/graphite dust grains. In particular, the West Complex shows an IRAC 4.5/8\\mic\\ ratio of 0.18 and a 3.6/5.7\\mic\\ ratio of 0.38. For this region, \\citet{blair2000} presented {\\it Hubble} FOS spectra that show a broad range of ionization stages as expected from shock heating, e.g.~O to O$^{3+}$ and N to N$^{4+}$, and elemental abundances consistent with LMC interstellar material. They derive a high electron temperature $T\\approx43,000$\\,K, a mean electron density of several thousand per \\ccm, and a preshock cloud density larger than 100\\ccm. Their shock models indicate that a shock speed of at least $\\sim$\\,200\\km\\ is needed to produce the observed high-ionization species.  Comparing this to shocked dust IR emission models calculated by \\citet{dwek1996}, we find that the measured IRAC color ratios are roughly reproduced by models with a preshock density of 100\\ccm\\ and a shock speed \\raisebox{0.5mm}{\\scriptsize $\\lesssim$}\\,400\\km, and models with a preshock density of 10\\ccm\\ and a range of shock speeds between 400 and 1000\\km. This suggests that part of the IRAC emission in the West Complex may be due to a population of shock-heated, small dust grains. However, it is likely that the measured IRAC colors represent a mix of dust continuum and ionic emission. The green channel 2 excess emission probably includes \\hi\\,Br$\\alpha$\\,4.05\\mic, since the knots appear in the H$\\alpha$ image from \\citet{morse1995} and the optical spectrum of the West Complex shows strong hydrogen recombination lines \\citep{blair2000}. The red channel 3 excess emission could also be due to \\feiif\\,5.34\\mic. If \\htwo\\ is present in those clouds, the IRAC 4.5\\mic\\ band covers the more highly excited lines, i.e.~the \\htwo\\,(0,0) S(8)-S(11) transitions, whereas the 5.7 and 8\\mic\\ bands cover \\htwo\\,(0,0) S(4)-S(7).  A different type of IRAC emission is mainly characterized by an excess in channel 2, i.e.~it appears as pure green emission in Fig.~\\ref{fig:n132dknots}, right panel. The faint green emission aligned along the southwestern rim is spatially coinciding with small clumps of gas seen in the {\\it HST} WFPC2 images \\citep{morse1996}. Judging solely from the IRAC band fluxes, it is difficult to determine the nature of the IR emission in these clumps. Unfortunately, none of them are covered by our IRS spectra. Many other isolated patches showing channel 2 excess emission do not have a counterpart in the narrow band {\\it HST} WFPC2 \\oiiif, \\siif, nor the H$\\alpha$ image from \\citet{morse1995}. However, the brightest green patches, one near the center of the ejecta, two blended patches near the northeastern edge of the remnant, and one outside the remnant perimeter near the northeastern corner of the image, have optical counterparts visible in archival, broad band {\\it HST} ACS data (proposal 12001, filters F475W, F550M, F850LP; see the related Hubble Heritage image, STScI-2005-30 news release). All counterparts are structured, extended sources with color variations. The object outside the remnant perimeter is most likely a background galaxy, whereas the nature of the objects within the perimeter of N132D remains unclear.   \\subsection{MIPS continuum emission from shock-heated grains} \\label{sec;MIPSdiscussion} We interpret the steeply rising continuum in the spectrum of the southeastern rim of N132D as thermal emission from dust grains that have been swept up and shock-heated by the supernova blast wave. Most of the MIPS\\,24\\mic\\ flux is then due to dust continuum flux, since the \\oivf\\,25.9\\mic\\ line flux is much weaker than the continuum in the wide MIPS\\,24\\mic\\ bandpass (cf.~Fig.~\\ref{fig:n132dspec+bbfit}). The association of the MIPS\\,24\\mic\\ flux with the blast wave along the remnant perimeter and not with the actual ejecta (cf.~Fig.~\\ref{fig:n132dmipsxray} and \\ref{fig:n132dknots}) confirms that the dust is part of the swept-up, ambient interstellar/circumstellar material. This is further supported by the appearance of the brightest MIPS emission along the southeastern rim, where the interstellar gas density is expected to increase due to the presence of a molecular cloud (cf.~CO contours in Fig.~\\ref{fig:n132d}, and Sect.~\\ref{sec:discussion_spatialvariation}). Generally, the MIPS\\,24\\mic\\ flux intensity follows the X-ray emission very well. Every peak in the MIPS\\,24\\mic\\ surface brightness is also an X-ray peak (cf.~Fig.~\\ref{fig:n132dmipsxray}), which demonstrates the connection between the X-ray emitting plasma and the heated dust grains.  A two-component modified-blackbody fit to the continuum yields temperatures of 110 and 58\\,K (cf.~Sect.~\\ref{sec;IRS}). These temperatures can be adopted as average dust temperatures if the grains are in thermal equilibrium, i.e.~balanced heating and cooling rates, and if the mean energy per grain heating event is much smaller than the thermal energy content of the grain at the equilibrium temperature \\citep[steady vs.~stochastic heating, cf.][ and references therein]{draine2001}. Under the conditions in the rim of N132D, i.e.~an average temperature $T\\sim8.4\\times10^6$\\,K and a mean preshock density \\raisebox{0.5mm}{\\scriptsize $\\gtrsim$}\\,3\\ccm\\ \\citep[cf.][]{hwang1993,morse1996}, gas-grain collisions probably dominate the grain heating \\citep{dwek1992}. If each gas-grain collision transfers an energy of the order of ${\\rm k}T\\approx1\\times10^{-16}$\\,J, i.e.~similar to the energy of a 0.6\\,keV X-ray photon, an individual heating event changes the temperature of a silicate/graphite grain with radius \\raisebox{0.5mm}{\\scriptsize $\\gtrsim$}\\,0.04\\mic\\ by less than 1\\,K. This estimate assumes spherical, homogenous grains with a density of $\\rm 2500\\,kg\\,m^{-3}$, and a specific heat capacity of $\\rm 185\\,J\\,kg^{-1}K^{-1}$ for silicate/graphite grains at $T\\sim100$\\,K. Small grains \\raisebox{0.5mm}{\\scriptsize $\\lesssim$}\\,0.04\\mic\\ in size can be heated substantially by a single heating event, and thus do not have a well-defined equilibrium temperature \\citep[cf.][]{li2001}. The fact that we observe two different average grain temperatures may indicate two distinct grain size populations and/or emission originating from separate physical regions with different heating conditions (see~Sect.~\\ref{sec;IRSgrainprocessing_discussion} for a more detailed discussion).  We can estimate the dust mass $M_{\\rm  dust}$ in those two components using \\citep{whitt2003} \\begin{equation}\\label{eqn;dustmass} M_{{\\rm{dust}}}  = \\frac{{4\\rho F_\\lambda  d^2 }}{{3B_\\lambda  (T_{{\\rm{dust}}} )}}\\left[ {\\frac{a}{{Q_\\lambda }}} \\right], \\end{equation} and adopting a density $\\rm \\rho=2500\\,kg\\,m^{-3}$ for silicate/graphite grains, a distance $d=50 \\rm \\,kpc$ for N132D, and an average ratio of grain radius $a$ over emissivity $Q_\\lambda$ of $1\\times10^{-5}\\rm m$. We evaluated the flux density $F_\\lambda$ and the Planck function $B_\\lambda$ for the fitted dust temperatures $T_{\\rm dust}$ at 23.7\\mic. Eq.~\\ref{eqn;dustmass} is an approximation assuming spherical dust grains of uniform size, composition, and in thermal equilibrium. We derived the dust masses in the entire southern rim (outlined by the green color in Fig.~\\ref{fig:n132dmipsxray}, left panel) by scaling the result from Eq.~\\ref{eqn;dustmass} by the MIPS\\,24\\mic\\ flux ratio of the southern rim and the IRS slit aperture, which gives $M_{{\\rm{dust,\\,58\\,K}}}\\sim0.05\\rm\\,M_\\odot$ and $M_{{\\rm{dust,\\,110\\,K}}}\\sim0.003\\rm\\,M_\\odot$. If we further scale this result to the total MIPS\\,24\\mic\\ flux of the whole remnant, the total swept-up ISM dust mass in N132D becomes about $\\rm 0.13\\,M_\\odot$. These numbers have considerable uncertainty, particularly the low temperature component, because of the assumptions made in Eq.~\\ref{eqn;dustmass} and the measurement uncertainties.  We can compare the derived dust masses with estimates of the swept-up gas mass from X-ray data, which give an approximate upper limit on the total amount of dust. \\citet{favat1997} fitted the BeppoSAX X-ray spectrum of N132D with a two-temperature non-equilibrium ionization model. Using the revised best-fit from \\citet{favat1998}, the total mass of the  X-ray emitting material is about 380--850\\msun, assuming a filling factor between 0.1 and 0.5. \\citet{hughe1998} applied Sedov models to fit {\\it ASCA} spectra of N132D, and derived a swept-up mass of about 650\\msun. Assuming a total swept-up gas mass of 600\\msun\\ for N132D, the amount of swept-up dust is $\\sim$\\,1.5\\msun, if the fraction of dust is 0.25 percent of the gas mass, i.e.~about one fourth of the Galactic value. This dust mass represents an approximate upper limit for the case of no dust destruction. Considering the factor-of-ten lower dust mass estimated directly from Eq.~\\ref{eqn;dustmass}, we can state that although some portion of the dust in N132D is surviving, as much as 90 percent might have been destroyed by the blast wave shock (cf.~Sect.~\\ref{sec;IRSgrainprocessing_discussion}).    \\subsection{Grain processing and destruction in N132D} \\label{sec;IRSgrainprocessing_discussion} \\label{sec:Dwek_IRmodel_spectra} Strong supernova shock waves compress, heat, and accelerate interstellar material, creating physical conditions that destroy interstellar dust grains. The grains are not instantly following the gas acceleration due to their inertia, but are gradually accelerated via gas-grain collisions, which cause a gradual grain erosion through sputtering \\citep[cf.][]{tiele1994}. Furthermore, the grain population develops velocity differences, and subsequent grain-grain collisions cause shattering and vaporization depending on the relative velocity and grain sizes \\citep{tiele1994,borko1995,jones1996}.  For N132D, \\citet{morse1996} have estimated a mean shock velocity of the X-ray emitting main blast wave of $\\sim$\\,800\\km, a mean preshock density of $\\sim$\\,3\\ccm, and a postshock temperature of $8.4\\times10^6$\\,K \\citep[cf.~also][]{hwang1993}, which is in good agreement with the Sedov model values from \\citet{hughe1998}. For such a high-velocity shock, the dominant grain destruction process is thermal sputtering \\citep{dwek1992,tiele1994,borko1995}. The lifetime $\\tau$ of silicate/graphite grains against this process was estimated by \\citet{draine1979}, and further refined by \\citet{tiele1994}, \\begin{equation}\\label{eqn;thermsput} \\tau  = \\left. {\\begin{array}{l} 0.79\\ {\\rm for\\ graphite}\\\\ 2.25\\ {\\rm for\\ silicate} \\end{array}} \\right\\}\\times 10^4 {\\rm{yr}}\\left( {\\frac{{{\\rm{cm}}^{{\\rm{ - 3}}} }}{{n_{\\rm{H}} }}} \\right)\\left( {\\frac{a}{{0.01{\\rm{\\mu m}}}}} \\right) \\end{equation} at a temperature of $8.4\\times10^6$\\,K. Thus, if the hydrogen density is $n_{\\rm H}=10\\ccm$, silicate/graphite grains with a radius $a=0.01\\mic$ are completely eroded after roughly 1500\\,yr. Likewise, 0.001\\mic\\ grains only survive 150\\,yr, whereas 0.1\\mic\\ grains would survive 15,000\\,yr.  \\citet{hughe1987}, \\citet{morse1996}, and \\citet{hughe1998} proposed that the supernova of N132D exploded most likely within a low-density cavity in order to reconcile the Sedov age with the dynamical age without invoking large explosion energies. This scenario would allow the blast wave to expand unhindered in the beginning until it encountered the denser material at the boundary of the cavity. If we assume an initial, average expansion velocity of $\\sim\\!10^4$\\km, the shock reached the cavity wall within about 800 years, supposing that the edge of the cavity lies close to the current MIPS\\,24\\mic\\ rim in the south of N132D and adopting the common center of the oxygen-rich ejecta as the expansion origin \\citep[cf.][]{morse1995}. Since the dynamical age of N132D is $\\sim$\\,2500\\,yr \\citep[cf.][]{morse1995}, we can infer that the shock has been interacting with the denser material for maybe 1700\\,yr. Hence, we would expect that most of the grains \\raisebox{0.5mm}{\\scriptsize ${\\lesssim}$}\\,0.01\\mic\\ that were {\\it initially} present are destroyed by now through thermal sputtering {\\it if} they were continuously residing in the hot plasma (cf.~Eq.~\\ref{eqn;thermsput}). Naturally, as the shock wave keeps expanding into new ISM material, the grain population is continually replenished.  In addition to thermal sputtering, shattering and vaporization remove some fraction of the large-size grains upon grain-grain collision. This is limited by the probability of a grain-grain encounter before being brought to rest with respect to the gas, which is of the order of the dust-to-gas mass ratio in the absence of betatron acceleration \\citep{draine1993,Draine2004}. Therefore, we conclude that thermal sputtering generally dominates the grain destruction in the fast shock of N132D, but shattering may redistribute the grain size distribution appreciably in denser clumps \\citep[cf.][]{jones1996}, where the shock speed may be much lower than the estimated mean value of $\\sim$\\,800\\km.  We compare our results to calculated model spectra from dust grains processed in fast, nonradiative shocks by \\citet{dwek1996}. Figure \\ref{fig:n132dspec_dwekmodels} shows the spectrum from the southeastern rim of N132D together with our modified-blackbody fit components and a series of calculated shock model spectra adapted from \\citet{dwek1996}. Their model assumes silicate/graphite grains in a power-law size distribution after \\citet{mathi1977} and grain destruction solely by sputtering without shattering/vaporization.  Overall, the scaled model spectra reproduce the observed spectrum of N132D remarkably well considering the model parameters are not fitted to the data. The applied scaling factors probably account for the filling factor of the emitting material, and were chosen empirically to match the flux of N132D. The main difference between the observed spectrum and the models is a lack of continuum emission at 6--9\\mic\\ and around 70\\mic, although the MIPS\\,70\\mic\\ flux measurement does not provide a very stringent constraint due to the background confusion. The long-wavelength emission may be due to high density material, e.g.~a swept-up dense knot, with a shock velocity lower than 400\\km. The lack of 6--9\\mic\\ flux in the models confirms that this emission is due to PAHs and not from a population of small, heated silicate/graphite grains. Note that the narrow feature near 11.5\\mic\\ in the calculated spectra is not the 11.3\\mic\\ PAH C-H bending mode but an IR active mode of crystalline graphite \\citep[cf.][]{drain1984}, which would likely be smoothed out in an imperfect polycrystalline sample \\citep{li2001}.  Our high-temperature modified-blackbody fit component is very similar to the 1500\\km/10\\ccm\\ shock model, which justifies that our continuum fit is physically meaningful, particularly for the PAH C-C-C emission in the 15--20\\mic\\ range (Fig.~\\ref{fig:n132dspec-bbfit}). Generally, however, a single preshock density and/or a single shock velocity are apparently not sufficient to explain the continuum of the observed spectrum. The medium into which the shock wave expands is clumpy \\citep[cf.~{\\it HST} WFPC2 images of N132D by][]{morse1996}, so we expect a combination of different densities and shock speeds.  \\begin{figure}[!htbp] \\epsscale{1.15} \\centering \\plotone{f10_color.eps} \\caption{N132D southeastern rim spectrum with modified-blackbody fit components (dashed lines, cf.~Fig.~\\ref{fig:n132dspec+bbfit}) and calculated shock models \\citep[][solid lines]{dwek1996}. The models are for 1500/10 (a), 1000/10 (b), 800/10 (c), and 400\\km/10\\ccm\\ (d) shock speed/preshock density, and are scaled by factors of 0.13, 0.14, 0.21, and 0.5, respectively.} \\label{fig:n132dspec_dwekmodels} \\end{figure}  \\subsubsection{Comparison with previous observations} \\label{sec:previousobservation} Figure~\\ref{fig:n132dspec_rim+knot+cloud+ngc} shows the spectrum of the galaxy NGC\\,7331 \\citep[cf.][]{smith2004} in comparison to our spectra for N132D. The spectrum of NGC\\,7331 bears a striking similarity to the spectrum of the molecular cloud/\\hii-region south of N132D, and also shows \\neiiif\\ and \\oivf/\\feiif\\ lines. Since it represents an average galactic ISM mixture, the resemblance to our spectra is not surprising. \\citet{smith2004} have fitted the 15--20\\mic\\ PAH features in this spectrum with two Drude profile components at 16.40 ($\\rm FWHM=0.34\\mic$) and 17.12\\mic\\ ($\\rm FWHM=0.96\\mic$), and measured a $16.4+17.1$ PAH flux of $1.5\\times10^{-7}$\\wms. This fit reproduces the overall shape of the emission in this region together with the \\htwo\\,S(1) component, but it does not take the blended features near 17.4 and 17.8\\mic\\ into account, which are also observed in other sources (cf.~Fig.~\\ref{fig:n132dspec_sources}). \\citet{smith2004} reported the strong 17.1\\mic\\ feature as a possible new PAH C-C-C bending mode detection. If we adopt Lorentzian components of narrower width, we find the same PAH components with similar fluxes and widths as seen in our spectra of the molecular cloud/\\hii-region near N132D (cf.~Table~\\ref{tab;gaussfit} and Fig.~\\ref{fig:n132dspec_sources}), and no need for a broad feature near 17.1\\mic. The 0.2--0.4\\mic\\ widths of these features are plausible for PAH molecules in the gas phase \\citep[][]{peete2004}. Our refitted, total PAH 15--20\\mic\\ flux in NGC\\,7331 is $1.1\\times10^{-7}$\\wms, which is a factor of 1.6 larger than observed in the southeastern rim of N132D, whereas the PAH band at 11.3\\mic\\ is about a factor 30 stronger in NGC\\,7331.  \\begin{figure}[!tb] \\epsscale{1.15} \\centering \\plotone{f11_color.eps} \\caption{Comparison of continuum subtracted 15--22\\mic\\ {\\it Spitzer} IRS emission spectra. All spectra are normalized to their 16.4\\mic\\ flux, and shifted along the flux axis for illustration purposes. We show our fit for the molecular cloud/H\\,{\\sc ii}-region south of N132D (online color version: red; cf.~Table~\\ref{tab;gaussfit}). IRAS\\,22308+5812 (compact H\\,{\\sc ii}, high-resolution mode, $R\\sim600$) has a unique, broad feature near 19.6\\mic, which also appears in {\\it ISO} SWS data presented by \\citet{vanke2000}, and weak emission longward of 20\\mic. See Table\\,\\ref{tab:pahlineratios} for references.} \\label{fig:n132dspec_sources} \\end{figure}  \\citet[][hereafter VHP]{vanke2000} have measured 15--20\\mic\\ PAH plateau strengths as well as the PAH bands at 6.2 and 11.3\\mic\\ for a variety of Galactic sources from {\\it ISO} SWS spectra. Table~\\ref{tab:pahlineratios} shows a comparison of different PAH band ratios from the {\\it ISO} results by \\citetalias{vanke2000} together with reanalyzed {\\it Spitzer} archival data and our results for N132D. A comparison of the absolute fluxes is hampered by the larger {\\it ISO} beam and different source distances, so we restrict ourselves to band ratios using \\wms\\ units instead. Note also that \\citetalias{vanke2000} measured the 16.4\\mic\\ PAH features with a local continuum, i.e.~just the excess portion of the feature on top of the 15--20\\mic\\ emission hump, whereas we generally use the mostly emission-free flux below 15 and above 22\\mic\\ to determine the continuum for fitting the 15--20\\mic\\ features. We analyzed {\\it Spitzer} archival calibration data for two sources in the \\citetalias{vanke2000} sample, IRAS\\,22308+5812 and NGC\\,7027, in order to test whether the results are comparable. We find that the measured PAH ratios for these two sources match within 30 percent, which is a reasonable agreement given the uncertainties due to the different beam sizes and continuum choices for the fitting.  \\citetalias{vanke2000} interpret the observed relative strengths of the broad plateaus and the narrower PAH bands in terms of variations in the size distribution, i.e.~the larger PAHs with a few thousand carbon atoms that give rise to the 15--20\\mic\\ emission probably originated from dense cloud cores that formed the massive stars in the \\hii-regions, whereas smaller PAHs with less than a few hundred carbon atoms give rise to strong 6.2--11.3\\mic\\ bands. They also tentatively attribute the discrete 16.4\\mic\\ band to \\mbox{C-C-C} bending modes of pendant benzene ring substructures in smaller PAH molecules \\citep[cf.~also][who propose PAH molecules containing pentagonal rings as the carrier of the 16.4\\mic\\ band]{mouto2000}.  \\begin{table}[!htbp] \\centering \\caption{PAH emission band ratios in different types of sources.}  \\begin{tabular}{lccccc} \\hline\\hline\\\\[-0.3cm] Source &                   Type    & $\\frac{{15-20\\mic}}{{11.3\\mic}}$ &  $\\frac{{16.4\\mic}}{{11.3\\mic}}$  & $\\frac{{15-20\\mic}}{{6.2\\mic}}$& $\\frac{{16.4\\mic}}{{6.2\\mic}}$ \\\\ % \\hline\\hline\\\\[-0.3cm] LkH$\\alpha$\\,234 \\htwo-ridge &            PDR &          0.2 &       0.08  &        0.1 &  0.04  \\\\%    &2.5\\\\  NGC\\,7331 (average) &                  Galaxy &        0.3 &       0.07  &        0.2 &   0.05 \\\\%    &4.3 \\\\  NGC\\,7023 &                  RN &          0.3 &       0.14  &        0.1 &    0.06 \\\\%   & 1.9\\\\  NGC\\,7027 (edge) &             PN &        0.4 &       0.04  &        1.2&   0.13  \\\\%   &9.6 \\\\  *NGC\\,7027 &                     PN &        0.5 &       0.04  &        0.7&  0.06   \\\\%   &11.8 \\\\  *CD-42$\\arcdeg$11721 &                    YSO &        0.7 &       0.10  &        0.6&  0.08   \\\\%    &7.1 \\\\  *IRAS 03260+3111 &                  HAeBe &        1.0 &       0.06  &        0.8 &  0.05  \\\\%    &17.0\\\\  *IRAS 23133+6050 &                   cH\\,{\\sc ii} &        2.2 &     $-$     &    1.5         &  $-$  \\\\%  & $-$ \\\\  *IRAS 22308+5812 &                   cH\\,{\\sc ii} &        2.3 &     $-$      &     1.8         &  $-$ \\\\%   & $-$ \\\\  IRAS 22308+5812 &                   cH\\,{\\sc ii} &        2.5 &    0.27         &        0.8  &  0.08 \\\\%   &9.5     \\\\  *IRAS 18317-0757 &                   cH\\,{\\sc ii} &        2.5 &   $-$        &    1.2         &  $-$ \\\\%   &$-$ \\\\  *Sh\\,2-106 IRS4 &            H\\,{\\sc ii} &        5.1 &     $-$       &      3.9         & $-$ \\\\%  &$-$\\\\  \\hline\\\\[-0.3cm]  N132D SE-rim &                     SNR &        7.1$\\pm$1.0 & $-$ & \\raisebox{0.5mm}{\\scriptsize ${>}$ }4 &   $-$  \\\\%   & $-$\\\\  N132D local backgr. &              &        0.7$\\pm$0.2    &    0.18         &     \\raisebox{0.5mm}{\\tiny ${\\gtrsim}$ }0.5     &  \\raisebox{0.5mm}{\\tiny ${\\gtrsim}$ }0.13 \\\\%  &4.1     \\\\   \\hline\\\\[-0.1cm] \\end{tabular} \\tablecomments{Sources denoted by an asterisk are {\\it ISO} SWS measurements taken from \\citetalias{vanke2000} and converted to \\wms. The other sources are {\\it Spitzer} archival data processed and fitted in the same manner than N132D (NGC\\,7331, see \\citet{smith2004}; LkH$\\alpha$\\,234, see \\citet{morri2004}; NGC\\,7023, see \\citet{werne2004}; NGC\\,7027 and IRAS\\,22308+5812, {\\it Spitzer} calibration observations). Dashes indicate that a discrete 16.4\\mic\\ PAH band is not detected in that source. PDR\\,=\\,photodissociation region; RN\\,=\\,reflection nebula; PN\\,=\\,planetary nebula; YSO\\,=\\,young stellar object; HAeBe\\,=\\,Herbig Ae/Be star; (c)H\\,{\\sc ii} \\,=\\,(compact)H\\,{\\sc ii}-region.} \\label{tab:pahlineratios} \\end{table}  Table~\\ref{tab:pahlineratios} is sorted to show the trend of rising PAH 15--20\\mic\\ to 11.3\\mic\\ ratio. For a diffuse ISM mix of starlight-heated, neutral and ionized PAHs, 15--20\\mic\\ features are efficiently produced by PAH particles in the 8--25\\,\\AA\\ size range (300 \\raisebox{0.5mm}{\\scriptsize ${\\lesssim}$} $ N_{\\rm C}$ \\raisebox{0.5mm}{\\scriptsize ${\\lesssim}$} $10^4$ carbon atoms). The 11.3\\mic\\ feature can be produced by particles as large as about 20\\,\\AA\\ or $ N_{\\rm C}\\approx4000$, but more efficiently by particles with radius $a\\,\\raisebox{0.5mm}{\\scriptsize ${\\lesssim}$}\\,10$\\,\\AA\\ or $N_{\\rm C}\\,\\raisebox{0.5mm}{\\scriptsize ${\\lesssim}$}\\,500$ \\citep{draine2006}. A direct, quantitative comparison of the observed PAH 15--20\\mic\\ to 11.3 and 6.2\\mic, or the 16.4\\mic\\ to 11.3 and 6.2\\mic\\ ratios with model calculations is difficult for a number of reasons. Foremost, there is a unique PAH size distribution, which differs from source to source. Secondly, individual, discrete PAH features may be due to specific substructures in PAH molecules of a certain size (cf.~interpretation of the 16.4\\mic\\ feature above). Finally, the specific local physical conditions affect the charge state, the amount of dehydrogenation, and the heating process of the PAH molecules.  Nevertheless, it is illustrative to note some general trends when comparing the observed PAH emission ratios in Table~\\ref{tab:pahlineratios} with calculated PAH band radiation efficiencies as function of size for a diffuse ISM mix of neutral/ionized PAHs \\citep[cf.][]{draine2006}. We approximate the predicted, total emitted power in the 15--20\\mic\\ region by scaling the calculated 17.4\\mic\\ band power with the total integrated absorption cross section per C atom in the 15--20\\mic\\ range divided by the cross section for the 17.4\\mic\\ band. This gives approximately a scale factor of 4, using the cross sections per C atom adopted by \\citet{draine2006}. The lowest observed 15--20/11.3, 15--20/6.2, 16.4/11.3, and 16.4/6.2\\mic\\ ratios in Table~\\ref{tab:pahlineratios} correspond then to the smallest PAHs considered in the calculations by \\citet{draine2006}, which contain about 20 carbon atoms. PAHs smaller than about 20--50 C atoms are photolytically unstable \\citep[cf.][]{allam1989,jochi1994,allai1996b,lepag2003}. Hence, the PAH size distribution of the sources near the top of Table~\\ref{tab:pahlineratios} is probably dominated by PAHs of that size. This is generally consistent with earlier estimates of PAH sizes derived from the 3.3 and 3.4\\mic\\ C-H stretching modes \\citep[][Table~3]{allam1989}, and from profile fitting of the PAH 3.3-11.3\\mic\\ bands with calculated spectra from a distribution of small PAHs \\citep{verst2001,pech2002}.  For the supernova remnant N132D, we expect that small grains/PAHs are rapidly destroyed by thermal sputtering (cf.~Sect.~\\ref{sec;IRSgrainprocessing_discussion}), i.e.~relatively larger grains/PAHs tend to dominate the size distribution, which is consistent with the high PAH 15--20\\mic\\ to 11.3\\mic\\ ratio and the absence of discrete 15--20\\mic\\ features. A PAH 15--20\\mic\\ to 11.3\\mic\\ ratio of $\\sim$\\,7 as observed in N132D corresponds to PAH sizes of about 20\\,\\AA\\ or 4000 C atoms, if we adopt the calculations by \\citet{draine2006} as a tentative guideline. This number is a rough estimate, since electron collisions rather than starlight dominate grain heating in N132D, changing the amount and frequency of energy deposition into grains compared to the model calculations. In addition, even large PAHs may be dehydrogenated in the hot plasma of N132D, which provides another rationale for the weak 11.3\\mic\\ feature and the resulting high 15--20/11.3\\mic\\ ratio. As a consequence, the size estimate from that ratio is probably an upper limit, since PAHs in N132D have most likely a higher degree of dehydrogenation than in the diffuse ISM. A lower limit estimate of the PAH sizes in N132D from the 15--20/6.2\\mic\\ ratio is about 500 C atoms, which is similar to the sizes of the very small grains/large PAHs proposed as carriers of the broad plateaus underlying the 6--14\\mic\\ PAH features \\citep[e.g.][and cf.~Sect.~\\ref{sec:discussion_PAHlines} ]{allam1989,peete2002,rapac2005}.  The lifetime of 500--4000 C atom grains, about 0.001--0.002\\mic\\ in size, is 80--160\\,yr according to Eq.~\\ref{eqn;thermsput}, which describes thermal sputtering of silicate/graphite grains in the forward shock of N132D. We presume that large PAHs/amorphous carbon particles behave similarly to graphite in terms of stability against thermal sputtering, which is a reasonable assumption since both consist mainly of aromatic carbon ring systems. Thus, it is plausible for large PAHs to survive thermal sputtering in N132D long enough to maintain a detectable population, particularly since both PAHs and grains will be replenished as the shock wave keeps expanding into new ISM material."
    },
    "0609/astro-ph0609305_arXiv.txt": {
        "abstract": "This paper explores the stability of an Earth-like planet orbiting a solar-mass star in the presence of a stellar companion using $\\sim 400,000$ numerical integrations.  Given the chaotic nature of the systems being considered, we perform a statistical analysis of the ensuing dynamics for $\\sim 500$ orbital configurations defined by the following set of orbital parameters: the companion mass $M_C$; the companion eccentricity $e$; the companion periastron $p$; and the planet's inclination angle $i$ relative to the stellar binary plane. Specifically, we generate a large sample of survival times ($\\tau_s$) for each orbital configuration through the numerical integration of $N \\gg 1$ equivalent experiments (e.g., with the same orbital parameters but randomly selected initial orbital phases). We then construct distributions of survival time using the variable $\\mu_s \\equiv$ log $\\tau_s$ (where $\\tau_s$ is in years) for each orbital configuration. The primary objective of this work is twofold.  First, we use the mean of the distributions to gain a better understanding of what orbital configurations, while unstable, have sufficiently long survival times to make them interesting to the study of planet habitability.  Second, we calculate the width, skew, and kurtosis of each $\\mu_s$ distribution and look for general features that may aid further understanding and numerical exploration of these chaotic systems. To leading order, most distributions are nearly Gaussian with a width $\\sigma \\sim 0.5$, although the longest-lived systems display substantial (non-Gaussian) tails. As a result, many independent realizations of these systems must be considered in order to characterize the survival time.  The situation is more complicated for orbital configurations with longer mean survival times, owing in part to the increasing importance of resonances. ",
        "introduction": "Recent (and ongoing) discoveries of exoplanetary systems (e.g., Butler et al. 1999; Marcy et al. 2001; Fisher et al. 2002; Tinney et al. 2002; McArthur er al. 2004) have shown that Sun-like stars are orbited by planets with a wide variety of orbital configurations.  Of course, current planetary searches are biased toward large bodies with short orbital periods, resulting in the discovery of planets with masses ranging from $\\sim 0.01 - 10 M_J$ and semi-major axis typically ranging from $\\sim 0.04 - 4$ AU.  It is expected that Earth-like planets will also form alongside their Jovian counterparts (e.g., Ruden 1999; Lissauer 1993), even in binary stellar systems (e.g., Whitmire et al. 1998; Quintana et al. 2002; Lissauer et al. 2004).  The discovery of exoplanetary systems has spurred a renewed interest in planetary dynamics, with a significant amount of attention being given to the stability of specific multi-planet systems (e.g., Laughlin \\& Adams 1999; S\\\"uli et al. 2005) and to the stability of (hypothesized) terrestrial planets located within a specific system's habitable zone (e.g., Noble et al. 2002; Asghari et al. 2004; Jones et al. 2005; Haghighipour 2006).  In addition, the identification of Jupiter-like planets in orbit around members of multiple star systems (e.g., Eggenberger et al. 2004 and references therein) has also provided an observational basis for studying planetary stability in binary systems (e.g., Harrington 1977; Pendleton \\& Black 1983; Rabl \\& Dvorak 1988; Holman \\& Wiegert 1999; David et al. 2003; Pilat-Lohinger et al. 2003; Musielak et al. 2005; Mudryk \\& Wu 2006). Of course, one must differentiate between stability within a habitable zone and the more general study of overall system stability, with the former imposing stricter conditions on a planet's allowed orbital motion.  Planetary orbits in binary systems can be found with a wide range of configurations.  The two most common orbital classes are P-type orbits, in which the planet's semi-major axis is larger than that of the binary so the planet orbits about both stars, and S-type orbits, in which the planet revolves around one of the stars with a semi-major axis smaller than that of the binary (e.g., Szebehely 1967, Dvorak et al.  1989, Pilat-Lohinger \\& Dvorak 2002). If the semi-major axis of the planetary orbit is sufficiently large (for P-type orbits) or sufficiently small (for S-type orbits), the orbital motion is stable and well-ordered.  However, in binary configurations for which the stellar bodies come sufficiently close to the planet, the motion can be chaotic and even unstable. One goal of this paper is to determine the regime of binary parameter space for which S-type planetary orbits are stable.  In the regime of parameter space near the border between stability and instability, the orbits tend to be chaotic.  A complete analysis of planetary stability in stellar binaries must reconcile the underlying chaotic nature of the systems being explored. While theoretical conditions can be used to determine whether a planetary system is unstable, it is still possible for such systems to last for vast spans of time.  Numerical work must then be used to delineate the boundaries of ``effectively'' stable parameter space. However, for a given orbital configuration, the survival time $\\tau_s$ of an unstable system (defined throughout this work as the number of years it takes for the planet to either be ejected or collide with either star) varies widely depending on the choices of initial orbital phases. Because the systems are chaotic, this variation is not smooth, i.e., small differences in the starting phase angles can lead to large differences in the resulting survival times.  As a result, the survival time for any set of binary properties can only be fully characterized in terms of a distribution of output measures. Further, as shown herein, the distribution of survival times has a substantial width.  Even though the systems are chaotic in the regime of interest, and hence display a wide distribution of survival times for effectively equivalent starting conditions, we stress that the distributions themselves are well-defined. As a result, the answer to the question -- How long does a planet survive in any given binary system? -- is a full distribution.  One goal of this work is to find such distributions for binary systems with intermediate semimajor axes $a \\approx 1 - 50$ AU.  A statistical analysis of the stability of an Earth-like planet orbiting a one solar-mass star in the presence of an outer-lying companion was recently performed by David et al. (2003; hereafter D03), leading to an estimate of the fraction ${\\cal F}_b$ of binary star systems that allow Earth-like planets to remain in the system over a time scale of 4.6 Gyr.  Specifically, D03 found that ${\\cal F}_b \\approx 0.5$ by calculating the survival time of an Earth-like planet (with an initial circular orbit of R = 1 AU) for a range of companion masses $M_C = 0.001 - 0.5 M_\\sun$, initial eccentricities $e$, and semi-major axis $a$.  For the parameter space explored by D03, the planet's survival time (for a specified companion mass) depends most sensitively on the companion's initial periastron distance $p = a(1-e)$, but spans over two decades when sampled over a range of semi-major axis values (with $e$ then set to give the desired value of $p$).  Nevertheless, the mean value of log $\\tau_s$ exhibits a clear exponential dependence on the periastron.  It is important to note that the range of survival times versus periastron presented in D03 results from both the chaotic nature of the three body problem and the sampling over different ($a, e$) pairs for a given periastron value (see \\S 5.2 for a more complete discussion).  This paper extends previous work on the stability of an Earth-like planet orbiting a Sun-like star in the presence of a stellar companion.  Although much of the previous work focused on coplanar systems (e.g., D03), this work explores the full range of inclination angles.  Specifically, we explore $\\sim 500$ different orbital configurations (defined through the choice of the companion mass $M_C$, eccentricity $e$, and periastron $p$, and the planet's inclination angle $i$ relative to the stellar orbital plane), and we broadly sample the orbital parameter space inhabited by most observed binary systems. Our main body of work explores $376$ different orbital configurations organized in four series (1 -- 4). An additional 136 orbital configurations with low inclinations, organized in three series (5 -- 7), are performed to further explore some of the detailed structure exhibited by the output measures of the first four series. We generate a large sample of survival times ($\\tau_s$) for each orbital configuration through the numerical integration of $N \\gg 1$ equivalent experiments (with the same basic orbital parameters but randomly selected initial orbital phases), performing a total of $\\sim 390,000$ separate orbital simulations; this sample of orbital simulations is thus an order of magnitude larger than our previous study (D03). This broad survey of parameter space provides a cleaner delineation of the orbital configurations that (while potentially unstable) have sufficiently long survival times to be interesting for planetary habitability (as noted above, however, stability within a planet's habitable zone imposes additional constraints on the planet's orbital motion).  Another focus of this new work is the characterization of the distributions of survival times.  As discussed earlier, the regime of parameter space near the stability/instability border is chaotic and the results must be described statistically.  Performing multiple realizations for each orbital configuration is a necessity.  This paper constructs distributions of survival times --- using the variable $\\mu_s \\equiv$ log $\\tau_s$, with $\\tau_s$ defined in years --- for each orbital configuration under consideration. The resulting distributions are then characterized by their mean, width, skew and kurtosis.  We also calculate the fraction of runs that lead to the planet's ejection from the system (and the fraction accreted by one of the stars). With the construction of these distributions, we can study how they vary over the regime of binary parameter space; these distributions also provide important guidance for future work (e.g., the width of the distribution determines how many independent realizations of the numerical problem are necessary to achieve a desired accuracy in estimating the mean value).  The paper is outlined as follows: We discuss the numerical method used in our work in \\S 2.  We present the results of our numerical work for series 1 -- 4 in \\S 3, and characterize the resulting distributions for these series in \\S 4.  We present the results of series 5 -- 7 in \\S 5, and use the results of these runs to explore certain aspects of the rich structure exhibited by these three body-systems. Specifically, we: 1) consider the effect of integer ratios of initial planet/companion orbital periods -- a necessary but not sufficient condition for resonance (see, e.g., Murray \\& Dermott 2000); 2) explore more fully the dependence of ejection time on eccentricity and on periastron; 3) characterize the fraction of ejection events; and 4) consider the stability exhibited by high inclination, low eccentricity orbits.  We present our conclusions in \\S 6. ",
        "conclusions": "This paper extends previous work on numerical simulations of Earth-like planets in binary systems.  Specifically, the survival time of an Earth-like planet orbiting a Sun-like star in the presence of an outer companion is determined for different orbital configurations, defined by the companion's mass, eccentricity, periastron, and the planet's inclination angle relative to the binary orbital plane.  Due to the chaotic nature of the systems being explored, the survival time for each orbital configuration was determined for $N$ equivalent realizations with randomly selected initial phase angles, allowing for a statistical analysis of the survival time.  We used $N = 10^3$ for the bulk of our exploration of parameter space, which focuses on four pairs of companion masses and eccentricities (series 1 -- 4), and $N = 10^2$ for three additional series of runs (5 -- 7).  In all, we performed $\\sim 400,000$ numerical integrations and explored $\\sim 500$ different orbital configurations.  Our results indicate that the values of survival time obtained for a given orbital configuration are log-normally distributed (to leading order).  We therefore use the logarithm of the survival time $\\mu_s \\equiv $ log $\\tau_s$ (where $\\tau_s$ is given in years) as our primary output measure.  We plot the dependence of $\\langle\\mu_s\\rangle$ on inclination angle and periastron value for each distribution in series 1 -- 4. These results confirm a weak dependence of survival time on inclination angle for cases with $i \\simless 40^o$, and the well known decrease in survival time at high ($i \\simgreat 40^o$) inclination angles.  A simple extrapolation of our results for low inclination angles indicates that while potentially Hill unstable, orbital configurations with a companion mass of $M_C = 0.1 \\msun$ and $M_C = 0.5 \\msun$ can nonetheless remain stable for $\\sim 5$ Gyrs when $p \\simgreat 4$ AU and $p \\simgreat 6$ AU, respectively.  Note that these estimates are conservative in that the extrapolation is carried out along the lower envelope of the (wide) range of survival times for a given periastron. These results are consistent with those of D03, and represent lower periastron values than those obtained through a Hill stability analysis (e.g., Gladman 1993; see also Barnes \\& Greenberg 2006), i.e., systems that are ultimately ``unstable'' according to analytic criteria can survive over the 4.6 Gyr age of the Solar System.  This work has important implications for searches for Earth-like planets (e.g., Terrestrial Planet Finder) in extrasolar systems, since a large fraction of stars are found in binaries. As summarized above, our numerical results show that all binaries with periastron greater than 6--7 AU allow Earth-like planets to be stable over the age of our Solar System. This finding, in conjunction with the observed distributions of binary parameters for solar-type primaries (see Duquennoy \\& Mayor 1991 and the compilation of D03), indicates that roughly half of all binaries are wide enough to contain stable Earths. In addition, binaries that are wide enough (with sufficiently large periaston) to allow for orbital stability of Earth-like planets are also wide enough to allow for the formation of terrestrial planets through the accumulation of planetesimals (e.g., Quintana et al. 2002; Quintana 2004; Marzari \\& Scholl 2000). Thus, at least half of the binary systems are habitable according to dynamical considerations.  Some of the extrasolar planets detected to date are found in binary systems, although these planets have masses comparable to that of Jupiter.  In the simulations performed here, the Earth-like planet acts like a test particle, but even a Jovian mass planet will act as a test particle in the potential of stellar bodies. As a result, one can directly apply our stability criteria to these systems: In order for a Jovian planet to remain stable over typical stellar ages of $\\sim5$ Gyr, the binary periastron $p_B$ must be wider than about $\\sim7$ times the semimajor axis $a_J$ of the Jovian planet, i.e., $p_B \\simgreat 7 a_J$.  In all of the binary systems with detected planets, the periastron satisfies this constraint.  A related question is whether or not an Earth-like planet can survive in a binary system that also contains a Jovian planet. Although the relevant 4-body simulations must be left for future work, a rough criterion can be formulated using our results to date: The Jovian planet must have a sufficently large periastron to allow the Eath-like planet to remain stable (and this constriant is roughly $p_J \\simgreat 2.5 - 3$ AU, e.g., D03) and the binary periastron $p_B$ must be large enough to allow the Jovian planet to be stable (roughly $p_B \\simgreat 7 p_J$).  The next major result of this work concerns the width of the distributions of survival time. We have calculated the $\\mu_s$ distributions over a broad range of orbital parameter space. It is well known that orbital systems can exhibit chaotic behavior in or near the unstable regime.  As a result, it is not possible to define a unique value of survival time for a given orbital configuration. Instead, the survival time can only be defined in terms of a distribution -- the mean and width of which then characterizes the underlying dynamics of the system.  For the orbital configurations explored in series 1 -- 4, the resulting $\\mu_s$ distributions exhibit Gaussian peaks with widths $\\sigma \\sim 0.5$ and generally display high-end tails when $\\langle\\mu_s\\rangle \\,\\simgreat 4$.  In this regime, the value of $\\langle\\mu_s\\rangle$ calculated via $N$ equivalent realizations of a given orbital configuration has a 68\\% probability of being within $0.05 \\,\\sqrt{(100/N)}$ of the true value, and a 99\\% probability of being within $0.15 \\, \\sqrt{(100/N)}$ of the true value.  The situation is more complex for orbital configurations whose $\\mu_s$ distributions have means in excess of $\\sim 4$.  While exhibiting peaks with widths $\\sim 0.5$, these distributions may also have significant tails, thereby leading to larger calculated total width values ($\\sim 0.7 - 2$).  Part of this increased complexity can be traced to the increasing importance of mean motion orbital resonance effects.  Additionally, the loss of a tangential component in the impulse imparted as the planet moves through closest approach for low eccentricity cases can also lead to fairly broad distributions, although the paucity of observed low eccentricity binary systems somewhat limits the importance of this latter effect. Nevertheless, our results indicate that the determination of survival times for long-lived systems poses a serious numerical challenge in that both the integration times must be longer and (because of the broader distributions) a larger number of equivalent realizations must be performed in order to get good statistical results.  The form of the distributions of survival times is not unexpected. Whenever a large number of independent variables play a role in determining a distribution, the result tends to exhibit a nearly gaussian form. In the limit of an infinite number of variables, the central limit theorem indicates that the resulting distribution will be gaussian, for any type of distribution for the input variables (e.g., Richtmyer 1978). Binary systems with planets contain an intermediate number of variables. Although the systems are sufficiently complicated to display interesting and complex behavior (Figs. 1 -- 18), the number of independent variables in not infinite. As a result, one expects the distribution of any composite variable (here the survival time for a given orbital configuration) to approach a gaussian form, but retain non-gaussian tails (since convergence is slowest in the tails). The result -- a gaussian distribution with tails -- is in fact what we find here.  As expected from the chaotic nature of the systems being explored, our output measures exhibit a rich amount of structure.  Specifically, we find that the presence of Kozai resonances result in a region of stability at $i \\sim 60^o$ when the companion orbit has a low eccentricity.  In addition, we find a significant amount of structure in the output measure $f_e$ - the fraction of events that lead to ejection for a given orbital configuration.  To our knowledge, this latter result has not received significant attention in the literature, and may warrant further study.  \\bigskip \\centerline{\\bf Acknowledgments}  We thank Matt Holman for useful discussions.  This work was supported at Xavier University by the Hauck Foundation, and by the National Science Foundation under Grant No. 0215836.  This work was supported at the University of Michigan by the Michigan Center for Theoretical Physics, and by NASA through the Terrestrial Planet Finder Mission (NNG04G190G) and the Astrophysics Theory Program (NNG04GK56G0). Finally, we thank the referee for a comprehensive set of comments that improved the manuscript.  \\newpage"
    },
    "0609/astro-ph0609769_arXiv.txt": {
        "abstract": "We present an exact solution of the equations for orbit determination of a two body system in a hyperbolic or parabolic motion. In solving this problem, we extend the method employed by Asada, Akasaka and Kasai (AAK) for a binary system in an elliptic orbit. The solutions applicable to each of elliptic, hyperbolic and parabolic orbits are obtained by the new approach, and they are all expressed in an explicit form, remarkably, only in terms of elementary functions. We show also that the solutions for an open orbit are recovered by making a suitable transformation of the AAK solution for an elliptic case. ",
        "introduction": "Two body problems are very classical in celestial mechanics and have been studied thoroughly since Kepler discovered the laws of motion of celestial objects (e.g., Aitken 1964, Goldstein 1980, Danby 1988, Roy 1988, Murray and Dermott 1999, Beutler 2004). The regular orbits of a system of two masses in Newtonian mechanics are of three types: ellipse, parabola and hyperbola. The latter two cases, in which the separation of two bodies will become infinite in the remote future, may be called open orbits. The singular orbit is a linear one corresponding to a head-on collision, which is extremely special. In this paper, we consider regular orbits mentioned above: for a binary system, the orbit determination may bring us some informations about its formation and evolution mechanism. In an open orbit, one may infer an impact parameter and an initial relative velocity of two masses, one of which may be ejected by some explosive mechanism such as a supernova or by three body scattering. For instance, some observations reveal that kicked pulsars move at unusually high speed (Anderson et al. 1975, Hobbs et al. 2005).  The orbit determination of {\\it visual double stars} was solved first by Savary in 1827, secondly by Encke 1832, thirdly by Herschel 1833 and by many authors including Kowalsky, Thiele and Innes (Aitken 1964 for a review on earlier works; for the state-of-the-art techniques, e.g, Eichhorn and Xu 1990, Catovic and Olevic 1992, Olevic and Cvetkovic 2004). Here, a visual binary is a system of two stars both of which can be seen. The relative vector from the primary star to the secondary is in an elliptic motion with a focus at the primary. This relative vector is observable because the two stars are seen. On the other hand, an {\\it astrometric binary} is a system of two objects where one object can be seen but the other cannot like a black hole or a very dim star. In this case, it is impossible to directly measure the relative vector connecting the two objects, because one end of the separation of the binary, namely the secondary, cannot be seen. The measures are made in the position of the primary with respect to unrelated reference objects (e.g., a quasar) whose proper motion is either negligible or known.  As a method to determine the orbital elements of a binary, an analytic solution in an explicit form has been found by Asada, Akasaka and Kasai (2004, henceforth AAK). This solution is given in a closed form by requiring neither iterative nor numerical methods. One may naturally seek an analytic method of orbit determination for open orbits. An extension for open orbits done earlier by Dommanget (1978) used the Thiele-Innes method, namely solved numerically the Kepler equation. As a result, the method by Dommanget does not provide an explicit solution in a closed form. Therefore, let us extend the explicit solution by AAK to open orbits. Then, we would face the following problem. AAK formalism uses a fact that the semimajor and semiminor axes of an ellipse divide it into quarters. This fact plays a crucial role in determining the position of the common center of mass on a celestial sphere; we should note here that the projected  common center of mass is not necessarily a focus of an apparent ellipse. The division into quarters is possible for neither a parabola nor a hyperbola.  The purpose of this paper is to generalize AAK approach so that we can treat an open orbit. This paper is organized as follows. Sec. 2 presents a generalized AAK formalism. In Sec. 3, the generalized approach is employed to obtain the method of orbit determination for a hyperbolic orbit. In Sec. 4, the formula for a parabolic orbit is presented. In Sec. 5, we recover these formulae by making a suitable transformation of that for an elliptic orbit with some limiting procedures. Sec. 6 is devoted to Conclusion. ",
        "conclusions": "The formulae for orbit determination of elliptic, hyperbolic and parabolic orbits are obtained in a unified manner by generalizing AAK approach, which originally needed a fact of the areal divisions by the semimajor and semiminor axes of an ellipse. We show also that the present formulae are recovered from AAK result by a suitable transformation among an ellipse, hyperbola and parabola."
    },
    "0609/astro-ph0609243_arXiv.txt": {
        "abstract": "We analyse the evolutionary history of galaxies formed in a hierarchical scenario consistent with the concordance $\\Lambda$-CDM model focusing on the study of the relation  between their chemical and dynamical  properties. Our simulations consistently describe the formation of the structure and its chemical enrichment within a cosmological context. Our results indicate that the luminosity-metallicity (LZR) and the stellar mass-metallicity (MZR) relations are naturally generated in a hierarchical scenario. Both relations are found to evolve with redshift. In the case of  the MZR, the estimated evolution is weaker than that deduced from observational works by approximately 0.10 dex. We also determine a characteristic stellar mass, $M_c \\approx 3 \\times 10^{10} M_{\\odot}$, which segregates the simulated galaxy population into two distinctive groups  and which remains unchanged since $z\\sim 3$, with a very weak evolution of its metallicity content. The value and role played by $M_c$ is consistent with the characteristic mass estimated from the SDSS galaxy survey by Kauffmann et al. (2004). Our findings suggest that systems with stellar masses smaller than $M_c$ are  responsible for the evolution of this relation at least from $ z\\approx 3$. Larger systems are stellar dominated and  have formed more than 50 per cent of their stars at $z \\ge 2$, showing very weak evolution since this epoch. We also found bimodal metallicity and age distributions from $z\\sim3$, which reflects the existence of two different galaxy populations. Although SN feedback may affect the properties of galaxies and help to shape the MZR, it is unlikely that it will significantly modify   $M_c$ since, from $z=3$ this stellar mass is found in systems with circular velocities larger than $100 \\kms$. ",
        "introduction": "\\label{intro}  Chemical abundances  could provide important clues on  the evolutionary history of galaxies (e.g., Freemann \\& Bland-Hawthorn 2002) since  they are the  result of the joint action of several physical processes such as Supernova (SN) feedback (Larson 1974; Larson \\& Dinerstein 1975; White \\& Rees 1978; Kauffmann \\& Charlot 1998; Somerville \\& Primack 1999; Cole et al. 2000; Springel \\& Hernquist 2003), gas inflow  (Pagel 1997), mergers and interactions,  among others. And, although it is not easy to disentangle the effect of each individual process, it may be possible to constrain galaxy formation models by studying how chemical and dynamical properties of galaxies evolve as a function of cosmic epoch (Pettini 2006) and confronting them with observations.   There are strong evidences for the existence of a good correlation between  the luminosities of local galaxies and their chemical abundances  which show that metallicities tend to increase with galaxy luminosity. In addition, recent studies have suggested that the LZR evolves with redshift so that, at a given metallicity, galaxies were brighter  in the past (Pettini et al. 2001; Kobulnicky et al. 2003; Kobulnicky \\& Kewley 2004; Lamareille et al. 2004; Maier et al. 2004; Shapley et al. 2004). An evolution in both the zero point and the slope is suggested, indicating a differential evolution in the chemical content of the systems as a function of luminosity.  In particular, in agreement with previous works, Kobulnicky \\& Kewley (2004), found that the largest change from $z \\sim 3$ is driven by  faint galaxies, while from $z \\sim 1$, they detected an increase in the metallicity level of $\\sim 0.14$ dex. Interestingly, Erb et al. (2006a) showed that the large uncertainties present in the determination of the LZR at high redshift, introduced mainly by the variations in the mass-to-light ratios, make it very difficult to draw a robust conclusion on its evolution.  Luminosities are usually used as a surrogate of stellar mass because of the difficulty in measuring the latter. However, it is now accepted that the correlation between metallicity and stellar mass is more fundamental. This  correlation  was first reported by Lequeux et al. (1979) in a study of irregular galaxies. Recently, Tremonti et al. (2004, hereafter T04) confirmed this correlation on a statistical basis by determining a tight correlation between gas-phase metallicity and stellar mass, extended over a factor of ten in metallicity and three decades in stellar mass, for a sample of $\\approx 30000$ local star-forming galaxies in  the Sloan Digital Sky Survey (SDSS). The stellar MZR shows a linear growth between  $10^{8.3} {\\rm M_{\\odot}} h^{-1} $ and $10^{10.35} {\\rm M_{\\odot}} h^{-1}$, flattening for larger stellar mass. T04 interpreted this behaviour as the efficient action of galactic winds over systems with masses lower than $\\sim 10^{10.35} {\\rm M_{\\odot}}h^{-1} $. This stellar mass  has been previously identified   as a characteristic mass for galaxy evolution (Kauffmann et al. 2004). Furthermore,  Gallazzi et al. (2005, hereafter G05) analysed stellar masses, light-weighted ages and stellar metallicities for a sample of galaxies of the SDSS finding similar trends  to those obtained by T04. G05 observed that both age and metallicity tend to increase with stellar mass, showing a rapid growth at intermediate masses and a gradual flattening  above  $M_*\\sim 10^{10.35} {\\rm M_{\\odot}} h^{-1}$. Recently, Savaglio et al. (2005) estimated a relative evolution of 0.10-0.15 dex at $z\\approx 0.7$ with respect to the local MZR of T04 while  Erb et al. (2006a) found a relative evolution of $\\approx 0.30$ dex for  the MZR of galaxies at $z\\approx 2.5$.  The understanding of the origin of the LZR and MZR within the context of the current cosmological paradigm can provide more stringent constrains for galaxy formation models since these correlations are produced by the intervening action of different physical processes. In this scenario, galaxies formed by the hierarchical aggregation of substructures which can affect the internal distribution of angular momentum and mass in galaxies, modifying the star formation rate and the metal production and distribution. SN energy feedback is also expected to play a main role by contributing to establish a self-regulated star formation activity and to trigger powerful outflows. Finally, the interaction between a galaxy and its environment (i.e. groups or clusters) may strip gas, starving galaxies and quenching the star formation activity and the synthetization of new chemical elements.  Numerical simulations which can describe the non-linear evolution of the matter and its chemical enrichment self-consistently have proved to be an adequate tool to tackle galaxy formation (Mosconi et al. 2001; Lia et al. 2001; Kawata \\& Gibson 2003). In particular,  Tissera, De Rossi \\& Scannapieco (2005) used chemical hydrodynamical simulations to show that galaxy-like systems in the concordance $\\Lambda$-CDM model followed the general trend of the observed  MZR, predicting an evolution in zero point and slope. The authors also claimed the need for strong SN galatic winds to decrease the fraction of metals locked into stars, specially in low mass systems. Although SN energy feedback has been regarded as a crucial physical mechanism to regulate the star formation activity in galaxies and to produce metal-loaded, gas outflows capable of enriching the intergalactic medium, its modelling within hydrodynamical simulations is still controversial.  For this reason, in this work we analyse in more detail  the origin of both the LZR and the MZR and study the properties of the simulated galaxies with the aim at improving our understanding of the evolution of both relations in a hierarchical scenario, focusing  on the role of dynamics.  A secondary goal of our work is to clarify where and how SN energy feedback is required. A full study including SN energy feedback will be carried out in the future when the SN model of Scannapieco et al. (2007, in preparation) gets fully tuned to be used in cosmological simulations.  This paper is organized as follows. Section 2 summarizes the main characteristics of the simulations. Section 3 provides a description of the simulated galaxies and an analysis of  their mean astrophysical properties. Section 4 and Section 5  discuss the simulated fundamental metallicity relations and their evolution. Section 6 investigates the origin of the characteristic stellar mass. And Section 7 summarizes our main results. ",
        "conclusions": "In this paper, we studied the dynamical, astrophysical and chemical properties of simulated galaxies in a hierarchical clustering scenario consistent with the concordance $\\Lambda$-CDM model. In this scenario, systems grew by the aggregation of smaller ones, generating correlations between luminosity, stellar mass, circular velocity and metallicity which are in general agreement with observations. Our main goal was to underpin the effect of the hierarchical assembly of galaxy building blocks on the metallicity fundamental relations. The non-linear formation of the structure makes complex the understanding of how star formation, gas cooling and accretion, mergers and environmental effects work together and of the impact they have on the determination of relations between dynamical and metallicity properties.  We found that the cooling and transformation of gas into stars inside substructures that grow hierarchically results in the determination of a characteristic mass which segregates systems dominated, on average,  by old, red and metal-rich stellar populations from those more gas-rich and dominated by young stars. The characteristic mass arises because the formation of large systems  is produced by the fusion of substructures dominated by old stellar populations with a small fraction of new born stars associated to the merger events. This behaviour generates  a change in the slope of the MZR since the outcomes of mergers are larger systems with approximately the same mean abundances. On the other hand, less massive systems tend to form their stars in a more passive way or by rich-gas mergers leading to a more strong correlation between stellar mass and metallicity. This characteristic stellar mass agrees with that obtained by Kauffmann et al. (2004) from the analysis of galaxies in the SDSS-DR2. Neverthless, the simulated evolution  of the MZR is milder than present observational evidence. Strong SN energy feedback could help to reconcile models and observations.  Our results suggest that this $M_{c}$ remains unchanged since $z=3$ playing  the role of segregating galaxies into these two distinctive populations. Systems with stellar masses smaller than $M_{c}$ are the ones experiencing the largest changes in their astrophysical properties since $z\\approx3$. These systems are responsible for the evolution of the luminosity-metallicity relation so that at a given metallicity, simulated systems are found to be around 3 magnitudes brighter at high redshift. Small simulated galaxies constitute a diverse population which has formed approximately 50 per cent of their stars since $z\\approx 1 $. Conversely,  large systems had half their stars formed at $z > 2$, on average.    Our results suggest that in the concordance cosmology, the hierarchical aggregation of the structure, which regulates the gas infall and the star formation activity in the absence of energy feedback, is a key process which can dynamically determine a characteristic stellar mass $M_c$. The simulated   $M_c$  plays a similar role to that of the characteristic mass obtained from observations. Certainly, SN feedback is also expected to play an important role in shaping the MZR as already shown by previous works such as Larson (1974) and Chiosi \\& Carraro (2002). However, it is accepted that the effects of SN feedback depend on the potential well of the systems. In hierarchical scenarios, this is an important aspect since the formation of galaxies may have involved the aggregation of substructures with a wide range of masses, each of them susceptible to be differentially affected  by SN energy feedback (depending on their virial masses). Then, to probe the effects of SN feedback, it is of upmost importance which  histories of formation are  assumed for galaxies. In this work, we provide  galaxy formation histories consistent with the concordance models. Within this context, our  findings suggest that since $z=3$, SN feedback  might have a weak effect on  $M_{c}$, owing to the fact that this characteristic mass corresponds to fast rotating systems even at $z=0$. However, this point will be adequately discuss in a forthcoming paper."
    },
    "0609/astro-ph0609525_arXiv.txt": {
        "abstract": "We present a photometric and spectroscopic study of four G--K dwarfs, namely HD~166, $\\epsilon$~Eri, $\\chi^1$~Ori and $\\kappa^1$~Cet.  In three cases, we find a clear spatial association between photospheric and chromospheric active regions.  For $\\chi^1$~Ori we do not find appreciable variations of photospheric temperature, and chromospheric H$\\alpha$ emission.  We applied a spot/plage model to the observed rotational modulation of temperature and flux to derive spot/plage parameters and to reconstruct a rough ``three-dimensional'' map of the outer atmosphere of $\\kappa^1$~Cet, HD~166 and $\\epsilon$~Eri. ",
        "introduction": "The atmospheres of main-sequence (MS) stars with spectral types later than F5 exhibit the effects of magnetic activity.  At ages of a few hundred Myr, MS stars have magnetic activity levels higher than the Sun but considerably lower than what is typically observed in close binary RS CVn systems or in BY Dra stars.  Systematic photometric monitoring of MS FGK stars has been pursued for more than three decades to study starspots in these stars (see, e.g., \\citealt{Radick83, Radick98}; \\citealt{Lockwood97}; \\citealt{Hen99}).  In particular, \\cite{Radick83} reported that two out of eleven stars monitored in Str\\\"omgren $uvby$ passbands showed light variations anti-correlated with the contemporaneous \\ion{Ca}{2} H\\&K S-index. \\cite{Lockwood97} found small-amplitude variations in $b$ and $y$ filters for about ten of 41 stars in their Lowell Observatory program, monitored for about eleven years, and reported photometric variability correlated with mean chromospheric activity.  In a more recent paper, \\cite{Radick98} analyzed the same dataset and found short-term variability behavior due to rotational modulation in at least 15 of the solar-like stars in their sample.   Recently, several studies have been conducted to analyze the photospheric and chromospheric active regions in young stars (\\citealt{Stra93}; \\citealt{Stou99}), but all the objects studied in these works are ultra-fast rotators (UFR) or pre-main sequence (PMS) stars.  In particular, \\cite{Stra93} found a marginal correlation between the starspot distribution and chromospheric inhomogeneities in LQ Hya, a rapidly rotating ($P_{\\rm rot}\\simeq1\\fd6$) single K2 V star, probably just arriving on the Zero Age Main Sequence (ZAMS).  However, they could not discriminate between plages and local velocity fields as the cause of the observed variations in the Full-Width Half-Maximum (FWHM) of the H$\\alpha$ line.  On the other hand, the two very active rapidly rotating Pleiades stars HII 686 and HII 3163 studied by \\cite{Stou99} did not display the maximum of H$\\alpha$ and \\ion{Ca}{2} Infrared Triplet (IRT) emission at the same phase of the spot transit.  A rotational modulation of the H$\\alpha$ emission with maximum nearly coincident with the minimum of the light curve has been instead detected by \\cite{Frasca97} in the active rapidly rotating binary TZ CrB.  Recently, \\cite{Frasca00} presented the first photospheric/chromospheric study of HD 206860, a solar-type star with an activity level intermediate between the Sun and very active PMS and UFR stars; they used \\ion{Ca}{2} H\\&K and H$\\alpha$ lines as chromospheric indicators and photometric observations as a photospheric indicator.  They found a clear rotational modulation in all the chromospheric and photospheric indicators, proving the presence of an uneven distribution of long-lived active regions and a spatial association between photospheric spots and chromospheric plages, as displayed in the Sun and in the most active RS CVn systems.  In this paper, we study the photospheric and chromospheric surface inhomogeneities in four young MS stars with activity levels intermediate between the Sun and the very active UFR stars.  We use light curves and/or temperature measurements as diagnostics of photospheric inhomogeneities and the H$\\alpha$ and \\ion{He}{1} D$_{3}$ lines as chromospheric diagnostics.  Only for $\\kappa^1$ Cet, for which we obtained contemporaneous photometric and temperature rotational modulations, are we able to determine a unique solution for the spot parameters (area and temperature) following the method described in \\cite{Frasca05}; for the other stars we give only a rough estimate of these spot parameters.  Moreover, we apply a spot/plage model to the photospheric and chromospheric rotational modulations in order to investigate the degree of spot/plage association in these mildly active stars, comparing the results with previous results obtained for very active binaries \\citep{Fra94,Frasca05} and single stars \\citep{Frasca00}. ",
        "conclusions": "We have analyzed the photospheric and chromospheric activity in four young, magnetically active solar-type stars, namely $\\chi^1$ Ori, $\\kappa^1$ Cet, HD~166 and $\\epsilon$ Eri.  The photospheric surface features have been recovered by means of the rotational modulation of luminosity and temperature as derived from the LDR method, while the chromospheric inhomogeneities have been studied from their excess H$\\alpha$ emission.  The H$\\alpha$ profiles can be reasonably well reproduced by means of only one Gaussian component, indicating the presence of plages spatially associated with the photospheric spots.  The temperature and light curves are always anti-correlated with the H$\\alpha$ emission modulation, confirming a close spatial association between spots and plages.  The only exception is $\\chi$1 Ori, for which we have not obtained clear rotational modulation of temperature and H$\\alpha$ curves.  The He line is always present in our spectra, but the $EW_{\\rm He}$ measurements show too much scatter to reveal any modulation with the rotational phase.  Thus, in general, the active regions in mildly active stars seem to have structures similar to solar active regions.  The spectroscopic measurements of HD~166 and $\\epsilon$ Eri span about five stellar rotations, but the spot/plage configuration seems to be largely unchanged, as observed in some other young solar analogues where the light curve remains stable for several rotations (\\citealt{Messina02}).  Moreover, from a simple spot/plage model analysis, we have derived the spot and plage parameters (temperature and area).  In the case of $\\kappa^1$ Cet, for which we had both simultaneous photometric and spectroscopic data, spot temperature and area have been uniquely determined.  The grid of the temperature solutions is flat caused by the small-amplitude of the temperature curve of this MS star, and this leads to great errors in the unique spot solution. For HD~166 and $\\epsilon$ Eri, we have also presented a rough estimate of these two parameters.  %  Finally, the temperature difference $\\Delta T$ between the quiet photosphere and spots, a key parameter tied to the blocking effect on convection produced by the intensification of the magnetic field, seems to be higher compared to that derived by \\cite{Frasca05} for the stars with lower gravity, such as IM Peg ($\\Delta T =$ 448 K) and HK Lac ($\\Delta T =$ 767 K).  On the other hand, the spot filling factors seem to be smaller than those obtained for these active stars.  Moreover, the chromospheric plages are larger than the associated spots, as observed in the Sun.  As a consequence, the increasing dissipation of magnetic energy with height above the photosphere, a characteristic of the Sun, seems to occur for these solar-type stars as well.  In future studies, we want first to extend the number of main sequence targets and then investigate the inhomogeneities at photospheric and chromospheric levels in pre-main sequence, rapidly rotating stars to separate gravitational effects from activity level effects."
    },
    "0609/astro-ph0609713_arXiv.txt": {
        "abstract": "We test the Secondary Infall Model (SIM) by direct comparison with the results of N-body simulations. Eight cluster-size and six galactic-size dark matter haloes have been selected at $z=0$ and re-simulated with high resolution. Based on their density profiles at the initial redshift, we compute their evolution by the SIM, assuming a simple prescription for the angular momentum. A comparison of the density profiles obtained by the SIM and the numerical experiments at $z=5$, $1$ and $0$ shows that, for most of the haloes at most epochs, the SIM reproduces the simulated mater distribution with a typical fractional deviation of less than $40$ per cent over more than six order of magnitudes in the density. It is also found that, within the SIM framework, most of the diversity in the shape of the density profiles at $z=0$ arises from the scatter in the primordial initial conditions rather than the scatter in the angular momentum distribution. A crude optimization shows that a similar degree of agreement is obtained for galactic and cluster haloes, but the former seem to require slightly higher amounts of angular momentum. Our main conclusion is that the SIM provides a viable dynamical model for predicting the structure and evolution of the density profile of dark matter haloes. ",
        "introduction": "\\label{sec:intr}  In the hierarchical clustering paradigm, cosmological structure evolves from a primordial density perturbation field via gravitational instability. One of the basic outcomes of the process is the formation of virialized objects, usually referred to as dark matter (DM) haloes. The problem of the collapse and virialization of DM haloes can be easily formulated, as it depends on a single, relatively simple, physical process: the dissipationless gravitational interaction of a system composed of a large number of point-like particles.  Yet, the long range and non-linear nature of the gravitational N-body problem has made it extremely elusive to rigorous analytical treatment. On a phenomenological level, however, tremendous progress has been achieved by means of numerical experiments. After roughly three decades of N-body simulations \\citep[starting with][among others]{Aarseth+79,vanAlbadavanGorkom77,White78}, this technique has reached the level at which different  algorithms and numerical codes have converged to yield an overall consistent description of the formation, evolution and internal structure of DM haloes \\citep{Knebe+00,Knebe+01_MLAPM}.  A major breakthrough was the suggestion by \\citet{NFW96,NFW97} (hereafter NFW) that the density profile of simulated DM haloes can be fitted by a simple analytical function, \\begin{equation} \\label{eq:nfw} \\rho(r) = \\frac{4\\rhos}{(r / \\rs)( 1 + r / \\rs)^2}, \\end{equation} where the two free parameters \\Rhos and \\Rs represent a characteristic density and radius of the halo. There is general consensus within the N-body community that the so-called NFW profile is able to provide a reasonably good fit to the numerical results in a variety of cold dark matter (CDM) cosmologies, although some doubts have been cast on the exact value of the logarithmic slope at the centre \\citep[the so-called density cusp, see e.g.][]{Moore+98,Moore+99,Ghigna+98,Ghigna+00,FukushigeMakino97,FukushigeMakino01,Fukushige+04,Power+03,Hayashi+04,Navarro+04}, as well as on the degree of `universality' of the fit, i.e. its dependence on the underlying cosmological model, the mass accretion history of the halo, or its environment \\citep[e.g.][]{JingSuto00,Klypin+01,Ricotti03,AvilaReese+05,Maulbetsch+_06}.  Observationally, the DM density profile must be indirectly inferred. In rich clusters of galaxies, estimates based on X-rays \\citep[e.g.][and references therein]{VoigtFabian06} or gravitational lensing \\citep[e.g.][]{Dahle03,Gavazzi+03,BartelmannMeneghetti04,DalalKeeton_03} seem to indicate that the mass distribution is indeed well described by equation~(\\ref{eq:nfw}). On galactic scales, observations are much more difficult to interpret. Dwarf spiral and low surface brightness galaxies are most likely dominated by DM in their centres, but the mass distribution inferred from rotation curves suggests a constant DM density core rather than a steep cusp  \\citep[e.g.][]{FloresPrimack94,Moore94}. Recent analyses show that observational data may actually be consistent with steeper profiles once the effects of inclination, non-circular orbits and triaxiality of the dark matter haloes are accounted for \\citep{Hayashi+04,HayashiNavarro_06}, but the controversy is still unresolved \\citep[e.g.][]{Gentile+04,deBlok05}.  It is therefore not clear whether the existence of a more or less `universal' density profile is supported by observational data, nor whether the possible disagreement is due to our lack of an understanding of the very complicated process of galaxy formation or it calls for an overall revision of the the standard model of cosmology. In any case, the qualitative shape of the density profile (shallower than isothermal in the inner parts and steeper in the outer), as well as (to some extent) its `universality' constitute a very robust prediction of the currently accepted DM structure formation paradigm.  Unfortunately, the NFW  fitting formula provides a good phenomenological description of the density profile of simulated DM haloes, but it does not provide a physical understanding of its origin. Ideally, one would have liked to have a self-consistent analytical model that enabled the calculation, from first principles only, of the density profile expected for a given cosmology, halo mass, formation time, environment, or whatever physically meaningful parameter that is found to play a relevant role.  No such model exists as yet, but we are now close to having one. Much of the analytical work on the formation of DM haloes is based on the seminal paper of \\citet{GunnGott72} on the dissipationless collapse of a spherical homogeneous perturbation in an otherwise expanding Friedmann universe. This was followed by \\citet{Gunn77}, who considered the collapse of a spherical inhomogeneous perturbation and the role of shell crossing. It also studied the case of the secondary infall, namely the late infall of shells onto an already collapsed and virialized perturbation, and showed that a self-similar solution can be found upon the use of adiabatic invariance.  \\citet{Gunn77} was followed by two different lines of research. On one hand, \\citet{FillmoreGoldreich84} and \\citet{Bertschinger85} independently found self-similar solutions of the collapse of scale-free spherical perturbations in an Einstein-de Sitter universe. On the other hand, \\citet{HoffmanShaham85} analyzed the dependence of the structure of proto-haloes on the primordial power spectrum and the way it affects the density profile of virialized DM haloes.  The basic predictions of the secondary infall model (SIM) were qualitatively confirmed by N-body simulations \\citep[e.g.][]{Quinn+86,Efstathiou+88,Crone+94}, which prompted further study and extension of the SIM, mostly focusing on the nature of the initial conditions, on the improvement of the dynamical model and on its cosmological implications \\citep[e.g.][]{RydenGunn87,Hoffman88,Ryden88_j,Ryden88_p,ZaroubiHoffman93,LokasHoffman00}. A common assumption made in these studies was that the halo particles are moving along radial orbits, and the generic result that emerges is that the inner density profile is roughly given by $\\rho \\propto r^{-2}$. Thus, a cuspy density profile is a generic outcome of the SIM. Actually, the predicted logarithmic slope is even steeper than the one measured in N-body experiments.  The next major improvement of the SIM was the introduction of non-radial motions, namely a distribution of angular momentum of individual particles, not necessarily implying, on average, a total angular momentum of the halo. A number of authors \\citep[e.g.][]{WhiteZaritsky92,Ryden93,AvilaReese+98} have pointed out that non-radial motions flatten the inner density profile. \\citet{Nusser01} extended the self-similar solutions to include a distribution of angular momentum, and several recent studies show that, by introducing enough angular momentum, the density profile predicted by the SIM can have a $\\rho \\propto r^{-1}$ density cusp \\citep{Hiotelis02,leDelliouHenriksen03,Ascasibar+04,Williams+04,Lu+06}.  Within the SIM framework, there are two key ingredients that determine the shape of the virialized density profile: the initial density profile of the proto-halo and the angular momentum distribution. The primordial density has been determined either by employing the statistical properties of Gaussian random fields \\citep{HoffmanShaham85,BBKS86,Hoffman88,Ascasibar+04} or by means of the mass accretion history of DM haloes \\citep{AvilaReese+98,NusserSheth99,Lu+06}. The angular momentum distribution has always been considered to be a free parameter, determined so as to yield best agreement with simulations. \\citet{Ascasibar+04} pushed the study of the SIM one step further by making a detailed comparison between the density profile predicted by the model with the actual numerically simulated DM haloes rather than the NFW fit. These authors assumed the initial density profile to be given by the average expectation around a local maximum of the primordial perturbation field, determined by the peak height and its smoothing scale. Optimizing over these two parameters and the angular momentum distribution, the resulting density profiles were in close agreement with the simulated ones, with an accuracy comparable to that of the best-fitting NFW profile.  The aim of the present paper is to go one step beyond \\citet{Ascasibar+04} in order to assess the validity of the SIM as a tool to understand the physical origin of the DM density profile. More precisely, we attempt to recover the density profile of actual simulated haloes, but instead of assuming a functional form for the initial conditions, we identify the primordial peak in the initial snapshot of the simulation and use the density profile around that point as an input for the SIM. We calculate the entire dynamical evolution of the halo upon assuming only one free parameter (the angular momentum distribution of the DM particles), and compare the density profile thus obtained with the full N-body simulation. We consider this as the ultimate test of the validity of the SIM, at least in the framework of the CDM-like cosmology.  The paper is structured as follows: The basic principles of the SIM are reviewed in \\S~\\ref{sec:sim}. Numerical simulations are described in \\S~\\ref{sec:nbody}, and the comparison of the SIM-calculated density profiles with the results of the full N-body simulation is presented in \\S~\\ref{sec:results}. A discussion and summary of the results is given in \\S~\\ref{sec:discu}. Technical details of the SIM and its implementation are given in Appendix~\\ref{sec:appSIM}. ",
        "conclusions": "\\label{sec:discu}   The main result of the present work is that the secondary infall model provides a valid theoretical framework for calculating the structure and evolution of DM haloes in an expanding universe, and that its predictions with respect to the density profiles are in close agreement with full N-body simulations. Comparing the SIM and simulated cumulative density profiles over more than six orders of magnitude, we find the typical discrepancy to be better than $40$ per cent. This level of agreement extends to the time evolution of the density profile and it is not limited to the final time snapshot only. Within the SIM framework, most of the diversity in the density profiles of DM haloes is contributed by the scatter in the primordial profiles rather than the scatter in the angular momentum.  Galactic-size haloes are less isolated objects than cluster haloes, in the sense that they have more similar-mass companions. The simulated galaxies have been selected from a typical filamentary region and naive reasoning would lead one to expect these objects to be less suitable for the application of the SIM compared with clusters. This is not the case. The agreement between the N-body and SIM density profile of galactic haloes is roughly as good as for clusters. The only difference is that the crude optimization over the angular momentum parameter, $\\eta$, finds that galaxies require slightly higher values of $\\eta$, consistent with the fact that the galactic haloes are more torqued and hence should have acquired more angular momentum. This in turn gives further support to the SIM in the sense that the value of $\\eta$ is not a mere free parameter but is related to the actual angular momentum of haloes.  The galactic-scale halo G4 deserves a special attention. The SIM-predicted profile with $\\eta=0.15$ provides a very bad fit to the simulated density profile at $z=0$. Yet, for the very high value of $\\eta=0.60$, a very good fit is obtained for the present epoch. Such a high value constitutes a remarkable exception, and no other object has a best-fitting $\\eta$ anywhere close to $0.60$. Visual inspection of the G4 halo reveals a major merger in progress at $z=0$, and the halo is clearly not in virial equilibrium. In spite of its dynamical state, the density profile of the halo is extremely well fitted by the SIM. Is such an agreement a mere coincidence, or does the SIM actually reproduce the internal structure the system? This intriguing question has led us to further investigate the structure of the G4 halo, and we found that it actually displays an extremely high amount of angular momentum due to the merger event.  Although the present work focuses on the density profile, it would be desirable to extend the formalism outlined in Appendix~\\ref{sec:appSIM} so as to compute their full dynamical structure of DM haloes; more precisely, the radial and tangential components of the velocity dispersion. Derived quantities, such as the angular momentum distribution or the anisotropy profile, can be readily calculated, and their dependence on the assumed value of $\\eta$ would be extremely helpful in validating our simple prescription for angular momentum quantitatively.  Moreover, one could use the upgraded formalism to address additional questions raised by numerical N-body simulations, such as the relation between the logarithmic slope of the density profile and the anisotropy of the velocity dispersion \\citep{HansenMoore06}, or the power-law radial dependence of the coarse-grained phase-space density \\citep{TaylorNavarro01,Ascasibar+04,Rasia+04}. This quantity is somewhat analogous to the `entropy' of the DM, formally defined so as to mimic the ideal gas thermal entropy, and it has indeed been found that they are closely related in adiabatic gasdynamical simulations \\citep{Faltenbacher_06}.  Eventually, the final goal would be to achieve a totally self-consistent prescription to evaluate the angular momentum distribution of DM haloes, which of course can only be achieved through an understanding of its physical origin. Although strict spherical symmetry would imply that the angular momentum of individual particles must be exactly conserved (and hence it should be imprinted by the initial conditions), several processes may be responsible for the conversion of radial into tangential motions in a more realistic case.  For instance, analytical and numerical studies have shown that a predominantly radial collapse is unstable to the growth of tangential motions, resulting in the so-called radial orbit instability \\citep[see e.g.][and references therein]{Huss+99}, and it has been conjectured \\citep{Barnes+06} that this mechanism could actually set the length scale \\Rs of DM haloes, which marks the transition region from almost isotropic orbits to radial ones.  A seemingly very different point of view is provided by cosmological simulations, which are assumed to incorporate all the physics relevant for the formation of DM haloes. Haloes in numerical experiments evolve through phases of violent `major merger' events in which objects of comparable masses merge to form a yet bigger halo. A detailed analysis shows that the NFW-like structure, and in particular the angular momentum distribution within $\\rs$, is exclusively determined by major mergers, and is hardly affected by the quiescent phases that follow \\citep[][and references therein]{SalvadorSole+05,RomanoDiaz+06}.  To summarize,  two very different frames of reference for DM halo formation end up predicting very similar density profiles. Assuming that this is not a coincidence, one should look for an underlying physical process that unites these two different approaches. Our results suggest that angular momentum might play an important role in that respect.  Yet, the present study points to another key element and this is the initial configuration of the haloes. Figure~\\ref{fig:div} shows the diversity in the density profile of the virialized objects caused by the variation in their initial density structure. It follows that the shape of the halo density profile cannot be discussed without an explicit reference to the nature of the initial conditions; a link that any complete theory of the collapse of DM haloes should incorporate. Although a full understanding is yet to be achieved, the study of the collapse and virialization of cosmic structures that commenced with \\citet{GunnGott72} has certainly provided us with many clues and a deep insight of the most relevant ingredients behind the process.  As a final note, we would like to suggest a potential practical application of the SIM as a complement (perhaps even substitute) of the multiple-mass technique. High-resolution re-simulations of individual DM haloes identified in low-resolution cosmological experiments require a significant amount of CPU resources, and they can be practically performed only for a small subset of all objects. We have shown, however, that the SIM provides a valid framework for calculating the dynamical evolution of DM haloes, and it takes about $10$~s to compute the density profile of a given object at a given epoch on a desktop PC. This gives rise to the interesting possibility of using the SIM to insert highly-resolved spherical haloes into large-scale, low-resolution cosmological simulations. In practice, one can use the low-resolution experiment to set up the initial conditions for all the objects and use the SIM to calculate their dynamical evolution, providing an efficient and relatively accurate way of generating extended catalogs of dynamically resolved clusters at a reasonable accuracy."
    },
    "0609/astro-ph0609239_arXiv.txt": {
        "abstract": "It has recently been suggested that the winds from Classical T Tauri stars in general, and the wind from TW Hya in particular, reaches temperatures of at least 300,000 K while maintaing a mass loss rate of $\\sim 10^{-11}$ \\Msol yr$^{-1}$ or larger.  If confirmed, this would place strong new requirements on wind launching and heating models.  We therefore re-examine spectra from the Space Telescope Imaging Spectrograph aboard the Hubble Space Telescope and spectra from the Far Ultraviolet Spectroscopic Explorer satellite in an effort to better constrain the maximum temperature in the wind of TW Hya.  We find clear evidence for a wind in the \\ion{C}{2} doublet at 1037 \\AA\\ and in the \\ion{C}{2} multiplet at 1335 \\AA.  We find no wind absorption  in the \\ion{C}{4} 1550 \\AA\\ doublet observed at the same time as the \\ion{C}{2} 1335 \\AA\\ line or in observations of \\ion{O}{6} observed simultaneously with the \\ion{C}{2} 1037 \\AA\\ line.  The presence or absence of \\ion{C}{3} wind absorption is ambiguous.  The clear lack of a wind in the \\ion{C}{4} line argues that the wind from TW Hya does not reach the 100,000 K characteristic formation temperature of this line.  We therefore argue that the available evidence suggests that the wind from TW Hya, and probably all classical T Tauri stars, reaches a maximum temperature in the range of 10,000 -- 30,000 K. ",
        "introduction": "Mass loss in the form of a wide-angle wind and/or collimated jet is a common feature observed in young, low mass stars that are still surrounded by accretion disks (for a review see Bally, Reipurth, and Davis 2006).  The material in these disks eventually accretes onto the central star, gets ejected in an outflow, or is incorporated into planets or other solar sytem-like bodies.  Understanding the processes through which young stars interact with and eventually disperse their disks is critical for understanding the rotational evolution of stars and the formation of planets.  Mass loss appears to be a necessary component in the accretion process both observationally as mentioned above and theoretically: angular momentum must be carried away in order to allow disk and stellar accretion to occur, and magnetized winds are efficient at removing angular momentum (e.g. Pudritz, Pelletier, and Gomez de Castro 1991, Shu et al. 1994).  Signatures of mass loss are observed primarily in optical, infrared, and radio wavelength spectral diagnostics.  Bipolar outflows are observed in millimeter wavelength CO lines  and other molecular tracers in the mm and sub-mm range (e.g. see review by Arce et al. 2006).   In the infrared, emission in H$_2$ (e.g. Gueth \\& Guilloteau 1999; Stanke, McCaughrean, \\& Zinnecker 2002) and [\\ion{Fe}{2}] (e.g.  Reipurth et al. 2000) lines are regularly seen in outflows from young stars.  At optical wavelengths, Herbig-Haro (HH) objects and their associated stellar jets are routinely imaged in narrow band filters centered on H$\\alpha$ or on a host of forbidden lines.  High resolution images of the central engine show that these jets originate from young stars surrounded by circumstellar disks, with the jet usually emerging perpendicular to the disk such as clearly seen in the case of HH 30 (Burrows et al. 1996, Ray et al. 1996).  Spectroscopically, winds and jets are diagnosed by the appearance of P-Cygni like line profiles in permitted lines such as H$\\alpha$ (and other Balmer lines) and Na D (e.g. Mundt 1984; Reipurth. Pedrosa, \\& Lago 1996; Alencar \\& Basri 2000), and in the profiles of forbidden lines which often show a strong asymmetry with substantial blue-shifted emission with little or no red-shifted emission (e.g. Edwards et al. 1987, Hamann 1994).  More recently, jets and winds from CTTSs have been detected at shorter wavelengths corresponding to higher temperature emissions.  Hartigan et al. (1999) detect several ultraviolet (UV) \\ion{Fe}{2} emission lines in the HH 47A bow shock, and several irradiated jets with UV signatures have been discovered recently (e.g. Reipurth et al. 1998, Bally \\& Reipurth 2001).  Gomez de Castro and Verdugo (2001) identify feature in the UV semiforbidden emission lines of \\ion{C}{3}] $\\lambda1908$ and \\ion{Si}{3}] $\\lambda1892$ which they associate with shocks at the base of the jet in a few CTTSs.  At X-ray energies, a number of jets have also been recently detected (e.g. Pravdo et al. 2001; Favata et al. 2002; Bally, Feigelson, \\& Reipurth 2003).  This X-ray emission forms in the fastest shocks associated with these stellar jets (see also Bally et al. 2006).  While these high excitation emissions appear associated with shocks within the jets from young stars, until recently it was generally thought that the stellar or disk wind at the base of these jets was only heated to a level which could produce features such as the the P-Cygni like profiles visible in the Balmer lines.  A hot wind component, possibly with an origin on the star as opposed to in the disk, has recently been proposed by a few authors.  Beristain, Edwards, and Kwan (2001) argue for a hot wind component in CTTSs based on the analysis of high resolution \\ion{He}{1} and \\ion{He}{2} line profiles observed in the optical.  Takami et al. (2002) and Edwards et al. (2003) find obvious blue-shifted absorption below the local continuum in the \\ion{He}{1} $\\lambda$10830 line in 6 young, low mass stars, again providing evidence for a hot wind component forming relatively near the star. Edwards et al. (2003) point out that the strength and ubiquity of this absorption indicates that the wind emanates over a large solid angle from the star: the \\ion{He}{1} $\\lambda$10830 line does not trace a highly collimated flow.  These results are reinforced by Edwards et al. (2006), who found subcontinuum \\ion{He}{1} $\\lambda$10830 wind absorption in 26 of 38 CTTSs. The $\\lambda5876$ and $\\lambda10830$ lines of \\ion{He}{1} require a strong ionizing flux or a high temperature for excitation, as their upper states lie $\\sim20$ eV above the ground level. If these lines are primarily collisionally excited, temperatures of $\\sim 25,000$ K are required to excite them (Athay 1965; Avrett, Vernazza, \\& Linsky 1976).  With photoionization followed by recombination and cascade, helium excitation can take place at local kinetic temperatures between 8000 and 15,000 K (Zirin 1975; Heasley, Mihalas, \\& Poland 1974; Wahlstrom \\& Carlsson 1994).   CTTSs in general (e.g. Feigelson et al. 2002), and TW Hya in particular (Kastner et al. 2002), are strong X-ray sources, so the detection of these \\ion{He}{1} lines in CTTS winds still leaves a wide range of possible temperatures present in these outflows.   Dupree et al. (2005) detect \\ion{He}{1}  $\\lambda10830$ absorption in the wind from TW Hya and T Tau.  They combine these observations with an analysis of the line profiles of \\ion{C}{3} $\\lambda977$ and \\ion{O}{6} $\\lambda1032$ observed with the Far Ultraviolet Spectroscopic Explorer ({\\it FUSE}) satellite to suggest that CTTSs possess continuous, smoothly accelerating winds which reach velocities of $\\sim 400$ \\kms\\ and temperatures of $\\sim 300,000$ K. These results are based primarily on the {\\it FUSE} data for TW Hya, since the data for T Tau is much lower in quality.  Dupree et al. (2005) find a minimum mass loss rate for the \\ion{O}{6} line of $2.3 \\times 10^{-11}$ \\Msol yr$^{-1}$.  Such a hot temperature for the wind from TW Hya is surprising, particularly given the high mass loss rate.  Alencar \\& Batalha (2002) and Herczeg et al. (2004) estimate a mass accretion rate onto TW Hya of $\\sim2 \\times 10^{-9}$ \\Msol\\ yr$^{-1}$.  Most estimates suggest that the mass loss rate from CTTSs is 0.1 -- 0.3 times that of the mass accretion rate (K\\\"onigl \\& Pudritz 2000, Shu et al.  1994), suggesting the mass loss rate from TW Hya may be as high as $5\\times 10^{-10}$ \\Msol\\ yr$^{-1}$. Such a high temperature, high mass loss rate wind puts strong constraints on the ultimate origin and heating of winds from CTTSs.  Current theories of mass loss from CTTSs primarily produce cold, magnetocentrifugally driven flows from the disks around these stars (e.g. K\\\"onigl \\& Pudritz 2000, Shu et al. 2000) which are then heated by a combination of ambipolar diffusion and X-rays to temperatures of $\\sim 10^4$ K (Shang et al. 2002; Shang, Li, \\& Hirano 2006).  Photoevaporation of the surface layer of the disk may also produce a low velocity ($\\sim 10$ \\kms) wind with a characteristic temperature of $10^4$ K as well (e.g. Matsuyama, Johnstone, \\& Hartmann 2003; Font et al. 2004).  Thus, if the Dupree et al. (2005) suggestion holds true, a totally new heating, and possibly driving, mechanism must be identified.  The analysis presented by Dupree et al. (2005) relies on the asymmetric {\\it shape} of the line profiles tracing the hot gas: no true absorption against a local continuum is detected.  Therefore, we re-examine the published (Herczeg et al. 2002, 2004; Dupree et al. 2005) high resolution spectra from the Space Telescope Imaging Spectrograph ({\\it STIS}) aboard the Hubble Space Telescope ({\\it HST}) and from the {\\it FUSE} satellite to look for firm wind signatures and confirm the high proposed temperature in the wind of TW Hya.  The {\\it STIS} and {\\it FUSE} bandpasses cover several lines with characteristic formation temperatures ranging from $\\sim 10^4$ K to $\\sim 300,000$ K.  We find firm evidence against the proposed wind in the \\ion{C}{4} and \\ion{O}{6} lines of TW Hya.  We therefore conclude that the wind in fact does not reach (at least at large optical depth as claimed) a temperature of $\\sim 100,000$ K, characteristic of the formation of the \\ion{C}{4} line.  In \\S 2 we describe the observational data.  Section 3 presents our analysis, and \\S 4 gives a discussion of our results. ",
        "conclusions": "The hot FUV emission lines from CTTSs are related to accretion, as is seen in correlations with mass accretion rate (Johns--Krull et al. 2000; Calvet et al. 2004).  The line profiles are often complex, and can show significant excess redshifted emission.  Such an asymmetric line profile is naturally expected for optically thin lines that are produced by hot gas in downflowing accretion columns.   Dupree et al. (2005) also suggest that narrow \\ion{O}{6} profiles in the lower-quality spectrum of T Tau could suggest wind absorption, but the narrow width is instead caused by strong absorption in interstellar H$_2$ lines (e.g., Roberge et al. 2001, Walter et al. 2003).  However, asymmetric or narrow profiles of warm emission lines, including \\ion{C}{2}-\\ion{C}{4}, \\ion{O}{6}, \\ion{He}{2}, \\ion{N}{5}, and \\ion{Si}{4} are insufficient to infer the presence of a wind from CTTSs. Detecting wind absorption from FUV spectra of CTTSs requires either measuring subcontinuum absorption or absorption of nearby atomic or H$_2$ lines.  We conclusively find that no \\ion{C}{4} wind absorption is detected in the STIS spectrum of TW Hya, based on several independent diagnostics. Since \\ion{C}{4} and \\ion{O}{6} trace gas at $T=100,000$ and $300,000$ K, respectively, the lack of any detectable \\ion{C}{4} in the wind means that no higher temperature \\ion{O}{6} gas will be detectable.  We confirm this conclusion by finding an absence of \\ion{O}{6} wind absorption in the {\\it FUSE} spectrum of TW Hya.  These findings emphasize the result that asymmetric line profiles with more red emission than blue as shown here do not necessarily diagnose the presence of a wind.  Indeed, Lamzin et al. (2004) argue that the extended blue wings of \\ion{C}{3} $\\lambda977$ \\AA, \\ion{C}{4} $\\lambda1550$ \\AA, and \\ion{O}{6} $\\lambda1032$ \\AA\\ arise in the magnetospheric infall around TW Hya.    The lack of detectable absorption by \\ion{C}{4} and the positive detection of \\ion{C}{2} wind absorption indicates that the abundance of \\ion{C}{4} in the wind is less than 1\\% of that of \\ion{C}{2}.  Based on the ionization equilibrium of C, this upper limit suggests a maximum wind temperature of $<50,000$ K.  The detection of optically thin wind absorption  in the \\ion{He}{1} $\\lambda$10830 \\AA\\ requires either some photoionization or temperatures of $>25,000$ K in the wind.  In either case, some \\ion{C}{3} wind absorption is expected.  Unfortunately the {\\it FUSE} spectrum of TW Hya at 977 \\AA\\ is not high enough quality to clearly detect the presence or absence of \\ion{C}{3} in the wind.    In addition to \\ion{C}{2} and \\ion{H}{1} discussed earlier, wind absorption from TW Hya is detected in \\ion{Al}{2}, \\ion{Si}{2}, \\ion{Mg}{2}, \\ion{O}{1}, and \\ion{N}{1}, but is not detected in \\ion{C}{1} or \\ion{Si}{1}.  Wind absorption is also prevalent from the CTTS RU Lupi (Herczeg et al. 2005), which has a much higher mass accretion rate than TW Hya, and therefore presumably a higher wind mass loss rate.  In addition to the same wind absorption lines as seen in TW Hya, the higher optical depth in the wind from RU Lupi allows us to detect \\ion{Cl}{1}, many \\ion{Fe}{2} lines, and \\ion{Si}{3}.  The similarity in wind absorption lines from RU Lupi and TW Hya suggests that the ionization of the wind can be similar for stars with mass loss rates that likely differ by at least a factor of 10.   The absence of \\ion{C}{1} absorption in the wind of both TW Hya and RU Lupi star suggests that the temperature of the wind is $>20,000$ K.  However, the presence of \\ion{O}{1}, \\ion{Cl}{1}, and \\ion{Fe}{2} in the wind requires gas at $T<20,000$ K. We therefore speculate that the winds from CTTSs have a kinetic temperature of $\\sim 10,000$ K, and that the ionization state is dominated by photoionization.  Elements with first ionization potential longward of the Lyman limit are photoionized.  A smaller amount of ionising photons may also irradiate the wind, which would explain prevalent wind absorption in the \\ion{He}{1} recombination line at 10830 \\AA\\ (Edwards et al. 2006), \\ion{N}{1} lines with lower levels excited to 2.4 eV, and the detection of \\ion{Si}{3} in the wind of RU Lupi (Herczeg et al. 2005).  No strong lines occur between 912--954 \\AA, which may explain why \\ion{Cl}{1} is not ionized in the wind of RU Lupi."
    },
    "0609/astro-ph0609180_arXiv.txt": {
        "abstract": "The cooling and reheating histories of dim isolated neutron stars(DINs) are discussed. Energy dissipation due to dipole spindown with ordinary and magnetar fields, and due to torques from a fallback disk are considered as alternative sources of reheating which would set the temperature of the neutron star after the initial cooling era. Cooling or thermal ages are related to the numbers and formation rates of the DINs and therefore to their relations with other isolated neutron star populations. Interaction with a fallback disk, higher multipole fields and activity of the neutron star are briefly discussed. ",
        "introduction": "\\label{intro} All \"magnificent\" seven ROSAT detected sources display X-ray fluxes consistent with thermal luminosities of the order of 10$^{31}$ - 10$^{32}$ ergs s$^{-1}$. In particular, RX J1856.5-3754, which is the only source in this class with a parallax distance determination and kinematic age, has luminosity  3 $\\times$10$^{31}$ ergs s$^{-1}$ at distance 120 pc. For RXJ1856.5-3754 and RX J 0720.4-3125 spectra can be fit with two blackbodies to cover X-ray and optical data. This is probably an indicative representation of the actual temperature modulation due to the magnetic fields. Moreover, the soft to hard luminosity ratio is the same, 0.5, in both sources, possibly indicating similar physics on the surface. The soft blackbody seems to cover the entire neutron star surface. This suggests that the entire surface luminosity is of order 10$^{31}$ - 10$^{32}$ ergs s$^{-1}$ also in the latter two sources. Are luminosities of 10$^{31}$ - 10$^{32}$ ergs s$^{-1}$ standard for these sources?\\\\ \\\\ Cooling provides the source of thermal X-ray luminosity for a young neutron star. Cooling proceeds first via neutrino emission from the interior of the neutron star, to give way to cooling by surface photon emission at an age of the order of  10$^{6}$ yrs. According to families of standard cooling curves (eg Tsuruta et al 2002), luminosities in the range of 10$^{31}$ - 10$^{32}$ ergs s$^{-1}$ correspond to the end of the neutrino cooling era and the transition to photon cooling.\\\\ \\\\ Why do we not  detect a few sources at age 10$^{5}$ yr and luminosity 10$^{33}$ ergs s$^{-1}$ ? The space density of 10$^{5}$ yr old sources is only 10 times less than the density of the 10$^{6}$ yr old sources.  With luminosities one order of magnitude larger according to the standard cooling curves, the detectable volume may not be significantly larger than it is for the 10$^{6}$ yr old sources, also considering that the detected neighbourhood being limited by absorption.\\\\ \\\\ We might expect an even larger number of older DINs, say at age 3 $\\times$ 10$^{6}$ yrs. While the cooling luminosity drops sharply at ages of a few 10$^{6}$ yrs such nearby sources would be above XMM and Chandra detection limits. The lower thermal luminosities and effective temperatures of older sources means the spectra are shifted to lower photon energies and therefore detection may be more difficult because of absorption. With only 7 sources,  distance and luminosity uncertainties and lack of reliable age estimates, it is at present too early to test whether the population of DINs is consistent with being powered by cooling luminosities alone. In particular it will be interesting to see if more statistics and age determinations yield sources at ages longer than 10$^{6}$ yrs without the drop-off in luminosity expected in the photon-cooling era . With typical distances of the order of 100 pc and if their galactic birthrate is of the order of the total neutron star birthrate, 10$^{-2}$ yr$^{-1}$, the DINs may well be a population dominated by sources older than 10$^{6}$ yrs, whose luminosity is supplied by something else after the drop-off in the cooling luminosity. ",
        "conclusions": "\\label{dis}  At typical separations of say 150 pc, the DINs must make up a galactic plane population of about 10 000. If the duration of the neutrino cooling epoch with L larger than 10$^{32}$ ergs s$^{-1}$ is taken to be their representative age, this age is 10$^6$ yrs. The rate of formation is therefore of the order of 10$^{-2}$ yr$^{-1}$ .  The DINs make up a substantial fraction of the supernova rate: they are a very abundant population of young neutron stars, comparable to isolated radio pulsars (and possibly RRATS).\\\\ \\\\ Dipole magnetic fields inferred from $\\dot{P}$ and P, as well as surface fields inferred from absorption features from proton cyclotron lines are about 6 $\\times$10$^{13}$ G. This is of the order of or above the quantum critical field, and only an order of magnitude less than the inferred dipole magnetic fields of AXPs and SGRs.  Periods are clustered in the same special narrow range as AXPs and SGRs, for 5 DINs observed periods range from 3.45 s to 11.37 s (Haberl 2005). This period clustering can be explained with fallback disks acting as a gyrostat around the neutron star, and determining its asymptotic equilibrium period (Alpar 2001). In this scheme DINs, AXPs and SGRs are supposed to be in a propeller interaction with the fallback disk.   Most numerous, after the no-disk radio pulsars, must be the very small mass disk cases, where the disk would have a long lifetime. These are proposed to be the DINs.\\\\ \\\\ All these classes of objects are in the upper right hand part of the P-$\\dot{P}$ diagram. All are close to but above the death line for radio pulsars. They share the P-$\\dot{P}$ neighbourhood with the high dipole field radio pulsars. So why do the DINs (AXPs and SGRs also) not function as radio pulsars? This is more of a question for the DINs as their inferred dipole surface fields are of the same order as the highest dipole field radio pulsars: the AXPs and SGRs might short out the pair creation in the polar cap more easily by virtue of fields that are one order of magnitude larger than those of the radio pulsars. Now there is a new enigma: RRATS (McLaughlin et al 2006), some of which are close to same corner of P-Pdot diagram. RRATS also may be similar to AXPs and SGRs in that they burst, though in the radio. What distinguishes between the different classes might indeed be the presence or absence, and nature of a fallback disk around the neutron star.  The $\\dot{M}$ history of the disk determines the equilibrium period to which the star approaches asymptotically.  For B = 10$^{12}$ - 10$^{13}$ G ,  a relatively wide range of $\\dot{M}$ gives a narrow range of equilibrium periods:   P$_{eq}$ scales with $ |\\dot{M}|^{- 4/7}$. An interesting example is that  if the dipole magnetic field on the surface is the quantum critical field, and $\\dot{M}$ is at the Eddington value, this leads to equilibrium periods of about 12 s. \\\\ \\\\ This approach is not in conflict with magnetar models (Thompson \\& Duncan 1995, Woods \\& Thompson 2006). The magnetar model functions on the basis of very strong magnetic fields anchored in the neutron star crust. These fields are likely to develop and release energy in local processes in the neutron star crust. The surface magnetar fields are likely to be concentrated in the higher multipoles, and the dipole magnetic field on the surface of magnetars may be of order or less than the quantum critical field? One might speculate that when the quantum critical field is reached on the surface, leading to pair creation, near surface currents and surface heating, magnetic structure in the higher multipoles will be amplified, at the expense of the dipole (global) surface field, which might be limited to a value near the quantum critical field. Arguments for surface multipole fields have been brought up in several connections (Zane \\& Turolla 2005; see also the contributions by Zane, by Tiengo and by Gaensler at this conference). \\\\ \\\\ In the context of fallback disk models and hybrid models with magnetar fields in the higher multipoles, the presence of fallback disk torques interacting with the dipole component of the magnetic field would lead to detectable thermal luminosities from the DINs if the disk lifetime extends beyond the 10$^6$ yr timescale of the cooling luminosity. As explored here, this is because the energy dissipation rate for propeller disk torques near equilibrium is stronger than energy dissipation rates supplied by ordinary or magnetar dipole spindown, and may be roughly constant throughout the disk lifetime, unlike the dipole spindown case for which the energy dissipation rate decays rapidly with a power law time dependence.\\\\"
    },
    "0609/astro-ph0609655_arXiv.txt": {
        "abstract": "The majority of the highest energy cosmic rays are thought to be electrically charged: protons or nuclei.  Charged particles experience angular deflections as they pass through galactic and extra-galactic magnetic fields.  As a consequence correlation of cosmic ray arrival directions with potential sources has proved to be difficult.  This situation is not helped by current data samples where the number of cosmic rays/source are typically $\\leq O(1)$.  Progress will be made when there are significantly larger data samples and perhaps with better catalogs of candidate sources.   This paper reports a search for correlations between the RXTE catalog of nearby active galactic nuclei, AGNs, and the published list of ultra-high energy cosmic rays from the AGASA experiment.  Although no statistically significant correlations were found, two correlations were observed between AGASA events and the most inclusive category of RXTE AGNs. ",
        "introduction": "\\label{introduction}  Perhaps the primary goal of all experiments studying the highest energy cosmic rays is to find the source of these particles.  While circumstantial evidence may favor one type of source over another, demonstration of a clear correlation between the direction of cosmic rays and their sources is arguably essential.  Unfortunately for electrically charged cosmic rays, galactic magnetic fields, and for the highest energy cosmic rays extra-galactic magnetic fields, cause angular deflections that can blur the correlation between cosmic ray arrival direction and source direction.  If the sources as viewed from the earth are extended\\,\\cite{waxman,cuoco} the problem is even more difficult.  Unless otherwise noted, for this paper we assume compact (point-like) sources for the highest energy cosmic rays.  If the angular blurring from magnetic fields is small\\,\\cite{dolag} ({\\it i.e.} not significantly greater than the experimental angular resolution) and/or for neutral primaries, then experiments should observe cosmic rays that cluster in arrival direction\\,\\cite{agasa_cluster,tinyakov}, and/or that correlate with potential astronomical ({\\it e.g.} BL Lac) sources\\,\\cite{tinyakov_BLLac,gorbunov_BLLac0,gorbunov_BLLac1,gorbunov_BLLac2,hires_BLLac}.  For nearby sources, where experiments should detect multiple cosmic rays/source, event clusters provide bounds on the cosmic ray source density\\,\\cite{dubovsky,blasi,kachelriess} potentially favoring one type of source for the highest energy cosmic rays over another. However at this time the situation is less than clear as some results\\,\\cite{finley,hires_cluster} question the significance of the reported clusters and/or some of the BL Lac correlations\\,\\cite{hires_BLLac,not_BLLacs}.  If deflections of charged cosmic rays by extra-galactic magnetic fields are not small\\,\\cite{sigl}, then lower energy, $E$, cosmic rays should experience the greatest angular deflections. Unfortunately small experiment data samples and a cosmic ray flux $\\propto E^{-3}$ have often caused studies to retain cosmic rays to energies, $E_{thresh}$, well below GZK\\,\\cite{gzk} energies\\,\\cite{agasa_cluster}. Furthermore deflections of the highest energy cosmic rays even by our galactic magnetic field can be substantial\\,\\cite{tinyakov_Bfield,kachelriess_Bfield,tanco}. As magnetic deflections scale proportional to the charge of the primary cosmic ray, nuclei in the cosmic rays may have significant deflections.  Although most searches have looked for clustering and/or source correlations on small angular scales, studies at larger angular scales have also found evidence for clustering and/or source correlations\\,\\cite{kachelriess_clusters,smialkowski,singh}. Certainly the angular scale of cosmic ray clusters and the magnitude, and thus relevance, of the deflections of ultra-high energy cosmic rays by magnetic fields is not universally agreed to at this time.  In the future, significantly larger data samples will allow analyses to increase $E_{thresh}$ while retaining the number of observed cosmic rays/source (for nearby sources) $\\geq O(1)$.   However another possibility is to exploit catalogs of candidate sources.  With a catalog of source directions, cosmic rays can be effectively correlated with sources even when if magnetic field deflections are ``not small'' and/or when the the number of observed cosmic rays per source is $<1$ allowing searches with existing data samples. That said, catalog based studies are limited by the completeness of the source catalog and the relevance (or not) of that class of astronomical source to the production of the highest energy cosmic rays.  Often conjectured astrophysical sources include gamma ray bursts, GRBs, and/or active galactic nuclei, AGNs\\,\\cite{BGG2002}.  This paper reports a search for correlations between a catalog of nearby AGNs\\,\\cite{rxte_catalog} and the published list of ultra-high energy cosmic rays from AGASA\\,\\cite{agasa_cluster}.  The components of our analysis are listed in Section\\,\\ref{section:components}.  Issues that relate to data and AGN selection are given in Section\\,\\ref{section:selection}.  The cosmic ray--AGN comparison results are given in Section\\,\\ref{section:comparison}. Section\\,\\ref{section:summary} summarizes this study. ",
        "conclusions": "\\label{section:summary}  We have searched for correlations between the published list of the highest energy events from the AGASA experiment\\,\\cite{agasa_cluster,AGASAwebsite} and the RXTE catalog of AGNs\\,\\cite{rxte_catalog}.  Two categories of RXTE AGNs were considered: {\\it all}-AGNs with RXTE 3-20keV intrinsic luminosities, $L_{3-20} \\geq 10^{41}$ ergs/s, and {\\it broadline}-AGNs with 2-8keV intrinsic luminosities, $L_{2-8} \\geq 10^{42}$ ergs/s motivated by the analysis of AGN evolution in Ref.\\,\\cite{barger}.  To retain RXTE source detection efficiencies $^>_{\\sim}50$\\%, source redshifts of $z\\leq0.005$ and $z\\leq0.03$ were required for the {\\it all}-AGN and {\\it broadline}-AGN categories respectively.  No correlations were observed between the AGASA events and the {\\it broadline}-AGN category of RXTE AGNs even though this category of AGN is most luminous in X-rays and even though the source density for this category of AGN is favored by some analyses\\,\\cite{blasi,kachelriess} of the highest energy cosmic rays.  In contrast, possible correlations were observed between AGASA events and the most inclusive, {\\it all}-AGN, category of RXTE AGNs.  We note that while not statistically conclusive, one of the nearby RXTE AGNs correlates with the AGASA/HiRes {\\it quartet} event cluster\\,\\cite{hires_cluster} and one correlates with one of the AGASA super-GZK events\\,\\cite{AGASAwebsite}.  Additional data would help confirm, or refute, the interesting possibility of highest energy cosmic ray--AGN correlations."
    },
    "0609/astro-ph0609324_arXiv.txt": {
        "abstract": "Convection and turbulence in stellar atmospheres have a significant effect on the emergent flux from late-type stars. The theoretical advancements in convection modelling over recent years have proved challenging for the observers to obtain measurements with sufficient precision and accuracy to allow discrimination between the various predictions.  An overview of the current observational techniques used to evaluate various convection theories is presented, including photometry, spectrophotometry, and spectroscopy. The results from these techniques are discussed, along with their successes and limitations.  The prospects for improved observations of stellar fluxes are also given. ",
        "introduction": "The gross properties of a star, such as broad-band colours and flux distributions, are significantly affected by the effects of convection in stars later than mid A-type. Consequently, our modelling of convection in stellar atmosphere models can significantly alter our interpretation of observed phenomena. By comparison with stars of known $T_{\\rm eff}$ and/or $\\log g$ (the fundamental stars), we can evaluate different treatments of convection in model atmosphere calculations. ",
        "conclusions": "The effects of convection on the stellar atmospheric structure can be successfully probed using a variety of observational diagnostics (\\cite[Smalley 2004]{SMA04}). The combination of photometric colours and Balmer-line profiles has given us a valuable insight into the nature of convection in stars. High quality observations that are currently available and those that will be in the near future, will enable further refinements in our theoretical models of convection and turbulence in stellar atmospheres."
    },
    "0609/astro-ph0609112_arXiv.txt": {
        "abstract": "% Microlensing of gravitationally lensed quasars by the stars in the foreground lens galaxy can be used to probe the nature of dark matter, to determine the mean stellar mass in the lens galaxy, and to measure the internal structure of quasar accretion disks.  Until recently, little progress has been made toward using microlensing for these purposes because of the difficulty in obtaining the necessary data and the lack of good analysis methods.  In the last few years, both problems have been solved.  In particular, Bayesian analysis methods provide a general approach to measuring quantities of physical interest and their uncertainties from microlensing light curves.  We discuss the data and the analysis methods and show preliminary results for all three astrophysical applications. ",
        "introduction": "%  Suppose I advertise to you a set of observations that simultaneously probe the nature of dark matter, the mean stellar mass in cosmologically distant galaxies, and resolves the internal structure of quasars.  This sounds like a marvelous scientific opportunity.  Yet despite having the capability of carrying out these observations for over a quarter of a century, there are few quantitative results.  The observations are the monitoring of gravitationally lensed quasars for the effects of microlensing by the stars in the lens galaxy.  In a gravitational lens, we see (usually) two or four images of a background quasar created by the deflection of light in the gravitational field of a foreground galaxy. If we then monitor the brightness of the quasar images, we will observe them to vary due to two effects.  First, if the luminosity of the quasar varies, we will see those variations in the individual images but with temporal shifts due to the different light travel times associated with each image. These time delays can be measured by cross correlating the light curves of the individual images, and they measure the surface mass density of the lens galaxy at the radius of the images from the center of the lens galaxy. Second, we will see uncorrelated variations in the brightness of the images due to {\\it microlensing} by the stars in the lens galaxy near each image of the quasar.   Our advertisement was for the scientific potential of microlensing.  Two basic problems have blocked delivering on the advertisement.  The first problem is sociological.  Time on optical telescopes is generally allocated by the night, and this sociological choice makes it nearly impossible to monitor gravitational lenses.  A scientific program to study variability in an interesting sample of lenses requires roughly 1~hour/night, {\\it every night}, on a 1--2m class telescope.   Since the time-allocation scheme does not allow for such studies,  monitoring lenses has been restricted to small numbers of lenses and short periods of time. The second problem is statistical.  While microlensing variability depends on the nature of dark matter, the mean stellar mass in the lens, and the internal structure of quasars, it is not amenable to simple analyses.   Even for the few lenses that were successfully monitored and showed microlensing variability, there was no good way to analyze the results without making compromising assumptions.  In this brief review, we outline the physics of microlensing in \\S2, the nature of the data in \\S3, and the Bayesian Monte Carlo method we have developed to analyze the data in \\S4.  In \\S5 we present results on the size of quasars, the fraction of mass in stars and the mean mass of the stars. In \\S6 we outline some of the physical, computational and statistical issues of our approach and how they might be addressed.  In \\S7 we discuss the future of the method.  Wambsganss (2006) provides a complete review of the microlensing field and Kochanek (2006) provides a review of strong lensing.  We use a flat $\\Omega_0=0.3$ cosmology and a a Hubble constant of $H_0 = 100 h$~km~s$^{-1}$~Mpc$^{-1}$.   \\begin{figure}[t] \\centerline{ \\psfig{figure=newmap1575a.ps,width=2.5in} \\psfig{figure=newmap1525a.ps,width=2.5in} } \\caption{ A typical magnification pattern for image A of Q2237+0305.  The left panel is smoothed to the source size needed to fit the optical light curves.  The right panel shows the same pattern for a source three times smaller.  Darker regions correspond to higher magnifications.  The peak magnifications drop from a factor of 10 (right panel) to a factor of 6 (left panel).  The center of the lens galaxy lies to the right. } \\label{fig:map} \\end{figure} ",
        "conclusions": "We have reached the point where we can begin to use microlensing as an industrial tool.  The data allows us to study problems -- stellar mass fractions, mean stellar masses and quasar structure -- that are virtually impossible to probe with other methods. Our main bottleneck now is time.  Time matters on two levels.  The first issue is simply observing time.  We have enough to monitor approximately 20 lenses well at one wavelength (R-band), and to obtain sparse coverage at other wavelengths (J,I,V and B-bands) and occasional data from HST and Chandra at X-ray and UV wavelengths blocked by the atmosphere.  The second issue is that the characteristic time scales for microlensing are long -- roughly a decade for any particular image.  For good statistical results we need to accumulate ``image monitoring centuries'' either by monitoring many lenses for shorter periods of time or a few lenses for very long periods of time.  Given that all the results to date are being obtained with roughly 10\\% of the time on a 1.3m telescope, there is enormous room to accelerate the scientific program.  All it requires is a few more queue-scheduled telescopes, hopefully in the 2-meter class.  Some of the most interesting problems require greater use of HST and Chandra because we expect the UV and X-ray emission of the accretion disk to come from the regions closest to the black hole.  In particular, the microlensing of wavelengths shortward of the 3500\\AA\\ atmospheric cutoff should show significant differences between the lower mass systems ($M_{BH} \\sim 10^8M_\\odot$) and the high mass systems ($M_{BH} \\sim 10^{9}M_\\odot$) because they correspond to regions still well outside the inner edge of the disk for the low mass systems but to regions very close to the inner edge for the high mass systems (see Fig.~\\ref{fig:dsize}).  The X-ray observations need to be extended from short observations sufficient to measure broad band variability, to long observations where we can examine the microlensing of the X-ray spectrum, particularly the Iron K$\\alpha$ lines.  Existing data show evidence for microlensing of the lines relative to the continuum (e.g. Dai et al. 2003), but better data is required to produce clean results."
    },
    "0609/astro-ph0609438_arXiv.txt": {
        "abstract": "The aim of this paper is to compare the two cases of an isolated neutron star, with a non-accreted crust, and that of an accreting neutron star, with an accreted crust, and try to estimate which one of the two would make a better source of gravitational waves. In order to do this we must evaluate the maximum ``mountain'' that the crust can sustain in these two cases. We first do this using the formalism of \\citet{ush} and find that the maximum quadrupole is very similar in the two cases, with the non-accreted crust sustaining a slightly larger mountain. We then develop a perturbation formalism for the problem, that allows us to drop the Cowling approximation and have more control over the boundaries. The use of this formalism confirms that there is not much difference between the two cases, but leads to results approximately one order of magnitude larger than those we obtain with the formalism of \\citet{ush}. ",
        "introduction": "The possibility that the crust of a rotating neutron star may sustain a ``mountain'', thus creating a source of gravitational waves,  has recently sparked much interest. This is mainly due to the suggestion by \\citet{bil}, echoing earlier ideas by \\citet{wag} and \\citet{pap}, that the gravitational wave torque due to such a mechanism may be dictating the spin equilibrium period for the LMXBs. A recent estimate of how large a quadrupole one may expect was given by \\citet{ush}, using the perturbation formalism of \\citet{mac}. From their work it appears that the quadrupole may be big enough to balance the accretion torque, if the breaking strain of the crust (a very poorly constrained parameter) is taken to be near the maximum of the possible values. Even though other gravitational wave mechanisms have been suggested for these systems (toroidal magnetic fields were suggested by \\citet{cutler} and r-modes by \\citet{rmode,rmode2}), and even though one can build models of the interaction between the accretion disc and the magnetic field that give the correct distribution of spin frequencies, as in \\cite{us}, the matter is still of great interest.  There are in fact many questions that can still be asked, the main one being if there is a difference between the ``mountain'' that one can build on the non-accreted crust of a rotating neutron star and the one that can be built on the accreted crust of a neutron star in a binary system (such as the LMXBs). We shall see that these two different scenarios can lead to significantly different compositions for the crust. The case of isolated neutron stars is particularly interesting as current gravitational wave detectors are beginning to put astrophysically  relevant limits on the ellipticities of known pulsars (\\cite{ligo}). In this work we will discuss this issue, and present a new scheme for calculating the maximum quadrupole which allows us to have better control of the boundary conditions at the base of the crust than in the scheme of \\citet{ush}. Another important point is that we do not adopt the Cowling approximation, which, as we will show, can lead to significant differences in the results. We intend to present a formalism in terms of perturbations of a spherically symmetric crust, without making the Cowling approximation  that was made in \\citet{ush}.   We will consider two cases for the crust core interface: in one case we will solve the hydrostatic equilibrium equations in the crust and in the core, and impose continuity of the tractions at the junction; in the other case we consider the core to remain unperturbed, and thus spherical. This is not physical but, as we shall see, allows us to estimate the errors associated with neglecting the perturbations of the core. The main point, however, is that if we leave the core unperturbed, this allows us to solve the TOV equations to obtain it's structure; we can thus obtain an estimate of how important relativistic effects may be. Such effects are relevant because, even though the Newtonian approximation is excellent at the low densities of the crust, the structure of the core can be highly relativistic and have an impact on global parameters such as the mass and radius of the star. We shall first develop the formalism for an $n=1$ polytrope and then apply it to the case of a realistic equation of state, for an accreted and a non-accreted crust. ",
        "conclusions": "We have reviewed the method of \\citet{ush} to determine the maximum deformation that a Neutron star crust can sustain. This procedure has then been applied to the case of an accreted and a non-accreted crust, in order to estimate whether one of the two would produce a better source of gravitational waves. It appears that the two kinds of crust are essentially equivalent, with the non-accreted crust being able to support a slightly larger quadrupole, as has been suggested in the literature (e.g. \\citet{sato}). It is important to note, however, that to obtain this result is has been necessary to extrapolate the composition of the crust in the high density regions at the base of the accreted crust. In order to confirm this result it would therefore be necessary to have a precise model for the composition in this region.  We have then examined a series of toy models, to determine what effect the Cowling approximation can have on the problem and, more crucially, if it is safe to assume that the shear modulus $\\mu\\longrightarrow 0$ at the base of the crust. We conclude that this is not the case, as the shear modulus is expected to vanish over a thin transition region close to the base of the crust, which can lead to large boundary terms. The Cowling approximation can also make a difference of factors of a few in the final result.  In order to drop the Cowling approximation, and have better control over the boundary conditions, we have developed a perturbation formalism and applied it to the case of an $n=1$ polytrope with a constant ratio $\\mu/\\rho$. The results of these tests have confirmed that the Cowling approximation used in \\citet{ush} can have a large impact. Our results, however, also suggest that it is not a bad approximation to consider an unperturbed core and perturb the crust, if the crust is sufficiently thick. This allows us to solve relativistic equations in the core, which can give corrections that are significant if we intend to use a realistic equation of state for the star. Having developed the formalism we apply it to the original problem of the accreted and non-accreted crusts, with a realistic model for $\\mu$, thus confirming that the two crusts can sustain very similar quadrupoles, the non accreted crust being able to sustain a slightly larger ``mountain''. What is surprising is the fact that the quadrupole found with this formalism is approximately an order of magnitude greater than the results obtained with the formalism of \\citet{ush}. Such a significant quadrupole, if it were confirmed, may allow us to put an upper limit on the breaking strain of the crust from observations. The problem, however, is not quite so simple. The calculations we have presented give a ``maximum'' quadrupole in as much as we have assumed all the strain to be in the $Y_{22}$ harmonic. But the stars we consider are rotating and must therefore be rotationally deformed, thus putting strain also in the $Y_{20}$ harmonic, as spin-up or spin-down moves the star away from its reference shape. This could bring the crust much closer to the yield point and significantly decrease the size of any additional ``mountain'' that it could sustain. Unfortunately, estimating exactly how much strain is built up by rotation is a complicated problem, as the unstrained configuration of the crust depends on the history of the crust, and there may have been many starquakes during the spin-down and subsequent spin-up phases. However, if a neutron star can sustain a quadrupole this large for a breaking strain of $\\sigma_{\\mbox{max}}=10^{-2}$, it would appear that even if the breaking strain is smaller there is still a possibility that the quadrupole may be large enough for gravitational waves to have a role in the spin evolution of the star.  Finally we should note that our approach was chosen in order to obtain a ``maximization'' argument, independent of the mechanism that produces the deformation, and to show that overall relativistic effects can be important. We have, therefore, integrated a set of source free equations, analogous to those of \\citet{mac}, with the deformation as a boundary condition. It would, however, be important to understand the physical processes that cause the deformation and  integrate the same equations with a source, such as the temperature or composition variations used in \\citet{ush}. Another important effect we are neglecting is that of the magnetic field. It is well known that the equilibrium shape of a star will not be spherical in the presence of a magnetic field and, even though crustal effects may be dominant early on during accretion, the process of magnetic burial could make the field locally strong enough for the magnetic deformations to become important. This effect was recently studied by \\citet{melatos} and \\citet{payne}, who were the first to perform globally self-consistent MHD calculations of magnetic mountains on accreting NS as gravitational wave sources. These calculations did not include crustal effects (as this would make the problem much less tractable), nevertheless it is important to note here that if one were to solve the full inhomogeneous problem, including the magnetic source terms (or terms due to other physical process we have neglected), one could obtain deformations larger than the crustal maximum we have calculated.  It would also be of some interest to consider the full problem in a relativistic formulation of elasticity, such as that of \\citet{lars}, as we have shown that the perturbations of the core can be important, especially in determining the dependence of the maximum quadrupole on the stellar parameters. We intend to turn our attention to these problems in the near future."
    },
    "0609/astro-ph0609262_arXiv.txt": {
        "abstract": "Since the introduction of the {\\sl Amati relation} (Amati et al. 2002), other largely promising techniques have appeared in an attempt to turn GRBs into reliable cosmological probes. These incluide the {\\sl Ghirlanda relation} (Ghirlanda et al. 2004) and the recently discovered {\\sl Firmani et al. relation} (Firmani et al. 2006a),  all of which take into account the most relevant physical properties of gamma-ray bursts (GRBs) as peak energy, jet openning angle, and both time lag and variability. Such discoveries hints at the long-sought {\\sl Holy Grail} of creating a cosmic ruler from  GRBs observables to be achievable. In this {\\sl Letter} we construct the Hubble diagram (HD) for the standard Friedmann-Lema\\^itre-Robertson-Walker (FLRW) cosmological model after enforcing it with the general relativistic energy conditions, heeding to investigate whether it still stands on as the leading scenario for cosmology in face of the distance modulus-redshift relation of a sample of GRBs that had their redshifts properly estimated and corrected upon applying on the data analysis the above quoted relations. Our $\\chi^2$ analysis support the view that FLRW plus the strong energy condition (SEC) is what better fits the GRB data. But this is not the whole story, since for a cosmological constant $\\Lambda \\neq 0$ the FLRW+SEC analysis with undefined $p=p(\\rho)$ suggests that $\\Lambda$ is not constant anyhow, because it does not follows the $p= \\omega \\rho$ HD, with $\\omega=-1$. Thenceforth, one concludes that either the cosmological constant does not exist at all, and consequently one is left with the FLRW+SEC which fits properly the GRBs HD and there is no late-time acceleration, or it does exist but should be time-varying. This all is contrary to current views based on SNIa observations that advogate for an actual constant $\\Lambda$ and an accelerating universe. In connection to the results above, this last argumentation would imply that either there is something wrong with the SNIa interpretation regarding cosmic late-time acceleration (Middleditch 2006) or we will have to move away from general relativity. We would better to think that general relativity is still the most correct theory of gravity. ",
        "introduction": "Gamma-ray bursts (GRBs) are the biggest explosions in the universe. The major breakthrough in understanding of GRBs came with the Beppo-SAX satellite discovery of  the X-ray afterglow of  GRB970228 (Costa et al. 1997). A swift follow-up of this event allows the subsequent detection of residual optical and radio emissions from this transient (van Paradijs et al. 1997; Frail et al. 1997). From these afterglows the first measurement of the redshift of a GRB was done, what definitely confirmed a long-standing suspicion that GRBs have cosmological origin. Such realization opens-up the possibility of using GRBs for a couple of fundamental researches in contemporary astrophysics and cosmology: a) as dust absorption-free tracers of the massive star formation history in the universe (Totani 1997; Pacsynski 1998; Wijers et al. 1998; Blain \\& Natarajan 2000, Lamb \\& Reichart 2000; Bromm \\& Loeb 2002, Lloyd-Ronning \\& Ramirez Ruiz 2002; Lloyd-Ronning, Fryer \\& Ramirez Ruiz 2002; Firmani et al. 2004; Yonetoku et al. 2004; Friedman \\& Bloom 2005; Schaefer 2006\\footnote{Regarding the recent claim by Schaefer (2006) on the exclusion of the cosmological constant $\\Lambda$ after confronting it with the GRB HD, a very interesting counter-argument was provided by Friedman (2006). He argues that the correct HD  should also show data calculated for a cosmology with $\\Lambda$, and that only a comparison of relative $\\chi^2$ can determine the favored cosmology. We agree with this statement, and as we demonstrate below that condition should be supplemented with the need for the energy conditions to be satisfied, too.}), and b) as standard candles and hence as cosmic probes with which one can follow back in time upto very high redshifts the expansion history of the universe (Schaefer 2003; Ghirlanda et al. 2004; Dai, Liang \\& Xu 2004; Firmani et al. 2005; Liang \\& Zhang 2005; Xu et al. 2005).  Unfortunately, to accomplish both of these key investigations is not an easy task since both face the problem of the small statistics of events with measured readshift ($z$) because, to determine $z$, is required deep-space optical/infrared or X-ray spectra. With regard to the item b), the large dispersion of the energetics of GRBs precludes of granting GRBs the status of cosmic rulers (Bloom, Frail \\& Kulkarni 2003). Nonetheless, a major effort is being pursued to overcome both the above mentioned drawbacks with significant achievements after realizing that there exist tight relations that connect GRB intrinsic energetics and/or luminosity with other physical observables (Ghirlanda et al. 2004; Firmani et al. 2006a). These relations become powerful distance indicators, and in events such supernovae (Woosley \\& Bloom 2006); for which $z$ can be inferred via an energetics-independent procedure, GRBs become standard candles for practical cosmographic studies (see Firmani et al. 2006a for a full accounting of this impressive progress).   Presently, after the first series of controversial statements on the viability of granting to GRBs the status of standard cosmic rulers (Bloom 2002; Schaefer 2002, 2003; Bloom, Frail \\& Kulkarni 2003; Friedman \\& Bloom 2005), a definite consensus appears to be arising, and the hope to have a {\\sl Holy Grail} to do cosmology upon GRBs is renascenting (Ghisellini \\& Lazzati 2004 (GGL2004); Friedman \\& Bloom 2005;  M\\\"ortsell \\& Sollerman 2005; Schaefer 2006). GGL2004 gave to the search for a {\\sl Holy Grail} a renascence after finding that the Amati et al. (2002) correlation becomes much tighter if one corrects $E_{\\rm iso}$ for the collimation factor of the GRB opening angle $\\theta_{\\rm j}$, which is the rest-frame time of the achromatic break in the light curve of an afterglow.\\footnote{It has been pointed out that the presence of a collimated break in GRB afterglow light curve is a strong evidence that the GRB emission is collimated into a cone of semi-aperture angle $\\theta_{\\rm j}$.} Thus the angle-corrected  $\\gamma$-energetics reads: $E_{\\gamma} = (1 - \\cos \\theta_{\\rm j}) E_{\\rm iso}$. The $E_{\\gamma} - E_{\\rm peak}$ correlation\\footnote{This is a power-law relation $E_{\\rm peak} - E_{\\gamma}$, where : $E_{\\rm peak} = \\tilde{\\kappa} \\left( \\frac{E_\\gamma} {E_0} \\right)^\\eta $, that relates the energy at the maximum ($E_{\\rm peak}$) of a given GRB with its total energy in $\\gamma$-rays ($E_{\\gamma}$), with the constant $E_0 = h^{-3/2}$ yielding the best-fit parameters $\\tilde{\\kappa}$ and $\\eta$ independent of the value of the Hubble parameter $h$.} have been used to put constraints on some cosmological parameters avoiding the {\\sl circularity problem} (Ghirlanda et al. 2004; Firmani et al. 2005). Besides, Firmani et al. (Firmani et al. 2006a) recently discovered new tigh correlation among prompt emission properties in long GRBs, which involves the isotropic peak luminosity, $L_{\\rm iso}$, the peak energy, $E_{\\rm peak}$, in the spectrum, $\\nu L_\\nu$, of the time-integrated prompt emission and the ``high signal\" time-scale, $T_{0.45}$, that was previously used to characterize the variability behaviour of GRBs. In the source rest-frame, this relation follows as: $L_{\\rm iso} = \\tilde{K} E_{\\rm peak}^{1.62} T_{0.45}^{ -0.49} $, with $\\tilde{K}$ a constant. It uses the GRB openning angle-corrected redshifts into $L_{\\rm iso}$.   Here we investigate whether the energy conditions (ECs) of general relativity (GR) when enforced on the standard  cosmological FLRW model can bring it a pace with the Hubble diagram (E. P. Hubble, 1929, Proc. Nat. Acad. Sci. USA, 15, 168) constructed from the sample of GRBs that have redshifts measured through spectroscopy or other methods (like photometry). For the present analysis we shall use a set of three samples of GRBs that had their redshifts estimated: a) Bloom, Frail \\& Kulkarni (2003; BFK2003), as shown in Table-1 and Figure 1a). b) The Schaefer (2006) GRB set analyzed with five methods, shown in Figure 1b). c) the $\\theta_{\\rm j}$-corrected GRBs sample, as from the discovery by Firmani et al. (2006a), plotted in Figure 1c).         \\begin{table} \\begin{center} \\caption{Selected sample of 24 GRB events with measured redshift, computed distance modulus and error bars, as tabulated by Bloom, Frail \\& Kulkarni (2003). \\label{tbl-1} } \\begin{tabular}{ccccc} \\tableline \\tableline GRB & Redshift & Dist.-Modulus & error & $\\theta_{\\rm j}$-Corrected\\footnotemark[1] \\\\ event & $z$ &  $\\mu(z)$ & $\\pm$ &  $\\bar{z}$\\\\ \\tableline 970508 & 0.8349 & 45.10 & 0.47 & -\\\\ 970828 & 0.9578 & 43.86 & 0.45 & 1.212\\\\ 971214 & 3.4180 & $<$47.83 & 0.37 & 4.801\\\\ 980613 & 1.0969 & $<$ 47.03 & 0.54 & - \\\\ 980703 & 0.9662 & 44.31 & 0.45 & 1.104\\\\ 990123 & 1.6004 & 44.00 & 0.53 & 1.358\\\\ 990510 & 1.6187 & 47.13 & 0.42 & 1.210\\\\ 990705 & 0.8424 & 43.98 & 0.43 & 1.109 \\\\ 990712 & 0.4331 & $<$42.13 & 0.33 & -\\\\ 991208 & 0.7055 & $>$43.35 & 0.86  & -\\\\ 991216 & 1.0200 & 44.02 & 0.83 & 0.823\\\\ 000131 & 4.5110 & $>$47.06 & 0.44 & 4.957\\\\ 000210 & 0.8463 & $<$43.90 & 0.29 & -\\\\ 000301C & 2.0335 & 46.30 & 0.40 & -\\\\ 000418 & 1.1181 & 42.97 & 1.43 & -\\\\ 000911 & 1.0585 & $>$44.00 & 0.46 & 1.173\\\\ 000926 & 2.0369 & 45.93 & 0.92 & -\\\\ 010222 & 1.4768 & 45.28 & 0.21 & -\\\\ 010921 & 0.4509 & 41.71 & 0.62 & -\\\\ 011121 & 0.3620 & $<$41.35 & 1.08 & -\\\\ 011211 & 2.1400 & 47.54 & 0.45 & -\\\\ 020405 & 0.6899 & 43.99 & 0.67 & -\\\\ 020813 & 1.2540 & 45.00 & 1.07 & -\\\\ 021004 & 2.3320 & 46.48 & 1.80 & -\\\\ \\tableline \\end{tabular} \\footnotetext[1]{Redshifts in this column were corrected by using the {\\sl Firmani et al. relation} } \\end{center} \\end{table} ",
        "conclusions": "In Figures 1a), 1b) and 1c), we present the main results of this investigation. First, after looking at Figures 1a) and 1b) one can verify that the upper DEC bound (dash-dotted line) matches perfectly the WEC $\\mu(z)$ prediction (open circle line). Besides, one can also realize that both the DEC upper and lower bounds are respected by the GRBs observations, that is, each of the bursts that have been observed has its representative point in the FLRW HD inside the limits of DEC1 and DEC2 in Figures 1a), 1b). Besides, there is no clear violation of either the WEC or the SEC. This results go almost contrary to conclusions achieved by Santos, Alcaniz \\& Rebou\\c cas (2006); who showed, after confronting the ECs with SNIa observations, that practically each of the ECs seems to have been violated in the recent past history of the cosmic evolution. We argue that as GRBs are almost free of any effect as absorption or the like, one is inclined to conclude that perhaps the astrophysics of SNIa has not been exhaustively understood (Middleditch 2006), and hence the interpretation of luminosity-dimmed SNIa as messengers of late-time acceleration should decidely be rethought, perhaps in the lines of Middleditch (2006).   Further, it is quite easy to see that the SEC prediction seems to fit more properly the GRBs HD  than any another of the ECs+FLRW. This is quite a surprising result. It seems to indicate, contrarily to most of the current expectations (see Firmani et al. 2006b), that the cosmological constant $\\Lambda$ (i.e. EoS with $\\omega =-1$) is not the driving force pushing out the universe in the late times since it ``naturally\" violates the SEC, as pointed out earlier. In fact, looking at Figure 1b), where Schaefer's (2006) 52 GRBs are plotted, this behavior seems to extend down to lower redshifts $z \\sim 0.1-0.5$. Thus our ECs+FLRW plus GRB HD seem to exclude any need for dark energy.  To bring our results into consistency, we have performed a $\\chi$-square analysis to compare the distance-modulus vs. redshift relation as predicted by each of the ECs+FLRW cosmological scenarios with the HD obtained from GRBs observations. Here we define  \\[\\chi^2 = \\sum_i \\frac{[\\mu^0_i(z) - \\mu^t_i(z)]^2} {\\sigma^2_{\\mu^0_{i}} } \\; , \\]  where respectively $ \\mu^0_i(z) $ and $ \\mu^t_i(z) $ define the measured and calculated (theoretical) value for the distance-modulus, and $\\sigma_{\\mu^0_{i}}$ represents the dispersion in the $\\mu^0_i(z)$ (see Table 1). Despite the $\\mu(z)$ uncertainties, the data clearly selects the SEC+FLRW as the best fit. Our calculations for the 24 GRBs from BFK2003 show that the average $\\chi^2$ for the WEC 13.32, for the SEC 4.864, and for the DEC 26.47. For the 52 Schaefer (2006) GRBs the $\\chi^2$ figures go as follows: for the WEC 12.17, for the SEC 1.36, and for the DEC 10.78. This therefore confirms what a visual inspection of Figures 1a) and 1b) allows to conclude. Hence, as the SEC implies that the universe is all the time decelerating irrespective of the type of curvature it may have, then one is forced to conclude that the cosmological constant is playing no role at all in the universe dynamics since it naturally violates the SEC. Nor it is being the fundamental field driving the purported universe late-time acceleration since there seems to be not such a transition down redshifts $z \\sim 0.5-1.0$. This in passing would imply that there is no need for the own dark energy field, as it appears to be ruled out by the GRBs HD. Besides, DE should comply with the ECs, and they (in particular the SEC) show that such a transition-to-acceleration at low redshifts seems to have not occurred.   Meanwhile, to elucidate any undisputable role of the cosmological constant we now impose the SEC $\\rho + 3p \\geq 0$ bound on the Friedmann's equations with $\\Lambda \\neq 0$ but unspecified $p=p(\\rho)$ EoS\\footnote{Notice that DE scenarios with $\\Lambda$ constant! have EoS $p = \\omega \\rho$, with $\\omega = -1$. This model clearly violates the SEC bound since SEC $\\longrightarrow p \\geq - 1/3 \\rho$. However, it provides a dynamics: $\\dot{R} \\geq H_0 R$, that follows the WEC bound given by the top curves of HDs in Figure 1.}, one obtains the dynamics (for $k=0$)  \\be \\frac{\\ddot{R}}{R} \\leq  \\frac{\\Lambda}{3}  \\; \\Longrightarrow \\left(\\frac{\\dot{R}}{R_0}\\right)^2 \\geq H_0^2 -  \\frac{\\Lambda}{3} \\left[ \\frac{z (z + 2)}{(1 + z)^2} \\right] \\; . \\label{lambda-neq-0} \\ee  (One can verify that for $\\Lambda = 0$ one recovers Eq.(\\ref{SEC-DYNAMIC}), above). Thus, the luminosity-distance becomes  \\be d_L \\leq  \\frac{(1 + z)}{\\left(H_0^2 - \\frac{\\Lambda}{3} \\left[ \\frac{z (z + 2)}{(1 + z)^2} \\right] \\right)^{1/2}} \\ln(1 + z) \\; , \\label{dl-lambda} \\ee  and therefore, the $\\mu(z)$-$z$ HD reads:  \\be {\\rm SEC + FLRW + \\Lambda \\; Hubble\\; diagram} \\atop \\; \\mu(z)-z \\;  \\Longrightarrow \\left\\{ \\mu(z) \\leq 5 \\log_{10} \\left( \\frac{(1 + z) \\; \\ln(1 + z) }{\\left(H_0^2 - \\frac{\\Lambda}{3} \\left[ \\frac{z (z + 2)}{(1 + z)^2} \\right] \\right)^{1/2}} \\right) + 25 \\label{mu-SEC} \\right. . \\ee  Putting $\\Lambda = 2.036 \\times 10^{-35}$ s$^{-2} = 0.214 {\\rm Gpc}^{-2}$, and $h=0.73$ one verifies that according to Figures 1a), 1b) and 1c) it gets close to the SEC limit of Eqs.(\\ref{hubble-law},\\ref{friedmann-equations}) without $\\Lambda$. The $\\chi^2$ analysis for this case gives 1.33 for the 52 GRBs and 4.36 for the 24 GRBs. It is a bit better than the SEC limit alone. However, the error bars do not allow to claim that there is any actual difference between the $\\chi^2$ value for SEC+FLRW+$\\Lambda$ and the sole for SEC+FLRW. Thus, this result can be interpreted in the lines of our discussion above as either indicating that there is no need for $\\Lambda$, or that $\\Lambda$ is time-variable. As the HD for the case $\\Lambda \\neq 0$ plus SEC+FLRW does not coincide with the one corresponding to the case $p= -\\rho$ (footnote {8}) then one is led to conclude that the cosmo-(i)logical constant does not exist at all or it is not constant either. Figures 1 illustrate the behavior of a $\\Lambda$+SEC+FLRW for different values of the Hubble parameter $h$.   In summary, despite the relatively large uncertainties still present in the luminosity distance of most GRBs and the small statistics of GRB with low redshifts, the approach introduced here appears to be more conclusive than an analysis standing on SNIa data alone. The main reason is that the sample of GRBs that we are using for the present analysis already reaches farther out, upto redshifts $z\\sim 7$ where the early-days cosmology was playing hard, while SNIa are naturally limited to distance scales no farther out than $z \\sim 1$. (Surely, the set of both lower (about $z \\sim 0.1-0.5$), and larger (upto $z \\sim 10-20$) redshifts will raise in the near future when more GRBs are to be observed by the SWIFT satellite).  On physical grounds, the outcome of this inedit test suggests that the FLRW model when faced with the GRB HD guarantees the prevalence of the WEC, SEC and DEC bounds, with the SEC being the better cosmological fit to the GRBs observational data. As SNIa are facing problems to go a pace with the ECs+FLRW, we think that it is the SNIa astrophysics that should be reviewed instead of general relativity (Middleditch 2006).  To the last, the GRBs HD has been built to demonstrate the reliability of GRBs as an independent, but nonetheless complementary tool to SNIa (and perhaps to CMB observations), to confront current cosmological models. Thus, everything now allows to forsee a promising future for a high-precision era in GRBs cosmology, a field which is by-now at its infancy (Lazzati et al. 2006), based on these novel relations that promise to become a major breakthrough in GRBs cosmology. This is the first paper of a thorough research of this key issue in current Cosmology. In forthcoming papers other cosmological models, specially those purporting modifications of gravity at very larger scales, will be analyzed in face of GRBs data to cross-check them for consistency with the HD of GRBs."
    },
    "0609/astro-ph0609748_arXiv.txt": {
        "abstract": "{Several studies have investigated the fractal and multifractal nature of magnetic features in the solar photosphere and its variation with the solar magnetic activity cycle.} {Here we extend those studies by examining the fractal geometry of bright magnetic features at higher atmospheric levels, specifically in the solar chromosphere. We analyze structures identified in CaIIK images obtained with the Precision Solar Photometric Telescopes (PSPTs) at Osservatorio Astronomico di Roma (OAR) and Mauna Loa Solar Observatory (MLSO).} {Fractal dimension estimates depend on the estimator employed, the quality of the images, and the structure identification techniques used. We examine both real and simulated data and employ two different perimeter-area estimators in order to understand the sensitivity of the deduced fractal properties to pixelization and image quality.} {The fractal dimension of bright 'magnetic' features in CaIIK images ranges between values of 1.2 and 1.7 for small and large structures respectively. This size dependency largely reflects the importance of image pixelization in the measurement of small objects. The fractal dimension of chromospheric features does not show any clear systematic variation with time over the period examined, the descending phase of solar cycle 23.} {These conclusions, and the analysis of both real and synthetic images on which they are based, are important in the interpretation of previously reported results.} ",
        "introduction": "The spatial distribution of the solar magnetic fields is very complex, depending on both the varying level of magnetic activity over the solar cycle and the height of observation in the solar atmosphere. This complexity likely results from both the dynamo process itself, which may occur on many different spatial scales, and the interaction of the field with convective motions as it emerges through the Sun's outer layers. The signatures of these processes have been investigated in previous works by fractal analyses of solar active regions, but the quantitative results obtained differ widely depending on the type of data and analysis techniques employed \\citep[e.g.][]{Jansen2003,McAteer2005}. Moreover, the lack of a unique definition of the fractal dimension itself often makes comparison of results difficult.  Among recent studies, one focused on the possible relationship between magnetic feature complexity and solar cycle phase \\citep{Meunier2004}. This study looked at a large sample of active region magnetograms acquired with the Michelson Doppler Interferometer (MDI) aboard the Solar and Heliospheric Observatory (SOHO) between April 1996 and June 2002. The measured fractal dimension increased with structure size (in agreement with \\citet{Meunier1999} and \\citet{nesme96}), showing a peculiar change in behavior near structures of area 550-800 ${\\rm Mm}^{2}$. A similar dependence on structure size was also found by \\citet{Jansen2003} in both high resolution photospheric magnetograms and magnetohydrodynamic (MHD) simulations. Additionally, \\citet{Meunier2004} investigated the relationship between the geometry of facular structures of different spatial scales and magnetic field intensity, flare activity, and solar cycle phase. She found that, while a region complexity generally increases with magnetic field intensity, there is no clear correlation with flare activity. Variations of fractal dimension with solar cycle were also reported, but their amplitudes and sign largely depended on object size and the associated magnetic field.   In order to investigate the complexity of magnetic features using observations representative of chromospheric heights, we have analyzed the fractal dimension of bright features identified in full-disk CaIIK images acquired by the Precision Solar Photometric Telescopes (PSPTs) at Osservatorio Astronomico di Roma (OAR) and Mauna Loa Solar Observatory (MLSO). The data analyzed span the past 6 years and thus allow investigations of variation with the solar cycle.  Several factors can influence fractal dimension estimation. Both image resolution and projection effect correlation, mass function, and perimeter-area estimators in studies of interstellar molecular clouds \\citep{Sanchez2005,Vogelaar1994}, and resolution and thresholding effects are important in fractal dimension estimation of flow patterns in field soil by box-counting methods \\citep{Baveye1998}. \\citet{Lawrence1996} studied similar effects in multifractal and fractal measures (box-counting, cluster dimension, threshold set) of solar magnetic active regions. The effect of structure selection technique was also investigated by \\citet{Meunier1999,Meunier2004}.  To link our findings directly to those of several recent solar studies \\citep[for example][]{Meunier2004,Jansen2003} we have employed the  perimeter-area relationship to define the fractal dimension of the identified features. To interpret our results, we have investigate the sensitivity of the deduced fractal dimension to the pixelization and resolution of the image and to the perimeter measure algorithm employed. In particular, we have determined how these factors influence the geometric properties deduced for objects of different sizes, and more generally have addressed the question of whether the perimeter-area relation is suitable to the study of the fractal and multifractal nature of solar magnetic features as a function of the solar cycle.  This paper is organized as follows. In the next section we describe the observations, data processing techniques, and geometrical measures employed. In section \\S3 we present the results obtained and in \\S4  compare them to those of previous efforts. In \\S5 we investigate the sensitivity of the deduced fractal dimension to the image resolution and the measurement techniques employed, by examining synthetic structures whose fractal properties are theoretically known: non-fractal objects, von Koch snowflakes, and those produced by fractional Brownian motion (fBm). The effect of atmospheric seeing is also investigated using PSPT images taken under variable seeing conditions. In section \\S6 we discuss our and previous results, in light of the conclusions drawn in \\S5. Finally, in \\S7 we summarize our work and conclude with a more general discussion on the validity of the methods. ",
        "conclusions": "We have analyzed the fractal dimension  of bright features identified in solar chromospheric CaII K images obtained during the last six years, corresponding to the descending phase of solar cycle 23. The results obtained are in general agreement with those reported in literature, in particular with those studies that have been carried out with similar fractal estimators (\\begin{em}d1\\end{em} and \\begin{em}D\\end{em}).  We have also investigated the effects of pixelization and resolution on the fractal dimension estimates obtained, studying these effects on real and simulated data, the latter including both non-fractal objects and objects whose fractal properties are well known, von Koch snowflakes and fBm images. We have shown that fractal dimension estimates suffer from pixelization errors at all scales, but errors are generally more important for areas less than  $\\sim500 - 1000$ pixel$^2$. Particularly, pixelization causes measured fractal dimensions to increase with object size in the case of both fractal and non-fractal objects. Our results thus indicate that the increase of fractal dimension with the feature size reported in literature by some previous analyses is likely  an effect of pixelization and image degradation, rather than a signature of an intrinsic multifractality of active regions. To reduce these effects, we thus suggest a restriction of the analyses to objects whose areas are larger then the quoted value.  Our analyses also showed that image degradation due to both seeing and instrumental effects smooths features edges making them appear rounder. The fractal dimension estimated is consequently lower than expected, even for large objects. Perhaps image degradation effects can be compensated for, as suggested in \\citet{Jansen2003}. Alternatively, a more careful estimate of these effects can be carried out by the analyses of images of known complexity objects (eg. fBms) convolved with realistic Point Spread Functions. We did not apply either of these compensations to our data and leave them for future investigation.  Finally we note that, as demonstrated by \\citet{Baveye1998} for the box counting method, pixelization effects can influence other fractal dimension estimators as well, and careful quantitative measure of the effects for each measure employed is essential to the interpretation of the results. We have not addressed this problem in this work and leave also this issue for future research."
    },
    "0609/astro-ph0609054_arXiv.txt": {
        "abstract": "Recent observations of the binary system \\linebreak LS~5039 with the High Energy Stereoscopic System \\linebreak  (H.E.S.S.) revealed that its Very High Energy (VHE) $\\gamma$-ray emission is modulated at the 3.9 days orbital period of the system. The bulk of the emission is largely confined to half of the orbit, peaking around the inferior conjunction epoch of the compact object. The flux modulation provides the first indication of $\\gamma$-ray absorption by pair production on the intense stellar photon field. This implies that the production region size must be not significantly greater than the gamma-gamma photosphere size ($\\sim 1\\,\\mathrm{AU}$), thus excluding the large scale collimated outflows or jets (extending out to $\\sim 1000\\,\\mathrm{AU}$). A hardening of the spectrum is also observed at the same  epoch between 0.2 and a few TeV which is unexpected under a pure absorption scenario and could rather arise from variation with phase in the maximum electron energy and/or the dominant VHE $\\gamma$-ray production mechanism. This first-time observation of modulated $\\gamma$-ray emission allows precise tests of the acceleration and emission models in binary systems. ",
        "introduction": "\\label{intro} In the commonly accepted paradigm, microquasars consist of a stellar mass black hole fed by a massive star. They can exhibit superluminous radio jets\\cite{Mirabel:Nature94}, and hints for the presence of an accretion disk. These {\\it scaled down} versions of Active Galactic Nuclei (AGN) also show, due to their much lower mass, variability on timescales shorter by several order of magnitudes (down to minutes or even seconds) that can be used to constrain the accretion and ejection scenarios (e.g. \\cite{Fender:1997}).  It has been suspected for a long time that these object could emit, through similar acceleration mechanisms as in AGN (leptonic or hadronic processes) high energy radiation up to the Very High Energy gamma-ray (VHE) domain, and that VHE radiation could give insight on the very central engine. The discovery  of VHE emission from LS~5039 by HESS\\cite{Aharonian:HESS_LS5039} and shortly after from LS~I~+$61^\\circ~ 303$ by MAGIC\\cite{Albert:MAGIC_LSI} confirmed this longstanding issue and established VHE astronomy as a powerful diagnostic probe of these objects.  LS~5039, identified in 1997\\cite{Motch:1997} as a massive X-ray binary system with faint radio emission\\cite{Marti:1997}, was resolved in 2000\\cite{Paredes:2000} into bipolar mildly relativistic radio jets ($v\\sim 0.2\\ c$) emanating from a central core, thus placing it into the {\\it microquasar class}. The detection of radio and variable X-ray emission\\cite{Bosch:2005} and its possible association with the EGRET source 3EG J1824-1514\\cite{Paredes:2000} suggests the presence of multi-GeV particles possibly accelerated in jets. LS~5039 is indeed the only object simultaneously detected in X-ray and radio in the field of view of the unidentified EGRET source\\cite{Ribo:Thesis}.  The binary system LS~5039 (Fig \\ref{fig:Cartoon}) consists of a massive 06.5V star in a $\\sim 3.9$ day mildly eccentric orbit ($e = 0.35$)\\cite{Casares:2005} around a compact object whose exact nature (black hole or neutron star) is still under debate. Under the assumption of pseudo-synchronization, Casares and collaborators\\cite{Casares:2005} constrain the compact object mass in the black hole range ($M_X = 3.7_{-1}^{+1.3} M_\\odot$) and obtain a low system inclination ($i \\sim 25^\\circ$), but a neutron star at higher inclination ($i \\sim 60^\\circ$) might still remain possible\\cite{Dubus:2006b}.   \\begin{figure} \\centering \\includegraphics[width=0.4\\textwidth]{Cartoon_Absorption_NoAxis.eps} \\caption{Orbital geometry of the binary system LS~5039 viewed from above and using the orbital parameters derived in\\cite{Casares:2005}. From \\cite{Aharonian:HESS_LS5039:2}. Shown are: phases ($\\phi$) of minimum ({\\it periastron}) and maximum ({\\it apastron}) binary separation; epoch of superior and inferior conjunctions occurring when the compact object and the star are aligned along the observer light-of-sight.} \\label{fig:Cartoon}       % \\end{figure} ",
        "conclusions": "New observations by HESS have established orbital modulation of the VHE $\\gamma$-ray flux and energy spectrum from the XRB LS~5039. The flux vs. orbital phase profile provides the first indication for $\\gamma$-ray absorption within an astrophysical source, suggesting that a large part of the VHE $\\gamma$-ray production region lies inside the pair absorption photosphere (within $\\sim 1\\,\\mathrm{AU}$) around the massive stellar companion. However, not all of the observed effects can be explained by absorption alone. A detailed study is now required to fully explain these new observations and understand the complex relationship between $\\gamma$-ray absorption and production processes within these binary systems."
    },
    "0609/astro-ph0609218_arXiv.txt": {
        "abstract": "If dark energy can be described as a perfect fluid, then, apart from its equation of state relating energy density and pressure, we should also especify the corresponding rest frame. Since dark energy is typically decoupled from the rest of components of the universe, in principle such a frame could be different from that of matter and radiation. In this work we consider the potential observable effects of the motion of dark energy and the possibility to measure  the dark energy velocity relative to matter. In particular we consider the modification of the usual interpretation of the CMB dipole and its implications for the determination of matter bulk flows on very large scales.  We also comment on the possible origin of a dark energy flow and its evolution in different models. ",
        "introduction": "The content of the universe can be appropriately described as a four-fluid system: baryons, radiation, dark matter and dark energy. Such fluids can be considered in good approximation to be decoupled from each other after recombination time. Each  fluid carries its own rest frame. Thus for instance, the CMB radiation, being highly homogeneous allows us to define its rest frame by means of the CMB dipole anisotropy. Indeed, in the usual interpretation \\cite{dipole}, the CMB dipole is due to the Doppler effect caused by the motion of the observer with respect to the last scattering surface. Therefore the radiation rest frame would be that in which an observer measures a vanishing dipole. This fact in turn has been used  to calculate the velocity of our Local Group with respect to the CMB radiation, just by substracting the velocity of the Sun with respect to the Milky Way and the Milky Way velocity with respect to the Local Group. The result is $v_{LG-CMB}=627\\pm 22$ km s$^{-1}$ in the direction $(l,b)=(276\\pm 3^\\circ,30\\pm 3^\\circ)$ in galactic coordinates.  The rest frame of (dark) matter is not so easily defined since we know that matter distribution is only homogeneous on very large scales. The presence of density inhomogeneities makes the velocity of a given matter volume to deviate from the pure Hubble flow. Such deviations $v$ are known as peculiar velocities and they are related to the density perturbations $\\delta(x)$ by this simple expression to first order in perturbation theory: \\begin{eqnarray} \\nabla \\cdot v=-H_0\\Omega_M^{0.6}\\delta(x) \\end{eqnarray} According to the Cosmological Principle, the universe is homogeneous on very large scales and the amplitude of  density perturbations on scales of size $R$ will decline as we take larger and larger values of $R$. This in turn means that the peculiar velocity of a matter volume of radius $R$ will also decline for larger $R$ (see Fig. 1), the matter rest frame eventually converging to the radiation frame on very large scales. In other words, matter and radiation would share a common rest frame according to the Cosmological Principle.  However as shown in Fig. 1, this theoretical framework is not conclusively confirmed by  observations. Indeed, in recent years several peculiar velocity surveys \\cite{surveys} have tried to determine the volume size at which the streaming motion of matter with respect to the CMB vanishes. In  the figure the results of different observations are compared with the {\\it rms} expected bulk velocity $V_b$ for  standard $\\Lambda$CDM model in a sphere of radius $R$. The results seem to agree with the theoretical expectations only at scales $R\\lsim 60h^{-1}$ Mpc. At larger scales, $R\\gsim 100h^{-1}$ Mpc, different data sets lead to different bulk velocities both in amplitude and direction. Moreover, there are indeed measurements in which  large matter volumes are moving at speeds $\\gsim 600$ km s$^{-1}$ with respect to the CMB frame, several standard deviations away from the theoretical predictions. These results have been argued to be affected by systematic errors in distance indicators, but if confirmed by future surveys, a revision of some of the underlying ideas in  Standard Cosmology would be required in order to understand the origin of such large flows.  \\begin{figure} \\includegraphics[height=.4\\textheight]{zarouibi.eps} \\caption{\\footnotesize Bulk  velocities vs. size $R$ for different peculiar velocity surveys \\cite{surveys} centered around the Local Group in the CMB frame. The dot labelled by COBE denotes the velocity of the Local Group measured by COBE. The solid line corresponds to the expected {\\it rms} velocity in the standard $\\Lambda$CDM cosmology, together with the 90\\% deviation in dashed lines. Figure from S. Zaroubi \\cite{zaroubi}} \\end{figure}  Concerning the dark energy rest frame, it is usually assummed that dark energy does not cluster on small scales and therefore can be considered as a highly homogeneous fluid which is almost decoupled from the rest of components, its only interaction being of gravitational nature. In such a case there is no reason to expect that its rest frame should necessarily agree with the matter/radiation frame and therefore an observational determination of the dark energy bulk velocity would be needed. In the following we will explore this problem, considering cosmological models with four decoupled homogeneous fluids and different rest frames \\cite{maroto}. ",
        "conclusions": ""
    },
    "0609/astro-ph0609732_arXiv.txt": {
        "abstract": "In many magnetized, dilute astrophysical plasmas, thermal conduction occurs almost exclusively parallel to  magnetic field lines.  In this case, the usual stability criterion for convective stability, the Schwarzschild criterion, which depends on entropy gradients, is modified.  In the magnetized long mean free path regime, instability occurs for small wavenumbers when $(\\partial P/\\partial z)(\\partial \\ln T / \\partial z) > 0$, which we refer to as the Balbus criterion.  We refer to the convective-type instability that results as the magnetothermal instability (MTI).  We use the equations of MHD with anisotropic electron heat conduction to numerically simulate the linear growth and nonlinear saturation of the MTI in plane-parallel atmospheres that are unstable according to the Balbus criterion.  The linear growth rates measured from the simulations are in excellent agreement with the weak field dispersion relation.  The addition of isotropic conduction, e.g. radiation, or strong magnetic fields can damp the growth of the MTI and affect the nonlinear regime. The instability saturates when the atmosphere becomes isothermal as the source of free energy is exhausted.  By maintaining a fixed temperature difference between the top and bottom boundaries of the simulation domain, sustained convective turbulence can be driven.  MTI-stable layers introduced by isotropic conduction are used to prevent the formation of unresolved, thermal boundary layers.  We find that the largest component of the time-averaged heat flux is due to advective motions as opposed to the actual thermal conduction itself. Finally, we explore the implications of this instability for a variety of astrophysical systems, such as neutron stars, the hot intracluster medium of galaxy clusters, and the structure of radiatively inefficient accretion flows. ",
        "introduction": "\\label{introduction} In many dilute, magnetized astrophysical plasmas, the electron mean free path between collisions can be many orders of magnitude larger than the ion gyroradius.  In this regime, the equations of ideal magnetohydrodynamics (MHD) that describe the fluid plasma must be supplemented with anisotropic transport terms for energy and momentum due to the near free-streaming motions of particles along magnetic field lines \\citep{brag65}.  Thermal conduction is dominated by the electrons compared to ions by a factor of the square root of the mass ratio.  In a non-rotating system, it is sufficient to neglect the ion viscosity \\citep{bal04}.  A fully collisionless treatment may be done using more complex closures such as \\citet{hp90}.  The implications of anisotropic transport terms on the overall dynamics of dilute astrophysical plasmas is only beginning to be explored \\citep{bal01, qdh02, sha03}.  One of the most remarkable results obtained thus far is that the convective stability criterion for a weakly magnetized dilute plasma in which anisotropic electron heat conduction occurs is drastically modified from the usual Schwarzschild criteria \\citep{bal00}. In particular, stratified atmospheres are unstable if they contain a {\\em temperature} (as opposed to {\\em entropy}) profile which is decreasing upward.  There are intriguing analogies between the stability properties of rotationally supported flows (where a weak magnetic field changes the stability criterion from a gradient of specific entropy to a gradient of angular velocity), and the convective stability of stratified atmospheres (where a weak magnetic field changes the stability criterion from a gradient of entropy to a gradient of temperature).  The former is a result of the magnetorotational instability \\citep[MRI;][]{bh98}.  The latter is a result of anisotropic heat conduction.  To emphasize the analogy, we will refer to  this new form of convective instability as the magnetothermal instability (MTI). The MTI may have profound implications for the structure and dynamics of many astrophysical systems.  In this paper, we use numerical methods to explore the nonlinear evolution and saturation of the MTI in two-dimensions.  We adopt an arbitrary vertical profile for a stratified atmosphere in which the entropy increases upward (and therefore is stable according to the Schwarzschild criterion), but in which the temperature is decreasing upwards (and therefore is unstable according to the Balbus criterion, $(\\partial P/\\partial z)(\\partial \\ln T / \\partial z) < 0$). We confirm the linear growth rates predicted by \\citet{bal00} for dynamically weak magnetic fields and numerically measure the growth rates for stronger fields.  We find that in the nonlinear regime vigorous convective turbulence results in efficient heat transport.  Full details are published in \\citet{ps05}, hereafter PS.  These results may have implications for stratified atmospheres where anisotropic transport may be present.  The most exciting potential application is to the hot X--ray emitting gas in the intracluster medium of galaxy clusters \\citep{pf05, mark98}.  The hot plasmas are magnetized and have very long mean free paths along the magnetic field lines; thus, they are a prime candidate for this instability. For example, the Hydra A cluster has $T \\approx 4.5$ keV and a density of $n \\approx 10^{-3} - 10^{-4} \\mathrm{cm}^{-3}$ giving a mean free path that's almost one-tenth of the virial radius. Other applications are to the atmospheres of neutron stars with moderate magnetic fields and radiatively inefficient accretion flows. ",
        "conclusions": "The most important conclusion of this work is that atmospheres with $dS/dZ < 0$ are not necessarily stable to convection.  In fact, dilute atmospheres with weak to moderate magnetic fields can be convectively unstable by the Balbus criterion resulting in an instability that we call the magnetothermal instability.  We have verified using MHD simulations that the measured linear growth rates agree with analytic WKB theory as predicted by \\citet{bal00}.  For adiabatic boundary conditions, we find the saturated state is an isothermal temperature profile, corresponding to the exhaustion of the free energy in the system.  For a driven instability with conducting boundary conditions, we find that the MTI efficiently transports heat, primarily by advective motions of the plasma in the vigorous convection that results.  The most promising application of this instability is to clusters of galaxies.  Structure formation calculations assuming $\\mathrm{\\Lambda}$CDM cosmologies predict monotonically decreasing temperature profiles of the intracluster gas \\citep{lok02}.  Observations of clusters with Chandra, such as Hydra A \\citep{dm02}, however, indicate essentially flat temperature profiles.  The intracluster medium is dilute, magnetized, and has a mean free path that could be as high as one-tenth of the cluster virial radius.  It may be that these temperature profiles are representative of the saturated state of the MTI.  This possibility will be explored in future work."
    },
    "0609/astro-ph0609504_arXiv.txt": {
        "abstract": "We  discussed prospects for directly detecting  circular polarization signal of gravitational wave background. We  found it is generally difficult to probe the  monopole mode of the signal due to broad directivity of gravitational wave detectors. But the dipole  $(l=1)$ and octupole $(l=3)$ modes of the signal can be measured in a  simple manner by combining outputs of two unaligned detectors, and we can dig them deeply under confusion  and detector noises.   Around $f\\sim 0.1$mHz LISA will provide ideal data streams to detect these anisotropic components whose magnitudes are as small as  $\\sim1$ percent of  the detector  noise level  in terms  of the non-dimensional  energy density $\\Omega_{GW}(f)$. ",
        "introduction": "As gravitational interaction is very weak,  significant efforts have been made to detect gravitational waves. But, on the other hand, we will be able to get rich information of the universe by observing gravitational waves that directly propagate to us with almost no absorption. Various astrophysical and cosmological models predict existence of stochastic gravitational wave background, and it is an interesting target for  gravitational wave astronomy \\cite{Allen:1996vm}. For its observational prospects, we need to understand how we characterize the background and what aspects we can uncover  with current  and future observational facilities.  One of such aspects is circular polarization that describes whether the background has asymmetry with respect to magnitudes of right- and left-handed waves. Circular polarization of gravitational wave background might be generated by helical turbulent motions (see {\\it e.g.} \\cite{Kahniashvili:2005qi} and references therein). Inflation scenario predicts gravitational wave background from quantum fluctuations during acceleration phase in the early universe, but asymmetry of left and right-handed waves can be produced with the gravitational Chern-Simon term that might be derived from string theory and might be related to creation of baryon number (see {\\it e.g.} \\cite{Lyth:2005jf} and references therein). Primordial gravitational wave background, including information of its circular polarization \\cite{Lue:1998mq},  can be indirectly studied with CMB measurement \\cite{Seljak:1996gy} at very low frequency regime $f\\sim 10^{-17}$Hz that is largely different from the regime directly accessible with gravitational wave detectors studied in this Letter. Gravitational wave background from Galactic binaries can be polarized, if orientation of their angular momentum have coherent distribution, such as, correlation with Galactic structure.  While  observational samples of local Galactic binaries do not favor such  correlation \\cite{inc}, this will also be an interesting target that can be directly studied with  LISA.  It is well known that observation of gravitational wave is intrinsically sensitive to its polarization state \\cite{Thorne_K:1987}. This is because we measure spatial expansion and contraction due to the wave, and polarization determines the direction of the oscillation perpendicular to its propagation direction. Another important nature of the observation is that we have to simultaneously deal with waves basically coming from all the directions, in contrast to observations with typical electromagnetic wave telescopes that have sharp directivity. Therefore, polarization information and directional information couple strongly in observational analysis of gravitational wave background. In general, orientation of a gravitational wave detector changes with time, and induced  modulation of data stream is useful to probe the polarization and directional information. For ground  based detectors, such as, LIGO, this is due to the daily rotation of the Earth \\cite{Allen:1996gp}. For a space mission like LISA, this change is determined by its orbital choice \\cite{lisa,Cornish:2001qi}. In addition,  we can also expect that several independent data streams of gravitational waves will be taken at the same time \\cite{Cutler:1997ta,Prince:2002hp}. In this Letter, in view of these observational characters,  we study how well we can extract information  of circular polarization of the background in a model independent manner about its origin. ",
        "conclusions": ""
    },
    "0609/astro-ph0609674_arXiv.txt": {
        "abstract": "We have calculated the H$\\alpha$ and Ca {\\sc ii} 8542 {\\AA} line profiles based on four different atmospheric models, including the effects of nonthermal electron beams with various energy fluxes. These two lines have different responses to thermal and nonthermal effects, and can be used to diagnose the thermal and nonthermal heating processes. We apply our method to an X-class flare that occurred on 2001 October 19.  We are able to identify quantitatively the heating effects during the flare eruption. We find that the nonthermal effects at the outer edge of the flare ribbon are more notable than that at the inner edge, while the temperature at the inner edge seems higher. On the other hand, the results show that nonthermal effects increase rapidly in the rise phase and decrease quickly in the decay phase, but the atmospheric temperature can still keep relatively high for some time after getting to its maximum. For the two kernels that we analyze, the maximum energy fluxes of the electron beams are $\\sim$ 10$^{10}$ and 10$^{11}$ ergs cm$^{-2}$ s$^{-1}$, respectively. However, the atmospheric temperatures are not so high, i.e., lower than or slightly higher than that of the weak flare model F1 at the two kernels. We discuss the implications of the results for two-ribbon flare models. ",
        "introduction": "Solar flares are one of the most significant active phenomena in the solar atmosphere. At present the widely accepted flare model based on magnetic reconnection is the CSHKP model, which was developed by \\citet{car64}, \\citet{Stu66}, \\citet{hira74} and \\citet{kopp76}. The model assumes that pre-existing closed magnetic loops in the corona are torn open by the force of the filament eruption as a result of magnetic instability. Subsequently, a current sheet is stretched out and energy stored in the magnetic field is released at the reconnection point. As a consequence, a cusp-shaped loop structure appears which was observed by {\\it Yohkoh} in soft X-rays  \\citep{tsuneta92} and by $Reuven$ $Ramaty$ $High$ $Energy$ $Solar$ $Spectroscopic$ $Imager$ $(RHESSI)$ in hard X-rays (HXRs) \\citep{sui03}. As successive magnetic reconnection goes on, the loop top HXR source \\citep{masuda94} rises \\citep{sui04,liu04} and the footpoint sources  separate from each other \\citep{fletcher01,liu04}. Recent observations reveal a more HXR source above the flare loop, called the coronal source \\citep{sui04}. The temperature in between the loop top and the coronal source is found to be higher than that in  lower and higher altitudes; therefore, this place is regarded as associated with the formation and development of a current sheet.   It is known that the primary energy release during a solar flare (the current sheet) is in the corona. The released energy results in a bulk heating of the plasma and acceleration of charged particles. Therefore, the chromosphere can be heated by either a conduction front or electron beam bombardment \\citep{brown73,can74,brown78,ems78}. Other energy transport processes, such as heating by an energetic proton beam \\citep{lin76,ems83,hen93} and soft X-ray irradiation \\citep{hen77,mac78,gan90,ber04} have also been invoked. Many studies have been devoted to the spatial distribution and temporal evolution of the thermal/nonthermal heating signatures. \\citet{can93} argued that energetic electrons favor to occur at the edge of high vertical currents. This is confirmed by \\citet{masuda01} that the spectrum tends to be harder at the outer edge of a ribbon than in the inner edge. \\citet{czay99} found that strong upflows, revealed by the blue shifts in EUV lines, appear at the outer edges of flare ribbons. \\citet{li04} showed that chromospheric downflows, revealed by the red asymmetries in  the H$\\alpha$ lines, are also the most obvious at the outer edges. These findings are basically consistent with the two-ribbon flare model, in which the outer edges map the footpoints of newly reconnected flare loops. The general scenario of flare evolution implies that the thermal/nonthermal heating processes can vary both spatially and temporally. Therefore, it is interesting to devise a method to diagnose the relative importances of thermal and nonthermal heating processes in a flare.  The purpose of this paper is to use two different chromospheric lines, H$\\alpha$ and Ca {\\sc ii} 8542 {\\AA} lines, to diagnose the heating processes in flares. The H$\\alpha$ line is the most observed and studied chromospheric line in solar flare spectroscopy. Theoretical calculations have shown how the H$\\alpha$ line varies with different flare parameters like the nonthermal electron flux, conduction flux, and coronal pressure \\citep{ric83,can84}. In particular, when considering the nonthermal excitation and ionization effects by electron beam, the H$\\alpha$ line can be enhanced significantly \\citep{fang93,kav02} . On the other hand, the Ca {\\sc ii} 8542 {\\AA} line is less sensitive to the nonthermal effects, though also enhanced to some extent. This line is more influenced by the coronal pressure and chromospheric temperature. The different responses of the H$\\alpha$ and Ca {\\sc ii} 8542 {\\AA} lines to thermal/nonthermal effects make it possible to diagnose the heating processes using these two lines together.  We make non-LTE calculations of the H$\\alpha$ and Ca {\\sc ii} 8542 {\\AA} line intensities for different atmospheric models and nonthermal electron beams. The results are used to diagnose the processes of a flare on 2001 October 19 in detail. The paper is organized as follows. The method for model calculations is given in \\S2. Theoretical results are presented in \\S3. \\S4 shows the diagnostics of the flare, followed by discussions and conclusions in \\S5. ",
        "conclusions": "In general, the heating mechanisms during a solar flare are complicated, i.e., a combination of different processes. Thermal conduction and nonthermal  electron beam bombardment are the two mostly discussed and important energy transport processes. Many studies are focused on solving this controversial and undetermined problem. \\citet{Saint05} deduced the thermal and nonthermal energies for a flare and found that they are of the same magnitude. However, in some other flares or in a specific phase or specific region of flares, it is possible that either thermal heating or nonthermal heating is dominant (e.g. \\citet{li05}; \\citet{ji04}). If seen  in chromospheric lines, the thermal component in the light curve is thought to be delayed by some time compared to the nonthermal component. This time delay is case-dependent.  Usually we cannot easily separate the thermal and nonthermal components from the light curve only. The detailed spectra provide a diagnostic tool to distinguish these two components.  Previous studies have shown that the nonthermal excitation and ionization effects caused by an electron beam bombardment have a great influence on chromospheric line profiles \\citep{fang93}. Since the H$\\alpha$ and Ca {\\sc ii} 8542 {\\AA} lines have different responses to thermal and nonthermal processes as discussed above,  we can diagnose different heating mechanisms using the observations of the two lines. In this work, we simply use the equivalent widths of the lines, which refrains us from invoking some unclear parameters like the macro-turbulent velocity. Of course, the HXR emission is closely related to nonthermal electrons; however, the thermal heating effects can better be studied through checking chromospheric lines.  In our analysis, we choose two lines, one of which is sensitive to thermal effects while another is sensitive to nonthermal effects, to diagnose these two heating processes. This method is tested to be practical and useful. Recently, \\citet{kar04} argued that the return current collision and excitation could significantly enhance the line radiation. If considering this effect, the deduced flux of the nonthermal electron beam should be somewhat lower. However, this does not influence our main conclusions.  The spatial distribution and temporal evolution of thermal and nonthermal effects during a solar flare are also complex. \\citet{li04} concluded that the  chromospheric downflow velocity tends to appear at the outer edge of  flare ribbons. This indicates that the nonthermal effects at the outer edge of flare ribbons are  the most significant. In the current work, by using a novel and different method, we also find that in the 2001 October 19 flare, the nonthermal effects at the outer edge are more distinct than that at the inner edge of flare ribbons. However, the chromospheric temperature at the inner edge seems higher than that at the outer part. The above results support the general scenario of flare development: successive magnetic reconnection occurs to form new flare loops whose footpoints are shown as the outer edge of flare ribbons. It is conceivable that in the newly formed loops, the nonthermal effects are more obvious than in the old loops. At the two kernels K2 and K3 in the 2001 October 19 flare, the electron beam fluxes increase quickly in the initial phase and decrease rapidly in the gradual phase while the thermal effects change gradually."
    },
    "0609/astro-ph0609452.txt": {
        "abstract": "We present high spatial resolution integral-field spectroscopy of 28 elliptical (E) and lenticular (S0) galaxies from the \\sauron\\/ representative survey obtained with the \\oasis\\/ spectrograph during its operation at the CFHT. These seeing-limited observations explore the central 8\\arcsec$\\times$10\\arcsec\\/ (typically one kiloparsec diameter) regions of these galaxies using a spatial sampling four times higher than \\sauron\\/ (0\\farcsec27 vs. 0\\farcsec94 spatial elements), resulting in almost a factor of two improvement in the median PSF. These data allow accurate study of the central regions to complement the large-scale view provided by \\sauron. Here we present the stellar and gas kinematics, stellar absorption-line strengths and nebular emission-line strengths for this sample. We also characterise the stellar velocity maps using the `kinemetry' technique, and derive maps of the luminosity-weighted stellar age, metallicity and abundance ratio via stellar population models. We give a brief review of the structures found in our maps, linking also to larger-scale structures measured with \\sauron. We present two previously unreported kinematically-decoupled components (KDCs) in the centres of NGC\\,3032 and NGC\\,4382. We compare the intrinsic size and luminosity-weighted stellar age of all the visible KDCs in the full \\sauron\\/ sample, and find two types of components: kiloparsec-scale KDCs, which are older than 8~Gyr, and are found in galaxies with little net rotation; and compact KDCs, which have intrinsic diameters of less than a few hundred parsec, show a range of stellar ages from 0.5 - 15~Gyr (with 5/6 younger than 5~Gyr), are found exclusively in fast-rotating galaxies, and are close to counter-rotating around the same axis as their host. Of the 7 galaxies in the \\sauron\\/ sample with integrated luminosity-weighted ages less than 5~Gyr, 5 show such compact KDCs, suggesting a link between counter-rotation and recent star-formation. We show that this may be due to a combination of small sample size at young ages, and an observational bias, since young KDCs are easier to detect than their older and/or co-rotating counterparts.\\looseness-2 ",
        "introduction": "In the era of the {\\it Hubble Space Telescope} ({\\it HST}), early-type galaxy nuclei (central 10-50~pc) have become a key component in the modern picture of galaxy evolution. Empirical relationships between the central luminosity profile and numerous large-scale galaxy properties, such as isophotal shape, rotational support, and stellar populations, strongly suggest that the formation and evolution of the central regions is inextricably linked to that of the galaxy as a whole \\citep[e.g.][]{faber97}. The ubiquitous presence of super-massive black holes (SMBHs) at the centres of galaxies \\citep[e.g.][]{ferrarese00,gebhardt00} is thought to be an important aspect of this relationship, controlling global star-formation and dictating the central density distribution \\citep[e.g.][]{silk98,granato04,binney04,dimatteo05}.  Most observational studies of early-type galaxy nuclei have so far been based only on {\\it HST} imaging, with a small number of objects also observed spectroscopically from space, either with a single aperture or a long slit. The limitations of these data become apparent when considering the complexity of structures found in the centres of these otherwise `boring' objects. Spectroscopy is required to establish dynamical black hole mass estimates, and it has been shown that a single long slit is grossly insufficient to constrain fully general dynamical models \\citep{verolme02,cappellari05}. To understand the stellar orbital structure of galaxy nuclei therefore requires information across a sizeable two-dimensional field, obtained with the highest possible spatial resolution. Additionally, {\\it HST} imaging reveals highly complex gas and dust distributions \\citep[e.g.][]{vdokkum95, tomita00,lauer05}, which are sensitive indicators of dynamical torques that are not detectable from the stellar distribution alone. Understanding the kinematic and chemical properties of such irregular gas distributions requires two-dimensional spectroscopy. {\\it HST} imaging studies also find broad-band colour gradients, indicative of distinct star-formation histories towards the very central regions \\citep[e.g.][]{carollo94b,carollo97a,lauer05,cote06}. However, to accurately quantify the populations requires spectroscopy on comparable spatial scales, of which there are rather few studies of early-type galaxy samples, due to the high signal-to-noise ratio required \\citep[e.g.][]{krajnovic04,sarzi05}. Moreover, these and comparable ground-based studies looking at central properties \\citep[e.g.][]{carollo94b,trager2000,thomas05} are restricted to either aperture measurements or long-slit profiles, making the association between population changes and other structural variations difficult to determine.  In response to these issues, we have initiated a program of high spatial resolution integral-field spectroscopy of early-type galaxy nuclei. We select our sample from galaxies observed as part of the \\sauron\\/ survey: a study of 72 representative nearby early-type galaxies and spiral bulges observed with \\sauron, our custom-built panoramic integral-field spectrograph mounted at the William Herschel Telescope (WHT) on La Palma \\citep[Paper I]{bacon01}. The aims of this survey and the observed galaxy sample are described in \\citet[Paper II]{dezeeuw02}. For the sub-sample of 48 E/S0s we present the stellar kinematic maps in \\citet[Paper III]{emsellem04}; the two-dimensional ionised gas properties in \\citet[Paper V]{sarzi06}; and maps of the stellar absorption line strengths in \\citet[Paper VI]{kuntschner06}. We refer the reader to these papers (and references therein) for discussion of the global properties of our sample galaxies.  To maximise the field of view, the spatial sampling of \\sauron\\/ was set to 0\\farcsec94$\\times$0\\farcsec94 (per lenslet) for the survey, therefore often undersampling the typical seeing at La Palma (0\\farcsec8). This does not affect the results for the main body of early-type galaxies, as they generally exhibit smoothly-varying structure on scales larger than the seeing. However, the central regions are poorly resolved in the \\sauron\\/ data, preventing accurate measurements of the central properties.  For this reason we have observed the centres of a subset of the \\sauron\\/ sample using the \\oasis\\/ integral-field spectrograph, mounted on the 3.6-m Canada-France-Hawaii Telescope (CFHT), Hawaii. This instrument allows optimal sampling of the natural seeing on Mauna Kea, resulting in a significant improvement in spatial resolution. To ensure a high degree of complementarity between the \\oasis\\/ data presented here and our \\sauron\\/ observations, we configured \\oasis\\/ to have an almost identical wavelength domain and spectral resolution. The \\oasis\\/ data can therefore be easily combined with \\sauron\\/ data, e.g. for constraining dynamical models \\citep[e.g.][]{shapiro06}. All data presented in this paper will be made available through the \\sauron\\/ web page.{\\footnote{http://www.strw.leidenuniv.nl/sauron}\\looseness=-2  The \\oasis\\/ observations themselves reveal a number of previously unresolved features in the centres of our sample galaxies. In particular, we find several galaxies which exhibit very young central stellar populations that can be directly associated with peculiar (misaligned) components in the stellar rotation fields. We explore this further by considering all objects in the \\sauron\\/ sample with clear kinematically decoupled components, and find a link between the size of these components and their stellar populations.  The layout of the paper is as follows: Section~\\ref{sec:obs} describes the instrument, the observations and the basic data reduction; Section~\\ref{sec:analysis} describes the application of our analysis tools used to derive the stellar kinematics, gas properties and line-strength measurements (including the application of stellar population models), and presents detailed comparisons with equivalent quantities from \\sauron; Section~\\ref{sec:results} presents the two-dimensional maps of all measured quantities, and discusses their main features; Section~\\ref{sec:discussion} presents a discussion of young stellar populations in our sample, and the connection with KDCs; and Section \\ref{sec:conclusions} concludes.  %% %%=================================================================== %%                  OBSERVATIONS AND DATA REDUCTION %%=================================================================== %% ",
        "conclusions": "\\label{sec:conclusions}  We have presented high spatial resolution integral field spectroscopy, using the \\oasis\\/ spectrograph, of a subset of 28 elliptical and lenticular galaxy centres selected from the \\sauron\\/ survey. From these observations, we derive the distribution and kinematics of stars and gas, as well as stellar absorption-line indices independently calibrated to the Lick system. This data set has a factor four higher spatial sampling and almost a factor two increase in median spatial resolution to that of \\sauron. We compare the results of our analysis to that of \\sauron\\/ spectra within equivalent integrated aperture spectra. The measured parameters, including velocity dispersion, emission line fluxes and stellar line-strength indices, are all consistent within the uncertainties, verifying our analysis.  Within the central regions probed by \\oasis, there is a rich variety of structure in the observed parameters, comparable to that observed on larger scales with \\sauron. In particular, we find a range of stellar dynamical substructure within the central kiloparsec of these objects, including central disks and counter-rotating cores, which we quantify using the technique of kinemetry. The ionised gas is similarly varied, being often dynamically decoupled from the stellar component, and varying flux ratios indicating a range of ionisation mechanisms, even within the same galaxy. We analyse the stellar populations by comparing the observed line-strength maps to model predictions, giving maps of luminosity-weighted stellar age, metallicity and abundance ratio. These maps reveal that the galaxies with the youngest global populations in our sample also have even younger centrally concentrated populations. These young nuclei may be explained by young nuclear star clusters or other recent circum-nuclear star formation as in H {\\small II} nuclei.\\looseness-2  The young nuclei show a propensity for harbouring compact kinematically decoupled stellar components, which can be associated directly with the young stars. Considering all galaxies in the \\sauron\\/ survey that show a clear kinematically decoupled stellar component, we explore the relationship between central stellar population and the size of the component. We find that, although the `classical' decoupled components show rather old stellar populations, there also exists a population of compact ($\\lsim 100$~pc) decoupled components which show a range of ages, from 15~Gyr to $\\lsim 0.5$~Gyr. Given the large difference in intrinsic size, we conclude that the two types of decoupled component are not directly related in the evolutionary sense, and that the compact variety are not necessarily an important feature {\\em by mass}, although the young stars within these decoupled components are largely responsible for the significant spread in integrated luminosity-weighted ages observed for these galaxies.  Our data suggest a link between counter-rotation and enhanced star-formation. By consideration of a simple model, we conclude that older KDCs of similar mass fractions are rather difficult to detect, and so there may be a sizeable number of low-mass counter-rotating components that we are missing. Likewise, co-rotating components of the same mass fraction leave a more subtle signature on the observed stellar kinematics, and so can only be detected when the KDC population is rather young. Our sample is therefore too small to reliably put constraints on the importance of counter-rotation to recent star formation.  These are the first observations in an ongoing follow-up of \\sauron\\/ survey galaxies using high-spatial resolution integral field spectroscopy to study the central regions of these objects in detail. The empirical results presented in this paper are a first step towards understanding the intrinsic structure of the central few hundred parsecs of these objects, and already reveal new insights. Future application of dynamical models, using both \\sauron\\/ and \\oasis\\/ data will substantiate this further.  %% %%==================================================================== %%                          ACKNOWLEDGEMENTS %%==================================================================== %%"
    },
    "0609/astro-ph0609168_arXiv.txt": {
        "abstract": "I review here some of the main observational features in microquasars and AGNs as sources of high-energy $\\gamma$-rays as well as some of the current models that try to explain the emission in the $\\gamma$-ray domain. ",
        "introduction": "Microquasars are X-ray binary stars  which exhibit relativistic jets, i.e., bipolar outflows of relativistic matter ejected perpendicularly to both sides of an accretion disc. The jets contain relativistic electrons that produce the synchrotron radiation detected at radio wavelength. These systems mimic many phenomena seen in quasars at smaller scales (Mirabel \\& Rodr\\'{\\i}guez 1999).  The recent detection of TeV $\\gamma$-ray emission coming from the microquasars LS 5039 and LS~I~+61~303 (Aharonian et al. 2005a; Albert et al. 2006a), using the ground-based Cherenkov telescopes HESS and MAGIC, is a remarkable result that consolidate microquasars as a new group of very high-energy (VHE) $\\gamma$-ray emitters. Moreover, this fact strengthes the analogy quasar/microquasar as it is well known that quasars are also sources of strong $\\gamma$-ray emission (Hartman et al. 1999). ",
        "conclusions": "Up to now there are 66 EGRET blazars and 11 TeV blazars detected. On the other hand, microquasars have consolidated as a new source population of the $\\gamma$-ray sky. As emitters of HE and VHE photons, they also mimic the high-energy behaviour of quasars. This fact puts microquasars among the most interesting sources in the Galaxy from the viewpoint of high-energy astrophysics.  Several kind of models predict that radio jets could be natural sites for the production of high-energy photons via both Compton scattering and proton proton collisions in microquasars.  Additional microquasars will very likely be detected soon with the Cherenkov telescopes and GLAST. This will bring more constraints to the physics of these systems and their relationship with AGNs."
    },
    "0609/astro-ph0609442_arXiv.txt": {
        "abstract": "The interaction of a fast wind with a spherical Asymptotic Giant Branch (AGB) wind is thought to be the basic mechanism for shaping Pre-Planetary Nebulae (PPN) and later Planetary Nebulae (PN). Due to the large speed of the fast wind, one expects extended X-ray emission from these objects, but X-ray emission has only been detected in a small fraction of PNs and only in one PPN. Using numerical simulations we investigate the constraints that can be set on the physical properties of the fast wind (speed, mass-flux, opening angle) in order to produce the observed X-ray emission properties of PPNs and PNs. We combine numerical hydrodynamical simulations including radiative cooling using the code {\\em FLASH} with calculations of the X-ray properties of the resulting expanding hot bubble using the atomic database {\\em ATOMDB}. In this first study, we compute X-ray fluxes and spectra using one-dimensional models. Comparing our results with analytical solutions, we find some agreements and many disagreements. In particular, we test the effect of different time histories of the fast wind on the X-ray emission and find that it is determined by the final stage of the time history during which the fast wind velocity has its largest value. The disagreements which are both qualitative and quantitative in nature argue for the necessity of using numerical simulations for understanding the X-ray properties of PNs. The X-ray luminosity in the 0.2 -- 10 keV range (1.24 -- 62 \\AA) covered by the CHANDRA/{\\em ACIS} instrument shows a very slow decrease with time over the range of evolutionary ages explored in our models (up to 3750 years). Furthermore, investigating the emission in other wavelengths ranges, we find that most of the luminosity emerges at longer wavelengths ($\\lambda \\gtrsim 140$ \\AA) from the cooler outer edge of the hot bubble. We apply our spherical models to the objects BD$+$30$^\\circ$3639 and NGC~40. We find that the model values of the X-ray temperature and luminosity for these objects are significantly higher than observed values and discuss several mechanisms for resolving the discrepancies. ",
        "introduction": "The shaping of Pre-planetary nebulae (PPN) and Planetary Nebulae (PN) is believed to result from the interaction of a fast, collimated post-AGB wind plowing into the slow, dense wind emitted during the AGB phase, followed by an isotropic tenuous wind during the PN phase \\citep{SaT98}. These interacting winds result in an expanding shell of shocked matter which forms the PPN and later the PN. Due to the large speed of the fast wind (from few x 100 km s$^{-1}$ up to 2000 km s$^{-1}$), one expects extended X-ray emission to be produced in PPNs (where the shaping mainly occurs) and PNs.  However, X-ray emission was only detected in 3 of 60 PNs observed with ROSAT, followed by a few more from CHANDRA and XMM \\citep[e.g.][]{GCG05}. In the case of PPNs, there is so far only one confirmed X-ray detection \\citep{SKF03}, although many have been observed with CHANDRA. X-ray emission should be present in principle in all of these systems -- hence the low detection rate implies that it is obviously below the detection limit.  The general problem of understanding the formation and shaping of PNs has been addressed analytically and numerically \\citep[see][and references therein]{BaF02}. In principle, one of the most direct probes of the fast wind and the interaction process which drives PN formation is the X-ray emission. It may help in answering several important questions -- e.g. what are the physical properties of the fast wind (speed, mass-flux, opening angle -- and the time histories of these parameters). Such studies can also help in identifying the nebular regions that are mainly responsible for the X-ray emission for comparison with recent X-ray images of PNs obtained with CHANDRA and XMM \\citep[e.g.][]{KSV00,KBB03,GCG05}.  X-ray emission as a probe was first exploited by \\citet{ZhP96}, who derived an analytical, self-similar, spherically symmetric model, taking into account the effects of thermal conductivity. Very recently, \\citet[][hereafter ASB06]{ASB06} presented new analytical models with radiative cooling, but without heat conduction. Although analytical modeling is potentially a powerful tool for understanding of the X-ray emission from PNs and its dependence on fundamental physical parameters, it necessitates making certain assumptions which may or may not be completely valid. For example, a major problem with such analytical models is, that they are not able to treat all the effects (hydrodynamics, radiative cooling, heat conduction) self-consistently. In this paper, we have therefore performed numerical simulations using the code {\\em FLASH} in which we vary the basic parameters of the fast and slow wind over an extensive parameter grid and compute the X-ray emission as a function of these parameters. A major goal of our work here is to test the main results from ASB06's analytical modeling and verify their general conclusions. We have therefore, as in ASB06, not included heat conduction in our treatment. However, we discuss the heat conduction process and how it may be inhibited by magnetic fields which are likely to be present in PNs. For example, the lack of spherical symmetry in most PNs has been explained by some authors as resulting from the presence of magnetic fields \\citep[e.g.][]{GLF05}. This paper, in which we have carried out spherically symmetric simulations of interacting winds, is a first in a series of papers in our quest to understand X-ray emission from PPNs and PNs.  The rest of the paper is organized as follows. In \\S \\ref{sec_model}, we describe our numerical model. In \\S \\ref{sec_res_1D}, we present the results of our one-dimensional simulations including the structure and X-ray properties of the expanding bubble and their dependency on the parameters of the fast wind. In \\S \\ref{sec_res_2D}, new effects from the computation of a spherically symmetrical model on a two-dimensional grid are shown. In \\S \\ref{sec_appl}, we apply our models to two PNs, BD$+$30$^\\circ$3639 and NGC~40, and examine several mechanisms for resolving the discrepancies between the models and observed values of the X-ray temperature $T_{\\rm x}$ and luminosity $L_{\\rm x}$ for these objects. We compare in detail our results with those of ASB06 in \\S \\ref{sec_diff} and discuss the still unsolved problem of fitting $T_{\\rm x}$ and $L_{\\rm x}$ in PNs in \\S \\ref{sec_fit}. Finally in \\S \\ref{sec_concl}, we present our conclusions. ",
        "conclusions": "\\label{sec_concl}  We computed the structure and evolution of, and X-ray emission from, numerical simulations of one-dimensional interacting winds models and compared the results with different analytical models. The kinematics and the evolution of our bubbles and the structure of the flows agrees well with analytical results. Comparing the X-ray properties of the hot bubbles in our models with the analytical results of ASB06, we find some agreements and many disagreements. The disagreements which are both qualitative and quantitative in nature argue for the necessity of using numerical simulations for understanding the X-ray properties of PNs.  We further investigated the luminosities in the energy ranges other than the 0.2 -- 10 keV range covered by {\\em ACIS}. We find that most of the X-ray flux may emerge at energies less than 0.09 keV (wavelengths above 137 \\AA), although we do not have reliable quantitative estimates of this flux from our models. At these wavelengths, commonly labeled as extreme and far UV, observations of PNs with FUSE (900 -- 1200 \\AA), EUVE (70 -- 760 \\AA), IUE (1150 -- 3200 \\AA) and GALEX for even longer wavelengths (1350 -- 2800 \\AA) would therefore be quite important. Unfortunately, photospheric emission from the hot central star easily blends with the emission from the PN. Instruments with high spatial resolution are needed.  We also investigated the X-ray spectra from our models, which are not addressed by the analytical studies. We showed that the shapes of the spectra are diagnostic of the fast wind velocity which determines the temperature of the hot bubble.  We applied our spherical model to the objects BD$+$30$^\\circ$3639 and NGC~40. We find in both cases that the simulated temperatures and luminosities of the hot bubble are much higher compared to the observations. While investigating these discrepancies, we identified several general issues which have to be considered for correctly modeling the X-ray emission of PNs.  The first one is the importance of an independent knowledge of the abundances -- especially in objects with [WC]-type central stars, where the abundances of the hot bubble (shocked fast wind) can be dramatically different from those in the dense shell (shocked slow wind). We could roughly reproduce the observed luminosity in BD$+$30$^\\circ$3639 by simply (i.e. in a non-self-consistent manner) removing several heavy elements including iron in the calculations of the X-ray properties. The treatment of cooling, however, has to be done self-consistently, as the cooling function itself is highly sensitive to abundances (e.g. iron).  A second issue is the possible dependence of $L_{\\rm x}$ and $T_{\\rm x}$ on the time history of the fast wind velocity. Noting that $T_{\\rm x}$ is most sensitive to $v_{\\rm f}$, we investigated whether $L_{\\rm x}$ and $T_{\\rm x}$ could be lowered by assuming that the value of $v_{\\rm f}$ was lower in the past, but found that such a time history will not help in explaining the observed low values of $L_{\\rm x}$ and $T_{\\rm x}$. The X-ray emission is determined by the final stage of the time history during which the fast wind velocity has its largest value. Although we have calculated only one case of wind evolution with a particular jump in the fast wind properties, the claim of ASB06, that the X-ray emission of the hot bubble is determined by the moderate fast wind and that the subsequent time history is unimportant, is not supported by our results.  We discussed the importance of heat conduction on lowering the temperature of the hot bubble, and its possible inhibition by magnetic fields. As the presence of magnetic fields is supported both by observations and theoretical considerations, heat conduction is likely inhibited. Even fields of the order of $\\mu G$ are sufficient to inhibit heat conduction across (and along) the field lines. Since the X-ray spectra strongly indicate a very large overabundance of carbon in the hot bubble (as expected in the fast wind from the [WC] central star of this object, but much higher than the carbon abundance of the dense shell), if heat conduction is operating in this object, it can not have evaporated a large amount of mass from the dense shell into the hot bubble. Alternatively, heat conduction is inhibited by magnetic fields in this object.  We also calculated a two-dimensional spherical model to examine the presence of instabilities and how they affect the structure of the hot bubble and the dense shell and the resulting X-ray properties. Due to these instabilities, some fraction of the total energy is redirected into the kinetic energy of non-radial motion, which leads to slower expansion velocities of the shell (by about 20 \\%) and a reduced luminosity between 0.2 -- 10 keV (by $\\sim 1.6$), although the temperature of the hot bubble does not change significantly.  From our detailed modeling, we conclude that the problem of hydrodynamical models for the formation of PNs producing values of $T_{\\rm x}$ and $L_{\\rm x}$ which are too high compared to observed values remains unsolved. The simulations in our study are the first step towards modeling PNs and PPNs detected in X-rays which show bipolar or asymmetric structures with varying degrees of collimation. Our preliminary results show that higher collimation leads to lower X-ray luminosities \\citep{StS06}. Detailed results will be presented in a future publication."
    },
    "0609/nucl-th0609024_arXiv.txt": {
        "abstract": "\\baselineskip 18.0pt The role of the isoscalar hyperon $\\Ld$ in probing the density dependence of the nuclear symmetry energy is studied in  multi-$\\Ld$ hypernuclei, hyperon-rich matter and neutron stars in relativistic  models. Relationships between the properties of three types of objects and the neutron thickness in $^{208}$Pb are established  with respect to the isoscalar-isovector coupling that modifies the density dependence of the symmetry energy. The exotic isotopes far from the neutron drip line can be stabilized by filling in  considerable $\\Ld$ hyperons.  The difference of the binding energy of multi-$\\Ld$ hypernuclei from different models is attributed to different symmetry energies.  The isovector potential together with the neutron thickness in  multi-$\\Ld$ hypernuclei investigated is very sensitive to the isoscalar-isovector coupling.  The large sensitivity of the $\\Ld$ hyperon fraction to the isoscalar-isovector coupling occurs at about  $2\\sim3\\r_0$ in beta equilibrated hyperon-rich matter. In neutron stars with hyperonization, an on-off effect with respect to the isoscalar-isovector coupling exists for the neutron star radius. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609660_arXiv.txt": {
        "abstract": "A sample of 66 galaxies from the catalog of Bettoni et al. (CISM) with anomalously high molecular-to-atomic hydrogen mass ratios $(M_{mol}/M_{HI}>2)$ is considered. The sample galaxies do not differ systematically from other galaxies in the catalog with the same morphological types, in terms of their photometric parameters, rotational velocities, dust contents, or the total mass of gas in comparison with galaxies of similar linear sizes and disk angular momentum. This suggests that the overabundance of $H_2$ is due to transition of $HI \\to H_2$. Galaxies with bars and active nuclei are found more frequently among galaxies which have $M_{mol}$ estimates in CISM. In a small fraction of galaxies, high $M_{mol}/M_{HI}$ ratios are caused by the overestimation of $M_{mol}$ due to a low conversion factor for the translation of {\\it CO}-line intensities into the number of $H_2$ molecules along the line of sight. It is argued that the \"molecularization\" of the bulk of the gas mass could be due 1) to the concentration of gas in the inner regions of the galactic disks, resulting to a high gas pressure and 2) to relatively low star-formation rate per unit mass of molecular gas which indeed takes place in galaxies with high $M_{mol}/M_{HI}$ ratios.   Astronomy Reports, 2006, Vol. 50, No. 8, p. 626--637. ",
        "introduction": "The densest and coolest component of the interstellar medium~--- molecular gas~--- is observed in all types of disk galaxies. In many cases, a mass of molecular gas inside galactic disks is comparable to a mass of neutral hydrogen ($HI$). The amount and spatial density of molecular gas reflects the specific properties of the evolution of the interstellar medium in galactic disks. The most massive gaseous formations~--- giant molecular clouds, which affect the dynamical evolution of the disk,--- are associated with molecular gas. However, more importantly, the mass and density of the molecular gas determine the star-formation rate and star-formation efficiency in galactic disks.  Despite the crucial role that molecular gas plays in galaxies, the factors determining its abundance relative to neutral gas remain poorly known. Most information on interstellar molecular gas has been obtained from measurements of {\\it CO}-line intensities~--- {\\it CO} has a very low excitation temperature. The conversion of the {\\it CO}-line luminosity into the total mass of gas is a separate problem. For galaxies with normal heavy-element abundances the conversion factor connecting the {\\it CO}-line intensity with the number of $H_2$ molecules along the line of sight, ${\\it X} = N(H_2)/I_{CO}$, is usually set equal to $(2-3)\\cdot 10^{20} \\quad mol/(K \\cdot km/s)$, however, the universality of this parameter remains a topic of discussion. Adopting a constant ${\\it X}$ value is, of course, a simplification, especially when galaxies with strongly differing mass, gas density, or star-formation intensity are compared. Various methods for measuring ${\\it X}$ and the dependence of ${\\it X}$ on the galactic luminosity show appreciable scatter, together with a systematic increase in ${\\it X}$ with decreasing luminosity (mass) of the galaxy and decreasing heavy-element abundance (\\cite{boselli}, \\cite{arimoto}). The conversion factors for dwarf galaxies and giant luminous galaxies may differ by an order of magnitude. We also expect ${\\it X}$ to be smaller for dense circumnuclear disks, where the molecular gas seems to be contained in both clouds and the diffuse intercloud medium (\\cite{downes}, \\cite{rosolovsky}).  Various authors (\\cite{elmegreen89}, \\cite{elmegreen93}, \\cite{wong}) have analyzed the conditions for molecular-to-atomic, $M_{mol}\\to M_{HI}$, transformations and vice-versa (see also references in these papers). The rate at which molecules are formed depends on both the density (or pressure) of a gas and on the density of {\\it UV} flux, which destroys molecules. According to Elmegreen (\\cite{elmegreen93}), the $M_{mol}/M_{HI}$ ratio is proportional to $P^{2.2}/j$, where $P$~is the gas pressure and ~$j$~is the {\\it UV} flux density. The gas compression in shocks can also play an important role in the \"molecularization\" of hydrogen~\\cite{kee}.  Since the gas pressure depends both on its surface density and on the volume density of the stellar disk, it decreases rapidly with galactocentric distance. Indeed, the azimuthally averaged density ratio $M_{mol}/M_{HI}$ follows the pressure in such a way that the contribution of $M_{mol}$ to the total mass of gas becomes small at a periphery of the galaxy~\\cite{wong}.  Available estimates show that the integrated ratio $M_{mol}/M_{HI}$ is appreciably lower than unity for most spiral galaxies. According to Casoli et al.~\\cite{casoli}, $M_{mol}/M_{HI}$ is, on the average, close to $0.2$. Boselli et al.~\\cite{boselli} found the average molecular-to-atomic gas mass ratio to be $M_{mol}/M_{HI} = 0.14$ for a sample of 266 galaxies. However, some spiral galaxies are known to have anomalously high $M_{mol}/M_{HI}$ values, exceeding two or even three (an order of magnitude higher than the average value). Our aim here is to analyze the specific features of such galaxies with a predominance of molecular gas and possible factors that could lead to such high $M_{mol}/M_{HI}$ estimates.  As our initial data, we used \"A new catalogue of ISM content of normal galaxies\" (CISM) of Bettoni et~al. \\cite{bettoni}. This catalog contains interstellar-medium data for almost two thousand \"normal\" galaxies, more than 300 of which have estimates for their molecular-gas mass. Normal galaxies are considered to be those that display no morphological distorsion (such as tidal tails or bridges) or strong perturbations of their circular velocities, although several galaxies among those included in the catalog can, nevertheless, be considered to be interacting galaxies (see below). The molecular-gas mass estimate is based on the adopted value of ${\\it X} = 2.3\\cdot 10^{20} \\quad mol/K \\cdot km/s$, taken to be the same for all galaxies. ",
        "conclusions": "Galaxies whose gas is predominantly in molecular form ($M_{mol}/M_{HI}>2$) occur among disk galaxies of all morphological types, although {\\it Sc~---Sd} types are among these objects. These galaxies differ only slightly from galaxies with lower relative mass of molecular gas in terms of their luminosities, surface brightnesses, rotational velocities, and dust contents per unit mass of gas. At the same time, galaxies that are rich in molecular gas show a tendency to have bars and active nuclei (to be LINERs), providing indirect evidence for a concentration of gas toward the nucleus.  The high $M_{mol}/M_{HI}$ ratios in the galaxy sample we have considered is not due to a loss of $HI$ or the acquisition of $H_2$, but instead is the result of the transformation of most of the $HI$ into~$H_2$, since the total gas mass, $M_{gas}$, remains normal and consistent with the specific angular momentum of the disk. However, the $M_{mol}$ values for some of the sample galaxies derived from {\\it CO}-line intensities are probably appreciably overestimated due to the use of a conversion factor that appears to be overestimated.  With only a few exceptions, galaxies with high $M_{mol}/M_{HI}$ ratios are characterized by moderate star-formation rates and comparatively high star-formation efficiencies, $SFR/M_{gas}$. However, the mean star-formation rate per unit mass of molecular gas in these galaxies is appreciably lower compared to galaxies with low $M_{mol}/M_{HI}$. This means that, on average, the gas spends more time in molecular form before it becomes atomic again.  As far as the integral properties of our sample galaxies are normal, their high content of molecular gas must be due to specific properties of the gas in local regions of the disk. The high fraction of molecular gas appears to be due to two interrelated factors. The first of these is a high density (pressure) of the gas due to its concentration in the inner region of the disk (or in spiral arms, if the galaxy possesses a high-contrast spiral structure). In this case, the local dust density should also be higher, facilitating the formation and preservation of molecules. We think that there are no galaxies with a deficit of gas in their central region among those with a predominance of molecular gas. The second possible factor is the long duration of the molecular state of the gas during the chain  $HI \\to H_2 \\to$~gravitationally unstable clouds~$\\to$~stars. Evidence for this is provided by the systematically lower star-formation rate per unit mass of molecular gas for the sample galaxies. If the gas is concentrated in the inner regions of the galaxy, this delay of its transformation into stars can be logically associated with the high angular velocity of rotation in the inner disk region, which enhances the gravitational stability of the gaseous disk against local perturbations, which, in turn, leads to the accumulation of gas and its subsequent transition into molecular form.  We note, however, that the concentration of molecular gas toward the center of the galaxy is a fairly common (although according to \\cite{boker}, it is less pronounced for galaxies of the latest morphological types). However, a predominance of molecular over atomic gas is a rare phenomenon, implying that a central concentration of a gas is essential, but no means a sufficient condition for such a predominance.  To understand better the causes for the transformation of  most of $HI$ into $H_2$ on a galactic scale a more complete picture of gas evolution  remains needed.  The work was supported by RFBR grant 04-02-16518."
    },
    "0609/astro-ph0609383_arXiv.txt": {
        "abstract": "{An investigation of the kinematics of a plasma stream rotating in the pulsar magnetosphere is presented. On the basis of an exact set of equations describing the behavior of the plasma stream, the increment of the instability is obtained, and the possible relevance of this approach for the understanding of the pulsar rotation energy pumping mechanism is discussed. ",
        "introduction": "One of the important stages in the development of pulsar radiation theory was the discovery that the rotation energy could transform into energy in the electrostatic field; this possibility stimulated the modeling of pulsar magnetospheres (see Ref. ~\\cite{de},~\\cite{gj}). In a typical model, the electrostatic field $\\bf E_0$, generated over the pulsar surface, has a finite projection along the dipole magnetic field $\\bf B_0$. The resulting electric force, $e E_0$ exceeds the gravitational force near the neutron star surface, and ends up uprooting charged particles from the surface layer. These particles are accelerated in the electric field and radiate $\\gamma$ photons because of the curvature radiation in the dipolar field. For photon energies greater than twice the electron rest mass ( $\\epsilon_\\gamma\\ge 2m_e c^2$),  the electron-positron pair production  ($\\gamma+{\\bf B_0}\\rightarrow e^+ +e^- +\\gamma'$) in the ambience of the magnetic field becomes possible. The particles, so produced, repeat the cycle; they are accelerated, generate $\\gamma$ radiation which creates more electron-positron pairs. This cascading leads to the build-up of a large flux of relativistic $e^+e^-$ pairs ; the process continues till the pair plasma screens the electric field $\\bf E_0$ as shown in Ref. ~\\cite{st},~\\cite{tad}.  The $e^+, e^-$ population in the pulsar atmosphere may be conveniently divided into three principal components: 1) the basic plasma mass (Bulk) with concentration $n_{pl}$, and a Lorentz factor $\\gamma_{pl}$, 2) a tail with concentration $n_t$, and Lorentz factor $\\gamma_t$, and 3) the remnant of the primary electron or ion beam (most likely electrons). Note that  cascading due to pair production is possible only in a highly relativistic scenario, because for nonrelativistic or even mildly relativistic velocities, the pair annihilation dominates pair production (see, for example, Ref. ~\\cite{dir},~\\cite{kom}). The reader is referred to Ref.~\\cite{kaz},~\\cite{mac} for a guide to processes that define and control the behavior of radiation in the pulsar magnetosphere.  The essence of this ``standard model'' is that the pulsar magnetosphere is permeated with a multi-component plasma. It is generally assumed that the energy contents of the three major components are of the same order: $n_{pl}\\gamma_{pl}\\approx n_t\\gamma_t\\approx n_b\\gamma_b$. Theoretical estimates for the density and the energy of the primary component are: $n_b\\approx (10^{11}-10^{12})\\mbox{cm}^{-3}$ ( also called the Goldreich Julian  concentration) and $\\gamma_b\\approx 10^6-10^7$. The origin of the pulsar radiation is supposed to be in the magnetosphere, and hence the energy contained in the $e^+e^-$ plasma must be enough to account for the observable radiation.  In these pulsar models, the measure of the region over which the electric field is nonzero is in great significance. The distance between the star surface and the region where it is screened out is called the vacuum gap, and may be estimated to be $(10^3-10^4)\\mbox{cm}$ for $E_0\\approx 10^7\\,$G(see for example Ref. ~\\cite{rud}). Unfortunately the  particle energy inventory accumulated within the gap is not sufficient to explain the observed visible radiation. Several mechanisms have been invoked to increase the gap size. The intermediate formation of positronium (electron-positron bound  state) which would hold back the decay, and lead to an increase in the gap was proposed in Ref.~\\cite{us}. According to Ref.~\\cite{aar}, the pulsar magnetic field lines, supposedly curved towards rotation, must be screwed even more; the twisted magnetic field lines will be rectified, and as a result will enlarge the gap. A different mode, the PFF (pair formation front) mechanism was introduced instead of the gap. Taking the heating of the stellar surface and the electron thermoemission into account leads to nonstationarity, and explains nonstationary behavior of pulsars (see Ref.~\\cite{kazm}).  In a series of papers, Muslimov and Tsigan have attempted to solve the gap problem via the general relativity route. Realizing  that in the vicinity of the rotating neutron star, the space-time is slightly curved, they work out the creation of the electric field in Kerr metric (see Ref.~\\cite{mus}). As a result, the gap size is somewhat enlarged but not enough to make a difference.  It seems that some additional source of energy or some new mechanism will be necessary to surmount the problem arising out of the ``insufficiency'' of the energy content in the vacuum gap. For instance, it seems perfectly plausible and possible that one could draw upon the pulsar rotation energy to augment the energy content of the magnetospheric $e^+e^-$ plasma; the rotation energy, for example,  could be transformed into   energy  associated with oscillations in the $e^+e^-$ plasma far from the pulsar surface. In this region the complicating effects of unremovable gravitation (having to use, the Kerr metric, for example) will not exist. This paper is an attempt to formulate and investigate the problem associated with the parametric pumping of plasma oscillations in the pulsar magnetosphere  taking into consideration the complicated nature of the  multicomponent electron-positron substrate.   Let us review a standard system to examine the plausibility of the proposed mechanism. For the Crab Nebula, the nebula radiation can be surely sustained by the rotational energy of the pulsar PSR~0531 located in the vicinity of the nebular center. The power of the nebula radiation exceeds the pulsar radiation power by two orders of magnitude (it is approximately $2\\cdot10^{38}\\,$erg/sec). The only source capable of providing such a prodigious  power is the slowdown of the pulsar rotation: $\\dot W\\approx I\\Omega\\dot \\Omega$, where $I$ is pulsar's moment of inertia with mass of order $(1.5-2.5)M_\\odot$ ($M_\\odot$ is solar mass), $\\Omega$ and $\\dot\\Omega$ are star's angular velocity and angular acceleration respectively.  This power is equal to $5\\cdot 10^{38}\\,$erg/sec.  The rotation energy loss rate is described by the  ratio $\\dot W/W\\simeq2\\dot\\Omega/\\Omega=2\\dot P/P$, where $P=2\\pi/\\Omega$ is the neutron star's rotation period, and $P$ and $\\dot{P}$ are measurable parameters for pulsars. Ratio $\\dot P/P$ for different pulsars ranges from $10^{-11}\\sec^{-1}$ (PSR 0531) to $10^{-18}\\sec^{-1}$ (PSR 1952+29). Thus, for ordinary plasma oscillations to be pumped by the pulsar rotation, their growth rate must not exceed $\\dot P/P$, the measurable rate of rotation energy loss.   A hydrodynamic approximation will be used to study the problem of rotation induced wave generation in plasmas. For relative simplicity one will assume that the  ${\\it e}^+e^-$ plasma  has two distinct energy ranges: the lower range (still highly relativistic) with $n_{pl}$, and $\\gamma_{pl}$, and the beam with $n_b$, and $\\gamma_b$. These flows propagate along the rotating monopole like magnetic field lines. ",
        "conclusions": "The purpose of this paper was to explore  the possibility of pumping the  rotational energy of the pulsar into the plasma. Considering a highly idealized system, the linear instability caused by rotation in a two component relativistic plasma embedded in a uniform magnetic field was examined. By using the hydrodynamic equations of motion, the continuity and the Poisson equation, it has been shown that the plasma waves can grow on the rotation energy with rather high growth rates. In fact the perturbation growth rate for distances $R/R_{cyl}\\geq 0.85$ is quite high in comparison with the supposed  instability rates we were seeking. If this mechanism is, indeed, operational, then  the only consistent scenario is that the linear stage with this growth rate is very short, and nonlinearities are turned in soon enough to considerably reduce the growth. Thus the need for a nonlinear theory is immediately and strongly indicated. Sooner or later we should consider this particular nonlinear effect, which will comprise one more step closer to the real scenario. There are two other shortcomings of this effort: a) the straight magnetic field lines has been examined, whereas real profiles are curved and b) only electrostatic waves have been considered. The preliminary results, however, unambiguously show that the energy content of the magneto spheric plasma can grow at the expense of the stellar rotational energy."
    },
    "0609/astro-ph0609510_arXiv.txt": {
        "abstract": "The extra-dimensional origin of dark matter is a fascinating and nowadays often discussed possibility. Here, we present the gamma-ray signatures that are expected from the self-annihilation of Kaluza-Klein dark matter particles. For comparison, we contrast this with the case of supersymmetry, where the neutralino annihilation spectra take a very different form. In both cases we find pronounced spectral signatures that could in principle be used to distinguish between these two types of dark matter candidates already with today's detector resolutions. ",
        "introduction": "While the existence of dark matter is by now widely accepted, with evidence coming from various independent observations over an impressive range of distance scales, its nature still remains unknown, providing a major challenge for cosmology and astroparticle physics. One of the most favourate candidates is certainly the supersymmetric neutralino, theoretically well motivated by considerations from particle physics (for reviews, see, e.g. \\cite{chirev}).  More recently, the proposal of an alternative dark matter candidate \\cite{CMS,STa} has received much attention. It arises naturally in models with universal extra dimensions (UED) \\cite{app}, where all standard model fields are allowed to propagate in a higher-dimensional bulk. After compactification of the internal space, these additional degrees of freedom appear as towers of new, heavy states in the effective four-dimensional theory. The lightest of these Kaluza-Klein particles (LKP) is usually stable due to KK parity, an internal ${Z}_2$ symmetry that appears as a remnant of higher-dimensional translational invariance. Taking into account radiative corrections to the KK masses, the LKP is expected to be well approximated by the $\\B$, the first KK excitation of the weak hypercharge boson \\cite{CMS}. Detailled relic density calculations \\cite{KM} show that it can account for the required dark matter density as determined by WMAP \\cite{wmap} if the compactification scale (and thus the $\\B$ mass) lies in the TeV range, the precise value depending on the exact mass spectrum and the resulting coannihilation channels. This is particularly interesting since it means that the LHC, which will probe compactification scales up to about 3 TeV \\cite{LHCbound}, will have access to practically the whole cosmologically relevant parameter space. The prospects for both direct and indirect LKP detection have been studied in some detail and generally found to be quite promising, at least for next generation's detectors \\cite{oldfrag,LKPdet,BBEGa,BBEGb}.  The underlying idea of indirect detection is always to look for exotic contributions to cosmic rays that might originate from  the  self-annihilation of dark matter particles. Unfortunately, total fluxes are notoriously difficult to predict due to the largely unknown distribution of dark matter. For gamma rays, in particular, these astrophysical uncertainties result in an overall \\emph{amplitude} that may vary over several orders of magnitude \\cite{FPS}. The \\emph{spectral form}, on the other hand, is unaffected and can be predicted with great accuracy from the underlying particle physics. It is thus of great interest to extract pronounced spectral signatures that could unambiguously identify dark matter contributions to astrophysical gamma-ray signals and maybe even allow to discriminate between different dark matter candidates. Here, we present two such signatures, a sharp cutoff at the dark matter particle's mass and the line signal from direct annihilation into gamma rays, and compare the situation for the LKP and the neutralino. With a massive vector versus a Majorana fermion particle, note that this actually entails a comparison between two different \\emph{types} of dark matter particles so that the results should not be restricted to the specific models studied here. ",
        "conclusions": "Absolut dark matter annihilation fluxes are hard to predict due to the large astrophysical uncertainties involved. For a possible claim of an indirect detection of dark matter to be convincing, it should therefore be grounded on unambiguous spectral signatures.  We have presented such signatures for the gamma-ray spectrum from both Kaluza-Klein and supersymmetric dark matter: we find a sharp cutoff at the dark matter particle's mass, with features at slightly lower energies that would allow to distinguish between these two scenarios already with the energy resolution of typical ACTs in operation (this situation is visualized in Fig.~\\ref{compfig}). Since photons with the highest energies are dominantly produced by final state radiation, we have thereby also demonstrated the importance of taking into account this type of radiative corrections when considering dark matter annihilations -- a view that has been strongly advocated also in \\cite{bir}.  The spectral features that we have reported here can be looked for wherever one expects enhanced dark matter densities, e.g.~in the direction of nearby dwarf galaxies or the galactic center \\cite{FPS}; individual dark matter clumps or intermediate mass black holes \\cite{imbh} in the galactic halo provide alternatives that are also well worth investigating. With the advance of new gamma-ray instruments of unprecedented size and energy resolution \\cite{newgamma}, these will be exciting possibilities to further explore."
    },
    "0609/astro-ph0609726_arXiv.txt": {
        "abstract": "The origin, evolution and role of magnetic fields in the production and shaping of proto-planetary  and planetary nebulae (PPNe, PNe) is a subject of active research.  Most PNe and PPNe are axisymmetric with many exhibiting highly collimated outflows, however, it is important to understand whether such structures can be generated by isolated  stars or require the presence of a binary companion.  Toward this end we study a dynamical, large-scale $\\alpha-\\Omega$ interface dynamo operating in a $3.0$ $M_\\odot$ Asymptotic Giant Branch star (AGB) in both an isolated setting and one in which a low-mass companion is embedded inside the envelope.  The back reaction of the fields on the shear is included and differential rotation and rotation deplete via turbulent dissipation and Poynting flux.  For the isolated star, the shear must be resupplied in order to sufficiently sustain the dynamo.  Furthermore, we investigate the energy requirements that convection must satisfy to accomplish this by analogy to the sun.  For the common envelope case, a robust dynamo results, unbinding the envelope under a range of conditions.  Two qualitatively different types of explosion may arise: (i) magnetically induced, possibly resulting in collimated bipolar outflows and (ii) thermally induced from turbulent dissipation, possibly resulting in quasi-spherical outflows.  A range of models is presented for a variety of companion masses. ",
        "introduction": "Most planetary and proto-planetary nebulae (PNe, PPNe) are highly aspherical and exhibit a diversified morphology of axisymmetric structures and/or collimated jets.  However, the production of such richly varied systems and shaping mechanisms, remains an open topic (\\cite{2002ARA&A..40..439B}).  A central question is whether a binary is required to produce such asymmetries or if an isolated Asymptotic Giant Branch star (AGB) is sufficient.  Recently, detection of magnetic fields in AGB stars (\\cite{2004MNRAS.348...34E}, \\cite{2004MNRAS.354..529B}) and the central stars of PNe (\\cite{2005A&A...432..273J}) has sustained interest in magnetic launching and collimation mechanisms.  Observational evidence of a magnetically collimated jet in an evolved AGB star (\\cite{2006Natur.440...58V}) further suggests that the magnetic field may play a dynamical role.  Single star magnetic outflow models have been proposed as mechanisms for producing and shaping PPNe and PNe (\\cite{1997ApJ...489..946P}, \\cite{EB2001}).  Whether such models can power and shape the PPNe is uncertain (\\cite{2006PASP..118..260S}).  In particular, envelope dynamos are expected to be short ($<100$ yrs) and drain differential rotation energy rapidly (\\cite{2004ASPC..313..401B}), making it unlikely that isolated stars can produce observed asymmetries.  If convection can resupply differential rotation energy, then an envelope dynamo in an isolated AGB star may be viable.  Anisotropic convection in the sun resupplies differential rotation through the $\\lambda$-effect (\\cite{1989QB523.R93......}, \\cite{2004muga.book.....R}), however it is not clear if such a mechanism operates in AGB stars.  Rather than single star models, the observed asymmetries may instead be the result of energy and angular momentum supplied by binary interactions.  This is supported by recent surveys suggesting that most, if not all PNe involve binary systems (\\cite{2004ApJ...602L..93D}, \\cite{2004ASPC..313..515S}, \\cite{2006astro.ph..6354M}, \\cite{2006A&A...452..257M}). Although there are many different types of binary interactions and outcomes, here we focus on common envelope (CE) evolution in the context of PNe progenitors (\\cite{1993PASP..105.1373I}).  A common envelope (CE) model in which low-mass ($<0.3$ $M_\\odot$) companions were embedded into the envelope of a $3.0$ $M_\\odot$ AGB star was investigated in \\cite{JN2006}.  The CE evolution is advantageous as it can supply angular momentum in an extremely short period ($<1$ yr) and produce a range of PPNe outflows.  Such a model would predict white dwarf + brown dwarf close binaries that survive the CE phase.  Recently, a WD+ BD binary in a close orbit has been detected (\\cite{2006Natur.442..543M}, \\cite{2006astro.ph..9366B}).  A separation distance of $0.65$ $R_\\odot$ indicates that the system incurred a common envelope phase in which the brown dwarf was responsible for ejecting the progenitor envelope.  This further motivates more detailed study of CE induced PNe.   In this paper, we reinvestigate the magnetic model presented in \\cite{EB2001} in more detail, in an effort to determine the viability of a single star dynamo.  We compare the results to a model in which the rotation profile is supplied by a CE phase as in \\cite{JN2006}.  In section 2, we review previous work, compare single star evolution to that resulting from a CE interaction and focus on the depth that the poloidal field can diffuse into the shear zone.  We present our interface dynamo in section 3, including the detailed back reaction of the fields on the shear, the generation of heat through turbulent dissipation and the spin down of the star due to Poynting flux.  In Section 4, we present the results of our model for three cases:  (i)  an isolated dynamo in which convection does not resupply differential rotation, (ii) an isolated dynamo in which a fraction of the convective energy resupplies differential rotation energy and (iii) a dynamo resulting from the in-spiral of a low-mass ($\\leq0.05$ $M_\\odot$) companion inside the stellar envelope.  We conclude in section 5. ",
        "conclusions": "Extraction of rotational energy is likely fundamental to the formation of multipolar PPNe and PNe.  Magnetic dynamos can play an intermediary role in facilitating the extraction of rotational energy.  Here, we have studied a large-scale, dynamical interface dynamo operating in the envelope of a $3.0$ $M_\\odot$ AGB star.  The back reaction of field amplification on the shear is included as are the drain of  rotation and differential rotation via both turbulent dissipation (thermally induced) and Poynting flux (magnetically induced).  Two different dynamos are studied: (i) that of an isolated star and (ii) that of a common envelope system in which the secondary is a low-mass companion ($<$ $0.05$ $M_\\odot$).  For the isolated case, we find that only when two conditions are met can the single star dynamo drive PPNe.  First, $\\Delta\\Omega$ must be re-seeded.  This may occur by analogy to the $\\lambda$-effect in the sun (\\cite{2004muga.book.....R}).  Secondly, $\\Omega$ must be constant.  This implies that the field is stored in the shear layer until the end of the AGB phase.  When these two stringent conditions are met, a small fraction of the energy cascade rate can provide the necessary shear energy to sustain the interface dynamo (in some cases as little as $0.1\\%$).  Not only is the dynamo maintained, but the time-integrated Poynting flux is large enough to overcome the envelope binding energy.    Whether or not isolated star dynamos can produce a PPNe, a binary interaction can do so more robustly for a wide range of cases.  Common envelope evolution is advantageous in several ways.  Energy and angular momentum are supplied very quickly ($<$ 1 yr) and often in less than or equal to a dynamo cycle period, allowing the dynamo to operate once the secondary has completed its in-spiral.  For our common envelope dynamo model, a range of companion masses can easily supply enough differential rotation energy to power either a dynamo driven jet or a thermally driven outflow.  A magnetically dominated explosion likely produces a collimated, poloidal outflow while that of a thermally induced explosion is expected to be more spherical.  We have highlighted some of the basic key issues of isolated and common envelope dynamos, however, more detailed research is needed in both areas.  The viability of anisotropic convection reseeding differential rotation must be determined.  If convection cannot reseed shear, then we are faced with the proposition that binary interactions are required to produce axisymmetric PNe.  In the CE case, the complex interaction between companion and envelope, multi-dimensional aspects of the dynamo, realistic rotation profiles of isolated stars and the inclusion of a wider array of secondary masses are just a few of the many problems which warrant future work.  Constraints on the turbulent diffusion coefficient as a function of radius also need to be determined.  The physics involved in transitioning from a dynamo to a fully active jet must be understood."
    },
    "0609/astro-ph0609040_arXiv.txt": {
        "abstract": "We consider fluid perturbations close to the ``phantom divide'' characterised by $p=-\\rho$ and discuss the conditions under which divergencies in the perturbations can be avoided. We find that the behaviour of the perturbations depends crucially on the prescription for the pressure perturbation $\\dep$. The pressure perturbation is usually defined using the dark energy rest-frame, but we show that this frame becomes unphysical at the divide. If the pressure perturbation is kept finite in any other frame, then the phantom divide can be crossed. Our findings are important for generalised fluid dark energy used in data analysis (since current cosmological data sets indicate that the dark energy is characterised by $p\\approx-\\rho$ so that $p<-\\rho$ cannot be excluded) as well as for any models crossing the phantom divide, like some modified gravity, coupled dark energy and braneworld models. We also illustrate the results by an explicit calculation for the ``Quintom'' case with two scalar fields. ",
        "introduction": "The discovery by the supernova surveys \\cite{sn1,sn2} that the expansion of the universe is currently accelerating came as a great surprise to cosmologists. Within the standard cosmological framework of a nearly isotropic and homogeneous universe and an evolution described by General Relativity, this behaviour requires a component with a negative pressure $p<-\\rho/3$, commonly dubbed {\\em dark energy}. However, it is very difficult to understand why the dark energy should appear at such a low energy scale, and why it should start to dominate the overall energy density just now. Explaining its nature is correspondingly regarded as one of the most important problems in observational cosmology.  Current limits on the equation of state parameter $w=p/\\rho$ of the dark energy seem to indicate that $p\\approx -\\rho$ \\cite{cora,wlim}, sometimes even that $p<-\\rho$ \\cite{phantom}, often called {\\em phantom energy} \\cite{caldwell}. Although there is no problem to consider $w<-1$ for the background evolution, there are apparent divergencies appearing in the perturbations when a model tries to cross the ``phantom divide'' $w=-1$\\cite{crossing}. Even though this region may be unphysical at the quantum level \\cite{cht,Cline:2003gs}, it is still important to be able to probe it, not least to test for alternative theories of gravity or higher dimensional models which can give rise to an effective phantom energy \\cite{para,nojiri,vli,vik2,kosh}. It would certainly be unwise to build a strong bias like $w\\geq-1$ into our analysis tools as long as experiments do not rule it out. In this paper we consider the evolution of the perturbations for models where $w$ crosses $-1$. We find that in many realistic cases the divergencies are only apparent and can be avoided. At the level of cosmological first order perturbation theory, there is no fundamental limitation that prevents an effective fluid from crossing the phantom divide.  The paper is organised as follows: We start with a short recapitulation of the first order perturbations in fluids, which also serves to define our notation. In section III we study the behaviour of barotropic fluids close to $p=-\\rho$, concluding that this class of fluids cannot cross the phantom divide self-consistently. We allow for non-adiabatic perturbations in section IV and show that the phantom divide can now be crossed as long as we define the pressure perturbation in a frame that stays physical. We then illustrate the results with the Quintom model of two scalar fields before presenting our conclusions. The appendices finally discuss the calculation of the effective Quintom perturbations in more detail. ",
        "conclusions": "As long as the cosmological data indicates the presence of a dark energy with an effective equation of state $p\\approx-\\rho$ it will be necessary to consider models with the same equation of state. In general, we will have to allow the equation of state to cross the phantom divide, $w=-1$. Even though such a fluid model may not be viable at the quantum level, it is possible that this behaviour is only apparent, or due to a modification of General Relativity or the existence of more spatial dimensions. In analysing the data we therefore have to be able to use general self-consistent models at the level of linear perturbation theory.  In this paper we have studied the behaviour of the perturbations in general perfect fluid models close to $w=-1$. We have shown that although models with purely adiabatic perturbations cannot cross $w=-1$ without violating important physical constraints (like causality or smallness of the perturbations), it is possible to rectify the situation by allowing for non-adiabatic sources of pressure perturbations. However, the parametrisation of $\\dep$ in terms of the rest-frame perturbations of the energy density cannot be used as this frame becomes unphysical at $w=-1$. By parameterising $\\dep$ instead in any other frame the divergencies are avoided.  We also computed all quantities in the Quintom model which provides an explicit example of the above mechanism. In this model, even though the propagation speed of sound waves remains finite and constant, the additional internal and relative pressure perturbations lead to an apparent sound speed which diverges. It is only the sum of all contributions to $\\dep$ which remains finite.  A more speculative conclusion is that it seems difficult for ``fundamental'' fields to cross $w=-1$ as their apparent rest-frame sound speed, defined through Eq.~(\\ref{eq:csrest}), would generally be the actual propagation speed of their perturbations, which must remain smaller than the speed of light. From the discussion in section \\ref{sec:non_ad} we learn that fields with a well defined rest-frame and sound speed have only two routes to phantom crossing: \\begin{itemize} \\item {\\em $dw/da=0$ at crossing:} This looks rather fine-tuned as the field needs to be aware of the presence of the barrier. On the other hand, a normal minimally coupled scalar field reaches $w=-1$ when $\\dot{\\phi}=0$, and it does so with zero slope. A scenario where phantom crossing of this kind is realised could be built by changing the sign in front of the kinetic term whenever $\\dot{\\phi}=0$. As an example, a cosine potential then leads to a ``phaxion'' scenario, see Fig. \\ref{fig:phaxion}. On average, the phaxion behaves like a cosmological constant, but oscillates around $w=-1$. However, it is unclear how to construct such a scenario in a covariant way. \\item {\\em vanishing sound speed:} A field with a vanishing sound speed can also avoid the logarithmic divergence of $\\dep$. This requires either a coupling between $1+w$ and $\\cs$ or else the sound speed might be zero at all times, see Fig.~\\ref{fig_cs0}. In the latter case the pressure of the field would need to remain close to $-\\rho$ at all times in order to prevent the small-scale perturbations from growing too quickly. Also, the sound speeds needs to be {\\em exactly} zero, otherwise the logarithmic divergence reappears. This may require a symmetry enforcing $\\cs=0$ as even a very small but non-zero sound speed would render this scenario inviable. \\end{itemize} In view of these difficulties, the Quintom family of models, although more complicated than a single field, may prove to be the conceptually simplest way to cross $w=-1$ with a model defined through an action. They also illustrate that the effective pressure perturbation may provide a kind of fingerprint of the mechanism behind the accelerated expansion of the universe if it can be measured.  \\begin{figure} \\epsfig{figure=figures/phax_b.eps,width=3.3in} \\caption{The phaxion: The sign in front of the kinetic term of a normal scalar field is flipped every time $\\dot{\\phi}$ passes through zero (so that $\\dot{\\phi}\\geq0$). As potential we use $V(\\phi)=1+\\cos(\\phi/m_0)$ which is bounded both from above and below. The field then moves to large values (top left graph), while $w$ oscillates around $-1$ (top right). The energy density $\\rho$ also oscillates so that the field mimics an effective cosmological constant (bottom left). As the phantom divide is crossed with zero slope, there is no divergence in the pressure perturbation (bottom right), but the time-dependent effective mass can lead to strong growth of the perturbations in spite of $\\cs=1$. \\label{fig:phaxion}} \\end{figure}  \\begin{figure} \\epsfig{figure=figures/cs0b.eps,width=3.3in} \\caption{The pressure perturbation $\\dep$ for three fluids with different rest-frame sound speed, $\\cs=0.1$ (top curve, red dash-dotted line), $\\cs=0.01$ (middle curve, blue dashed line) and $\\cs=0$ (lowest, solid line). Only the last case does not have a logarithmic divergence in the pressure perturbation at the phantom divide. \\label{fig_cs0}} \\end{figure}  Concerning the use of phantom-crossing fluids for the purpose of data analysis, it is straightforward to implement them by just avoiding the usual parametrisation of the pressure perturbations in terms of the rest-frame sound speed. However, there does not seem to be a canonical way to choose the pressure perturbations if one is not allowed to use the dark energy rest-frame. The aim for the far future will be to directly measure the pressure perturbations of the dark energy in order to gain insight into the physical origin of the phenomenon. For now, we have to ensure that the definition of $\\dep$ does not lead to unphysical situations, while preserving the usual parametrisation in terms of the rest-frame sound speed as much as possible far away from the phantom divide. Maybe the simplest way out is to regularise the adiabatic sound speed $\\ca$ which appears in Eq.~(\\ref{eq:dp_rest}) because of the gauge transformation into the rest-frame. While any finite choice of $\\dep$ is a physically acceptable choice, it is preferable to modify Eq.~(\\ref{eq:dp_rest}) in a minimal way so that the usual interpretation of $\\cs$ is preserved for $w\\neq-1$. We propose to use \\be \\tilde{c}_a^2 = w - \\frac{\\dot{w} (1+w)}{3 \\HH [(1+w)^2+\\lambda]} \\label{eq:intca} \\ee where $\\lambda$ is a tuneable parameter which determines how close to $w=-1$ the regularisation kicks in. A value of $\\lambda\\approx 1/1000$ should work reasonably well, as shown in Fig.~\\ref{fig_interpolation}. In this case $\\dep$ is well-defined and there are no divergencies appearing in the perturbation equations (\\ref{eq:pert_cs_d}) and (\\ref{eq:pert_cs_v}). Although the differences in $\\dep$ for the different choices of $\\lambda$ in Fig.~\\ref{fig_interpolation} look important, we have to remember that, firstly, the perturbations in the dark energy are normally small (especially close to the phantom divide) and that secondly they are only communicated to the other fluids via the gravitational potential $\\psi$. As in Fig.~\\ref{fig-2b}, we find also here that the dark energy perturbations are subdominant. Computing the CMB power spectrum for $\\lambda=1/100$ and $\\lambda=10^{-4}$ we find a relative difference in the $C_\\ell$ of less than 1\\% on large scales, which rapidly drops to $10^{-6}$ by $\\ell=50$, well below cosmic variance. We expect detectable differences only for strongly clustering dark energy with a sound speed that is close to zero or even negative. Although it is reassuring that directly measureable quantities are not sensitive to the precise value of $\\lambda$ in Eq.~(\\ref{eq:intca}), this also shows that it will be very difficult to distinguish between different models of dark energy if $w\\approx-1$.  \\begin{figure} \\epsfig{figure=figures/interpol.eps,width=3.3in} \\caption{The pressure perturbation for a fluid example with $\\cs=0.1$. The red dash-dotted curve shows the standard case with the logarithmic divergence at $w=-1$. The blue (dashed), black (solid) and green (dotted) curve use a regularised adiabatic sound speed $\\tilde{c}_a^2$ which does not diverge. The black curve with $\\lambda=1/1000$ provides a reasonable fit. Larger values smooth too much, while smaller values start to follow the divergence and exhibit a temporary instability in the solution for $V$. \\label{fig_interpolation}} \\end{figure}"
    },
    "0609/astro-ph0609089_arXiv.txt": {
        "abstract": "We report on new \\chandra\\ exploratory observations of six candidate Type~2 quasars at \\hbox{$z$=0.49--0.73} selected among the most \\oiii\\ luminous emitters from the Sloan Digital Sky Survey (SDSS). Under the assumption that \\oiii\\ is a proxy for the intrinsic luminosity of the central source, their predicted rest-frame \\xray\\ luminosities are $L_{2-10keV}\\approx10^{45}$~\\lum. For two of the targets, the photon statistics are good enough to allow for basic \\xray\\ spectral analyses, which indicate the presence of intrinsic absorption ($\\approx$~10$^{22-23}$~cm$^{-2}$) and luminous \\xray\\ emission ($L_{\\rm X}\\gtrsim10^{44}$~\\lum). Of the remaining four targets, two are detected with only a few (3--6) \\xray\\ counts, and two are undetected by \\chandra. If these four sources have the large intrinsic \\xray\\ luminosities predicted by the \\oiii\\ emission, then their nuclei must be heavily obscured ($N_{\\rm H}>$~few~$\\times10^{23}$~cm$^{-2}$) and some might be Compton thick ($N_{\\rm H}>1.5\\times10^{24}$~cm$^{-2}$). We also present the results for two Type~2 quasar candidates serendipitously lying in the fields of the \\chandra\\ targets, and provide an up-to-date compilation of the \\xray\\ properties of eight additional SDSS Type~2 quasars from archival \\chandra\\ and \\xmm\\ observations (five with moderate-quality \\xray\\ data). The combined sample of 16 SDSS Type~2 quasars (10 \\xray\\ detections) provides further evidence that a considerable fraction of optically selected Type~2 quasars are obscured in the \\xray\\ band (at least all the objects with moderate-quality \\xray\\ spectra), lending further support to the findings presented in Vignali, Alexander and Comastri (2004a) and unification schemes of Active Galactic Nuclei, and confirms the reliability of \\oiii\\ emission in predicting the \\xray\\ emission in obscured quasars. ",
        "introduction": "The quest for the identification of Type~2 quasars, the high-luminosity, long-sought after ``big cousins\" of local Seyfert~2 galaxies (i.e., characterized by high-ionization, narrow emission lines and a lack of broad emission lines in the rest-frame optical/ultraviolet spectra), has been the topic of numerous investigations over the past few years. Although there is not a unique definition of Type~2 quasars in the \\xray\\ band, generally the most accepted view is that these sources should be the luminous and obscured ($L_{\\rm X}>10^{44}$~\\lum\\ and $N_{\\rm H}>10^{22}$~cm$^{-2}$) Active Galactic Nuclei (AGN) population predicted by unification schemes of AGN (e.g., Antonucci 1993) and required by many synthesis models of the \\xray\\ background (XRB; e.g., Comastri et al. 2001; Gilli, Salvati \\& Hasinger 2001). For example, in the most recent version of the synthesis model of Gilli, Comastri \\& Hasinger (2006), the contribution of Type~2 quasars to the XRB is estimated to be $\\approx$~15~per cent.  Over the past 4--5 years, deep \\xray\\ surveys (such as those conducted in the \\chandra\\ deep fields; e.g., Giacconi et al. 2002; Alexander et al. 2003; see Brandt \\& Hasinger 2005 for a recent review) performed with \\chandra\\ and \\xmm\\ have resolved a large fraction of the XRB ($\\approx$~80~per~cent in the 2--10~keV band; e.g., Bauer et al. 2004; Hickox \\& Markevitch 2006; but see also Worsley et al. 2005) and detected many high-redshift Type~2 quasar candidates (e.g., Crawford et al. 2001, 2002; Mainieri et al. 2002; Barger et al. 2003). Optical and infrared spectroscopy is critical to determine the redshifts of these sources and confirm that these candidates are ``genuine'' Type~2 quasars (i.e., matching both the optical and \\xray\\ definitions). However, the majority of these objects are optically faint (most likely because they lie at high redshifts) and are therefore challenging targets to obtain spectroscopic redshifts. As a consequence, although considerable efforts have been made, the number of spectroscopically confirmed Type~2 quasars from these surveys is comparatively small (e.g., Norman et al. 2002; Stern et al. 2002; Fiore et al. 2003; Perola et al. 2004; Gandhi et al. 2004; Caccianiga et al. 2004; Szokoly et al. 2004; Mainieri et al. 2005a; Maiolino et al. 2006), despite the fact that as many as 15--20~per~cent of luminous obscured quasars are predicted to be observed in the deepest \\xray\\ surveys (Alexander et al. 2001; Mainieri et al. 2005b).  The goal of our project is to select optically luminous Type~2 quasar candidates identified on the basis of high \\oiii5007\\AA\\ luminosity from the ground-based, large-area Sloan Digital Sky Survey (SDSS; York et al. 2000) spectroscopic data and then follow up these sources with \\xray\\ observations to confirm whether or not they are luminous and obscured also in the \\xray\\ band. The advantage of this approach, compared to the study of Type~2 quasars found in moderately deep and ultra-deep \\xray\\ surveys, is that at the optical magnitude limit of the SDSS, Type~2 quasars are generally easy to study at both optical and \\xray\\ wavelengths. In addition, \\xray\\ studies of optically selected Type~2 quasars can explore whether selection in the optical and \\xray\\ bands identifies sources with intrinsically different properties. Due to observational limitations (e.g., small-size samples, low percentages of spectroscopic identifications, limited regions of the sky being surveyed), such studies have been challenging or impossible in the past.  In this paper we extend the work presented by Vignali, Alexander and Comastri (2004a, hereafter V04; 2004b) and focus on the study of the \\xray\\ properties of Type~2 quasar candidates originally selected from the SDSS in the redshift range $\\approx$~0.3--0.8. These sources are classified on the basis of high-excitation, narrow emission lines without underlying broad components and with line ratios characteristic of non-stellar ionizing radiation (Zakamska et al. 2003; see Zakamska et al. 2004, 2005 and 2006 for observations of these objects at other wavelengths). From the analysis of archival \\rosat\\ and \\xmm\\ observations, V04 were able to place constraints in a statistical sense on the \\xray\\ emission of 17 of these objects (three \\xray\\ detections and 14 upper limits). Assuming the correlation between the \\oiii\\ and the \\hbox{2--10~keV} flux found for Seyfert~2 galaxies (Mulchaey et al. 1994) and adopting the parameterization reported in $\\S$3.2 of Collinge \\& Brandt (2000), V04 found that at least 47~per~cent of the observed sample showed indications of \\xray\\ absorption, including the four highest luminosity sources with predicted intrinsic luminosities of $\\approx$~10$^{45}$~\\lum.  Here we present the exploratory \\chandra\\ observations for six additional SDSS Type~2 quasar candidates and provide an up-to-date list of the \\xray\\ properties of Type~2 quasar candidates for which archival \\chandra\\ and \\xmm\\ data are available (see also Ptak et al. 2006, hereafter P06, where pointed and archival \\xray\\ observations of a sample of mostly bright SDSS Type~2 quasar candidates are presented). In total we have 16 sources (10 \\xray\\ detections) with available \\chandra\\ and \\xmm\\ data at present. The choice of \\chandra\\ and \\xmm\\ is motivated by their broad-band coverage (thus minimizing the uncertainties due to extrapolations as in the case of \\rosat\\ data), relatively large field-of-view, spectroscopic capabilities, and good sensitivity to relatively \\xray\\ faint sources.  Hereafter we adopt the ``concordance'' (WMAP) cosmology ($H_{0}$=70~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{\\rm M}$=0.3, and $\\Omega_{\\Lambda}$=0.7; Spergel et al. 2003). ",
        "conclusions": ""
    },
    "0609/astro-ph0609276_arXiv.txt": {
        "abstract": "{} { Analysis and interpretation of the Her X-1 X-ray light curve obtained with the ASM onbord RXTE over the period 1996 February to 2004 September. } { We construct the averaged X-ray light curves by means of adding up light curves corresponding to different 35\\,day cycles. A numerical model is introduced to explain the properties of the averaged light curves. } { We report that the features found previously in the averaged X-ray lightcurve are confirmed by the new RXTE/ASM data. In particular, anomalous dips and post-eclipse recoveries in two successive orbits in the short-on state are found to be prominent and stable details of the light curve. We argue that a change of the tilt of the accretion disk over the 35\\,d period is necessary to account for these features and show that our numerical model can explain such a behavior of the disk and reproduce the observed details of the light curve. } {} ",
        "introduction": "Her X-1/HZ Her is a close binary system with a 1.7\\,d orbital period consisting of an accretion powered 1.24\\,s X-ray pulsar \\citep{Giacconi73, Tananbaum72} and an optical companion -- a main sequence star of spectral type A/F \\citep{Crampton74}. The X-ray flux of the source shows a $\\sim$35\\,d periodicity which is thought to be due to a counter-orbitally precessing tilted accretion disk around the neutron star \\citep{GerBoy76}.  The main feature of the X-ray light curve is the alternation between {\\em on} (high X-ray flux) and {\\em off} (low X-ray flux) states. The 35\\,d cycle contains two {\\em on} states -- the {\\em main-on} ($\\sim$7 orbital periods) and the {\\em short-on} ($\\sim$5 orbital cycles) -- separated by $\\sim$4-5 orbital cycles. The maximum X-ray flux of the main-on is $\\sim$5 times higher than that in the maximum of the short-on. The {\\em off} states are usually explained by periodic obscuration of the X-ray source by the disk which most likely has a twisted form \\citep{GerBoy76}.  According to the generally accepted model \\citep[see e.g.][ and references therein]{SchandlMeyer94,Kuster05} the main-on state starts when the outer disk rim opens the line of sight to the source. Subsequently, at the end of the main-on state the inner part of the disk (that may be surrounded by a hot rarefied corona) covers the source from the observer. This picture seems to be confirmed by spectral analysis. The opening of the X-ray source is accompanied by notable spectral changes including evidence for the presence of a strong absorption, whereas the decrease in X-ray intensity occurs more slowly and without appreciable spectral changes (\\citealt{Giacconi73, Kuster05}, see however \\citealt{Oosterbroek00}).  The transition from the {\\em off} to the main-on state is usually called the {\\em turn-on} of the source. An interesting fact is that turn-ons occur near orbital phases $\\phi_{\\rm orb} \\simeq$ 0.2 or $0.7$. This behavior has been explained by \\citet{LevJer82} and \\citet{Katz82}. These authors assumed that the disk undergoes a wobbling with half synodal period due to tidal torques. At $\\phi_{\\rm orb} \\simeq 0.2$ and $0.7$ the angle between the disk plane and the observer's line of sight changes most rapidly.  From the data described above it follows that the duration of the 35\\,d cycle is a half-integer number of orbital cycles. Observations show that in most cases the duration of 35\\,d cycles is 20, 20.5 or 21 orbital cycles \\citep{Staubert83,Staubert06b}.  Another interesting feature of the X-ray light curve is a sudden decrease in flux accompanied by significant spectral changes -- the {\\em X-ray dips}. The dips were observed many times by different instruments \\citep[][ and references therein]{Shakura99}. X-ray dips can be separated into three groups: {\\em pre-eclipse dips}, which are observed in the first several orbits after X-ray turn-on, and march from a position close to the eclipse toward earlier orbital phase in successive orbits; {\\em anomalous dips}, which are observed at $\\phi_{\\rm orb} = 0.45 - 0.65$ and {\\em post-eclipse recoveries}, which are occasionally observed as a short delay (up to a few hours) of the egress from the X-ray eclipse in the first orbit after turn-on.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{asmx1.ps}} \\caption{\\small Her X-1 ASM light curve sample.} \\label{asm} \\end{figure}   Several physical models have been proposed to explain the origin and properties of the X-ray dips. For example, in many models \\citep{CrosaBoynton80, Gorecki82, Leahy97, Schandl96} the dips are caused by thickening of the outer rim of the accretion disk due to interaction with the accretion stream; \\citet{VrtilekHalpern85} suggested that the dips are produced by the accretion stream above and below the disk \\citep[see also][]{BochkarevKoritskaia89}. However, these models made various assumptions about the disk parameters (for example, in the original model by Crosa and Boynton the disk size was uncomfortably large to explain the X-ray dip marching behaviour).    Another model for X-ray dips was proposed by \\citet{Shakura99} In this model pre-eclipse dips occur when the accretion stream crosses the observer's line of sight before entering the disk. This can happen only if the stream is non-coplanar to the systems orbital plane. The reason for the stream to move out of the orbital plane is non-uniform X-ray heating of the optical stars atmosphere by the X-ray source, which produces a temperature gradient near the inner Lagrange point. The non-uniformity of the heating comes from the partial shadowing of the optical star surface by the accretion disk. Furthermore, such a stream forces the outer parts of the accretion disk to be tilted with respect to the orbital plane. The tidal torques cause the disk to precess in the direction opposite to the orbital motion. Due to tidal torques and the dynamical action of the accretion stream the outer parts of the disk develop a notable wobbling (nutational) motion twice the synodal period. For an observer, the X-ray source can be screened by the outer parts of the disk for some time during the first orbit after the X-ray turn-on. This causes anomalous dips and post-eclipse recoveries.  The analysis of Her X-1 observations obtained with the {\\em All Sky Monitor} (ASM) on board of the {\\em Rossi X-ray Timing Explorer} (RXTE) satellite \\citep[see][]{Bradt93} over the period 1996 February to 2004 September allowed to improve the statistics of the averaged X-ray light curve of Her X-1 with respect to previous analyses \\citep{Shakura98a,ScoLea99,StiBoy04}. In addition to well-known features (pre-eclipse dips and anomalous dips during the first orbit after X-ray turn-on) the averaged X-ray light curve shows that anomalous dips and post-eclipse recoveries are present for {\\em two} successive orbits after the turn-on in the short-on state. These details have already been reported by \\citet{Shakura98a}. Now using about two times more data we confirm these features to be stable details of the light curve (they are also seen in one RXTE/PCA pointed observation of the source during the short-on state \\citep{InamBaykal05}).  In this work we present averaged X-ray light curves of Her X-1 obtained from the analysis of ASM observations. Based on this data we modify the model developed by \\citet{Shakura99} in order to reproduce the details of the averaged light curve which were left unexplained  previously. We show that the accretion disk needs to change its tilt during the 35\\,d cycle from $\\sim 20^\\circ$ in the main-on to $\\sim 4^\\circ$ in the short-on. ",
        "conclusions": "The following main results have been obtained from analysing and modeling RXTE/ASM X-ray light curves of Her~X-1. \\begin{enumerate} \\item The shape of the averaged X-ray light curves is determined more accurately than was possible previously \\citep[e.g. by ][]{Shakura98a}. This has allowed us to confirm the significance and stability of the features in the averaged short-on light curve which have been found previously: the anomalous dips in two successive orbits after the beginning of the short-on state with a post-eclipse recovery after the first eclipse. \\item We argue that the changing of the tilt of the accretion disk with 35\\,d phase is necessary to account for the appearance of these features and the observed duration of the main-on and short-on states. \\item We show that the joint action of tidal forces and the dynamical interaction of the gas stream with the outer parts of the accretion disk leading to its wobbling can be responsible for the observed behavior of the disk. \\end{enumerate}"
    },
    "0609/astro-ph0609795_arXiv.txt": {
        "abstract": "We derive expressions, in terms of ``polar shapelets'', for the image distortion operations associated with weak gravitational lensing. Shear causes galaxy shapes to become elongated, and is sensitive to the second derivative of the projected gravitational potential  along their line of sight; flexion bends galaxy shapes into arcs, and is sensitive to the third derivative. Polar shapelets provide a natural representation, in which both shear and flexion transformations are compact. Through this tool, we understand progress in several weak lensing methods. We then exploit various symmetries of shapelets to construct a range of shear estimators with useful properties. Through an analogous investigation, we also explore several flexion estimators. In particular, some of the estimators can be measured simultaneously and independently for every galaxy, and will provide unique checks for systematics in future weak lensing analyses. Using simulated images from the Shear TEsting Programme (STEP), we show that we can recover input shears with no significant bias. A complete software package to parametrize astronomical images in terms of polar shapelets, and to perform a full weak lensing analysis, is available on the world wide web. ",
        "introduction": "\\label{intro}  Weak gravitational lensing is a powerful method to map the distribution of mass in the Universe, regardless of its nature or state (for reviews see Mellier 1999; Bartelmann \\& Schneider 2001; Refregier 2003). The apparent shapes of background galaxies become distorted as their light travels near mass concentrations along their line of sight to the Earth. The well-known shearing of galaxies, in which intrinsically circular sources would be seen as elongated ellipses, is induced by an amount proportional to the second derivative of the projected foreground gravitational potential. Such distortion has been measured around individual galaxy clusters (\\eg\\ Wittman \\etal\\ 2001; Wittman \\etal\\ 2003; Bacon \\& Taylor 2003; Brada\\v{c} \\etal\\ 2005; Wittman \\etal\\ 2006) and, in a statistical fashion, by large scale structure (recent measurements include Massey \\etal\\ 2004; Van Waerbeke \\etal\\ 2005; Heymans \\etal\\ 2005; Jarvis \\etal\\ 2006; Hoekstra \\etal\\ 2006; Semboloni \\etal\\ 2006; Hetterscheidt \\etal\\ 2006; Schrabback \\etal\\ 2006; Kitching \\etal\\ 2007; Massey \\etal\\ 2007b).  A higher-order effect, known as ``flexion'', is also emerging as a probe of the distribution of mass on small scales, and particularly in the inner cores of galaxy clusters (Goldberg \\& Natarajan 2002; Irwin \\& Schmakova 2003; Goldberg \\& Bacon 2005; Bacon \\etal\\ 2006; Irwin \\& Schmakova 2006; Okura, Umetsu \\& Futamase 2006; Goldberg \\& Leonard 2007). Variation in the shear signal across the width of a background galaxy causes bending in its apparent shape. This is the next term in a lensing expansion that leads towards the formation of an arc, as in strong lensing. The flexion is sensitive to the third derivative of the projected gravitational potential.  Precise image analysis techniques are required to detect weak gravitational lensing, because the shapes of galaxies are changed by the effect by only a few percent. In fact, the lensing contribution to the shape is about an order of magnitude smaller than the dispersion of galaxies' intrinsic morphologies and the spurious distortions introduced by typical imperfections in telescopes. The widely used shear measurement method by Kaiser, Squires \\& Broadhurst (KSB, 1995) has been successful in many contexts, but contains several documented shortcomings: it is found to be insufficiently accurate to measure shears with a desired accuracy of less than 1\\% (\\cf\\ Bacon \\etal\\ 2001; Erben \\etal\\ 2001; Heymans \\etal\\ 2005 (STEP1); Massey \\etal\\ 2007a (STEP2)), and it is mathematically ill-defined for realistic Point Spread Functions (\\cf\\ Kaiser 2000; Kuijken 1999; Hirata \\& Seljak 2003).  Several new shear measurement methods are being developed, to fully exploit future space-based weak lensing surveys with HST or the proposed SNAP, DUNE or JDEM missions, and ground-based wide-field surveys such as those with Megacam, CTIO DES, VISTA darkCAM, Pan-STARRS and LSST. A review of the various shear measurement methods is found in STEP2, along with their division into ``active'' and ``passive'' categories. Active techniques work by modelling galaxies as intrinsically circular, then shearing the models until they most closely match the observed ellipticities. Passive methods work by measuring the apparent ellipticities of objects as well as higher order shape statistics, which are used to calibrate the ellipticities.  Flexion measurement methods are still in relative infancy. Initial attempts to mathematically describe the flexion distortion (Goldberg \\& Natarajan 2002; Irwin \\& Shmakova 2003) were formidably complicated. A passive estimator has been constructed by Okura, Umetsu \\& Futamase (2006), and further expanded by Goldberg \\& Leonard (2007). A completely different, probabilistic approach is taken by Irwin \\& Shmakova (2006) and Irwin, Shmakova \\& Anderson (2006). However, several important features in these approaches remain to be developed, and they remain mathematically complex; it is therefore desirable to find a formalism which allows maximum physical insight into the problem. An advance towards this was made by Goldberg \\& Bacon (2005), who related flexion to the formalism of Cartesian shapelets (Refregier 2002; Refregier \\& Bacon 2002). Shapelets contain all the mechanics necessary to deconvolve galaxies and flexion estimators from the effects of a PSF. The active method of Goldberg \\& Bacon (2005) and Goldberg \\& Leonard (2007) has been used to successfully detect the flexion signal. The mathematics has a simpler form, although it is still not as elegant as possible.  Here, we present the image manipulations of lensing theory in terms of the ``polar shapelets'' formalism (Refregier 2003; Massey \\& Refregier 2005). This suggests a complete, orthonormal set of basis functions into which any galaxy shapes can be decomposed. It also provides a neat way to deconvolve arbitrary galaxy shapes from an arbitrarily complicated PSF, so we can set out under the assumption that this problem is solved. Polar shapelets then provide a natural representation for both shear and flexion operations, with simple mathematical forms that yield transparent physical interpretation. The complex number approach used throughout polar shapelets matches very conveniently with the complex ellipticity notation of Blandford (1991) now ubiquitous in shear literature, and with the complex formalism of flexion developed by Bacon \\etal\\ (2006). A complete software package to decompose images into polar shapelets is available from the shapelets web site\\footnote{{\\tt http://www.astro.caltech.edu/$\\sim$rjm/shapelets}}.  We then exploit the inherent symmetries of polar shapelets to explore a comprehensive range of passive measurement methods for both shear and flexion. To create a shear or flexion estimator, we simply need to find a combination of shapelet coefficients that has the desired properties under each transformation.  We generally keep the estimators as close as possible to linear in the image, to minimise both noise and bias in the final result. The shapelet methodology resembles a continuation of the KSB method to higher order. However, the inclusion of higher order shape information, and a complete parametrization of galaxy morphology, provides several new opportunities to improve on KSB, and to remove its instabilities. Some of the shear and flexion estimators that we describe are also independent, and can be obtained simultaneously for each galaxy. These will provide invaluable new cross-checks for systematics in the data analysis, which are unique to this method, and can also be combined to increase the overall ratio of signal to noise.  As we shall discuss, one of the shear estimators has already been proved highly successful in a blind test on simulated images containing an applied shear, as part of the STEP programme (Heymans \\etal\\ 2006; Massey \\etal\\ 2007a). We defer detailed testing of the remainder until the next STEP cycle.  This paper is organised as follows. In \\S\\ref{lensing} we describe the shapelet decomposition and the action of weak gravitational lensing in shapelet space. In \\S\\ref{shear_estimators} we derive several possible weak shear estimators, and discuss the performance of a key estimator on the simulated STEP images. In \\S\\ref{flexion_estimators} we derive several possible weak flexion estimators. We conclude in \\S\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  We have described the mechanics of weak gravitational lensing in terms of ``polar shapelets'' (Refregier 2003; Massey \\& Refregier 2005). This is a set of basis functions that can be used to parametrize images, in a way that appears convenient for both weak shear and flexion measeurement. The symmetries of polar shapelets are well-matched to those of lensing. For example, the complex notation of polar shapelet coefficients, where their modulus represents the amount of power, and their phase represents their orientation, mirrors that commonly used in the literature to define a complex ellipticity. In addition, polar shapelets concisely encapsulate the ideas of weak flexion that have been recently developed.  The symmetries inherent to the polar shapelet formalism provide useful insight into the parallel symmetries of weak lensing distortions. We have exploited this relation to construct estimators that are able to simultaneously extract both the weak shear and flexion signal from the observed shapes of distant galaxies. We attempt to bypass some of the limitations of KSB and other shear measurement methods that were reviewed in STEP2. Adopting the classification scheme from that programme, we briefly discussed the recasting of alternative, ``active'' shear and flexion estimators into the polar shapelet formalism, and more comprehensively explored the options available for ``passive'' shear and flexion estimators.  The Gaussian-weighted shear estimator $\\tilde{\\gamma}_{\\rm Gaussian}$ recovers old methods like KSB and RRG, but in a compact complex notation. The unweighted shear estimator $\\tilde{\\gamma}_{\\rm unweighted}$ takes advantage of the noise-free shapelet reconstructions to diagonalise the shear susceptibility tensor into a scalar quantity. This particular quantity happens to be easily derivable in real space as well. The simplification of the shear susceptibility is completed with $\\tilde{\\gamma}_{\\rm shear-invariant}$. With this, the shear susceptibility is simply the object's flux: a robustly measured quantity, and one that does not change to first order during shear. The growing complexity of these shear estimators solves progressively more of the four issues highlighted with previous-generation shear measurement methods. However, they also converge more slowly, and require high-$n$ coefficients to be  available. The later estimators may consequently be accessible only on galaxy images with higher signal-to-noise. The best estimator to use (which may not even be the same for an entire population) may therefore depend on the flux of an object, or the nature of residual problems found in any particular dataset. One final, particularly interesting alternative option is the estimator $\\tilde{\\gamma}_{\\rm profile}$ that can reduce the rms shear noise due to intrinsic galaxy ellipticities, by exploting additional information about each galaxy's radial profile.  Analogous estimators for flexion also exist for most of these options.  Interestingly, our method permits several independent shear estimators and several flexion estimators to be obtained simultaneously for each galaxy. Because we calculate them following a {\\it single} PSF deconvolution or non-linear iteration, their agreement (or otherwise) will provide a stringent new test on the PSF modelling and for other residual systematic effects. Such tests are unique to our approach, as the interpretation of analogous active shear estimators would be hindered by the need to perform a separate PSF deconvolution for each estimator, and possibly removing any shared defects.  We have demonstrated the performance of one of our key shear estimators via blind tests that were part of the Shear TEsting Programme (STEP; Massey \\etal, 2007a). We shall compare the practical performance of our remaining shear estimators in a future round of STEP simulations. We are also implementing an option to input a known flexion signal in the image simulation suite of Massey \\etal\\ (2004). We are planning a smaller-scale, flexion version of STEP, to calibrate the performance of emerging weak flexion estimators, including the ones presented in this work as well as others presented elsewhere. In the mean time, a complete IDL software package capable of performing the shapelet image decomposition, including the weak lensing manipulation and analysis described in this paper, is available from the shapelets web site {\\tt http://www.astro.caltech.edu/$\\sim$rjm/shapelets}."
    },
    "0609/astro-ph0609106_arXiv.txt": {
        "abstract": "{ Lithium and zirconium abundances (the latter taken as representative for s-process enrichment) are  determined for a large sample of massive Galactic O-rich AGB stars,  for which high  resolution optical spectroscopy has been obtained (R$\\sim$40,000$-$50,000). This is done by computing synthetic spectra based on classical hydrostatic  model atmospheres for cool stars using extensive line lists.  The results obtained are discussed in  the framework of ``hot bottom burning'' (HBB) and nucleosynthesis models. The complete sample is studied attending to various observational properties such as the position of the stars  in the IRAS two-colour diagram ([12] $-$ [25] {\\it vs} [25] $-$ [60]), Galactic distribution, expansion velocity (derived from the OH maser emission) and period of variability (when  available). We conclude that a considerable fraction of the sources observed  are actually massive AGB stars ($M>3$--4 $M_{\\odot}$)  experiencing HBB, as deduced from the strong Li overabundances found. A comparison of our results with similar studies carried out in the past for the Magellanic Clouds (MCs) reveals that, in  contrast to MC AGB stars, our Galactic sample does  not show any indication of s-process element enrichment. The differences observed are explained as a consequence of metallicity effects. Finally, we discuss the  results obtained in the framework of stellar evolution by  comparing our results with the data available in the literature for Galactic post-AGB stars and PNe. ",
        "introduction": "The asymptotic giant branch (AGB) is formed by stars with initial masses in the range between 0.8 and 8 $M_{\\odot}$ in a late stage of their evolution. The internal structure of an AGB star consists of an electron-degenerate C--O core surrounded by a shell of He. During most of the time H burning is the main source of energy for the AGB star but, occasionally, the inner He shell ignites in a ``thermal pulse'' and, eventually, the byproducts of He burning may reach the outer layers of the atmosphere in a ``dredge-up'' of processed material (the so-called third dredge-up). Repeated thermal pulses add  carbon to the stellar surface.  As a consequence, originally O-rich AGB stars  can turn into C-rich AGB stars (C/O $>$ 1). This scenario  would explain the observed spectral sequence M--MS--S--SC--C in AGB stars (e.g.\\ Mowlavi 1999).  Another important characteristic of AGB stars is the presence of neutron-rich elements (s-elements such as Sr, Y, Zr, Ba, La, Nd, Tc, etc.) in their atmospheres. These species are formed by slow-neutron captures in the intershell region. After a dredge-up episode, according to theoretical models, protons can be partially mixed in the  $^{12}$C- and $^{4}$He-rich region between the hydrogen- and helium-burning shells (e.g.\\ Straniero et al.\\ 1995, 1997; Mowlavi 2002; Lattanzio 2003), and can react with $^{12}$C to form $^{13}$C during the interpulse phase. The $^{13}$C($\\alpha$,n)$^{16}$O reaction releases neutrons, which are captured by iron nuclei and other heavy  elements, forming s-elements that can later  be dredged up to the stellar surface in the next thermal pulse (Straniero et al.\\ 1995, 1997; Busso et al.\\ 2001). $^{22}$Ne is another neutron source which needs to be considered. In this case the activation takes place during the convective thermal pulses (e.g.\\ Straniero et al.\\ 1995, 2000; Vaglio et al.\\ 1999; Gallino et al.\\ 2000). $^{22}$Ne can be formed from $^{14}$N and $^{14}$N is formed by the CNO cycle in the hydrogen shell. The $^{22}$Ne($\\alpha$,n)$^{25}$Mg neutron source requires higher temperatures and occurs typically at higher neutron densities than the $^{13}$C($\\alpha$,n)$^{16}$O reaction. Thus, a different s-element pattern is expected depending on the  dominant neutron source. According to the most recent models the $^{13}$C($\\alpha$,n)$^{16}$O reaction is the preferred neutron source for masses around 1--3 $M_{\\odot}$, while for more massive stars (i.e.\\ M $\\gtrsim$ 3--4 $M_\\odot$)  neutrons are thought to be mainly released through the $^{22}$Ne($\\alpha$,n)$^{25}$Mg reaction  (see, for example, Busso, Gallino \\& Wasserburg 1999 and Lattanzio \\& Lugaro 2005 for a recent review). In the literature, there is  strong evidence that most Galactic AGB stars enriched in s-process elements have masses around 1--3 $M_{\\odot}$, where the $^{13}$C($\\alpha$,n)$^{16}$O reaction is the neutron donor (e.g.\\ Lambert et al.\\ 1995; Abia et al.\\ 2001). Unfortunately, a confrontation of the predictions made by these models with observations of more massive AGB stars in our Galaxy is not yet available.  In the case of the more massive O-rich AGB stars ($M > 3$--4 $M_\\odot$), the convective envelope can penetrate the H-burning shell activating the so-called ``hot bottom burning'' (hereafter, HBB). HBB takes place when the temperature at the base of the convective envelope is hot enough ($T \\geq 2\\times 10 ^{7}$ K) and $^{12}$C can be converted into $^{13}$C and $^{14}$N through the CN cycle (Sackmann \\& Boothroyd 1992; Wood, Bessel \\& Fox 1983). HBB models (Sackmann \\& Boothroyd 1992;  D'Antona \\& Mazzitelli 1996; Mazzitelli, D'Antona \\& Ventura 1999) also predict the production of  the short-lived $^{7}$Li by the chain $^{3}$He($\\alpha$,$\\gamma$)$^{7}$Be (e$^{-}$,$v$)$^{7}$Li, through the so-called ``$^{7}$Be transport mechanism'' (Cameron \\& Fowler 1971). One of the predictions of these models is that lithium should be detectable, at least for some time, on the stellar surface.  The activation of HBB in massive O-rich AGB stars is supported by  studies performed of AGB stars in the Magellanic Clouds (hereafter, MCs) (Plez, Smith \\& Lambert 1993; Smith \\& Lambert 1989, 1990a; Smith et al.\\ 1995) which show a lack of high luminosity C stars beyond $M_{\\rm bol} \\sim -6$. Instead, the more luminous AGB stars in the MCs are O-rich. The detection of strong Li overabundances together with strong s-element enhancement in these luminous AGB stars in the LMC and in the SMC is the signature that these stars are indeed HBB AGB stars that have undergone a series of thermal pulses and dredge-up episodes in their recent past. The HBB nature of these stars is also confirmed by the detection of a very small \\Crat\\ ratio ($\\simeq 3$--4), expected only when HBB is active (Mazzitelli, D'Antona \\& Ventura 1999).  Unfortunately, current HBB theoretical models have been tested almost  exclusively using the results obtained from the study of the more massive AGB stars in the MCs but have never been applied to Galactic sources, mainly because of the lack of observations  available but also because of the  inaccurate information on distances within our Galaxy.  Indeed, only a handful of Li-rich stars have been found in our Galaxy so far (e.g.\\ Abia et al.\\ 1991, 1993; Boffin et al.\\ 1993) and, unlike those detected in the Magellanic Clouds, they are not so luminous ($-6 \\leq M_{\\rm bol} \\leq - 3.5$). Most of them are low mass C-rich AGB stars (Abia \\& Isern 1996, 1997, 2000) and intermediate mass S- and SC-stars (Abia \\& Wallerstein 1998) and not O-rich M-type stars. A few are less luminous red giant branch (RGB) stars at the bump (Charbonnel 2005). Some of these AGB stars are among the most Li-rich stars in our Galaxy, the so-called ``super Li-rich'' AGB stars (Abia et al.\\ 1991). Under these conditions, however, HBB is not expected to be active and the Li production is not well understood. Note that HBB is expected to be active only in the most massive (and luminous) AGB stars (from $\\sim$4 to 7 $M_{\\odot}$) (Mazzitelli, D'Antona \\& Ventura 1999), which should not be C-rich, but O-rich. Actually, the best candidates in our Galaxy  are the so-called {\\it OH/IR stars},  luminous O-rich AGB stars extremely bright in the infrared, showing a characteristic double-peaked OH maser emission at 1612 MHz. These stars are also known to be very long period variables (LPVs), sometimes with periods of more than 500 days and large amplitudes of up to 2 bolometric magnitudes. However, they experience very strong mass loss rates (up to several times 10$^{-5}$ $M_{\\odot}$ yr$^{-1}$) at the end of the AGB and most of them appear heavily obscured  by thick circumstellar envelopes, making optical observations very difficult (e.g. Kastner et al. 1993; Jim\\'enez-Esteban et al. 2005a; 2006a). Thus, no information yet exists on their lithium abundances and/or possible s-process element enrichment.  In this paper we present the results obtained from a wide observational programme based on high resolution optical spectroscopy of a carefully selected sample of Galactic AGB stars thought from their observational properties to be massive. The criteria followed to select the sources in the sample are explained in Section 2. The high resolution spectroscopic observations made in the optical and the data reduction process  are described in Section 3, while the main results obtained  are presented in Section 4. In this section we also show how spectral  synthesis techniques were applied to derive the atmospheric parameters   as well as the lithium and zirconium abundances of the stars in our sample. A discussion of these results in the context of HBB models, nucleosynthesis models and stellar evolution can be found in Section 5. Finally, the conclusions derived from this work are given in Section 6. ",
        "conclusions": "As a result of a high resolution spectroscopic survey carried out in the optical for a large sample (102) of massive Galactic O-rich AGB stars, we have detected the Li I resonance line at 6708 \\AA~in 25 stars, while 32 stars did not display this Li I line. The rest of the stars observed (45 stars) were heavily obscured by their thick circumstellar envelopes and, as such, they were too red or not found at optical wavelengths.  We have classified the sample into three subgroups on the basis of their OH expansion velocity and variability period. From their relative distribution in the IRAS two--colour diagram [12] $-$ [25] vs [25] $-$ [60] and Galactic distribution, we conclude that they must represent populations of Galactic O-rich AGB stars with different progenitor masses.  Combining MARCS model atmospheres and synthetic spectroscopy with extensive line lists we have been able to derive the fundamental stellar parameters ($T_{\\rm eff}$, log $g$, C/O, $\\xi$, etc.) and obtain the Li and Zr abundances in those stars for which an optical spectrum was obtained. Our chemical abundance analysis shows that half of stars show Li overabundances in the range between {log $\\varepsilon$(Li)} $\\sim$ $+$0.5 and $+$3.0 dex. This is interpreted as a signature of the activation of HBB, confirming that they are actually massive AGB stars ($M >$ 4 $M_\\odot$). However, these stars do not show any zirconium enhancement (taken as  representative for the s-process element enrichment), indicating that they behave differently from the more massive (and luminous) AGB stars in the MCs.  Assuming that the variability period and the OH expansion velocity can be taken as distance-independent mass indicators and by comparing our results with the theoretical predictions of HBB and nucleosynthesis models we conclude that the O-rich AGB stars in our sample are all considerably massive ($M >$ 3 $M_\\odot$) even those non-detected in lithium, but only the more massive ones ($M >$ 4 $M_\\odot$; with periods longer than 400 days and {$v_{\\rm exp}$(OH)} $>$ 6 km s$^{-1}$) experience HBB. The lack of lithium enrichment found in some of the sources  thought to belong to the group of more massive AGB stars in the sample is explained as the consequence of the timescale of the Li production phase being of the order of or slightly smaller than the interpulse time ($\\sim$10$^{4}$ years) or, alternatively, by the $^{3}$He exhaustion in the envelope  interrupting the Li production, in agreement with the predictions made by the models.  The group of stars displaying the most extreme observational properties (with periods sometimes longer than 700 days and {$v_{\\rm exp}$(OH)} $>$ 12 km s$^{-1}$) are believed to represent the more massive AGB stars in our Galaxy. Unfortunately, these stars are strongly obscured by their thick circumstellar envelopes and we could not carry out any type of analysis in the optical range. This very high mass population should also experience HBB and show strong Li enhancement but no s-process element overabundances. We propose that the HBB status of these obscured stars can be determined by measuring their \\Crat\\ ratios in the near-infrared, which can also be used as a HBB indicator.  Our results suggest that the dramatically different abundance pattern found in AGB stars belonging to the MCs and to our Galaxy can be explained in terms of the different metallicity conditions under which these stars evolved. This constitutes strong observational evidence that the chemical evolution of massive AGB stars is strongly modulated by metallicity.  Globally considered, our results can also be used to establish evolutionary connections between TP-AGB stars, post-AGB stars and PNe. In particular, we suggest that the HBB AGB stars identified in our sample must be the precursors of N-rich type I PNe."
    },
    "0609/astro-ph0609330_arXiv.txt": {
        "abstract": "In the quest of understanding how star formation occurs and propagates in the low metallicity environment of the Small Magellanic Cloud (SMC), we acquired deep F555W ($\\sim$V), and F814W ($\\sim$I) HST/ACS images of the young and massive star forming region NGC 346. These images and their photometric analysis provide us with a snapshot of the star formation history of the region. We find evidence for star formation extending from $\\approx$10~Gyr in the past until $\\approx 150\\, {\\rm Myr}$ in the field of the SMC. The youngest stellar population ($\\sim 3\\pm 1\\, {\\rm Myr}$) is associated with the NGC~346 cluster. It includes a rich component of low mass pre-main sequence stars mainly concentrated in a number of sub--clusters, spatially co--located with CO clumps previously detected by Rubio et~al.\\ (2000). Within our analysis uncertainties, these sub--clusters appear coeval with each other. The most massive stars appear concentrated in the central sub--clusters, indicating possible mass segregation. A number of embedded clusters are also observed. This finding, combined with the overall wealth of dust and gas, could imply that star formation is still active. An intermediate age star cluster, BS90, formed $\\sim 4.3\\pm 0.1\\, {\\rm Gyr}$ ago, is also present in the region. Thus, this region of the SMC has supported star formation with varying levels of intensity over much of the cosmic time. ",
        "introduction": "The Small Magellanic Cloud (SMC) is an excellent laboratory to investigate the star formation (SF) processes and the associated chemical evolution in dwarf galaxies. Its current sub-solar chemical abundance (Z=0.004) makes it the best local counterpart to the large majority of dwarf irregular (dIrr) and Blue Compact Dwarf (BCD) galaxies, whose characteristics may be similar to those in the primordial universe. Because the SMC is part of a triple system, its star formation can be triggered by interactions with the Milky Way (MW) and the Large Magellanic Cloud (LMC). This kind of interactions might have been very frequent in the past, if galaxies are building up via hierarchical mergers.  The SMC proximity allows us to resolve into single stars the youngest and most compact star clusters, down to the sub solar mass regime. Young star clusters have been extensively studied in the MW and in the LMC, and a direct comparison of the characteristics of these clusters with those of the SMC allows us to address fundamental issues such as the universality of the initial mass function (IMF), and the occurrence of primordial mass segregation within star forming regions.  In this paper we present an in--depth study of the stellar content of NGC~346 ($\\alpha_{2000}=$ 00:59:05.2; $\\delta_{2000}=-$72:10:28). This star cluster is an outstanding benchmark for star formation studies. It is an extremely young, \\citep[$\\sim$3~Myr,][]{bouret03} moderately compact cluster, that excites the largest and brightest H{\\sc II} region - N66 - in the SMC \\citep{relano02}. N66 has a diameter of about $420\\arcsec$, which, at the distance of the SMC \\citep[60.6~Kpc, ][]{hilditch05} corresponds to $\\sim 123\\, {\\rm pc}$.  NGC~346 lies in a very active and interesting region: at a distance of 2~arc--minutes to the East (35~pc in projection) of the center of the cluster, we find the massive Luminous Blue Variable star HD~5980. N66 contains also at least two known supernova remnants (SNR): SNR~0057-7233 \\citep{ye91}, located to the Southwest of N66, and SNR~0056-7226. SNR~0057-7226 is found in proximity to HD~5980, although far-ultraviolet absorption studies have determined that HD~5980 is actually behind the SNR. High--resolution X--ray observations of N66 \\citep{naze02} show a $130\\arcsec \\times 130\\arcsec$ region of extended thermal emission corresponding to the SNR, with an emission peak at its center. SNR~0057--7226 has a relatively uniform surface brightness and does not show any clear temperature gradient from center to rim. More details on the properties of the SNRs and H\\,{\\sc ii} regions around NGC~346  are given by \\citet{reid06}.  Both H$\\alpha$ and O{\\sc vi} observations imply the presence of a strong shock in N66, associated with SNR~0057-7226, as the blast wave encounters the denser material of the ionized nebula \\citep{danfort03}. The dynamical properties of the region inferred from the emission O{\\sc vi} lines observed near NGC~346 suggest the existence of a super--bubble, caused by a supernova explosion \\citep{hoopes02}.  NGC~346 is relatively faint in the X--rays, with a total luminosity of $L_X\\sim1.5\\times 10^{34}\\, {\\rm erg s^{-1}}$ in the 0.3--10.0~KeV energy range. Most of this emission seems correlated with the location of the brightest stars in the core of the cluster \\citep{naze02}. Many other X-ray sources are detected in the field of N66. Optical counterparts of these sources are B--type stars, suggesting that many of them may be X--ray binaries \\citep{naze03}.  NGC~346 contains a major fraction of the O~stars known in the entire SMC \\citep{walborn78, walborn86, niemela86, massey89}. The bright end of its stellar population ($V\\la 19.5\\, {\\rm mag}$) has been well investigated in the past twenty years. Evidence of very--early spectral type stars was found \\citep{niemela86,walborn86,massey05,heap06}, indicating the presence of stars as young as $\\sim 1\\, {\\rm Myr}$, and an IMF slope ($\\Gamma = d\\log \\xi (\\log m)/d \\log m =-1.9$, where $\\xi (\\log m)$ is the number of stars born per unit logarithmic mass interval per unit area), close to the slope for massive stars observed both in the LMC and in the solar neighborhood (Massey, Parker, \\& Garmany  1989, see also Chabrier 2005 for a review). Sequential star formation, starting from the Southwest of the association towards the center of the cluster was suggested by \\citet{massey89}.  From the spectral analysis of some dwarf O-type stars, the derived mean metallicity of the cluster is $Z/Z_{\\odot}=0.2 \\pm 0.1$ \\citep{haser98, bouret03}, while the estimated age is $t\\sim 3 \\times 10^6\\, {\\rm yr}$. In a more recent work, from the analysis of 21 bright stars, \\citet{mokiem06} identified two different stellar components: though the majority of the stars are 1--3~Myr old, some of the investigated stars seem somewhat older (3--5~Myr). No relation between the age and the spatial distribution was found of the stars.  Extended CO clouds \\citep{rubio00} are found within NGC~346. Ground-based narrow-band images show evidence of dust and compact embedded clusters surrounding the central cluster, suggesting the presence of a second stellar generation possibly triggered by the outflows of the central ionizing cluster, as observed in 30~Doradus and N11 \\citep{walborn92,walborn02}.  A rich population of pre--Main Sequence (pre--MS) stars has been recently discovered by Nota et~al.\\ (2006--- hereafter Paper~I) in NGC~346. These stars have likely formed within NGC~346, $\\simeq 3-5\\, {\\rm Myr}$ ago. They are in the mass range $3-0.6\\, {\\rm M_\\odot}$, and appear mostly concentrated in the main cluster and in the knots of molecular gas identified by \\citet{rubio00}.  Our ultimate objective is to study how star formation occurred and propagated in this region, by characterizing its stellar content, luminosity and mass function. In this paper we present the photometric data, and a qualitative description of the stellar population in NGC~346. We discuss the properties and the spatial distribution of the stellar populations identified, and we propose a preliminary assessment of the star formation history in the region. The paper is organized as follows: a description of the observation and the data reduction procedure is presented in Section~2; in Section~3 we present the color--magnitude diagrams (CMDs), the description of the stellar populations identified and of their spatial distribution. In Section~4 we discuss the rich BS90 star cluster in the SMC foreground of NGC~346. The results of this paper are discussed in Section~5. ",
        "conclusions": "Our analysis of the stellar content of N66/NGC~346 provides a snapshot of the star formation history of the entire region. From the analysis of the photometric data, we can conclude that different star formation episodes took place.  The oldest stellar population observed belongs to the SMC field. The morphology of the CMD suggests that, in the field, a major episode of SF occurred approximately between 3 and 5~Gyr ago, but stars with ages up to at least 10~Gyr are also present. We noted also that at most recent epochs the SF activity in the field has been significantly lower, with a possible moderate enhancement $\\sim$150~Myr ago.  An intermediate--old age star cluster is located $\\sim$23~pc in projection from the center of NGC~346. A comparison between the CMDs of this old cluster and that of the SMC field (see Fig.~\\ref{f:rgbs}) indicates that the stars in the old association formed in an almost instantaneous burst of star formation $\\sim$$4.3\\pm 0.1\\, {\\rm Gyr}$ ago. The discovery of this intermediate--old age star cluster further supports the hypothesis suggested by \\citet{gardiner92},  \\citet{dolphin01}, and \\citet{ mccumber05} that, in the SMC, both the field and the star cluster populations have likely formed in a quasi--continuous mode.  The youngest stellar population belongs to NGC~346. Padua isochrones indicate that NGC~346 is $3\\pm 1\\, {\\rm Myr}$ old. The MS of the youngest population abruptly interrupts at ${\\rm F555W}\\simeq 21$. At ${\\rm F555W}>21$ we identify hundreds of stars whose colors and magnitude are consistent with those predicted for pre--MS stars (see also Paper~I).  The high spatial resolution of our observations reveals that the youngest stellar population is not uniformly distributed within the ionized nebula: we identified at least 15 Sub--clusters, which differ in size and stellar content. Within the uncertainties due to the comparison with isochrones, the sub--clusters are likely coeval with each other (see Tab.~\\ref{t:sc}). However a relatively older sub--cluster (Sc--16, with an age $15\\pm 2.5\\, {\\rm Myr}$) is located at the Northeast periphery of our data. This sub--cluster is likely not related to the star forming episode that originated NGC~346.  Recent spectroscopic investigations of NGC~346 \\citep{heap06,mokiem06} confirm that the majority of the stars associated with the cluster have an age of 3~Myr, even if stars as young as 1~Myr are also present \\citep[but see the discussion in ][]{massey05}. No relationship between the age of the stars and their distance from the cluster center is identified. Due to the age uncertainties quoted by the authors and the lower precision of the isochrones comparison technique, our results are in good agreement.  The sub--clusters coincide with the clumps of molecular gas identified by \\citet{rubio00}. Many of them, especially towards the Southeast, are close to obscure Bok--globules, which can contain still embedded Class~I and 0 young stellar objects that would not be visible in the wavelength range covered by our data. The presence of diffuse gas and dust around these sub--clusters suggests that NGC~346 has likely not exhausted its fuel and star formation is possibly still ongoing.  All along the dust ``ridge'' of N66 it is possible to observe many dense nebular knots and dust pillars oriented toward the central clusters, that are similar in shape to those identified in 30~Doradus by \\citet{walborn02}. The presence of these structures strongly supports the hypothesis that the stellar winds, developed from the most massive stars in NGC~346, are probably triggering a secondary episode of star formation at the periphery of N66.  The most massive stars are located in the sub--clusters in the center of the nebula. We also find that sub--clusters located toward the North--Northeast are generally smaller and constituted mainly by pre--MS stars. These clusters coincide with the highest CO velocity clumps analyzed by \\citet{rubio00}. This seems to indicate that the UV radiation and the stellar winds from the stars in Sc--1, 2, and 3 are not isotropic, but they are more efficient in disrupting the molecular cloud toward the Northwest, making the star formation less efficient than toward the Southwest. Alternatively, it is possible that the interaction with SNR 0057--7233 is preventing the normal expansion of N66 along this direction.  NGC~346 is probably the result of the collapse and subsequent fragmentation of the initial giant molecular cloud in multiple ``seeds'' of star formation. The fact that the observed sub--clusters appear almost coeval, and that many of them appear to be connected by arcs of dust and gas, seems to be a good observational match to the conditions predicted by the hierarchical fragmentation of a turbulent molecular cloud model \\citep{klessen00,bonnell02,bonnell03}.  According to this model the fragmentation of the cloud is due to supersonic turbulent motions present in the gas. The turbulence induces the formation of shocks in the gas, and produce filamentary structures \\citep{bate03}. The chaotic nature of the turbulence increases locally the density in the filamentary structures. When regions of high density become self--gravitating, they start to collapse to form stars. Simulations show that star formation occur simultaneously at several different location in the cloud \\citep{bonnell03}, as appears to be the case for NGC~346.  Further UV and IR photometric and spectroscopic observations are necessary to better establish the mass and the ages of the stars, to verify if small differences in age are present among the sub--clusters, and to confirm the presence of ongoing star formation within N66. Furthermore, an analysis of the gas dynamics within N66 can confirm if the parent cloud underwent a hierarchical fragmentation: the formation of shocks in the gas, due to the initial supersonic turbulence, rapidly removes kinetic energy from the gas \\citep{ostriker01}, and we expect to find very low gas velocities around the sub--clusters."
    },
    "0609/astro-ph0609610_arXiv.txt": {
        "abstract": "The asymptotic giant branch (AGB) phase of stellar evolution is common to most stars of low and intermediate mass. Most of the carbon and nitrogen in the Universe is produced by AGB stars. The final fate of the AGB envelopes are represented by planetary nebulae (PN). By studying PN abundances and compare them with the yields of stellar evolution is possible to quantify carbon and nitrogen production, and to study cosmic recycling in galactic and Magellanic Cloud populations. In this paper we present the latest results in PN chemical abundance analysis and their implication to the chemical evolution of the galaxy and the Magellanic Clouds, with particular attention to carbon abundance, available only thanks to ultraviolet spectroscopy. ",
        "introduction": "Low- and intermediate-mass stars (M$_{\\rm MS} \\sim 0.8 - 8$ M$_{\\odot}$, where M$_{\\rm MS}$ is the main sequence mass) and their ejecta are excellent probes of stellar populations in galaxies, and have been detected in all types of galaxies as well as in the intracluster space. Stars in this mass range constitute a major component (by mass) of the stellar material in the Universe, thus a correct understanding of their evolution in different environments has the potential to advance many astrophysical fields. Furthermore, these stars play a fundamental role in cosmic recycling, being major contributors to the carbon and nitrogen abundances for the next generation of stars. Stars in this mass range go through the asymptotic giant branch (AGB) phase, and through their evolution they enrich the interstellar medium with helium, carbon, nitrogen, and other elements. Typically the abundances of oxygen, argon, and neon are unaffected during AGB evolution in most stars, thus these elements can be used to probe the environment at the time of progenitor formation.   The AGB phase is characterized by the periodicity of nuclear burning phases and dredge-up phases, and, as a result, the products of stellar evolution are carried to the stellar outer layers. Planetary nebulae (PN) are the stellar envelopes of AGB stars, ejected at the tip of the AGB, carrying the products of stellar evolution to the interstellar medium. Thus by analyzing PN abundances one can constraint the AGB evolutionary models, and the environment at the time of the formation of the AGB progenitors. Since the models of stellar evolution predict that the chemical yields form AGB stars depend on initial stellar mass and metallicity, large samples of PN and in different metallicity environments should be compared with evolutionary models to obtain the best constraints. Since PN are ubiquitous and bright, and easily recognizable due to their unique spectra, they are ideally suited to probe AGB evolution and populations in nearby galaxies. In this paper we present the state of the art of planetary nebulae abundance analysis in the Galaxy and the Magellanic Clouds, in particular the impact of model constraining through carbon abundance, as derived from ultraviolet spectroscopy. We review large and homogeneous data sets that we used for our analysis, both for the Galaxy ($\\S$2) and the Magellanic Clouds ($\\S$3). We also present the comparison with the corresponding stellar evolution models. A look at the galactic disk population and the gradients of elements in the galactic disk is shown in $\\S$4. In $\\S$5 we give a summary of the results, and describe the future challenges of this field.    \\begin{figure}[ht!] \\centerline{\\includegraphics[width=28pc]{fig1_prague.epsi}} \\caption{Log (N/O) versus log (O/H) for galactic disk PN. Open symbols: non type I PN; filled symbols: type I PN. Lines: models [7] for z=0.0126, 0.0159, 0.02, 0.0252, and 0.00317 respectively from right to left.} \\end{figure} ",
        "conclusions": "Stellar evolution of low- and intermediate-mass stars is well probed by studying AGB stars and their remnants. In this paper we have shown how diagnostic abundance plots of PN can be used to constraint stellar evolution through the AGB. Carbon abundances offer a further challenge, since the carbon emission lines correspond to the UV part of the spectra. The Magellanic Cloud PN allow the comparison with stellar evolution models at lower metallicities. We have shown that the bulge oxygen and nitrogen abundances are higher, on average, than those of the galactic disk, suggesting the existence of two distinct stellar populations. We have also shown that the galactic disk population presents a very large spread in environment metallicity (or metallicity at the formation of the progenitors), and that galactic disk PN have an effective ON cycle, especially at low metallicity. This result is somehow strengthened by the analysis of LMC and SMC PN, where the destruction of oxygen in favor of nitrogen appears clear for SMC PN of type I. The N/O versus C/O (and versus He/H) diagnostic plots indicate that the type I PN are the remnant of the evolution of more massive AGB stars. For the Magellanic Cloud PN the diagrams are sparsely populated, and new data of carbon abundances are essential for model constraining at low metallicity. Galactic gradients of oxygen and neon have been found to be shallow, in agreement with other recent and independent PN gradient studies.  The major future challenge in the field of AGB and post-AGB evolution is twofold: on the theoretical side, models should take into account the fraction of AGB stars that are within a binary system, and account for yields of binary evolution; on the observational side, the acquisition of larger, homogeneous sets of UV spectra to determine carbon and neon abundances in galactic and local group PN would be essential for progress in this field. The study of gradients could be extended also to other local group galaxies. Finally, PN have been observed also outside galaxies, in the intracluster medium. A better characterization of these extragalactic sources would be essential to constraint the mass of dark halos in galaxies [18], and to determine whether the evolution of these peculiar PN is different [19] from that of galactic PN.        \\begin{figure}[ht!] \\centerline{\\includegraphics[width=28pc]{fig4_prague.epsi}} \\caption{Galactic disk PN metallicity gradients, for oxygen and neon.} \\end{figure}    \\noindent {\\it Acknowledgements: Thanks to the Organizers for a very interesting meeting, and to Katia Cunha, Eva Villaver, Bruce Balick and Dick Shaw for their contribution to the science presented here.}"
    },
    "0609/astro-ph0609426_arXiv.txt": {
        "abstract": "LSB galaxies have low metallicities, diffuse stellar disks, and massive HI disks. We have detected molecular gas in two giant LSB galaxies, UGC~6614 and F568-6. A millimeter continuum source has been detected in UGC~6614 as well. At centimeter wavelengths we have detected and mapped the continuum emission from the giant LSB galaxy 1300+0144. The emission is extended about the nucleus and is most likely originating from the AGN in the galaxy. The HI gas distribution and velocity field in 1300+0144 was also mapped. The HI disk extends well beyond the optical disk and appears lopsided in the intensity maps. ",
        "introduction": "We searched for CO(1--0) emission from seven LSB galaxies using the BIMA array. Three of the galaxies are large spirals; UGC~5709, UGC~6614 and F568-6 (Malin~2). The remaining four are relatively smaller galaxies with irregular disks; they are NGC~5585, UGC~4115, UGC~5209 and F583-1. The BIMA maps suggested the presence of CO(1--0) emission in two galaxies, UGC~6614 and F568-6. Using the IRAM 30m telescope we detected molecular gas in the disks of both galaxies and in the nucleus of F568-6. These observations show that molecular gas in LSB galaxies is present in their disks as well as their nuclei, but the distribution is localized over isolated regions and hence difficult to detect. We have also detected millimeter continuum emission from the nucleus of UGC~6614. The emission has a flat spectrum in the 1.4 to 110~GHz range and is from the AGN in the galaxy \\cite{Das06}.  \\vspace{-5mm} ",
        "conclusions": ""
    },
    "0609/astro-ph0609083_arXiv.txt": {
        "abstract": "We present an observation of the very late thermal pulse object V4334~Sgr (Sakurai's Object) with the Infrared Spectrometer (IRS) on the Spitzer Space Telescope. The emission from 5-38\\mic\\ is dominated by the still-cooling dust shell. A number of features are seen in absorption against the dust shell, which we attribute to HCN and polyyne molecules. We use these features to determine the $^{12}$C/$^{13}$C ratio for the absorbing gas to be $\\sim3.2^{+3.2}_{-1.6}$; this implies that, despite the H-content of the molecules, the hydrocarbon-bearing gas must have originated in material produced in the very late thermal pulse. We see no evidence of emission lines, despite the recently-reported optical and radio observations that suggest the effective temperature of the stellar remnant is rising. ",
        "introduction": "\\begin{figure*} \\setlength{\\unitlength}{1cm} \\begin{center} \\leavevmode \\begin{picture}(5.0,4.5) \\put(0.0,4.0){\\special{psfile=sak_mn1.eps hoffset=-170 voffset=38 hscale=28 vscale=28 angle=-90}} \\put(0.0,4.0){\\special{psfile=sak_mn2.eps hoffset=70 voffset=38 hscale=28 vscale=28 angle=-90}} \\end{picture} \\caption[]{{(a)}: The Spitzer IRS spectrum of Sakurai's Object in 2005. The curve is a {\\sc dusty} fit to the data; see text for details. {(b):} Detail in the spectrum in the wavelength range 12.5-16.5\\mic, plotted as $\\lambda^2\\,f_\\lambda$ to emphasize the depression (which is indicated by horizontal line) around 13\\mic; much of the scatter is real, arising in ro-vibrational bands of the identified molecules.} \\end{center} \\label{SED} \\end{figure*}  V4334~Sgr (Sakurai's Object) is a low-mass ($\\sim0.5$\\Msun) star that is retracing its post-Asymptotic Giant Branch (AGB) evolution along the Hertzsprung-Russell diagram following a very late thermal pulse (VLTP) \\cite{herwig}. The star was discovered in early 1996 \\cite{nakano}, and was found to be at the centre of a faint planetary nebula (PN) of diameter $40''$ (e.g. Pollacco 1999). It first appeared as an F supergiant, possibly with a hot dust shell \\cite{duerbeck96}.  Analysis of the optical spectrum for mid-1996 by Asplund et al. \\shortcite{asplund99} revealed that Sakurai's Object had undergone a transformation to a hydrogen-deficient, carbon-rich star and that the elemental abundances of a range of species were close to those found in the R~CrB stars. They also used the C$_2$ (1-0) and (0-1) Swan bands at 4740\\AA\\ and 5635\\AA\\ respectively to determine that the $^{12}$C/$^{13}$C ratio was in the range 1.5-5. Pavlenko et al. \\shortcite{pavlenko04} used the first-overtone CO bands to determine that the $^{12}$C/$^{13}$C ratio was $4\\pm1$. This is substantially less than the value in (for example) the Solar System ($^{12}$C/$^{13}$C = 70; Mathis 2000) and likely arises from second stage CNO cycling following He burning in the VLTP \\cite{asplund99}.  In 1998, Sakurai's Object was completely obscured by an optically thick dust shell from which (as of December 2005) it has still not emerged. Indeed analysis of the infrared (IR) spectrum showed that dust emission has been making an increasing contribution to the IR emission since 1997 April \\cite{pavlenko02}. Since the major dust event of 1998, the 1-5\\mic\\ spectrum is black body-like and is that of the ejected dust shell. Analysis of the post-1998 IR observations has shown that the dust is carbon, primarily in amorphous form \\cite{tyne1,eyres}. The 1-5\\mic\\ spectra during the period 1999-2001 were consistent with massloss that increased from $\\sim\\Mdot{0.5}{5}$ to $\\sim\\Mdot{1.2}{5}$ for a gas-to-dust ratio 200, appropriate for solar abundances (note that a lower gas-to-dust ratio of 75 was suggested by Evans et al. \\shortcite{evans-jcmt} for the hydrogen-deficient, carbon-rich wind of Sakurai's Object). The maximum grain radius (in a grain size distribution $n(a)\\,da \\propto a^{-q}\\,da$, with $q\\simeq3$) increased by a factor $\\sim3$ \\cite{tyne1}.  Sub-millimetre observations \\cite{evans-jcmt} show that the flux density at 450\\mic\\ and 850\\mic\\ was still rising in 2003, indicating that the massloss was being maintained; the spectral energy distribution (SED) resembled a black-body at 360~K.  Observations of \\fion{N}{ii}$\\lambda6548,6583$ and \\fion{O}{ii}$\\lambda7325$ in the spectrum of Sakurai's Object in 2002 \\cite{kerber} hinted that the stellar temperature had increased from $\\sim5\\,500$~K in 1997 to $\\gtsimeq20\\,000$~K. The reheating seems to be confirmed by recent observations at 8.6~GHz with the VLA \\cite{hajduk}, which detected free-free emission from the newly-ionized gas.  We describe here observations of Sakurai's Object with the Infrared Spectrometer (IRS) \\cite{houck} on the Spitzer Space Telescope \\cite{werner}.  \\vspace{-5mm} ",
        "conclusions": "We have reported an observation of V4334~Sgr (Sakurai's Object) with the Spitzer IRS. We find that the IR SED is dominated by the optically thick dust shell that still enshrouds the stellar remnant. The dust shell ``photosphere'' has a temperature $\\sim200$~K, while the temperature of the inner dust shell is inferred to be $\\sim400$~K. There is evidence that the condensation of new dust has ceased, and that the character of the dust has changed since 2001.  There are several prominent absorption features, which we attribute to HCN, acetylene and polyynes. The $^{12}$C/$^{13}$C ratio is $\\sim3$; this is substantially smaller than the Solar System value ($\\sim70$) and is consistent with (a)~the CNO cycling following the VLTP, and (b)~the value obtained from the IR first overtone CO bands.  We see no evidence for the presence of emission lines (such as \\fion{Ne}{ii}~$\\lambda12.8$\\mic, \\fion{Ne}{iii}~$\\lambda15.5$\\mic) that would be a signature of the anticipated increased $T_{\\rm eff}$ of the star."
    },
    "0609/astro-ph0609560_arXiv.txt": {
        "abstract": "The Centaurus cluster ($z=0.0104$) was observed with the X-ray Imaging Spectrometer (XIS) onboard the Suzaku X-ray satellite in three pointings, one centered on the cluster core and the other two offset by $\\pm 8\\arcmin$ in declination. To search for possible bulk motions of the intracluster medium, the central energy of He-like Fe-K line (at a rest-frame energy of 6.7~keV) was examined to look for a positional dependence. Over spatial scales of 50~kpc to 140~kpc around the cluster core, the central line energy was found to be constant within the calibration error of 15~eV.  The 90\\% upper limit on the line-of-sight velocity difference is $|\\Delta v|<1400~{\\rm km\\,s^{-1}}$, giving a tighter constraint than previous measurements.  The significant velocity gradients inferred from a previous Chandra study were not detected between two pairs of rectangular regions near the cluster core. These results suggest that the bulk velocity does not largely exceed the thermal velocity of the gas in the central region of the Centaurus cluster.  The mean redshift of the intracluster medium was determined to be $0.0097$, in agreement with the optical redshift of the cluster within the calibration uncertainty.  Implications of the present results for the estimation of the cluster mass are briefly discussed. ",
        "introduction": "Clusters of galaxies are the most massive  gravitationally bound structures in the Universe, and are thought to grow into larger systems through complex interactions between smaller systems. Signatures of such merging events may manifest themselves in non-Gaussian velocity distributions of member galaxies, temperature and density inhomogeneities in the intracluster medium (ICM), and bulk motions of the ICM. Indeed, numerical simulations predict the existence of bulk ICM flows with a substantial fraction of the virial velocity ($\\sim 1000~{\\rm km\\,s^{-1}}$ for rich clusters), lasting several Gyr after each merger event (e.g., \\cite{Roettinger_etal_1996, Norman_Bryan_1999}). Therefore, measurements of the ICM velocity structure provide very important information with which to understand the cluster dynamics.  If the ICM has a significant bulk velocity compared to its thermal velocity, the associated non-thermal pressure would endanger the assumption of hydrostatic ICM equilibrium, which is usually assumed in deriving the total gravitating mass in a cluster.  For example, the factor 2--3 discrepancy between the X-ray and lensing mass estimations in some objects (e.g., \\cite{Ota_etal_2004, Hattori_etal_1999}) could be due to this effect.  Therefore, the presence/absence of ICM bulk motion has a great impact on the cluster studies, including their mass estimates and cosmological applications.  We are thus encouraged to observationally constrain the ICM motion. Such an attempt may be carried out by examining X-ray emission lines from heavy elements for Doppler shifts and broadenings (e.g., \\cite{Dupke_Bregman_2001}), or by utilizing kinetic Sunyaev-Zel'dovich effects in the radio wavebands (e.g., \\cite{Sunyaev_etal_2003}). For the former method, X-ray spectroscopy of the emission lines from heavy ions, particularly the $6.7$~keV K-lines from Helium-like iron atoms, provides the most sensitive way of studying such ICM motions.  By observing X-ray spectra in the 0.5--8.5 keV range from the core regions of the Centaurus cluster (Abell~3526) with the Chandra ACIS, \\citet{Dupke_Bregman_2006} recently reported on a large velocity gradient of $2400\\pm1000~{\\rm km\\,s^{-1}}$ over a spatial scale of 100~kpc, reconfirming their previous ASCA measurements \\citep{Dupke_Bregman_2001}.  However, their results might still be subject to residual uncertainties arising from temporal and intrachip gain variability, requiring an independent cross confirmation.  Since a bulk velocity of $1000~{\\rm km\\,s^{-1}}$ along the line of sight translates to a shift in the He-like Fe-K line energy of $\\Delta E = 22$~eV, we need not only sufficient photon statistics and good energy resolution, but also precise instrumental gain calibrations with an accuracy better than $\\sim 20$ eV (or $\\sim 0.3\\%$).  The Centaurus cluster (Abell~3526), at the redshift of $z_{\\rm cl}=0.0104$, is a nearby X-ray bright cluster of galaxies, hosting the cD galaxy NGC~4696.  It has been extensively studied at various wavelengths.  \\citet{Lucey_etal_1986} found a bimodal velocity distribution of galaxies, corresponding to the main cluster (Cen~30) centered on NGC~4696, and a subcluster (Cen~45) centered on NGC~4709, with a velocity difference of $\\sim 1500~{\\rm km\\,s^{-1}}$.  They hence argued that Cen~30 is currently accreting Cen~45.  The X-ray measured ICM temperature is $\\sim 3.8$~keV in outer regions, and decreases toward the center to $\\sim 1.7$~keV (e.g., \\cite{Ikebe_etal_1999, Makishima_etal_2001, Takahashi_etal_2004}). Although the central part of this cluster is relatively relaxed (e.g., \\cite{Allen_Fabian_1994}), at larger radii the temperature maps derived with ASCA reveal a hot spot coincident with Cen~45 \\citep{Churazov_etal_1999, Furusho_etal_2001}, supporting the idea that there is a past or on-going merger episode.  The central region is very rich in metals, with a large abundance gradient \\citep{Fukazawa_etal_1994, Ikebe_etal_1999, Graham_etal_2006, Sanders_Fabian_2006}.  Emitting very intense Fe-K emission lines, the central region of the Centaurus cluster provides an excellent opportunity for a study of ICM bulk-motion.  In this paper, we present an analysis of the bulk velocity of the gas over the central 240 kpc of the Centaurus cluster, conducted with the X-ray Imaging Spectrometer (XIS; \\cite{Koyama_etal_2006}) on board the Suzaku satellite \\citep{Mitsuda_etal_2006}.  The XIS has an excellent sensitivity at the iron K-line energies, and a very low background. In addition, the current accuracy of the XIS in-orbit calibration is already high enough to enable us to make meaningful searches for possible ICM bulk motions.  The Suzaku X-ray telescope (XRT; \\cite{Serlemitsos_etal_2006}) has a spatial resolution of $\\sim 2\\arcmin$ (Half Power Diameter; HPD), allowing us to study spatial velocity variations on reasonable scales ($\\sim$a few arcminutes). We therefore test the previously claimed large velocity gradient near the cluster center.  Throughout this paper we adopt $\\Omega_{\\rm M}=0.3$, $\\Omega_{\\Lambda}=0.7$, and $H_0=70~{\\rm km\\, s^{-1}\\,Mpc^{-1}}$; $1\\arcmin$ corresponds to 13.5~kpc at the cluster redshift of $z_{\\rm cl}=0.0104$.  The quoted statistical errors refer to 68\\% confidence ranges, unless otherwise stated. ",
        "conclusions": "From the careful analysis of the high-quality Fe-K line data obtained with the Suzaku XIS, we found no evidence of significant spatial gradients in the Fe-K line center energies inside the XIS fields of view.  The derived upper limits apparently exclude the large velocity difference previously claimed by \\citet{Dupke_Bregman_2006}.  In \\S\\ref{subsec:comparison_chandra}, we further investigate this issue by analyzing the same regions as studied by \\citet{Dupke_Bregman_2006}.  In \\S\\ref{subsec:implications}, implications of the present results are discussed, followed by \\S\\ref{subsec:turbulent} where we constrain Doppler broadening of the Fe-K lines in search for turbulent motions of the ICM.   \\subsection{Comparisons with previous results} \\label{subsec:comparison_chandra}  A large velocity difference by $\\sim 2500~{\\rm km\\,s^{-1}}$ was reported in the core of the Centaurus cluster based on the previous ASCA and Chandra observations \\citep{Dupke_Bregman_2001, Dupke_Bregman_2005, Dupke_Bregman_2006}.  Specifically, by comparing the redshifts measured with the Chandra/ACIS-S3 spectra in various regions near the cluster center, \\citet{Dupke_Bregman_2006} obtained (1) a maximum velocity gradient of $\\Delta v=2900\\pm700~{\\rm km\\,s^{-1}}$ between a pair of rectangular regions, $2\\arcmin.3\\times2\\arcmin.7$ (PEAK-) and $1\\arcmin.5\\times2\\arcmin.5$ (PEAK+), and (2) another large difference by $\\Delta v=2400\\pm1000~{\\rm km\\,s^{-1}}$ between two $2\\arcmin.4\\times 3\\arcmin$ regions (R-2 and R+3).  We show these four regions on our figure~\\ref{fig2}.  Since the report by \\citet{Dupke_Bregman_2006} apparently disagrees with our results, we examined the present Suzaku data for the above two cases, by accumulating spectra from the same sky regions on the three FI chips.  (As seen in figure~\\ref{fig2}, the PEAK- and PEAK+ regions are located near cells \\#35 and \\#18, and R-2 and R+3 are near cells \\#43 and \\#21 in our analysis.)  After co-adding the central and offset data, the total exposure time is 183~ks for PEAK- and R-2, and 188~ks for PEAK+ and R+3; these are about 5 times longer than the Chandra exposure.  \\begin{table*} \\begin{center} \\caption{Comparisons of velocity differences measured with Suzaku and Chandra} \\label{tab2} \\begin{tabular}{llll}\\hline\\hline Region           & Suzaku/XIS & Chandra/ACIS \\\\ \\cline{2-2} \\cline{3-3} & $\\Delta v$~(${\\rm km\\,s^{-1}}$) \\footnotemark[*] & $\\Delta v$~(${\\rm km\\,s^{-1}}$)\\footnotemark[$\\dagger$] \\\\ \\hline PEAK-, PEAK+ &  $-660 \\pm 390\\,(\\pm 660)$ & $2900\\pm700$\\\\ R-2, R+3            &  $-540 \\pm 360\\,(\\pm 660)$ & $2400\\pm1000$ \\\\ \\hline \\end{tabular} \\end{center} \\footnotemark[*] The velocity difference derived from the XIS spectra. The $68$\\% statistical errors and (the 68\\% systematic errors) are quoted.  \\footnotemark[$\\dagger$] The velocity difference and the $1\\sigma$ error derived from the Chandra ACIS spectra \\citep{Dupke_Bregman_2006}. \\end{table*}  As to (1), the model fitting to the XIS spectra has given $z_{-}=0.0101\\pm0.0005\\,(\\pm 0.0015)$, and $z_{+}=0.0079\\pm0.0012\\,(\\pm 0.0015)$, resulting in a negative value for the redshift difference $\\Delta z \\equiv z_+ - z_- = -0.0022 \\pm 0.0013\\,(\\pm 0.0022)$, or the velocity difference, $\\Delta v = c\\Delta z= -660 \\pm 390 \\,(\\pm 660)~{\\rm km\\,s^{-1}}$.  Because PEAK- and PEAK+ are close to each other, the intrachip gain difference between them is considered to be smaller than our canonical value, $\\pm 470~{\\rm km\\,s^{-1}} $.  Nevertheless, to be conservative, we multiplied this by $\\sqrt{2}$, to obtain $\\pm 660~{\\rm km\\,s^{-1}}$ as our systematic error on the velocity difference.  The above results then imply that the velocity difference between PEAK- and PEAK+ does not significantly exceed the calibration uncertainty of $\\pm 660~{\\rm km\\,s^{-1}}$.   If we further sum the systematic and statistical errors in quadrature, the 90\\% confidence range becomes $-1890~{\\rm km\\,s^{-1}} < \\Delta v < +570~{\\rm km\\,s^{-1}}$, which is inconsistent with the Chandra measurement of $2900\\pm 700~{\\rm km\\,s^{-1}}$.  The results are summarized in table~\\ref{tab2}.  As to (2), a similar analysis of the Suzaku XIS data has given the redshifts of the R-2 and R+3 regions as $z_{\\rm R-2}=0.0107\\pm0.0006\\,(\\pm 0.0015)$ and $z_{\\rm R+3}=0.0089\\pm0.0010\\,(\\pm 0.0015)$, respectively.  Therefore, the redshift and velocity differences become $\\Delta z = z_{\\rm R+3}-z_{\\rm R-2} = -0.0018\\pm 0.0012\\,(\\pm 0.0022)$ and $\\Delta v= -540 \\pm 360 \\,(\\pm 660)~{\\rm km\\,s^{-1}}$, respectively (table~\\ref{tab2}).  The allowed 90\\% range then becomes $-1730~{\\rm km\\,s^{-1}} < \\Delta v < +650~{\\rm km\\,s^{-1}}$, which again excludes the Chandra measurement of $\\Delta v$ between R-2 and R+3.   Then, can we reconcile the apparently conflicting Chandra (plus ASCA) and Suzaku results?  One possibility is that the much wider PSF of the Suzaku XRT, compared to that of Chandra, has diluted the velocity gradient.  To examine the effect, we again executed the {\\tt xissim} simulation of the XIS spectra for PEAK- and PEAK+, simply assuming that the gas bulk motion exists only in those two regions and that they have the velocity difference close to the 90\\% lower limit of the Chandra result, $\\Delta v=+1700~{\\rm km\\,s^{-1}}$; other regions in the XIS field of view are assumed to have no bulk motion, with $z=0.0104$.  The model fit to the two simulated spectra indicates that the Suzaku XIS should observe a velocity difference of $\\Delta v=+690~{\\rm km\\,s^{-1}}$, as a result of the telescope resolution which reduce the velocity difference by a factor of $\\sim 2.5$.  Therefore, the Chandra measurement may be brought to be marginally consistent with the present Suzaku upper limit.  As to (2), the same simulation indicates that the 90\\% lower-limit value of $\\Delta v=+760~{\\rm km\\,s^{-1}}$ between R-2 and R+3 based on the Chandra report will be reduced to $\\Delta v =+300~{\\rm km\\,s^{-1}}$ in the XIS data after smeared by the PSF. This is consistent with the 90\\% upper limit set by Suzaku, $+650~{\\rm km\\,s^{-1}}$.  In this way, the positive $\\Delta v$ reported by \\citet{Dupke_Bregman_2006} for the two pairs of regions can be tentatively made consistent with the negative $\\Delta v$ we measured. However, this was possible by invoking rather extreme assumptions; namely, to employ the Chandra lower limit and the Suzaku upper limit on $\\Delta v$, and to assume that the ICM bulk motion is localized to the small regions of $\\sim 25$~kpc.  A comparison of the present results with the reported ASCA measurement \\citep{Dupke_Bregman_2001} may give additional clues.  The Suzaku XRT \\citep{Serlemitsos_etal_2006} is an improved version over its predecessor onboard ASCA.  Similarly, the Suzaku XIS is an improved instrument based on the SIS experiment onboard ASCA.  As compared with the XRT+SIS combination onboard ASCA, the XRT+XIS of Suzaku hence has a $\\sim1.5$ times better PSF, a $\\sim4.5$ times larger effective area (at 7 keV; two SIS sensors versus three XIS cameras).  Furthermore, the energy scale of the Suzaku XIS has been better calibrated \\citep{Koyama_etal_2006} than the ASCA SIS which was the first single-photon detection CCD used in cosmic X-ray observations. Therefore, any Doppler shifts in the Fe-K line detectable by ASCA would be detected by Suzaku with a much higher significance (unless the exposure is too short).  From these considerations, we conclude that it is rather difficult to reconcile the the present results with the Chandra (plus ASCA) measurements.  Additional information will be obtained by analyzing a longer Chandra dataset on the Centaurus cluster. \\citep{Fabian_etal_2005}. \\footnote{Some of the present authors actually analyzed this data set, but did not detect any significant velocity differences between the R-2 and R+3 regions identified by \\citet{Dupke_Bregman_2006}.}  \\subsection{Implications of the derived upper limits} \\label{subsec:implications}  Utilizing the Fe-K$\\alpha$ line, we obtained a 90\\% upper limit on the velocity difference of $|\\Delta v| < 1400~{\\rm km\\,s^{-1}}$ between any pair of regions within the XIS field of view, separated typically by $\\sim 50$~kpc or more.  This gives a tighter constraint on the ICM bulk motion than was suggested by previous reports.  Since the ion sound velocity in the ICM is $s\\equiv(\\gamma P/\\rho)^{1/2}=(\\gamma kT/\\mu m_p)^{1/2} \\sim 880 ~{\\rm km\\,s^{-1}}$ for a temperature of $kT=3$~keV, our results imply that the line-of-sight velocity of the ICM bulk motion, if any, does not largely exceed the sound velocity. In other words, the ram pressure associated with the ICM bulk motion is at most comparable to the thermal ICM pressure, thus approximately validating the assumption of hydrostatic equilibrium in calculating the cluster mass.  Of course, we can probe the gas motion only along the line of sight, and hence we need to assume 3-dimensional velocity structures in order to quantitatively utilize the present upper limits.  As an elementary exercise, below we evaluate the upper limit on rotational motion assuming for simplicity that the gas is rigidly rotating at a typical circular velocity of $\\sigma_{r} \\sim |\\Delta v|/2$, and derive its quantitative effect on cluster mass estimation. Assuming an approximate spherical symmetry, and including the centrifugal force, the balance against the gravitational pull at a radius $r$ on the rotational equatorial plane then becomes \\begin{equation} -\\frac{GM(r)\\rho_{\\rm gas}}{r^2} = \\frac{\\partial P_{\\rm gas} }{\\partial r} - \\frac{f\\rho_{\\rm gas}\\sigma_r^2}{r}, \\label{eq1} \\end{equation} where $M(r)$ is the cluster mass interior to radius $r$, $\\rho_{\\rm gas}$ is the gas density, $P_{\\rm gas}=\\rho_{\\rm gas} kT (\\mu m_p)^{-1}$ is the ICM pressure, and $f$ is the fraction of the ICM that is rotating ($0\\leq f \\leq1$).  The gas density profile of the Centaurus cluster roughly follows $\\rho_{\\rm gas}\\propto r^{-1}$ within the central $r\\lesssim 10\\arcmin$ region, assuming the double-$\\beta$ model \\citep{Ikebe_etal_1999, Mohr_etal_1999}. Together with the assumption $\\sigma_r \\propto r$, this yields \\begin{equation} f\\rho_{\\rm gas}\\sigma_r^2 \\propto r, \\end{equation} and hence the second term on the right-hand side of equation~\\ref{eq1} can be rewritten as \\begin{equation} \\frac{f\\rho_{\\rm gas}\\sigma_r^2}{r} = \\frac{\\partial (f \\rho_{\\rm gas}\\sigma_r^2)}{\\partial r}. \\end{equation} Equation~\\ref{eq1} then becomes \\begin{eqnarray} -\\frac{GM(r)}{r^2} & = & \\frac{1}{\\rho_{\\rm gas}} \\frac{\\partial P_{\\rm gas} }{\\partial r} + \\frac{1}{\\rho_{\\rm gas}} \\frac{\\partial (f \\rho_{\\rm gas}\\sigma_r^2)}{\\partial r}\\nonumber \\\\ & =& \\frac{1}{\\rho_{\\rm gas}} \\frac{\\partial}{\\partial r}P_{\\rm gas} (1+f \\beta_r), \\end{eqnarray} where \\begin{equation} \\beta_r \\equiv \\frac{\\mu m_p \\sigma_r^2}{ kT} \\sim 1.07 \\left(\\frac{\\mu}{0.63}\\right) \\left(\\frac{\\sigma_r}{700~{\\rm km\\,s^{-1}}}\\right)^2 \\left(\\frac{kT}{3~{\\rm keV}}\\right)^{-1} \\label{eq4} \\end{equation} is the ratio of the kinetic energy due to rotational motion to the thermal energy.  Thus, the total mass should be higher than the hydrostatic mass by a factor of $(1+ f\\beta_r)$, because of the presence of additional pressure support.  At a typical radius of $\\sim 100$~kpc, the rotational velocity is $\\sigma_r < |\\Delta v|/2 = 700~{\\rm km\\,s^{-1}}$, yielding $(1 + f\\beta_r)\\lesssim 2$.  Considering an inclination angle $i$ of the rotation axis and the mean $\\sin{i}$ factor, $2/\\pi$, the rotational velocity is $\\sigma_r < |\\Delta v|/2\\sin{i} = 1100~{\\rm km\\,s^{-1}}$ and $(1 + f\\beta_r)\\lesssim 3.6$.  Therefore cluster mass estimation with the assumption of hydrostatic equilibrium is justified within a factor of $\\sim 2-3.6$.  In order to put a stronger constraint on the cluster mass to compare with other techniques (particularly gravitational lensing), a two-fold improvement on the energy-scale calibration is necessary.  \\subsection{Turbulent broadening of iron lines}\\label{subsec:turbulent}  The velocity structure in the cluster can also be studied by analyzing the X-ray line widths for turbulent Doppler broadening of the X-ray emission lines.  The iron ions provide the best tracers of this effect, because the broadening due to a turbulent velocity, $\\sigma _{\\rm turb}$, on the order of $\\sim 300~{\\rm km\\,s^{-1}}$, can exceed their thermal broadening ($\\sim 2$~eV) by several times \\citep{Inogamov_Sunyaev_2003}.  However, considering the energy resolution of the FI CCDs, 160~eV (FWHM) or 70~eV (Gaussian $1\\sigma$), what can be done with the present data is at most to examine the observed Fe-K line width for possible broadening beyond the instrumental energy resolution and its intrinsic line width ($\\sigma_0\\sim 28$~eV for $kT=3$~keV).  In order to carry out this search under the highest statistics, we again fitted the XIS0+XIS2+XIS3 spectrum from the central $18\\arcmin\\times18\\arcmin$ region (analyzed in \\S\\ref{subsec:analysis_large}) with the APEC plus three Gaussian model.  Here we assumed that the gas has no bulk velocity, and fixed the centroid of the He-like K$\\alpha$ line to $6.677(1+z_{\\rm cl})^{-1}$~keV.  As a result, the Gaussian width has been obtained as $\\sigma = 47\\pm2$~eV.  Since the degradation of the energy resolution with time is not currently implemented in the XIS response matrix \\citep{Koyama_etal_2006}, the fitted Gaussian width is expressed by $\\sigma^2 = \\sigma_{\\rm inst}^2 + \\sigma_{0}^2 + \\sigma_{\\rm turb}^2$, where $\\sigma_{\\rm inst}$ is the width of the Mn-K$\\alpha$ calibration line as mentioned before, and the average of the three FI CCDs gives $\\sigma_{\\rm inst}=40\\pm4$~eV.  Thus, the measured width is consistent with that expected from the instrumental and intrinsic widths, $(\\sigma_{\\rm inst}^2 + \\sigma_{\\rm 0}^2)^{1/2}=49\\pm2$~eV, where the intrinsic width $\\sigma_{\\rm 0}$ includes also the thermal Doppler although it is negligibly small ($\\sim 2-3$~eV) in the energy range of interest.  The turbulent broadening is then loosely constrained as $\\sigma_{\\rm turb}< 32~{\\rm eV}\\sim 1400~{\\rm km\\,s^{-1}}$ (90\\% confidence).  Although $\\sigma_{\\rm 0}$ significantly increases below a temperature of 2~keV, it is within the range of $26~{\\rm eV} < \\sigma_0< 30~{\\rm eV}$ as long as $2~{\\rm keV}<kT<4~{\\rm keV}$, and does not significantly affect the above discussion.  The derived limit on $ \\sigma_{\\rm turb}$ is consistent with a stronger constraint, $< 800~{\\rm km\\,s^{-1}}$, derived by an argument that strong turbulence would smear out, via diffusion, the observed abundance gradient in the central 100~kpc of the cluster \\citep{Graham_etal_2006}.  Observations of resonance scattering of the X-ray emission lines will provide another way to constrain the ICM turbulent motion.  This is, however, beyond the scope of the present paper."
    },
    "0609/astro-ph0609757.txt": {
        "abstract": "The black-hole X-ray binary transient GRO J1655-40 underwent an outburst beginning in early 2005. We present the results of our multi-wavelength observational campaign to study the early outburst spectral and temporal evolution, which combines data from X-ray ({\\it RXTE}, {\\it INTEGRAL}), radio ({\\it VLA}) and optical ({\\it ROTSE}, {\\it SMARTS}) instruments. During the reported period the source left quiescence and went through  four major accreting black hole states: low-hard, hard intermediate, soft intermediate and high-soft. We investigated dipping behavior in the {\\it RXTE} band and compare our results to the 1996--1997 case, when the source was predominantly in the high-soft state, finding significant differences. We consider the evolution of the low frequency quasi-periodic oscillations and find that the frequency strongly correlates with the spectral characteristics, before shutting off prior to the transition to the high-soft state. We model the broad-band high-energy spectrum in the context of empirical models, as well as more physically motivated thermal and bulk-motion Comptonization and Compton reflection models. {\\it RXTE} and {\\it INTEGRAL} data together support a statistically significant high energy cut-off in the energy spectrum at $\\approx 100-200$ keV during the low-hard state. The {\\it RXTE} data alone also show it very significantly during the transition, but cannot see one in the high-soft state spectra. We consider radio, optical and X-ray connections in the context of possible synchrotron and synchrotron self-Compton origins of X-ray emission in low-hard and intermediate states. In this outburst of GRO J1655--40, the radio flux does not rise strongly with the X-ray flux. ",
        "introduction": "X-ray Novae, also known as soft X-ray transients, are a subclass of the low-mass X-ray  binaries (LMXBs) for which prolonged periods of quiescence are occasionally interrupted by dramatic, accretion-powered  optical, UV, and X-ray outbursts \\citep{csl97,mr04}. A majority of these binaries have been determined to contain compact primaries with mass functions exceeding the nominal 3-solar-mass limit for neutron stars; indeed, most of the known stellar mass black holes  are associated with this class of objects. These outbursts are frequently accompanied by radio emission which is generally associated with collimated outflows. The radio emission is transient in nature, and a wide range of behavior has been documented.  Conventional classification of black hole (BH) spectral states includes five states, namely: {\\it quiescent state}, {\\it low-hard state} (LHS), {\\it intermediate state} (IS), {\\it high-soft} (or {\\it thermally dominant}) {\\it state} (HSS) and {\\it very high} (or {\\it steep power law}) {\\it state}. \\citet{mr04} recently reviewed these states. They advocate the names in parentheses, as being more accurately descriptive. However, we follow tradition in the use of the abbreviation HSS. The LHS energy spectrum is dominated by a hard (spectral index 1.4--1.6) power law. A thermal component is either very small or not seen at all. Strong variability (up to 40 \\%) in the form of band-limited noise is observed in this state, often accompanied by quasi-periodic oscillations (QPO) seen in the Fourier Power Density Spectrum (PDS) as narrow peaks. Stable radio emission with a flat spectrum  has been observed during the LHSs of a set of BH X-ray sources. It has been associated with steady relativistic jets and outflows \\citep{gallo}. BH transients are usually observed in the LHS when their luminosity is less than 5 \\% of the Eddington luminosity. The HSS, in contrast to the LHS, is dominated by a thermal component with a characteristic temperature of $\\approx$ 0.5--1.0 keV, which is attributed to bright emission from an optically thick accretion disk. In the HSS the power law index is steeper ($>$2.0) and variability is suppressed to not more than several percent (which is dominated by red noise, i.e. $P(f)\\sim 1/f$). The LHS and the HSS are the most well observed and documented BH states.  During transitions between them, a BH source enters an IS, where the source may exhibit diverse behavior with mixed LHS and HSS properties. \\citet{hb05} considered the phenomenology of the IS in detail and introduced two substates, which they call {\\it hard intermediate state} (HIMS) and {\\it soft intermediate state} (SIMS) \\citep[See also][]{bel05,bel06}. For the HIMS the energy spectrum is softer than in the LHS and the thermal component is more pronounced. The PDS is still dominated by band-limited noise with characteristic frequencies higher than those observed for the LHS. Radio emission is also observed in the HIMS, but with a slightly steeper spectrum. In the SIMS the disk component starts to dominate the energy spectrum and further softening is observed. The total rms variability drops sharply. The PDS is a sum of band-limited and power law components, or is sometimes even more complicated in shape. This subdivision of IS into HIMS and SIMS is consistent with the behavior of GRO J1655-40 during the hard-to-soft transition that we observed, so we adopt this terminology when we want to refer to a particular intermediate substate. In our data we identify LHS, HIMS, SIMS, and HSS periods.  The X-ray binary GRO~J1655--40 is a well known example of BH X-ray transient, having undergone several major outbursts within the the last 12 years. It was discovered by the BATSE instrument onboard the Compton Gamma-Ray Observatory in mid 1994 \\citep{zh94,har95}. The secondary star being relatively bright, the binary parameters are exceptionally well determined amongst LMXBs. %\\citet{OB97}, using The most recent optical  photometry of \\citet{greene01} led to a BH mass estimate of $6.3\\pm0.5 M_\\odot$. \\citet{hr95} inferred a distance to GRO~J1655--40 of $3.2\\pm0.2$ kpc from radio jet analysis. \\citet{mir02} discussed in detail different methods of distance estimates and concluded that for GRO~J1655--40 the upper limit is 3.5 kpc and the lower limit is 0.9 kpc. %GRO~J1655--40 was discovered by the BATSE instrument onboard the %Compton Gamma-Ray Observatory in mid 1994 \\citep{zh94,har95}. The light curve of the discovery outburst was irregular, in the sense that it deviated substantially from the often seen \u0093fast-rise exponential decay\u0094 form \\citep{csl97}.  It was instead characterized by several distinct peaks, with subsequent sporadic, lower amplitude outburst activity continuing into 1995. These initial events were correlated with the contemporaneous radio observations. Most notable was the discovery of radio jets which exhibited apparent superluminal motion \\citep{har95,hr95,tin95}. The later, low-amplitude events however, were apparently not associated with additional plasma ejections leading to the type of radio emission initially seen \\citep[e.g.][]{tav96}.   In April 1996, after a short period of quiescence, a new outburst appeared in the data of the {\\it Rossi X-ray Timing Explorer} ({\\it RXTE}) All-Sky Monitor \\citep{rem96}. This event also consisted of two giant flares which lasted 100 and 222  days. Although the radio behavior was far less dramatic, extensive {\\it RXTE} coverage allowed this outburst to be studied in greater detail than the 1994--1995 events. However, pointed {\\it RXTE} observations started when the source was already in the HSS and the state transition was not covered. The 2005 outburst was recognized at a very early stage and observers at different wavelengths were relatively prepared and successful at responding relatively quickly. In this paper we present the analysis of the first multi-frequency observations of an initial  hard-to-soft state transition for a GRO~J1655--40 outburst. The data comprise one the most detailed multi-frequency observations of the initial hard-to-soft transition in a BH transient [see \\citet{broc02} for a review of the observed state transitions in BH sources].   On February 17, 2005 (MJD 53419) {\\it RXTE} Proportional Counter Array (PCA) bulge scan observations detected emission consistent with coming from the position of GRO~J1655--40 and strongly suggested a new outburst \\citep{mar05}. The flux in the 2-10 keV band was 4.1 +/- 0.6 mCrab. Near infra-red (NIR) J photometry on February 20.39 UT found the source 0.5 mag brighter than at its brightest in quiescence \\citep{atel417} and the NIR activity was confirmed on February 21.3-21.4 \\citep{atel418}. NRAO {\\it Very Large Array Radio Observatory} ({\\it VLA}) observations on February 20.6 found that it was again a detectable radio source \\citep{atel419}. A campaign of daily pointed {\\it RXTE} observations started on February 21, 2005 (MJD 53423).  These observations and the PCA bulge scans triggered our  program of Target of Opportunity (ToO) observations with INTernational Gamma-Ray Astrophysical Laboratory ({\\it INTEGRAL}), the goal of which was to characterize the early-rise phase of X-ray nova events. Subsequently the source brightened and triggered other programs for observing bright transients.  The observations we consider in this paper span a 25 day period starting on 21 February 2005. We have utilized the radio flux-history from an ongoing program of transient monitoring at the {\\it VLA}. The outburst was also monitored by the {\\it Swift} observatory with a series of pointed observations \\citep{broc06}, the first of which fell on the initial LHS within the time frame of the data reported here. The source was observed extensively until the outburst faded, by  {\\it RXTE} \\citep{ho05a,ho05b,ho05c} and other missions.   In this paper we report on the results of our multi-wavelength observational campaign, beginning with the outburst detection and ending on March 16, 2005 (MJD 53445) after the source completed the hard-to-soft state transition. The data are described in \\S 2. Our analysis of the {\\it INTEGRAL} and {\\it RXTE} observations is presented in \\S 3. This includes light curve analysis as well as our spectral model fitting. We discuss dipping behavior in \\S 4. Interpretation and conclusion are presented in \\S 5. ",
        "conclusions": "}  We have studied  multi-frequency data spanning the early stage of the 2005 outburst of GRO J1655-40. The  empirical behavior observed includes: spectral transition from the LHS through the HIMS and SIMS to the  HSS, accompanied by the expected rise in the thermal component and a steepening of power law spectral index; QPO appearance and its frequency rise throughout the LHS and the intermediate states and its turnoff upon entering the HSS; a  high-energy cut off in the LHS and IS; intensity dips during the LHS which are consistent with photo-electric absorption and correlated differently with orbital phase than the  HSS dips observed by K00 for the 1996--1997 outburst.  We observe rising radio emission  during the LHS, a declining radio flux in the HIMS and its turn off in the SIMS. The radio data from MJD\\,53433 are unique in showing a pronounced peak at 5\\,GHz.  X-ray binaries in the LHS typically produce a very stable, flat-spectrum radio source, as seen here in the plateau stage.  A 5\\,GHz peak is highly unusual, and suggests an expanding, optically-thick source, typical of the early stages of jet ejection. The radio coverage allows only poor limits on the ejection date (later than MJD\\,53429, earlier than 53433), but it is tempting to associate this ejection with the abrupt change from a steady to a sharply brightening X-ray corona.  In this context it is also interesting that during the X-ray exponential rise, the radio flux was comparatively steady, with a marginal steepening of the  spectrum.  While the latter is not convincing on any given day, the radio spectrum is consistently flatter (harder) during the plateau stage, and steeper (softer) during the exponential rise.  Taken together with the decoupling between the X-ray and the radio {\\it flux}, it seems likely that the radio emission here has a different physical origin from that during the plateau.  The plateau seems much like the standard LHS, in the ratio of X-ray to radio fluxes, and the flat radio spectrum, suggestive of a compact core. The radio emission during the HIMS may instead be a more extended jet, though it is far from obvious why that jet should be so stable (in flux density and spectrum).  Currently, there are two competing theories explaining the origin of the X-ray power law spectrum in BH sources. In the first, the ``canonical'' model, the power law slope is produced by Comptonization of the soft disk photons in the extended optically thin corona, located above the accretion disk. But \\citet[][and references therein]{mar03}, argue that the power law component in the LHS can be attributed to a combination of the optically thin high energy tail of the synchrotron spectrum from the radio jet base and synchrotron self-Compton emission adding the highest energy component. In the framework of the  model, the synchrotron jet base subsumes the role of the corona in Comptonization models \\citep[e.g.][]{nowak05,markoff05}. However, the lack of radio and X-ray correlation in our data  make it difficult in the case of GRO~J1655--40 to associate the base of a compact jet with the Compton-scattering corona which creates the power law X-ray emission. \\citet{gallo} found evidence for radio and X-ray correlation for several sources, though they also saw that in some cases the radio was quenched at some point in the X-ray rise. It is possible that the radio emitting region is larger than the X-ray corona, and not dominated by the X-ray emitting part. This would mean that the situation may be more complex. Recently the jet-base model has been subject to other observational challenges. \\citet{mac06} argued that if the LHS spectrum is due to the jet, which is radiatively inefficient in nature, the transition to the HSS should be accompanied by a sharp change in luminosity. While we observe a rather quick rise in the thermal component during the SIMS around MJD 53440 (see Figure \\ref{flux_vs_time}), the corresponding power law flux does not drop, but jumps up together with black body component before it starts to decay smoothly. Overall, our analysis does not validate the jet-base model; however some presence of synchrotron and synchrotron self-Compton emission can not be ruled out.  It is apparent that the system is far from a steady state configuration during the LHS and the state transition.  Most probably, the accretion disk, which is quiescent prior to the outburst, gradually gets involved in outburst activity by some instability propagating either  from the innermost parts of the disk to the outside or from the outer disk to the inside\\citep[e.g.][]{can95}. Information from the QPO phenomena can give an important clue towards distinguishing between the two possible directions of the heat front propagation. For example, in the case of shock-wave or front propagation in the disk one could expect the X-ray flux to be modulated at the frequency of the disk Keplerian rotation at the radius of the shock location. If this process is indeed responsible for the observed flux oscillations, then the outward propagation scenario is not correct and the shock is propagating inward from larger radius, as advocated by \\citet{ch05}. However, there is evidence for a closer  coupling of the QPO phenomena with the power law component than  with the disk component, so that the accretion disk shock wave oscillation model for the oscillations is not compelling. QPO rms variability is found to to increase with energy \\citep[see][and references therein]{klis95}. \\citet{vig03} found a close correlation of power law spectral index with QPO centroid frequency for  BH sources GRS~1915+105, GRO~J1655--40, XTE~J1550--564 and 4U~1630--47. For GRO~J1655--40, Vignarca et al. used data from the 1996--1997 outburst. Figure \\ref{qpo_alpha} shows that in the 2005 outburst rise the spectral index was correlated with QPO frequency above about 0.8 Hz. Both the power law and the black body (seed photon source) fluxes were correlated with the frequency, while the radio flux is not (see Figure \\ref{flux_corr}). In the BMC model, the power law flux and the black body are related. However, the spectral index is independent of QPO until a significant rise in flux has occurred. Qualitatively, these correlations are consistent with the BMC model and do not exhibit the same indications of a jet-base model shown in some other sources by \\citet{mig05}. There are other models that could relate this QPO frequency to the fluxes, such as that of the occultation of the power law component by a precessing disk ring \\citep{schn06}. The dynamical consistency of such a ring and the connection to the spectrum of the disk component should be further investigated.  The location of dips has been discussed in terms of the impact of the accretion stream on the disk \\citep[e.g.][]{bis05}. The location may be a function of the rate of accretion flow. The structure of the dips may put limits on the size of the emission region, as has been done for neutron star low mass binaries \\citep[e.g.][]{chur01}. Detailed study of this is beyond the scope of this paper.  In summary, the 2005 outburst of GRO~J1655--40 was observed intensely with X-ray instruments, with very interesting and diagnostic results. During the early phases there was good radio coverage, and less thorough, but nevertheless useful optical coverage.  The X-ray data allow spectral fits to generic Comptonization and power law plus reflection models with a cutoff, multicolored disk and fluorescence. If the corona is really the base of a relativistic outflow, the radio flux does not exhibit the proportionality to the X-ray flux that is seen in some other BH sources, nor a simple anti-correlation of the flux between the power law and the disk components of the spectrum. Low frequency QPO seem to suggest that their origin, presumably at an interface between an optically thick and a coronal component, is moving inwards in radius, as the frequency is positively correlated with the luminosity in the power law, the luminosity in a disk component (and inversely with its inner radius), while the radio flux decreases as the QPO reaches its maximum. The spectra of the power law component are limited by a high energy cutoff, although we are unable to study this aspect in detail. While our study in this paper extends only to when the HSS began, the observations that followed will address many other points, including the degree of similarity of the decay of this transient to its rise.  We wish to thank J. Homan and collaborators for the use of the data from their proposal and  E. Kuulkers and the ESA INTEGRAL Science Operation Center (ISOC) for providing the on-line Galactic Bulge Monitoring database."
    },
    "0609/astro-ph0609206_arXiv.txt": {
        "abstract": "We present submillimeter observations of dark clouds that are part of the Spitzer Legacy Program, From Molecular Cores to Planet-Forming Disks (c2d).  We used the Submillimetre Common User's Bolometer Array to map the regions observed by Spitzer by the c2d program to create a census of dense molecular cores including data from the infrared to the submillimeter.  In this paper, we present the basic data from these observations: maps, fluxes, and source attributes.  We also show data for an object just outside the Perseus cloud that was serendipitously observed in our program. We propose that this object is a newly discovered, evolved protostar.  Finally, these data are publicly available at http://peggysue.as.utexas.edu/SIRTF/DATA/. ",
        "introduction": "In recent years, significant progress has been made toward a theory for the formation of low mass stars, which are the likely sites for the formation of planetary systems.  Maps of submillimeter dust continuum emission are excellent tracers of the column density and mass. Dust continuum emission has been used to map both large clouds (e.g. \\citealt{mot98}; \\citealt{joh04}; \\citealt{enoch06}) and smaller cores (e.g., \\citealt{shirley00}; \\citealt{visser02}; \\citealt{young03}; \\citealt{kirk05}).  Studies of small cores provide complementary advantages to less-biased surveys of large clouds. By focusing on regions already known to have dense gas, we can study their properties in more detail. The large scale surveys can in turn check the biases that are introduced. This paper presents data on dust continuum emission toward a large number of small cores.  For a full analysis, data at submillimeter wavelengths should be combined with data through the infrared. The Spitzer Legacy program, From Molecular Cores to Planet-Forming Disks (c2d), described by \\citet{evans03} has obtained extremely sensitive observations from 3.6 to 70 \\micron\\ toward 105 small cores \\citep{huard06}. In order to obtain the full benefit of this database, it is crucial to add data at longer wavelengths where the dust emission is optically thin.  We have mapped 38 small cores, which were initially included in the c2d program, using the Submillimetre Common-User Bolometer Array (SCUBA) at both 850 and 450 \\micron.  Because of time constraints, 15 of these cores were dropped from the c2d program, though some were covered by the large cloud map of Perseus (\\citealt{joergensen06}, \\citealt{rebull06}).  Together with other complementary data at optical and near-infrared wavelengths, our observations provide a comprehensive catalog of multi-wavelength data for a sample of small molecular cores.  The goal of this paper is to present the SCUBA data in a uniform way, consistent with other papers presenting submillimeter continuum data on the c2d cores and clouds. We leave detailed analysis to later papers that combine these data with the c2d data. ",
        "conclusions": "We have presented SCUBA observations of low-mass cores that were initially proposed to be observed with the Spitzer Space Telescope as a part of the c2d Legacy Program.  We describe these observations and the data reduction and analysis.  In addition, we show the maps, source fluxes, and source sizes.  The core mass has been calculated from the 850 $\\mu$m fluxes, and we analyze how the mass varies with source size and shape.  We find several small clouds with multiple, distinct cores, but, for most regions, only one of the cores is detected by IRAS.  Finally, we present inadvertent observations of an object near Perseus that may be a newly discovered protostar.  All data presented here are publicly available as FITS images at http://peggysue.as.utexas.edu/SIRTF/DATA/."
    },
    "0609/astro-ph0609801_arXiv.txt": {
        "abstract": "We examine the growth of eccentricities of a population of particles with initially circular orbits around a central massive body. Successive encounters between pairs of particles increase the eccentricities in the disk on average.  As long as the epicyclic motions of the particles are small compared to the shearing motion between Keplerian orbits, there is no preferred scale for the eccentricities. The simplification due to this self-similarity allows us to find an analytic form for the distribution function; full numerical integrations of a disk with 200 planetesimals verify our analytical self-similar distribution.  The shape of this non-equilibrium profile is identical to the equilibrium profile of a shear-dominated population whose mutual excitations are balanced by dynamical friction or Epstein gas drag. ",
        "introduction": "Modern computational power allows the simultaneous integration of the orbits of increasingly numerous particles. Much of the planet formation process, however, involves particle numbers that exceed the limits of computational efficiency.  This limitation is often circumvented with a statistical approach. By monitoring the gravitational interactions of the particles in a time-averaged sense, various properties of the particle population can be calculated without a full $N$-body simulation.  \\citet{CS06}, hereafter Paper I, motivate a Boltzmann equation to describe the evolution of the eccentricity distribution of an ensemble of particles in which the relative motion between any two interacting particles is dominated by the shearing motion of close circular orbits.  Such a regime of orbital eccentricities is called shear-dominated.  The solution of their equation provides a simple analytic expression for the equilibrium eccentricity distribution that results when dynamical friction can balance the mutual interactions of the particles; the analytic expression matches results from numerical simulations remarkably well.  In this letter we derive analytically the non-equilibrium distribution function of interacting shear-dominated particles in the absence of dynamical friction. \\S \\ref{secReview} reviews the construction of the Boltzmann equation. In \\S \\ref{secAnalytical}, we show that the distribution function behaves self-similarly, and the shape of the non-equilibrium distribution function is identical to the equilibrium distribution of Paper I. \\S \\ref{secNumerical} corroborates this result with numerical simulations.  Conclusions follow in \\S \\ref{secConclusions}. ",
        "conclusions": "\\label{secConclusions}  We have written a time-dependent Boltzmann equation that describes the eccentricity distribution function of a population of orbiting particles under the influence of their mutual excitations in the shear-dominated regime. Reasoning that the distribution function of eccentricities should behave self-similarly, that is, retain a constant profile while its normalization and scaling depend on time, we have reduced the Boltzmann equation to a form that has already been shown to have an analytic solution.  Numerical experiments confirm the self-similarity and the analytic solution.  Although we have only considered disks of a single particle size, the formalism above applies trivially to disks with mass distributions. In fact, the characteristic eccentricity, equation \\ref{eqEStarofT}, depends only on the total surface mass density of the disk. As long as bodies of every part of the mass spectrum are in the shear-dominated regime, the eccentricities of all bodies are drawn from the same distribution.  This is a consequence of the fact that gravitational acceleration is mass-independent.  In contrast, the dynamical friction of Paper I depends on the size of each particle.  The equilibrium distributions in that case do differ for each mass group.  Since $e_c(t)$ is an increasing function of time, the condition of shear-dominated dynamics will be violated eventually. While the disk is shear-dominated however, most of the disk bodies have eccentricities of about $e_c(t)$.  The mean eccentricity, $\\int\\int e f(e,t) d^2 \\vec e$, is formally infinite; in reality the mean depends logarithmically on the maximum eccentricity achievable from one interaction.  Higher moments of the distribution, such as $\\langle e^2 \\rangle$, are dominated by the bodies with the maximum eccentricity.  The random kinetic energy of the disk bodies, for example, is then set by the few bodies with the highest eccentricities regardless of the value of $e_c(t)$.  That the shape of the distribution function is identical to the distribution function of Paper I is ultimately not surprising.  In both scenarios the bodies in question excite their orbital parameters via the same shear-dominated viscous stirring mechanism. If dynamical friction is acting on these bodies, their eccentricities decrease with time proportionally to their magnitude.  An equilibrium between excitations and this damping produces a characteristic eccentricity around which the eccentricities of all bodies are distributed. Without an agent of dynamical friction, the typical eccentricity of a body in the disk, $e_{\\rm typical} \\sim e_c(t)$, grows with time. However, the ratio of the eccentricity of a particle that has not interacted recently, $e$, to that typical eccentricity shrinks proportionally to itself:  \\begin{equation} \\dot x \\equiv \\frac{d}{dt} \\left( \\frac{e}{e_c(t)} \\right) \\sim - x . \\end{equation}  \\noindent This is formally equivalent to the damping provided by the dynamical friction of Paper I."
    },
    "0609/astro-ph0609030_arXiv.txt": {
        "abstract": "{We present multi-frequency radio observations of a sample of $z\\sim2$ obscured (type-2) quasars in the Spitzer extragalactic First Look Survey area. We combine the public data at 1.4 GHz, used in the selection of these sources, with new observations at 610 MHz (GMRT) and at 4.9 GHz (VLA). We find the sample includes sources with steep, flat  and gigahertz-peaked spectra. There are no strong correlations between the presence or absence of emission lines in the optical spectra and the radio spectral properties of the sample. However, there are no secure flat-spectrum type-2 quasars with narrow emission lines which would be problematic for unified schemes.  Most of the population have straight radio spectra with spectral index $\\alpha \\sim 1$ as is expected for developed, potentially FRI-like, jets in which continous injection of relativistic electrons is accompanied by inverse-Compton losses against the cosmic microwave background.    } ",
        "introduction": "The `zoo' of active galactic nuclei (AGN) includes many different objects with signs of non-stellar activity, powered by a supermassive black hole at the centre of a host galaxy \\citep{1984ARA&A..22..471R}. The different observed properties are usually explained via variations in the rate of accretion of this central black hole, the black hole mass, and the orientation of the axes of symmetry with respect to the observer \\citep[e.g.][]{1993ARA&A..31..473A}.  Obscured (type-2) quasars are intrinsically the same as unobscured (type-1) quasars, except for the orientation, where the dusty torus around the central black hole blocks the line of sight to the bright accretion disk and the broad-line region. Only the narrow-line region is then visible to the observer. Radio-loud type-2 quasars, in the form of narrow emission line radio galaxies, have been known for a long time. For the more radio-luminous galaxies \\citep[the FRII class of ][] {1974MNRAS.167P..31F}, convincing evidence exists for their obscured quasar nuclei \\citep[e.g.][] {2005A&A...442L..39H,Ogle:2006cu}.  However, not all FRII radio-galaxies are powerful accreters \\citep{Ogle:2006cu}, and amongst the less radio-luminous FRI class few quasars were found \\citep [see ] [ for an exception] {2001ApJ...562L...5B}, although recent radio imaging of quasars has challenged this (Heywood, Blundell \\& Rawlings, MNRAS, submitted). Amongst the radio-quiet population, the type-2 quasars had remained elusive until recently.    The search for high-redshift type-2 quasars has been carried out mainly in the X-rays \\citep [see][for a review] {2005ARA&A..43..827B} and, more recently, in the mid-infrared \\citep [e.g.][]{2004ApJS..154..166L,2005Natur.436..666M}.  The sample of \\citet {2005Natur.436..666M,2006MNRAS.370.1479M} consists of $z\\sim2$ radio-intermediate type-2 quasars: quasars with radio luminosities below the FRI/FRII break, but above that would be expected from star-formation alone. These were selected in the Spitzer extragalactic First Look Survey (FLS) by demanding objects to have a 24-$\\mu$m flux density $S_{24}>300~\\mu$Jy, a 3.6-$\\mu$m flux density $S_{3.6}\\leq45~\\mu$Jy, and a 1.4 GHz flux density in the range $350~\\mu$Jy $\\leq S_{1.4} \\leq$ 2 mJy.  \\citet {2006MNRAS.370.1479M} found secure optical redshifts (from narrow lines) for $\\sim50$\\% of their sample with the remaining spectra typically completely blank: they discuss the possibility of their blank optical spectra being due to a dusty host galaxy obscuring all the central activity, even the narrow emission lines. Similar type-2s with blank optical spectra are found in the samples of \\citet {2004ApJS..154..166L} and \\citet {2005A&A...440L...5L}.       \\begin{table*} \\begin{center} \\begin{tabular}{lllcrrrrrr} \\hline \\hline Name & RA & Dec &   $S_{4.9 \\rm GHz}$ &  $S_{1.4 \\rm GHz}$ & $S_{610 \\rm MHz}$ & $\\alpha^{4.9}_{1.4}$ & $\\alpha^{1.4}_{610}$ & $z_{\\rm spec}$ & Spec \\\\  & (J2000)  &  &  \\mc{1}{c}{/ $\\mu $Jy} & \\mc{1}{c}{/ $\\mu $Jy} & \\mc{1}{c}{/ $\\mu $Jy} & &  \\\\ \\hline AMS01 & 17 13 11.17 & +59 55 51.5  & 164 $\\pm$ 52 & 490  $\\pm$ 31 & 1167 $\\pm$ 165 & 0.87$^{+0.35}_{-0.27}$ & 1.04$^{+0.13}_{-0.13}$ & - &1 \\\\ AMS02 & 17 13 15.88 & +60 02 34.2  & 408 $\\pm$ 52  & 1184 $\\pm$ 55 & 2985 $\\pm$ 152 & 0.85$^{+0.15}_{-0.13}$ & 1.11$^{+0.08}_{-0.08}$ & - &1 \\\\ AMS03 & 17 13 40.19 & +59 27 45.8  &  365 $\\pm$ 53  & 1986 $\\pm$ 87 & 5180 $\\pm$ 191 & 1.35$^{+0.16}_{-0.14}$ & 1.15$^{+0.07}_{-0.07}$ & 2.698 &2 \\\\ AMS04A$^{\\dag}$ & 17 13 40.62 & +59 49 17.1  &  135 $\\pm$ 53  & 536  $\\pm$ 72 & 1092 $\\pm$ 77 & 1.10$^{+0.50}_{-0.38}$ & 0.86$^{+0.26}_{-0.24}$ & 1.782 &2 \\\\ AMS04B$^{\\ddag}$ & 17 13 40.62 & +59 49 17.1  &  232 $\\pm$ 106  & 867  $\\pm$ 101 & 1519 $\\pm$ 122 & 1.05$^{+0.58}_{-0.40}$ & 0.68$^{+0.24}_{-0.23}$ & 1.782 &2 \\\\ AMS05 & 17 13 42.77 & +59 39 20.2  &  443 $\\pm$ 51  & 1038 $\\pm$ 49 & 1214 $\\pm$ 164 & 0.68$^{+0.14}_{-0.13}$ & 0.19$^{+0.13}_{-0.13}$ & 2.017 &2 \\\\ AMS06 & 17 13 43.91 & +59 57 14.6  & 191$\\pm$53   & 444  $\\pm$ 31 & 1115 $\\pm$ 153 &0.67$^{+0.31}_{-0.25}$ & 1.11$^{+0.14}_{-0.14}$ & 1.76  &1,3 \\\\ AMS07 & 17 14 02.25 & +59 48 28.8  &  239 $\\pm$ 51  & 354  $\\pm$ 27 & 578  $\\pm$ 151 & 0.31$^{+0.25}_{-0.22}$ & 0.59$^{+0.20}_{-0.20}$ & - &1 \\\\ AMS08 & 17 14 29.67 & +59 32 33.5  & $<$718   & 655  $\\pm$ 36 & 889  $\\pm$ 151 &$>$-0.07 & 0.37$^{+0.21}_{-0.22}$ & 1.979 &2 \\\\ AMS09 & 17 14 34.87 & +58 56 46.4  &  174 $\\pm$ 51  & 426  $\\pm$ 29 & 1118 $\\pm$ 190 & 0.71$^{+0.33}_{-0.26}$ & 1.16$^{+0.18}_{-0.18}$ & - &1 \\\\ AMS10 & 17 16 20.08 & +59 40 26.5  & 501$\\pm$ 36  & 1645 $\\pm$ 73 & 4590 $\\pm$ ~93 & 1.04$^{+0.13}_{-0.12}$ & 1.24$^{+0.74}_{-0.07}$ & - &1 \\\\ AMS11 & 17 18 21.33 & +59 40 27.1 & $<$104  & 356  $\\pm$ 29 & 628  $\\pm$ ~89 &$>$0.98 & 0.68$^{+0.21}_{-0.21}$ & - &1 \\\\ AMS12 & 17 18 22.65 & +59 01 54.3  &  305 $\\pm$ 83  & 946  $\\pm$ 45 & 2440 $\\pm$ 194 & 0.90$^{+0.19}_{-0.17}$ & 1.14$^{+0.10}_{-0.94}$ & 2.767 &2 \\\\ AMS13 & 17 18 44.40 & +59 20 00.8  &  495 $\\pm$ 80  & 1888 $\\pm$ 83 & 4559 $\\pm$ 135 & 1.07$^{+0.18}_{-0.16}$ & 1.06$^{+0.07}_{-0.72}$ & 1.974 &2,3 \\\\ AMS14 & 17 18 45.47 & +58 51 22.5  &  158 $\\pm$ 52  & 469  $\\pm$ 29 & 1099 $\\pm$ 196 & 0.87$^{+0.37}_{-0.28}$ & 1.03$^{+0.14}_{-0.14}$ & 1.794 &2 \\\\ AMS15 & 17 18 56.93 & +59 03 25.0  & 326 $\\pm$ 53  & 440  $\\pm$ 29 & 570  $\\pm$ 174 & 0.24$^{+0.19}_{-0.17}$ & 0.31$^{+0.23}_{-0.25}$ &  &4 \\\\ AMS16 & 17 19 42.07 & +58 47 08.9  &  $<$108  & 390  $\\pm$ 60 & $<$170   & $>$1.02 &$<$-1.00 & 4.169 &2 \\\\ AMS17 & 17 20 45.17 & +58 52 21.3  &  265 $\\pm$ 53  & 615  $\\pm$ 34 & 1354 $\\pm$ 226 & 0.67$^{+0.22}_{-0.19}$ & 0.95$^{+0.18}_{-0.19}$ & 3.137 &2 \\\\ AMS18 & 17 20 46.32 & +60 02 29.6  &  212 $\\pm$ 53  & 390  $\\pm$ 29 & 502  $\\pm$ 113 & 0.49$^{+0.29}_{-0.24}$ & 0.30$^{+0.25}_{-0.26}$ & 1.017 &2 \\\\ AMS19 &  17 20 48.00 & +59 43 20.7  &  506 $\\pm$ 52  & 822  $\\pm$ 41 & 958  $\\pm$ 152 & 0.50$^{+0.18}_{-0.16}$ & 0.18$^{+0.17}_{-0.17}$ & 2.25 &1,3 \\\\ AMS20 & 17 20 59.10 & +59 17 50.5  & 1391 $\\pm$53  & 1268 $\\pm$ 58 & 674  $\\pm$ 187 & -0.07$^{+0.07}_{-0.07}$ &-0.76$^{+0.32}_{-0.40}$ & - &1 \\\\ AMS21 & 17 21 20.09 & +59 03 48.6  & 145 $\\pm$ 53  & 449  $\\pm$ 29 & 1019 $\\pm$ 169 & 0.91$^{+0.42}_{-0.30}$ & 0.99$^{+0.16}_{-0.16}$ & - &1 \\\\ \\hline \\hline \\end{tabular} \\caption{ \\noindent Positions, redshifts and radio data for the entire sample. The integrated fluxes with errors are from the AIPS tasks IMFIT, SAD and JMFIT (for 4.9 GHz, 1.4 GHz and 610 MHz respectively). The last column states the spectroscopic properties: 1 blank in optical spectrum, 2 shows narrow lines ($\\lesssim 2000$ km s$^{-1}$) in the optical spectrum, 3 has a redshift determined from Spitzer-IRS, 4 has not been observed by optical spectroscopy. The optical redshifts obtained using the William Herschel Telescope-ISIS instrument are from \\citet {2006MNRAS.370.1479M}, while the Spitzer-IRS redshifts are from \\citet {2005ApJ...628..604Y} or \\citet {2006ApJ...638..613W}. $^{\\dag}$The properties of AMS04 without the flux density of the source at 17 13 41.20 +59 49 24.0 which might be an extended component of AMS04. $^{\\ddag}$The properties of AMS04 including the flux density of the possible extended component. } \\label{tab:table} \\end{center} \\end{table*}   Radio observations can be used to test whether the type-2 quasars are obscured by the torus. The first expectation is that radio jets, if present, are at large angles to the line-of-sight and steep radio spectral indices are expected. If the jets in some of the type-2 quasars were recently triggered, then GHz-peaked radio spectra, characteristic of young jets \\citep {1998PASP..110..493O} might be observed. Classic type-2 quasars (presumably those showing narrow-lines) are not expected to have the jet close to the observer's line of sight and exhibit the flat radio spectra seen in some type-1 quasars. Therefore, the suggestion by Miller, Rawlings \\& Saunders (1993) and Falcke, Malkan \\& Biermann (1995) that the radio-intermediate quasar population is comprised of Doppler-boosted flat-spectrum radio-quiet quasars would pose serious difficulties for the narrow-line type-2 quasars in this sample. However, if some quasars are obscured by dust independent of the obscuration, as suggested by \\citet {2006MNRAS.370.1479M}, then the jets could be close to the line of sight, and flat radio spectra might be observed. This predicts that a small fraction of the sources with no  narrow lines should have flat radio spectra, although this fraction will be boosted by a radio selection criterion.  \\vspace{-0.1cm} In this letter we report the results of observing the entire Mart\\'\\i nez-Sansigre et al. sample in the FLS at 610~MHz and 4.9 GHz, combined with public data at 1.4 GHz, and the conclusions on orientation or age that can be inferred.  Throughout this paper we adopt a $\\Lambda$CDM cosmology with the following parameters: $h = H_{0} / (100 ~ \\rm km ~ s^{-1} ~ Mpc^{-1}) = 0.7$; $\\Omega_{\\rm m} = 0.3$; $\\Omega_{\\Lambda} = 0.7$. ",
        "conclusions": ""
    },
    "0609/astro-ph0609389_arXiv.txt": {
        "abstract": " ",
        "introduction": "For the plasma physicist, the solar corona offers an outstanding example of a space plasma, and surely one that deserves a lifetime of study. Not only can we observe the solar corona on scales of a few hundred kilometers and monitor its changes in the course of seconds to minutes, we also have a wide range of detailed diagnostics at our disposal that provide immediate access to the prevalent physical processes.  Yet, solar physics offers a rich field of unsolved plasma-physical problems. How is the coronal plasma continuously heated to $> 10^6$~K? How and where are high-energy particles episodically accelerated? What is the internal dynamics of plasma in magnetic loops? What is the initial trigger of a coronal flare? How does the corona link to the solar wind, and how and where is the latter accelerated? How is plasma transported into the solar corona?  Why, then, study stellar coronae that remain spatially unresolved in X-rays and are only marginally resolved at radio wavelengths, objects that require exposure times of several hours before approximate measurements of the  ensemble of plasma structures can be obtained?   \\begin{table}[t!] \\caption{Symbols and units used throughout the text} \\label{symbols} \\begin{tabular}{ll} \\hline\\noalign{\\smallskip} Symbol, acronym\t          \t    & Explanation\t       \\\\ \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} $R_*$ or $R$                                & Stellar radius [cm]\t\t      \\\\ $R_{\\odot}$             \t            & Solar radius [$7\\times 10^{10}$~cm]\t\t      \\\\ $c$                                         & Speed of light [$3\\times 10^{10}$~cm~s$^{-1}$]\t\t      \\\\ $k$                                         & Boltzmann constant [$1.38\\times 10^{-16}$~erg~cm$^{3}~$K$^{-1}$]\t\t      \\\\ $m_e$                                       & Mass of electron [$9.1\\times 10^{-28}$~g]\t\t      \\\\ $d$             \t                    & Distance [pc]\t\t      \\\\ $p$             \t                    & Pressure [dyne~cm$^{-2}$]\t\t      \\\\ $L$             \t                    & Coronal loop semi-length [cm]\t\t      \\\\ $T$             \t        \t    & Coronal electron temperature [K]\t  \\\\ $T_b$             \t        \t    & Radio brightness temperature [K]\t  \\\\ $n, N$                                      & non-thermal electron density ($N$: integrated in energy)       \\\\ $N$                                         & Number of flares       \\\\ $n_e$             \t        \t    & Electron density [cm$^{-3}$]\t  \\\\ $n_{\\rm H}$             \t            & Hydrogen density [cm$^{-3}$]\t  \\\\ $B$             \t        \t    & Magnetic field strength [G]\t  \\\\ $f$             \t        \t    & Surface filling factor [\\%]\t  \\\\ $\\Gamma$             \t        \t    & Loop area expansion factor (apex to base)\t  \\\\ $E$             \t                    & Energy  [erg]\t\t      \\\\ $L_R$             \t                    & Radio luminosity \t[erg~s$^{-1}$~Hz$^{-1}$]\t\t      \\\\ $L_X$             \t                    & X-ray luminosity \t[erg~s$^{-1}$]\t\t      \\\\ $L_{\\rm bol}$             \t            & Stellar bolometric luminosity \t[erg~s$^{-1}$]\t\t      \\\\ $\\Lambda = \\Lambda_0T^{\\chi}$               & Cooling function  [erg~s$^{-1}$~cm$^{3}$]   \\\\ $\\tau$                                      & (Decay) time scales, also: optical depth   \\\\ $\\nu$                                       & Radio frequency    \\\\ $\\nu_p$, $\\omega_p$                         & Plasma frequency, angular  plasma frequency \\\\ $\\nu_c$, $\\Omega_c$                         & Gyrofrequency, angular  gyrofrequency \\\\ $\\kappa$                                    & Absorption coefficient    \\\\ $\\gamma$                                    & Lorentz factor \\\\ $L_B$                                       & Magnetic scale height \\\\ HRD          \t                \t    & Hertzsprung-Russell Diagram\t\t\t \\\\ EM             \t                \t    & Emission Measure  \t\t     \\\\ $Q$, DEM             \t        \t    & Differential Emission Measure  Distribution\t\t      \\\\ EMD             \t        \t    & (discretized, binned) Emission Measure Distribution      \\\\ \\noalign{\\smallskip}\\hline \\end{tabular} \\end{table}   There are many reasons. In the context of the {\\it solar-stellar connection,} stellar X-ray astronomy has introduced a range of stellar rotation periods, gravities, masses, and ages into the debate on the magnetic dynamo. Coronal magnetic structures and heating  mechanisms may vary together with variations of these parameters. Parameter studies could provide valuable insight for constraining relevant theory. Different topologies and sizes of magnetic field structures lead to different wind mass-loss rates, and this will regulate the stellar spin-down rates differently.  Including stars into the big picture of coronal research has also widened our view of coronal plasma physics. While solar coronal plasma resides typically at $(1-5)\\times 10^6$~K with temporary excursions to $\\approx 20$~MK during large flares, much higher temperatures are found  on some active stars, with steady plasma temperatures of several tens of MK and flare peaks beyond 100~MK. Energy release in stellar flares involves up to $10^5$ times more thermal energy than in solar flares, and pressures that are not encountered in the solar corona.  The present chapter provides a ``stellar astronomer's view'' of magnetic coronae. It cannot replace the knowledge of detailed physical processes in the solar corona, nor do we expect to find entirely new physical concepts from  the study of stellar coronae without guidance by solar physics. However, as I hope to show in the following sections, stellar astronomy has provided some unexpected and systematic trends that may well help understand systematics of coronal behavior across a large range of stellar parameters. In this sense, the goal of this chapter is to present an overview of the basic, observed stellar-coronal phenomena and their interpretation analogous to solar physics, rather than the derivation of the basic plasma-physical mechanisms themselves. For a deeper understanding of the latter, I must refer the reader to appropriate textbooks and lectures on solar coronal physics. ",
        "conclusions": ""
    },
    "0609/astro-ph0609340_arXiv.txt": {
        "abstract": "We report on an ambitious multi-site campaign aimed at detecting stellar variability, particularly solar-like oscillations, in the red giant stars in the open cluster M67 (NGC 2682). During the six-week observing run, which comprised 164 telescope nights, we used nine 0.6-m to 2.1-m class telescopes located around the world to obtain uninterrupted time-series photometry. We outline here the data acquisition and reduction, with emphasis on the optimisation of the signal-to-noise of the low amplitude (50--500$\\,\\mu$mag) solar-like oscillations. This includes a new and efficient method for obtaining the linearity profile of the CCD response at ultra high precision ($\\sim10\\,$parts per million). The noise in the final time series is $0.50\\,\\rmn{mmag}$ per minute integration for the best site, while the noise in the Fourier spectrum of all sites combined is 20$\\,\\mu$mag. In addition to the red giant stars, this data set proves to be very valuable for studying high-amplitude variable stars such as eclipsing binaries, W UMa systems and $\\delta\\,$Scuti stars. ",
        "introduction": "Asteroseismology of stellar clusters is potentially a powerful tool. The assumption of a common age, distance, and chemical composition provides stringent constraints on each cluster member, which significantly improves the asteroseismic output \\citep{GoughNovotny93}. Hence, detecting  oscillations in cluster stars in a range of evolutionary states holds promise of providing new tests of stellar evolution theory. Driven by this great potential, several studies have been aimed at detecting solar-like oscillations in the open cluster M67 \\citep{GillilandBrown88,Gilliland91,GillilandBrown92b,Gilliland93} and in the globular cluster M4 (Frandsen et al. in prep.). The most ambitious campaign was reported by \\citet{Gilliland93}, who used seven 2.5-m to 5-m class telescopes during one week in a global photometric network to target 11 turn-off stars in M67. Despite these efforts, they did not claim unambiguous detection of oscillations. However, one of their conclusions was that oscillations should be detectable in the more evolved red giant stars due to higher expected oscillation amplitudes. However, the oscillation periods of up to several hours and expected frequency separations of a few $10^{-6}\\,$Hertz would require a time base of roughly one month on these stars. Recent month-long studies using single- or dual-site high-precision radial velocity measurements ($\\sigma\\sim2\\,$m/s) on bright field stars have clearly demonstrated that solar-like oscillations are present in red giant stars \\citep{Frandsen02,Barban04,Ridder06}. However, due to non-continuous coverage these data suffered badly from aliasing in the Fourier spectrum, which complicated the detailed frequency analysis. Multi-site or space observations are therefore required \\citep{Stello06}. Such observations will hopefully soon become available for red giant stars in the field from the MOST, COROT and Kepler missions. However, after the cancellation of the ESA Eddington mission, no current or planned space project will measure stellar oscillations in cluster stars. Using high-precision spectrographs from ground to measure radial velocities in red giants is not possible due to the lack of a global network that can achieve high-precision velocities on an ensemble of relatively faint cluster stars. Hence, the only feasible approach is ground-based photometry.     In this investigation we again target M67 using multi-site photometric observations. Unlike the previous studies on this cluster, our primary targets are the red giant stars (see Fig.~\\ref{fig0}). Extrapolating the $L/M$-scaling relation \\citep{KjeldsenBedding95} predicts the amplitude of these stars to be in the range 50--500$\\mu$mag. Although very low, these amplitudes are significantly higher than for the turn-off stars targeted by, e.g., \\citet{Gilliland93}. \\begin{figure} \\includegraphics[width=84mm]{fig0.eps} \\caption{Colour-Magnitude diagram of the open cluster M67 (photometry by \\citet{Montgomery93}). The red giant target stars are indicated with filled symbols. The numbers  correspond to those indicated in  Fig.~\\ref{fig2}. The solid line is an isochrone ($(m-M)=9.7\\,\\rmn{mag}$, $\\rmn{Age}=4.0\\,\\rmn{Gyr}$, $Z=0.0198$ and $Y=0.2734$) from the BaSTI database \\citep{Pietrinferni04}.} \\label{fig0} \\end{figure} In addition to the red giants, more than 300 cluster stars were observed during the campaign. Many are high-amplitude variables such as W UMa systems, $\\delta\\,$Scuti stars and eclipsing binaries, some of which are already known. This campaign provides a unique data set to investigate these stars as well (Bruntt et al., in prep.).  With emphasis on the low-amplitude red giant stars, the main purpose of this paper is to describe the optimisation of time-series data towards achieving the highest possible signal-to-noise in the Fourier spectrum (in amplitude). Further discussion on the extraction of p-modes from the Fourier spectra of these stars will be presented by Stello et al. (in prep.).    \\section[]{Observations} We observed the open cluster M67 from 6 January to 17 February 2004 using nine telescopes (0.6-m to 2.1-m class) in a global multi-site network. The sites were distributed in longitude to allow continuous time series photometry during the six-week observing program. We were allocated 164 nights of telescope time which, due to bad weather, yielded about 100 clear nights (see Fig.~\\ref{fig1} and Table~\\ref{tab1}). In the first 18 days we observed 34\\% of the time and in the following three weeks the coverage was 80\\%. For the entire campaign (43 days) the coverage was 56\\%.  \\begin{figure} \\includegraphics[width=84mm]{fig1s.eps} \\caption{Time series of star No. 10 for all sites (after removing outliers and correcting for colour-dependent extinction; see Sect.~\\ref{colour}). The site abbreviations are explained in Table~\\ref{tab1}.} \\label{fig1} \\end{figure} The participating telescopes and detectors had different properties and the data sets are therefore rather diverse in terms of field-of-view (FOV), cadence and noise properties. A summary of the observations and the instrument characteristics for each site is given in Table~\\ref{tab1}. We indicate the smallest and largest FOV in Fig.~\\ref{fig2}. The red giant stars are indicated as well. %  Observations at each telescope were planned to optimise the signal-to-noise for solar-like oscillations in the red giant stars. We did that by calculating both the noise and the expected oscillation amplitudes (the signal) in different photometric filters. The amplitudes were estimated from \\citet[~Eq.~5]{KjeldsenBedding95} and the noise was estimated by photon counting statistics. These calculations showed that the Johnson B and V filters were favourable, but an on-site test was required to establish which of these was superior at each telescope. The observers therefore chose filters based on an initial test at the beginning of the first night. All sites except Kitt Peak chose the V filter. No phase change is seen in the solar oscillations between observations obtained in different filters over the range 400--700 nm \\citep{Jimenez99}. We therefore expect the same adiabatic behaviour for high-order solar-like oscillations in other stars as well. Hence, combining data based on different filters can be done after a simple rescaling of the amplitude, and corresponding adjustment of the weights to preserve the signal-to-noise.  To obtain a noise level which was essentially limited by scintillation and photon noise, it was important to avoid drift on the CCD of the stellar field. The aim was to have each star confined within a few pixels. Not all sites had autoguiding systems and as a result we found very different drift characteristics from site to site (see Fig.~\\ref{fig1a}). The images were defocused to obtain a higher duty cycle but we avoided crowding.  \\begin{table*} \\centering \\begin{minipage}{157mm} \\caption{Summary of observations.} \\begin{tabular}{@{}lrrrrrrrrrrr@{}} \\hline Site\\footnote{Site abbreviations and observers are: SSO$_1$ (Wide Field Imager at Siding Spring Observatory, Australia, Z.E.D., D.S., A.P.J. and L.L.K.); SSO$_2$ (Imager at Siding Spring, A.P.J., J.N. and D.S.); SOAO (Sobaeksan Optical Astronomy Observatory, Korea, S.-L.K., J.-A.L. and C.-U.L.); SAAO (South Africa Astronomical Observatory, T.A. and H.R.J.); RCC (Ritchey-Chr\\'etien-Coud\\'e at Piszk\\'estet\\H{o}, Konkoly Observatory, Hungary, J.N.); Sch (Schmidt at Piszk\\'estet\\H{o}, Z.C. and R.S.); LaS (La Silla Observatory, Chile, H.B.); LOAO (Mt. Lemmon Optical Astronomy Observatory, Arizona, Y.B.K. and J.-R.K.); Kitt (Kitt Peak National Observatory, Arizona, R.L.G.); Lag (Mt. Laguna Observatory, California, C.S. and M.Y.B.).} & Telescope & Filter & FOV    &  Image     & Observing & N$_{\\rmn{exp}}$ & Median  & Exp. & Duty & \\#nights & \\#nights \\\\%& Duty \\\\ & aperture  &        &        &  scale     &  time     &                 & cadence & time & cycle&  alloc.  & good     \\\\%& cycle\\\\ \\hline SSO$_1$& 1.0m & V & 14\\farcm0 & 0\\farcs38/pix &  43.0h & 2205 &  62s &  30s & 48\\%& 16 &  9 \\\\%& 56\\%\\\\ SSO$_2$& 1.0m & V & 14\\farcm0 & 0\\farcs60/pix &  51.0h & 1157 & 144s &  30s & 21\\%& 17 & 10 \\\\%& 59\\%\\\\ SOAO   & 0.6m & V & 20\\farcm5 & 0\\farcs60/pix &  33.6h &  467 & 240s & 120s & 50\\%& 17 &  7 \\\\%& 41\\%\\\\ SAAO   & 1.0m & V &  6\\farcm0 & 0\\farcs31/pix & 112.9h & 2595 & 149s &  80s & 54\\%& 28 & 22 \\\\%& 79\\%\\\\ RCC    & 1.0m & V &  7\\farcm0 & 0\\farcs29/pix &  23.6h &  722 &  89s &  70s & 79\\%& 14 &  5 \\\\%& 36\\%\\\\ Sch    & 0.6m & V & 17\\farcm0 & 1\\farcs10/pix &  31.3h & 1584 &  53s &  35s & 66\\%& 16 &  6 \\\\%& 38\\%\\\\ LaS    & 1.5m & V & 13\\farcm5 & 0\\farcs39/pix & 109.0h & 3945 &  90s &  24s & 27\\%& 22 & 18 \\\\%& 82\\%\\\\ LOAO   & 1.0m & V & 22\\farcm5 & 0\\farcs66/pix &  41.8h & 2886 &  46s &  12s & 26\\%& 14 &  7 \\\\%& 50\\%\\\\ Kitt   & 2.1m & B & 10\\farcm0 & 0\\farcs30/pix &  75.5h & 1563 & 172s &  52s & 30\\%& 11 &  9 \\\\%& 82\\%\\\\ Lag    & 1.0m & V & 14\\farcm0 & 0\\farcs41/pix &  46.4h & 1320 & 114s &  20s & 18\\%&  9 &  7 \\\\%& 78\\%\\\\ \\hline Total  &      &   &           &               & 561.1h &18444 &      &      &     &164 &100 \\\\%& 61\\%\\\\ \\hline \\vspace{-1cm} \\label{tab1} \\end{tabular} \\end{minipage} \\end{table*}  \\begin{figure*} \\includegraphics[width=120mm]{fig2s.eps} \\caption{M67 field-of-view (FOV). SAAO with the smallest FOV and LOAO with the largest are indicated with dashed squares. The red giant stars are marked with circles. Source: STScI Digitized Sky Survey.} \\label{fig2} \\end{figure*} \\begin{figure} \\includegraphics[width=84mm]{fig1as.eps} \\vspace{-1.5cm} \\caption{The position on the CCD of star No. 10 relative to a reference frame. The insets show the inner 20 by 20 pixels. Although autoguiding is good at LaS, instrument rotation introduced drift during the observing run for stars far from the rotation axis.} \\label{fig1a} \\end{figure}   The exposure time at each site was adjusted to have star No. 4 just below the saturation limit, which provided a safety margin for the large group of clump stars that were $0.3\\,$mag fainter (see Fig.~\\ref{fig0}). However, the two brighter stars (No. 3 and 11) were therefore often saturated. Due to their expected long oscillation periods (Stello et al., in prep.), on time scales similar to typical instrumental drift, the results on these stars were likely to provide only limited scientific output. Keeping these two stars above the saturation limit resulted in lower noise for the stars at the base of the red giant branch (stars 6, 13, 12, 14) which were more likely to produce useful results.   \\section[]{Calibration} We calibrated each CCD image using four steps: \\begin{enumerate} \\item overscan subtraction (not all CCDs had an overscan region), \\item subtraction of bias (the bias levels were stable enough to use a single master bias image for each CCD), \\item correction for non-linearity (see Sect.~\\ref{linearity}), \\item flat-fielding to correct for pixel-to-pixel variations in the quantum efficiency (we used one master flat field for each CCD; for Kitt Peak and the RCC this was based on dome flats, while we used sky flats for all other sites). \\end{enumerate} We found that the dark current was negligible compared to the read-out-noise for all sites and it was therefore ignored. These four steps were standard except the non-linearity correction, which is described in more detail in the following section.  \\subsection{CCD linearity calibrations}\\label{linearity}  In this project a few target stars were relatively close to the CCD saturation limit, at flux levels for which the non-linear response of the CCD gain could be significant. It is important to correct for these gain variations to attain the high photometric precision required by this project.  Linearity at high flux levels was investigated for all CCDs using a ``classical'' linearity test. We used an approach similar to that described by \\citet{Gilliland93}. The method measures relative variations in the CCD amplifier gain, rather than absolute calibration in terms of $e^-$/ADU. Flat-field images were taken sequentially with increasing exposure time, interleaved with reference images; e.g. $3\\,$s, $10\\,$s, $3\\,$s, $10\\,$s,$\\cdots$, $3\\,$s, then followed by $3\\,$s, $20\\,$s, $3\\,$s, $20\\,$s,$\\cdots$, $3\\,$s, until the final series, in which the longer exposures were almost saturated. The reference images allowed instabilities of the light source to be measured and removed. In some cases, however, the flat-field lamp varied on time scales too short to be sampled and a correction could therefore not be made. The mean counts in the flat-field image, scaled according to the exposure time, were plotted versus the mean counts. Images with the same exposure time were grouped to form a single point, with an uncertainty equal to the group rms. In Fig.~\\ref{fig2a} we show the results of the linearity tests for all CCDs. Saturation occurred at $65536\\,$ADU except for the Schmidt, where it was at $16384\\,$ADU. \\begin{figure} \\includegraphics[width=84mm]{fig2a.eps} \\caption{Classical linearity tests for all CCDs. For SSO$_1$, SSO$_2$ and SAAO two separate tests have been merged. Error bars are plotted as $3\\times$rms to make them visible. The solid lines are polynomial fits used to correct the data for non-linearity. } \\label{fig2a} \\end{figure} From Fig.~\\ref{fig2a} (upper right panel) we see that non-linearity can introduce variability in the stellar time series up to several percent in non-photometric conditions (variable atmospheric transparency). Using ensemble photometry (Sect.~\\ref{ensemble}) will, however, reduce this effect if the ensemble stars are roughly of equal colour and luminosity. We decided to correct for non-linearity for all sites except the Schmidt, which did not show measurable non-linear effects, and Laguna, where we only had measurements of the gain variations up to approximately 30000 ADU, which was significantly lower than the intensity levels of most target stars. For LOAO, the data did not justify a description of the gain variation to be higher than a first-order polynomial fit, although a few points could indicate that higher-order features were present. The linearity calibration of the data from La Silla was based on a method described in Sect.~\\ref{newmethod}, hence no fit was applied to the data shown in Fig.~\\ref{fig2a}.  We obtained calibration images at La Silla for a new and more elegant method for determining the linearity properties. This method provides a much more precise determination of the CCD gain variations, which we will compare with the classical method in the next section.  \\subsubsection{Ultra-high-precision method}\\label{newmethod} The basic concept of the method described in this section was first outlined by \\citet{Baldry99} and \\citet{Knudsen00}, and was developed into a fully applicable method by \\citet{Stello02}. Like the classical method, this method measures relative variations in the CCD amplifier gain using flat-field images of different exposure times, but it differs by using flat fields that have a strong gradient e.g. by using a grism. Each flat field showed a large smooth variation in light level from approximately the bias level to a significant fraction of the saturation limit, with  the longest exposure reaching saturation (see Fig.~\\ref{fig3}). \\begin{figure} \\includegraphics[width=84mm]{fig3.eps} \\caption{Linearity flat field of the longest exposure, $410\\,$s (see text). The flat plateau at high row numbers is due to the digital saturation of the CCD.} \\label{fig3} \\end{figure} The advantage of this method is that the effect from instabilities in the light source used to obtain the flat fields is very small, because we are sampling a large range (in the longest exposures the entire range) of the CCD gain response in a single exposure. The resulting measurement  precision of the gain variations is several orders of magnitude better than the classical method. Further, this method requires relatively few images to achieve high precision, making it very efficient.  We obtained spectral flats using the DFOSC spectrograph on the Danish 1.54-m telescope (La Silla). Light variation in one direction across the CCD was achieved using a grism to disperse the light from the slit illuminated by an internal telescope calibration lamp (Fig.~\\ref{fig3}). A series of images were acquired in the following way: $3\\times30\\,$s, $3\\times130\\,$s, $3\\times30\\,$s, $3\\times250\\,$s, $3\\times30\\,$s, $3\\times370\\,$s, $3\\times30\\,$s, $3\\times390\\,$s, $3\\times30\\,$s, $3\\times410\\,$s, $3\\times30\\,$s. The control exposures of $30\\,$s enabled long-term drift in the flat field-lamp to be removed. Although this improves the precision, it is not critical. After subtraction of overscan and bias, we corrected for the long-term drift of the flat-field lamp and made an average (master) flat field for each exposure time. We then collapsed each master flat field by averaging in one direction to obtain one dimensional \\textit{intensity curves}, as shown in Fig.~\\ref{fig4}. \\begin{figure} \\includegraphics[width=84mm]{fig4.eps} \\caption{Two intensity curves of different exposure times, $410\\,$s and $130\\,$s (see text). The flat part of the $410\\,$s exposure at high row number is due to saturation of the CCD.} \\label{fig4} \\end{figure}  Due to non-linear effects in the CCD, the intensity we measure in the $n$th row is: \\begin{eqnarray} I(n) &=& T_{\\rmn{exp}} \\cdot S(n) \\cdot g(I(n)), \\label{row_intens} \\end{eqnarray} where $T_{\\rmn{exp}}$ is the effective exposure time in seconds (after correcting for dead-time of the shutter and short-term fluctuations in the light source), $g(I(n))$ is the CCD amplifier gain as a function of the measured intensity and $S$ defines the shape of the intensity curves, hence $S(n)$ expresses the intensity in the $n$th row from a 1-s exposure. Although the actual gain variations are a function of the received flux, we assume the CCD amplifier has a well-defined output signal for every input signal.  From two intensity curves with different exposure times, say $410\\,$s and $130\\,$s, we constructed the relative curve \\begin{equation} { I_{410\\rmn{s}}(n) \\over I_{130\\rmn{s}}(n) } = { 410\\,\\rmn{s} \\over 130\\,\\rmn{s} } \\cdot { g(I_{410\\rmn{s}}(n)) \\over g(I_{130\\rmn{s}}(n)) }, \\label{rel_gain} \\end{equation} \\noindent where we have corrected for shutter dead-time and verified that the shape $S$ was stable. In Fig.~\\ref{fig6} (top panel) we show the curve $I_{410\\rmn{s}}(n)/I_{130\\rmn{s}}(n)$ versus $I_{410\\rmn{s}}(n)$, which we call a \\textit{gain-ratio curve}. At each intensity, $I_1$, this curve shows the gain ratio, $g(I_1)/g(I_2)$, where $I_2=I_1\\cdot130\\,\\rmn{s}/410\\,\\rmn{s}$. Thus, finding $g(I)$ for any intensity level requires inversion of the gain-ratio curve in an iterative process.  \\begin{figure} \\includegraphics[width=84mm]{fig6.eps} \\caption{\\textbf{Top panel:} Gain-ratio curve based on two intensity curves of different exposure time ($410\\,$s and $130\\,$s). \\textbf{Bottom panel:} Inversion from smoothed gain-ratio curve $R_1(I)$ (dotted) to final gain-curve $g(I)$ (solid). For comparison, the measurements from the classical method (Fig.~\\ref{fig2a}, LaS) and their error bars ($1\\times$ rms) are indicated. The uncertainty on $g(I)$ is of the order $10\\,$ppm. The smoothing of $R_1(I)$ introduces a systematic error in $g(I)$ of roughly 1\\% near the saturation limit because the curve is steep.} \\label{fig6} \\end{figure} We started out using a smoothed version of the measured gain-ratio curve, which we denoted $R_1(I)$, as a first estimate for the actual underlying gain curve $g(I)$ (see Fig.~\\ref{fig6}, bottom panel). Then, assuming $g_1(I)=R_1(I)$, we computed a new gain-ratio curve, $R_2(I)=g_1(I)/g_1(I\\cdot 130\\,\\rmn{s}/410\\,\\rmn{s})$ for all $I$. The new estimate for the gain was then corrected by the relative deviation between $R_{1}$ and $R_{2}$ according to $g_2(I)={ R_{1} \\over R_{2} }g_1(I)$, etc. This iterative process stopped when $R_{i}$ matched the measured $R_{1}$ and the corresponding $g_i(I)$ was the desired gain curve. Before smoothing the measured gain-ratio curve (Fig.~\\ref{fig6}, top panel) we removed the noise at low intensities (0--1000 ADU) by replacing it with a linear fit to the data points from 1000 to 10000 ADU. In the example shown, the gain-ratio curve is very similar to the final gain curve. However, this is not a feature of the method, but is due to the gain characteristics of the particular CCD amplifier.  To ensure that all features in the gain were detected, we examined gain-ratio curves based on different exposure-time ratios. If a feature, say a bump in the CCD gain, is periodic for increasing intensity and hence repeated at all pairs of intensities ($I_1,I_2$) related as $I_2=I_1\\cdot 410\\,\\rmn{s}/130\\,\\rmn{s}$, it will not show up in the gain-ratio curve based on $410\\,$s and $130\\,$s exposures (or any combination with the same exposure-time ratio). Our gain curves based on flat fields with different exposure-time ratios, all showed an excellent match within the errors.  In Fig.~\\ref{fig6} (bottom panel) we compare our new method with results from the classical method for La Silla (Fig.~\\ref{fig2a}). The two methods are in  agreement with each other, but the series of flat-field images for our new method is significantly faster to obtain, provides the relative gain for all intensities and has a precision more than 100 times better. However, it requires temporally stable but spatially variable illumination of the CCD (e.g. spectral flat fields) which is not possible at every telescope. ",
        "conclusions": "We have collected 100 telescope nights of photometric multi-site data of the open stellar cluster M67 over a six-week period. The focus of this paper was the discussion of our approach towards achieving the highest signal-to-noise ratio for the very low amplitude (50--500$\\,\\mu$mag) solar-like oscillations in the red giant stars. This included our careful reduction of the CCD images to obtain the lowest possible noise in the time-series data, and our weighting scheme to reduce the noise level in the Fourier spectrum.   We have obtained a point-to-point scatter in the time-series photometry down to about 1$\\,$mmag for most sites, while the largest participating telescope reached 0.5$\\,$mmag (Fig.~\\ref{fig9}). These values are similar to those from previous campaigns on M67 by \\citet{GillilandBrown88}, \\citet{Gilliland91} and \\citet{GillilandBrown92a} which all used telescopes of similar size but for shorter time spans (maximum two weeks). Comparison of our best nights with known noise terms demonstrates that the attained point-to-point scatter is consistent with irreducible terms dominated by photon and scintillation noise for all sites but one (RCC), which shows extra noise of unknown origin (Fig.~\\ref{fig9}). With these scatter values, our weighting scheme provided a mean noise level in the Fourier spectra of approximately 20$\\,\\mu$mag (in amplitude), which would allow us to detect solar-like oscillations in the red giant stars with $\\rmn{S/N}\\gid4$ assuming $L/M$-scaling."
    },
    "0609/astro-ph0609176_arXiv.txt": {
        "abstract": "We performed spectral line observations of \\cojto, \\tcojoz, and \\ceojoz\\ and polarimetric observations in the $\\lambda=1.3 \\textrm{ mm}$ continuum and \\cojto\\ toward a multiple protostar system, L1448 IRS 3, in the Perseus molecular complex at a distance of $\\sim 250 \\textrm{ pc}$, using the BIMA array. In the $\\lambda=1.3 \\textrm{ mm}$ continuum, two sources (IRS 3A and 3B) were clearly detected with estimated envelope masses of $0.21$ and $1.15 \\textrm{ \\msun}$, and one source (IRS 3C) was marginally detected with an upper mass limit of $0.03 \\textrm{ \\msun}$. In \\cojto, we revealed two outflows originating from IRS 3A and 3B. The masses, mean number densities, momentums, and kinetic energies of outflow lobes were estimated. Based on those estimates and outflow features, we concluded that the two outflows are interacting and that the IRS 3A outflow is nearly perpendicular to the line of sight. In addition, we estimated the velocity, inclination, and opening of the IRS 3B outflow using Bayesian statistics. When the opening angle is $\\sim 20\\arcdeg$, we constrain the velocity to $\\sim 45 \\textrm{ \\kms}$ and the inclination angle to $\\sim 57\\arcdeg$. Linear polarization was detected in both the $\\lambda=1.3 \\textrm{ mm}$ continuum and \\cojto. The linear polarization in the continuum shows a magnetic field at the central source (IRS 3B) perpendicular to the outflow direction, and the linear polarization in the \\cojto\\ was detected in the outflow regions, parallel or perpendicular to the outflow direction. Moreover, we comprehensively discuss whether the binary system of IRS 3A and 3B is gravitationally bound, based on the velocity differences detected in \\tcojoz\\ and \\ceojoz\\ observations and on the outflow features. The specific angular momentum of the system was estimated as $\\sim 3\\times 10^{20} \\textrm{ cm}^2 \\textrm{ s}^{-1}$, comparable to the values obtained from previous studies on binaries and molecular clouds in Taurus. ",
        "introduction": "\\label{sec_intro}  L1448 is a dark cloud located approximately one degree southwest from NGC 1333 in the Perseus cloud complex at a distance of $\\sim 250$ pc \\citep[e.g.][]{enoch2006}. Three infrared sources were observed by the Infrared Astronomical Satellite (IRAS) and denoted as IRS 1, IRS 2, and IRS 3 by \\citet{bachiller1986}. Due to its brightness in the IRAS bands, IRS 3 has been focused on more than the others. Meanwhile, \\citet{bachiller1990, bachiller1995} revealed a well-collimated, large outflow originating from L1448-mm, located 70$''$ southeast of L1448 IRS 3. In fact, IRS 3 is overlapped with the blueshifted lobe of the outflow. \\citet{curiel1990} detected the L1448-mm continuum source and two sources separated by $7''$ in the L1448 IRS 3 region, using the Very Large Array (VLA) at $\\lambda=6$ and 2 cm wavelengths. They named the sources L1448 C (for Center) and L1448 N(A) and N(B) (for North) for L1448-mm and the two sources in IRS 3, respectively. \\citet{terebey1997} detected another source 20$''$ northwest of the two IRS 3 sources at $\\lambda=2.7$ mm continuum, using the Owens Valley Radio Observatory (OVRO) millimeter interferometer. \\citet{looney2000} resolved all three sources in IRS 3 with high resolution Berkeley Illinois Maryland Association (BIMA) observations at $\\lambda=2.7$ mm continuum. In this paper, we call these sources L1448-mm and L1448 IRS 3A, 3B, and 3C, using IAU nomenclature.  In addition to the multiplicity of young sources in L1448 IRS 3, multiple outflows in the region have been suggested in previous studies. \\citet{bachiller1990, bachiller1995} suggested an outflow originating from IRS 3B, since they detected a redshifted component in the blueshifted lobe of the mm-source outflow and since IRS 3B is the brightest source at mm wavelengths. In contrast, \\citet{curiel1990} suggested that the outflow was driven by IRS 3A, based on a spectral index of 0.2, similar to thermal jet models \\citep[e.g.][]{reynolds1986}, and coincident H$_2$O maser observations. \\citet{davis1995} and \\citet{eisloffel2000} presented H$_2$ emission images showing shocked regions. They suggested up to three outflows in the IRS 3 region from collimated features in the images. Several Herbig-Haro objects were also detected and explained as outflows driven by the three IRS 3 sources \\citep{bally1997}. \\citet{wolfchase2000} suggested two outflows from IRS 3A and 3B, based on previous studies and their large-scale maps of CO $J=1\\rightarrow0$\\ emission. Moreover, \\citet{girart2001} reported a well collimated redshifted component detected along a line of position angle $110 \\arcdeg$ from IRS 3B in SiO $J=2\\rightarrow1$ emission. These previous studies, however, could not clearly show the outflow features relative to the sources due to a lack of angular resolution and the inherent complexity of the region.  Theoretical studies have suggested that magnetic fields play key roles in the outflows of protostars as well as star formation itself. Observations of magnetic morphology and strength are possible using the Zeeman effect \\citep[e.g.][]{crutcher1999} and linear polarization of dust emission and spectral lines \\citep[e.g.][]{girart1999}. \\citet{crutcher2004} recently estimated the magnetic strength as well as morphology, from linear polarization of dust emission in three prestellar cores (starless cores) using the Chandrasekhar-Fermi method \\citep{chandrasekhar1953}. In the case of magnetic fields related to protostars (Class 0 sources) with outflows, \\citet{girart1999} detected linear polarization in the $\\lambda=1.3$ mm dust emission and CO $J=2\\rightarrow1$\\ spectral line. Their study was unique for low mass protostars with outflows.  In this paper, we present polarimetric observations in the $\\lambda=1.3$ mm continuum and CO $J=2\\rightarrow1$\\ showing two outflows and magnetic fields in the L1448 IRS 3 region, using the BIMA array. Moreover, \\tco~$J=1\\rightarrow0$ and \\ceo~$J=1\\rightarrow0$ observation results are presented. We argue that the binary system of L1448 IRS 3A and 3B is gravitationally bound, using these \\tco~and \\ceo~observations. ",
        "conclusions": "We present CO $J=2\\rightarrow1$, $^{13}$CO $J=1\\rightarrow0$\\ and C$^{18}$O $J=1\\rightarrow0$\\ observations, and $\\lambda=1.3$ mm continuum and CO $J=2\\rightarrow1$\\ polarimetric observations. IRS 3A and 3B are distinctly detected with mass estimates of $0.21$ and $1.15 \\textrm{ \\msun}$ respectively at $\\lambda=1.3$ mm, but IRS 3C is marginally detected with upper mass limit of $0.03 \\textrm{ \\msun}$ (\\S \\ref{sec_cont}). The ambient velocities of IRS 3A, 3B, and 3C are estimated as $5.5$, $4.5$, and $3.5 \\textrm{ \\kms}$, respectively, from \\tco~$J=1\\rightarrow0$\\ and \\ceo~$J=1\\rightarrow0$\\ channel maps (\\S \\ref{sec_13coc18o}). The two close sources, IRS 3A and 3B, have a velocity difference less than $1 \\textrm{ \\kms}$; the difference is a kinematical constraint on a gravitationally bound binary system. Moreover, we estimated the specific angular momentum of the binary system as $\\sim 3\\times10^{20}~\\textrm{cm}^2~\\textrm{s}^{-1}$, similar to the upper limit of binaries and to the lower limit of molecular clouds in Taurus (\\S \\ref{sec_binary}).  We present CO $J=2\\rightarrow1$\\ observations showing two outflows, one each from IRS 3A and 3B (\\S \\ref{sec_outflows}). The outflow driven by IRS 3A has $\\textrm{PA} = 155 \\arcdeg$ and is nearly perpendicular to the line of sight, while the outflow by IRS 3B has $\\textrm{PA}=105 \\arcdeg$. In addition, we posit that the two outflows are interacting in the southern lobe of the IRS 3A outflow and the redshifted lobe of the IRS 3B outflow, based on a comparison of the kinetic energies of lobes. Coupled with the fact that the linear polarization detected in CO $J=2\\rightarrow1$\\ is likely to come from the IRS 3B outflow, IRS 3B is closer than IRS 3A. We also detected that the IRS 3B outflow has accelerating and non-accelerating features in the velocity-position diagram, which is interpreted as either a precessing outflow or a ``trumpet'' outflow. Assuming a ``trumpet'' outflow of IRS 3B rather than its precession, we estimated the velocity, inclination angle, and the opening angle, using Bayesian statistics. The velocity and the inclination angle are constrained between $100$ and $40 \\textrm{ \\kms}$ and between $75\\arcdeg$ and $50\\arcdeg$, respectively, as the opening angle between $8\\arcdeg$ and $26 \\arcdeg$. Furthermore, using an opening angle of $\\sim 20\\arcdeg$ from Spitzer Space Telescope observations, the velocity and the inclination angle of the IRS 3B outflow are $\\sim 45 \\textrm{ \\kms}$ and $\\sim 57\\arcdeg$.  Linear polarization in both the $\\lambda=1.3 \\textrm{ mm}$ continuum and CO $J=2\\rightarrow1$\\ spectral line is marginally detected around the center and outflow of IRS 3B, respectively (\\S \\ref{sec_magfields}). The dust emission polarization gives a magnetic field perpendicular to the outflow, which may be a toroidal magnetic field parallel to the circumstellar disk. In contrast, the spectral line polarization suggests either a perpendicular or a parallel magnetic field to the IRS 3B outflow. To determine the relation between the magnetic field and the outflow direction, more sensitive polarimetric observations are required.  The L1448 IRS 3 is an excellent region to study star formation. The IRS 3A outflow, nearly perpendicular to the line of sight, enables the study of the disk and outflow structures in protostar systems more easily, since it shows the profile projected in the plane of the sky. In addition, the binary system of IRS 3A and 3B having two outflows in quite different directions gives an opportunity to study the interaction between two sources as well as constrain binary system formation itself. Finally, more sensitive polarimetric observations will provide clues on the connection between outflows and magnetic fields."
    },
    "0609/astro-ph0609448_arXiv.txt": {
        "abstract": "We present five recycled pulsars discovered during a 21-cm survey of approximately 4,150 deg$^2$ between 15$\\arcdeg$ and 30$\\arcdeg$ from the galactic plane using the Parkes radio telescope.  One new pulsar, PSR~J1528$-$3146, has a 61\\,ms spin period and a massive white dwarf companion. Like many recycled pulsars with heavy companions, the orbital eccentricity is relatively high ($\\sim$0.0002), consistent with evolutionary models that predict less time for circularization. The four remaining pulsars have short spin periods ($3\\,{\\rm ms} < P < 6\\,{\\rm ms}$); three of these have probable white dwarf binary companions and one (PSR J2010$-$1323) is isolated. PSR J1600$-$3053 is relatively bright for its dispersion measure of 52.3\\,pc\\,cm$^{-3}$ and promises good timing precision thanks to an intrinsically narrow feature in its pulse profile, resolvable through coherent dedispersion.  In this survey, the recycled pulsar discovery rate was one per four days of telescope time or one per 600 deg$^2$ of sky.  The variability of these sources implies that there are more millisecond pulsars that might be found by repeating this survey. ",
        "introduction": "\\label{sec:intro}  Most pulsars are thought to descend from massive stars in the disk of the galaxy, and as a result, the bulk of pulsar search efforts have historically been concentrated near the galactic plane where the majority of pulsars, with their fairly short observable lifetimes, reside \\citep[e.g.][]{sstd86,clj+92,jlm+92,mlc+01}.  Early surveys along the plane met with limited success when searching for recycled pulsars.  The value of finding new recycled pulsars comes from their contribution to our understanding of the binary evolution processes \\citep{bv91}, tests of general relativity \\citep{vbb+01}, and their use in a millisecond pulsar (MSP) timing array for the detection of low frequency gravitational waves \\citep{fb90}.  The pioneering survey of \\cite{wol91a} led to the discovery of two fascinating systems in a small area at high galactic latitude.  This success spawned several surveys at high galactic latitudes, which yielded a large number of recycled pulsars \\citep{bhl+94,fcwa95,lnl+95,nll+95,cnt96}.  Simulations of the low-luminosity galactic recycled pulsar population point to a more isotropic distribution than is expected for the more luminous long-period pulsars that are less subject to the deleterious effects of the interstellar medium \\citep{jb91}.  This difference is because recycled pulsars have had a longer time to migrate away from their birthplace in the galactic disk and their shorter periods strongly limit the distance to which they can be detected in the electron-rich galactic plane.  These facts, combined with the relative insensitivity to dispersion and lower sky temperatures afforded by high-frequency observations, suggested that a 21\\,cm survey for pulsars away from the galactic plane would be extremely productive \\citep{tbms98}.  The success of the Swinburne Intermediate Latitude Pulsar Survey demonstrated the validity of this approach, discovering eight recycled pulsars in $\\sim 2950$\\,deg$^2$ at galactic latitudes between 5$^\\circ$ and 15$^\\circ$ \\citep{eb01, eb01b, ebvb01}. Recent re-processing of the Parkes multibeam survey which concentrated on the galactic plane ($\\left| b \\right| < 5^\\circ$) has also yielded millisecond and recycled pulsars in large numbers \\citep{fsk+04}.  In this paper we describe five recycled pulsars discovered in this extension of the Swinburne Intermediate Latitude survey.  In section \\ref{sec:survey} we describe the survey parameters before discussing the five pulsars individually in section 3, giving their timing solutions (where possible), pulse profiles and derived parameters. In section 4 we discuss our results. ",
        "conclusions": "This survey has once again demonstrated the efficacy of the Parkes 64\\,m telescope's multibeam receiver for discovering recycled pulsars. The recycled pulsar discovery rate of one per 600 deg$^2$ is comparable to that of Edwards et al. (2001).  However, only one new recycled pulsar was discovered more than 25$\\arcdeg$ and less than 30$\\arcdeg$ from the plane, compared to the five between 5$\\arcdeg$ and 10$\\arcdeg$ from the plane, which may indicate that the density of MSPs detectable with this backend and telescope combination drops off considerably when more than 25$\\arcdeg$ from the galactic plane.  The ratio of the number of recycled pulsars still in binaries compared to those that are now single puts a limit on the fraction of systems where recycling goes out of control, resulting in a solitary pulsar either because the pulsar gets too close to the companion and ablates it with its strong wind, or tidally disrupts the companion completely. In this survey 6 of the 7 recycled pulsars are in binary systems, which in itself suggests that in the majority of cases recycling is not fatal to the donor star. Thanks to the Parkes multibeam surveys the disk population of millisecond pulsars is now becoming large enough for us to search for trends that might explain why some millisecond pulsars are single and others are not, and which parameters are correlated. The recycled pulsars found in this survey have reinforced the previously observed trend of pulsars with high-mass companions having longer spin periods and higher eccentricities than those with low-mass companions.  Large-scale uniform surveys allow us to model the underlying populations' properties, including the $z$-scale height of MSPs. There are 11 field MSPs (defined for our purposes as having $P < 20$\\,ms) with $|z| > 0.5$\\,kpc in the ATNF Pulsar Catalogue\\footnote{http://www.atnf.csiro.au/research/pulsar/psrcat/} \\citep{mht+05}, with $DM$-based distances provided by the NE2001 model \\citep{cl02} for pulsars without direct distance estimates.  Of these pulsars, 4 were discovered by the Parkes 70\\,cm survey \\citep{mld+96, lml+98}, and none by the Swinburne surveys (though we would have expected these surveys to have similar sensitivities to high-$|z|$ pulsars).  The Arecibo surveys that are responsible for most of the rest are difficult to model.  To investigate the implications of the large-area Parkes surveys we simulated a population of 100,000 MSPs with a wide range in luminosity and a scale height of 500\\,pc.  Our simulation predicted that half of all discoveries should have $|z| > 0.5$\\,kpc for the Parkes 70cm and the Swinburne surveys. In reality, only about 13\\% of the MSPs detected have $|z| >$ 0.5\\,kpc, which either suggests that 0.5 kpc is an overestimate of their scale height, or our modeling of their luminosities is incorrect.  We note that somewhat different results are obtained using the older \\citet{tc93} Galactic electron model to estimate distances.  Based on this model, there are 17 MSPs with $|z| > 0.5$\\,kpc, 5 of which were discovered in the Parkes 70\\,cm survey and 5 of which were discovered by the Swinburne surveys.  This suggests that values of the $z$ scale height from previous population studies based on the older distance model, such as the $0.65^{+0.16}_{-0.12}$\\,kpc Gaussian scale height of \\citet{cc97}, will be revised downward by the use of the new Galactic electron model.  Again restricting ourselves to disk pulsars with $P < 20$\\,ms, we find 15 isolated MSPs and 44 in binary systems.  The distribution of height above the Galactic plane is shown for these two populations in Figure \\ref{fig:zhist}.  The RMS deviation from $z = 0$ is $240 \\pm 40$\\,pc for isolated MSPs and $450 \\pm 50$\\,pc for binary MSPs (if we consider only pulsars with $P < 10$\\,ms the isolated sample is unchanged, while the binary sample drops to 37 objects with RMS in $z$ of $470 \\pm 60$\\,pc).  These results are similar to those obtained by \\citet{lkn+06} using a sample of 9 isolated and 20 binary pulsars with $P < 10$\\,ms.  The median distance to the isolated and binary MSPs in our sample is 960\\,pc and 1240\\,pc respectively, regardless of the period cutoff used.  The difference between these two populations is less than that noted by \\citet{lkn+06}, who found more than a factor of two difference in their sample.  This result may suggest that the luminosity distributions of isolated and binary MSPs are not as different as previously thought, with implications for the birth velocity distributions required by the different observed scale heights.  A flux monitoring campaign on a large population of MSPs including those discovered here should enable direct investigation of the tantalising scale height---luminosity question for isolated and binary MSPs. If the isolated MSPs really do have lower luminosities we would expect them to have had a different accretion history than those still in binary systems.  If it is the tight binary MSPs that evolve into isolated MSPs through either ablation or tidal destruction of their companions, then we might expect different underlying velocity distributions and therefore different scale heights for the isolated and binary populations."
    },
    "0609/astro-ph0609481_arXiv.txt": {
        "abstract": "\\noindent The present matter density of the Universe, while highly inhomogeneous on small scales, displays approximate homogeneity on large scales. We propose that whereas it is justified to use the Friedmann-Lema\\^ itre-Robertson-Walker (FLRW) line element (which describes an exactly homogeneous and isotropic universe) as a template to construct luminosity distances in order to compare observations with theory, the \\emph{evolution} of the scale factor in such a construction must be governed not by the standard Einstein equations for the FLRW metric, but by the modified Friedmann equations derived by Buchert \\cite{buch1,buch4} in the context of spatial averaging in Cosmology. Furthermore, we argue that this scale factor, defined in the spatially averaged cosmology, will correspond to the effective FLRW metric \\emph{provided} the size of the averaging domain coincides with the scale at which cosmological homogeneity arises. This allows us, in principle, to compare predictions of a spatially averaged cosmology with observations, in the standard manner, for instance by computing the luminosity distance versus red-shift relation. The predictions of the spatially averaged cosmology would in general differ from standard FLRW cosmology, because the scale-factor now obeys the modified FLRW equations. This could help determine, by comparing with observations, whether or not cosmological inhomogeneities are an alternative explanation for the observed cosmic acceleration. ",
        "introduction": "The goal of Cosmology is to describe the evolution history of the Universe as a whole, tracing it as far back as possible from the present epoch. This seemingly daunting task is made tractable by the fact that the observed matter density in the Universe appears to be approximately homogeneous when viewed on the largest scales. Further, observations of the Cosmic Microwave Background (CMB) radiation reveal it to be isotropic (after subtracting a dipole contribution due to our local motion) down to the level of one part in $10^5$. The twin conditions of homogeneity and isotropy, if valid on all length scales, would immensely simplify the form of the metric and matter tensors used in the General Relativistic description of Cosmology. It can be shown \\cite{FRWref} that on purely geometric grounds, the most general metric describing a homogeneous and isotropic universe takes on the form of the Friedmann-Lema\\^ itre-Robertson-Walker (FLRW) line element. The standard method of proceeding is to assume that the geometry of the universe can be described by a background metric with small perturbations, where the background is exactly homogeneous and isotropic and is hence of the FLRW form. (For a detailed discussion of the assumptions involved in the standard Cosmology, see Ref. \\cite{ras}.)  This approach is adopted not only for the early universe, for example in describing Big Bang nucleosynthesis and the dynamics of the CMB anisotropies, but is also assumed to be valid at recent epochs. In particular, the luminosity distances constructed to compare data from Type Ia Supernovae (SNe) with theory, assume a metric of the FLRW form with its dynamics given by the Einstein equations for a homogeneous and isotropic universe. We know, however, that the matter content of the Universe today is \\emph{not} homogeneous on small scales - one finds inhomogeneity at the scale of individual galaxies for example; and there is evidence \\cite{voids} that voids of order $30h^{-1}$ Mpc\\footnote{Here $h$ is the conventional parameter which appears in the definition of the Hubble constant $H_0=100h$ kms$^{-1}$Mpc$^{-1}$.} in diameter account for $40$-$50\\%$ of the observed Universe. There is also evidence for voids 3-5 times this size \\cite{voids2}, as well as local voids on smaller scales \\cite{voids3}. The evidence for homogeneity of the matter distribution in the Universe only appears at scales larger than about $100h^{-1}$ Mpc or so \\cite{homog}. It is then reasonable to assume that the metric describing our Universe must also reflect the inhomogeneity of the matter on the corresponding scales. One may now ask whether it is at all justified to use the luminosity distances constructed using the FLRW metric with its associated dynamics to compare theory with observational data from SNe. In this paper we argue that the use of such luminosity distances is indeed justified, provided we are only interested in using such constructions in probing length scales larger than a certain minimum scale which is fixed by requiring that the assumption of homogeneity be valid (a lower bound on such a scale could be $100h^{-1}$ Mpc for example). More importantly, we propose that the \\emph{dynamics} of this (effective or template) macroscopic metric must be governed not by the standard Einstein equations applied to the FLRW metric, but by the effective equations obtained after spatial averaging on constant time slices, as developed by Buchert \\cite{buch1,buch4}.  An outstanding open issue has been the following: how does one compare the predictions of the inhomogeneous, spatially averaged cosmology with observations, and should one be attempting an interpretation in terms of spatially averaged variables in the first place? For instance, one could study light propagation and determine the luminosity distance versus redshift (\\dlz) relation which best fits the observations, within an \\emph{inhomogeneous} model of the Universe. Following the important and interesting early work of Dyer and Roeder, where the angular-size and luminosity distances were obtained for beams propagating in \"lumpy\" FLRW universes \\cite{dyer-roeder}, detailed investigations have been carried out in \\cite{inhomdlz}.  Such a program has also received considerable attention in the context of the spherically symmetric inhomogeneous Lema\\^itre-Tolman-Bondi (LTB) models \\cite{LTB,cel}. However, it is not clear to us that this is necessarily the best way to proceed, at least as far as distant SNe are concerned. A realistic computation of light propagation in an inhomogeneous geometry seems difficult at best, and while the use of simplified toy models such as the LTB solutions in this endeavor may be justified on the grounds that we may inhabit a locally underdense region \\cite{voids2}, it is not clear that the assumptions justifying the use of the toy model remain valid at large scales (see also Ref. \\cite{cel} for a recent review of the LTB approach). And after all, the concordance model infers the presence of a dark energy by comparing the predictions of the standard homogeneous and isotropic FLRW cosmology against observations \\cite{riess}. If one decides to interpret the FLRW cosmology as a template against which observations may be compared, then in order to decide whether cosmic inhomogeneities could be an alternative to the concordance model, one must construct a homogeneous and isotropic cosmology, after a suitable smoothing procedure, and compare its predictions with observations such as the \\dlz\\ relation for Type Ia SNe.  One should note that the effects of averaging within the standard perturbed FLRW framework of Cosmology, have been studied by several authors (see, e.g., Refs. \\cite{hirata,gruzinov}), and it has been argued that these effects must necessarily be small. The papers in Refs. \\cite{hirata,gruzinov} however, do not necessarily address the entire problem. For example, Ref. \\cite{hirata} deals with only super-horizon fluctuations, whereas work done by others (see, e.g, Ref. \\cite{kolb}) indicates that there may be nontrivial contributions due to averaging from \\emph{sub}horizon fluctuations where the matter perturbations have gone nonlinear. And while the authors of Ref. \\cite{gruzinov} construct a cosmological model with negligible effects of averaging, there do exist toy models of inhomogeneous spacetimes in which these effects are large (see, e.g., Ref. \\cite{aptp}).  Furthermore, the possibility of \\emph{nonperturbative} effects of averaging has not, to our knowledge, been satisfactorily addressed, much less ruled out. It is important to realize that current $N$-body simulations for example, which are based on a Newtonian approximation of gravity, consistently miss possible relativistic effects of averaging. This is because the modifications due to averaging in a Newtonian setting can be shown to reduce to boundary terms in general \\cite{ehlers}, and to vanish in the case of spherical symmetry \\cite{kerscher}. These terms therefore do not contribute in simulations which either set periodic boundary conditions on the simulation box, or use a spherically symmetric simulation region. As we will see, the ``backreaction'' effects in Buchert's framework (i.e., the modifications to the Cosmological equations due to averaging) are coupled to the average of the spatial curvature, this being a completely relativistic effect. Nontrivial consequences of this coupling deserve more attention than they have received in the context of $N$-body simulations (see Ref. \\cite{kerscher}).  We also note that an argument for describing the Universe using FLRW plus Newtonian perturbations (e.g., as given in Ref. \\cite{wald}), is based on the fact that it is Newtonian perturbations \\emph{and} a negative pressure component, which describe the Universe very well. In other words, one argues that the effect of averaging inhomogeneities must be small, but at the cost of adding an ill-understood Dark Energy component in the Universe. On the other hand one should allow for the following possibility (as of now neither confirmed nor ruled out): at an initial epoch, say at the end of inflation, the perturbations in the homogeneous metric would give a small ``backreaction'' as a result of averaging. As a result of its coupling with the curvature, this backreaction would grow and become significant around the time of structure formation, when matter inhomogeneities on small scales become nonlinear. The backreaction would also nontrivially affect the growth of the metric perturbations. Such a possibility could only be tested within a fully nonperturbative treatment of the evolution of inhomogeneities.  The goal of this paper is modest. We argue that the scale factor, defined in the spatially averaged cosmology via the volume of a domain, should correspond to an effective FLRW metric, provided the size of the averaging domain coincides with the scale at which cosmological homogeneity arises. This allows us, in principle, to compare predictions of a spatially averaged cosmology with observations, in the standard manner, for instance by computing the \\dlz\\ relation. (However, see Ref. \\cite{buchmorph} for an alternative approach to matching the spatially averaged dynamics with observations.) The predictions of the spatially averaged cosmology would in general differ from standard FLRW cosmology, because the scale-factor now obeys modified Einstein equations. We will demonstrate that consistent models of spatially averaged FLRW templates exist which can, at least in principle, explain observations such as those of the \\dlz\\ relation for Supernovae. \\label{intro}   In a nutshell, our new results in this paper are that we propose a template metric which takes into account the kinematical effects of inhomogeneities in the Universe. To implement this idea, we first suggest that in Buchert's spatial averaging scheme, the size of the averaging domain should coincide with the cosmological scale at which homogeneity sets in. We then argue that, as a first step towards predicting the relation between inhomogeneities and observations, the expansion factor $a_{\\cal D}(t)$ defined by Buchert should be identified with the scale-factor $a(t)$ in standard FLRW cosmology. The template metric is then the FLRW metric, with $a_{\\cal D}(t)$ playing the role of $a(t)$. We then show, with the help of a few toy examples, that the luminosity distance - redshift relation in a model Universe with an accelerating $a_{\\cal D}(t)$ can be computed in the standard manner. To the best of our knowledge, such a template metric has not been proposed in earlier work by other researchers. ",
        "conclusions": "We have argued that for the purposes of Cosmological study, an explicit coarse graining of the matter density in the Universe must be performed using the formalism of spatial averaging developed by Buchert. Due to the observed large scale homogeneity of the matter density and the isotropy of the CMB radiation, the notionally smoothed out template metric obtained after this coarse graining must be described by a FLRW line element, with the scale factor being given by Buchert's functional $\\aD(t)$. The procedure for constructing luminosity distances is therefore identical to that in the standard approach, except that the \\emph{evolution} of the effective scale factor must be given by the modified FLRW equations \\eqref{avg7} supplemented by the integrability condition \\eqref{avg9}. (The scale factor and its evolution will be further modified by including perfect fluid sources with non-zero pressure \\cite{buch4} and the curvature backreaction and volume effects \\cite{buchRG}.) Such a construction of luminosity distances must be interpreted as a convenient way of comparing observations with theoretical models, while retaining the essence of the physical inhomogeneous geometry. It will be an interesting project to compare such luminosity distances in realistic backreaction scenarios with current Supernovae observations in an attempt to constrain the magnitude and evolution of the backreaction.  We end by discussing some additional issues concerning an explicit smoothing of the metric which we avoided in the present treatment, and also some potential problems with our construction. Firstly, we return to the issue of holding the comoving range $R_{hom}$, and in particular the smoothing scale $L$, constant. This is justified only so long as the condition $L \\ll L_{\\rm Hubble}(z)$ is satisfied, where $L_{\\rm Hubble}(z)$ is the appropriate comoving scale associated with the observable Universe at redshift $z$. For example, if we choose a scale $L$ which satisfies this condition at the present epoch, the same scale will not in general be valid at the last scattering epoch of the CMB. However, it seems reasonable that for a small enough redshift range, the condition $L \\ll L_{\\rm Hubble}(z)$ can be satisfied with a single constant $L$. For epochs sufficiently far back in the past, the scale $L$ would have to be reduced, and the construction of the \\dlz\\ relation would no longer be valid.  Further, a change in the size of the averaging domain would in general affect both the scale factor \\aD\\ and the ``constant'' \\uD{k}\\ in a non-trivial way. In particular, any change in \\uD{k}\\ would render the effective FLRW metric \\eqref{smop1} meaningless. Also, a change in \\uD{k}\\ need not necessarily be caused by a change in the averaging domain size. It seems likely \\cite{buch-corresp} that an explicit smoothing of the geometry using Buchert and Carfora's technique will yield a 3-space of different spatially constant curvatures at different times, this difference being independent of any change in the averaging domain size with time and being possibly incommensurate with our assumption of $\\Cal{R}_{\\rm eff}\\propto\\aD^{-2}$. This could be interpreted in terms of a change in \\uD{k}\\ with time, and in such a case, one would be dealing with a \\emph{different} FLRW template at different times. The effects of such a construction are not clear at present. Clearly this issue needs more careful study, and we hope to address this in the future. Of course, if \\uD{k}\\ is made a function of time, care must be taken that the matter content is now so chosen that Einstein's equations continue to be satisfied, as pointed out by R\\\"{a}s\\\"{a}nen \\cite{r1} in response to \\cite{Hou}.  Also, as pointed out by R\\\"{a}s\\\"{a}nen \\cite{ras}, the use of a homogeneous and isotropic FLRW template will necessarily ignore physical effects such as a non-trivial shear in the underlying geometry. A potential problem with our template construction is that effects of the non-trivial scalar curvature $^{(3)}\\Cal{R}(\\vec{x},t)$ on light propagation are also simplified to a proportionality to $\\aD^{-2}$. Although the full scalar curvature will show up in the evolution of $\\aD(t)$ via its spatial average \\avg{\\Cal{R}}, it is possible that this construction may inaccurately model the real Universe. A study of the evolution of the effective spatial curvature as mentioned in the preceding paragraph may help resolve this issue.  There is also the related issue of degeneracy, namely that several underlying inhomogeneous models may reproduce the same backreaction effects when averaged. For example, it was shown \\cite{aptp} in the context of the LTB solutions, that models with a vanishing spatial curvature (the so called marginally bound models) in the inhomogeneous geometry also have a vanishing backreaction \\uD{\\Cal{Q}} after averaging. One is now faced with the situation wherein interpreting say luminosity distance observations within the inhomogeneous geometry may distinguish two inhomogeneous models, but interpreting the data \\emph{after} averaging would not make this distinction and would hence perhaps rule out \\emph{both} models. This issue has been recently highlighted by Enqvist and Mattsson (the last paper in Ref. \\cite{LTB}) who show that a marginally bound LTB model can fit the data from Type Ia SNe without any spatial averaging. Since the standard FLRW matter dominated model is also (a special, homogeneous case of) a marginal LTB model which \\emph{doesn't} fit the SNe data \\cite{riess}, the formalism after spatial averaging would rule out both models by demanding a non-trivial backreaction instead. However, since the LTB (and in fact any inhomogeneous) models themselves face a similar degeneracy issue (for example several different LTB models can currently claim to fit the SNe data \\cite{LTB}), this matter cannot be considered as settled, and also deserves further attention.  Lastly, we note that Buchert's averaging scheme itself has been criticized in the literature (see, e.g., Ref. \\cite{wald}) on the grounds that the effects of averaging in a noncovariant manner as done in this scheme, are likely to be gauge artifacts. It should be noted that within second-order perturbation theory, gauge invariance of the averaged equations has been demonstrated in \\cite{kolb} and \\cite{li-schwarz}. It is also possible to approach the problem of determining modifications to the standard Cosmological equations via a covariant method such as Zalaletdinov's Macroscopic Gravity \\cite{zala} for example. Such an attempt has in fact been made recently and has led to effective Cosmological equations very similar to those of Buchert (see Ref. \\cite{spatlim}). The proposal in the present paper however, is attractive to us due to its simplicity in implementation, and may be treated as a phenomenological model of the more rigorous equations derived in Ref. \\cite{spatlim}. \\label{discuss}"
    },
    "0609/astro-ph0609354_arXiv.txt": {
        "abstract": "{To constrain different cosmological models and especially the equation of state parameter $w$ of dark energy, the phases and amplitudes of baryonic acoustic oscillations in the galaxy power spectrum can be used as cosmological probes.} {It will be shown that the phases as well as the amplitudes of baryonic acoustic oscillations can be extracted out of a galaxy power spectrum, including various observational effects like growth suppression, redshift space distortions, and galaxy biasing. The phases of baryonic acoustic oscillations are used as a standard ruler for a cosmological test to constrain $w$ over a large redshift range with a minimum of assumptions.} {A non-oscillating phenomenological fitting function is used to extract the oscillatory part of the galaxy power spectrum and to disentangle phase information from amplitude information. The method (FITEX) is tested with simulated data of the Hubble Volume Simulation to include redshift space effects, non-linear structure growth and biasing.  A cosmological test is introduced, which compares the extracted oscillations against a theoretical model template to derive constraints on the $w$ parameter.} {The phenomenological fitting function is found to model the various distortions of the galaxy power spectrum to the sub-percent level. The various distortions only boost the amplitude of the oscillations. The theoretical template is accurate enough to test for small deviations in the phases of the oscillations, resulting from different $w$ values. A cosmological test, using the baryonic acoustic oscillations as a standard ruler, is able to constrain $w$ in a robust way. These constraints are based solely on the phase of these oscillations, assuming no errors of the value at the end of the Compton drag epoch, as the standard ruler. The $w$ constraints are expected to further improve when the full two dimensional redshift space effects and the amplitudes of the baryonic acoustic oscillations are taken into account.} ",
        "introduction": "Acoustic oscillations as observed in the temperature anisotropies of the cosmic microwave background radiation (CMB) are traditionally used to constrain the values of certain cosmological parameters. The intrinsic amplitudes and locations of the oscillations are determined by the densities and pressures of the various energy components in the very early and hot Universe. The oscillations are furthermore modified by their subsequent geometric projection onto the present hypersphere, where they are observed. It turns out that both intrinsic CMB oscillations and their geometric projection are mainly shaped by the physical properties at high redshifts. Excluding the very large scales where cosmic variance makes analysis quite difficult, the observed oscillations in the CMB are highly degenerate against certain changes of the energy density $\\rho_{\\rm DE}$ and the equation of state parameter, $w=p_{\\rm DE}/\\rho_{\\rm DE}c^2$, of the dark energy, dominating at low redshifts (see also Caldwell, Dave \\& Steinhardt 1998).  However, related so-called baryonic acoustic oscillations (BAOs) are expected to be observable at lower redshifts in the matter power spectrum. Therefore, the phases of the BAOs can be used as cosmic rulers in a similar manner as the CMB oscillations, but now at much smaller redshifts (Eisenstein, Hu \\& Tegmark 1998). Here, the observed phases are affected by the late-time geometry and thus by the value of $w$, the prime cosmological parameter of the present investigation.  The BAOs are classical Doppler peaks in the density distribution of matter (e.g. Hu \\& Sugiyama 1996). They are triggered in the hot Universe by oscillatory velocity patterns of the baryons. These velocity oscillations are generated, in the same environment as the CMB density oscillations, by sound waves on scales where radiation pressure could stabilize the fluid against gravitational collapse. At later times, this oscillatory velocity pattern of the baryon field produces kinematically a new field of matter density fluctuations in form of quasi-regular matter oscillations. The amplitudes of these BAOs are determined by the ratio of baryonic matter to the overall matter density.  The BAOs thus constitute a quasi-regular pattern of oscillating substructures superposed with small amplitudes onto a general density field which fluctuates irregularly with several 100 times higher amplitudes.  The imprints of the BAOs on galaxy and clusters distributions, gravitational shear maps etc. can provide at least in principle a clean cosmic ruler for precise tests of the $w$ parameter. The characteristic scale, $s$, of these imprints is the comoving distance, sound waves can travel during the epoch where baryons and photons are strongly coupled through Compton and Thomson scattering (Compton drag). Cosmological tests based on the resulting sound horizon, $s$, as a metric ruler are expected to have only very small systematic errors because the phases of the BAOs located at scales small compared to $s$ are solely determined by well-understood physical processes. A practical problem is, however, to separate the oscillatory BAO modes from the irregular fluctuation field with sufficient accuracy.  One might object that the small amplitudes of the BAOs can easily be washed out especially by structure growth which can mix perturbation modes with adjacent wavenumbers $k$. This is certainly true for oscillations in the non-linear regime (Meiksin, White \\& Peacock 1999). On larger scales, however, recent observations of 2dF galaxies at redshifts $z<0.3$ clearly show several BAOs in the galaxy power spectrum (Cole et al. 2005). In addition, the sequence of BAOs in $k$-space is projected into a single wiggle in the space domain. The corresponding excess galaxy correlation was in fact observed in the two-point spatial correlation function of the luminous red SDSS galaxies at $z<0.5$ (Eisenstein et al. 2005). In addition to these observational indications, results from recent numerical simulations suggest that within certain redshift and scale ranges, stable BAOs could exist as useful probes for cosmological investigations (e.g. Seo \\& Eisenstein 2005, Springel et al. 2005, Jeong \\& Komatsu 2006). However, several critical issues like the exact behavior of baryons during structure growth on BAO-scales can only be discussed with much larger simulations and a better understanding of galaxy formation.  Applications of BAOs for cosmological tests of $w$ are confronted with the following situation. Present observational constraints of $w$ using different combinations of CMB, galaxy clusters, galaxies, and gravitational lensing data are all found to be consistent with $w=-1$, i.e., the value for the cosmological constant. The $1\\sigma$ error of the $w$ values, derived from partially dependent observations, is around 10-20\\% (for a recent review see, e.g., Schuecker 2005). To improve current estimates, we thus have to extract the BAOs and to measure their phases quite accurately in the presence of non-linear and scale-dependent effects. Therefore, even tiny systematic errors in the analysis on percent levels can have severe consequences for the observational accuracy of $w$.  The basic aim of the present paper is to describe a new, simple, and robust method to extract BAOs from (non)linear, scale and redshift-dependent biased, redshift-space galaxy power spectra under realistic lightcone survey conditions. For brevity, we call the method `fit and extract' (FITEX). The method has the potential to reduce many of the abovementioned sources of systematic errors. Our starting point is to use as few assumptions as possible for the extraction of the BAOs from a complex power spectrum. For reasons which will become clear later, we do the separation by fitting a flexible non-oscillating function to an observable which is directly related to the transfer function and not to the power spectrum itself. The crucial point is to show that BAOs extracted in this simple manner are still governed by simple physical processes, and that FITEX stays robust even under complex survey conditions. We believe that such model-independent approaches are quite important, especially in light of the fact that at least in the near future we cannot expect to model the abovementioned non-linear and scale-dependent effects with accuracies on the sub-percent level.  Our computations are in several cases optimized to galaxy surveys covering typically several 100 square degrees on the sky in the redshift range between $z=2-4$. This choice is motivated by the Hobby Eberly Dark Energy Experiment (HETDEX), which is planned to measure the $w$ parameter with several million Ly-$\\alpha$ emitting galaxies at these redshifts (Hill \\& McQueen. 2004, Gebhardt et al. in prep.). We see the results of the present paper as a useful contribution to a realistic error forecast for this important project.  We organized the paper as follows. In Sect.\\,\\ref{BAOS} we motivate the basic idea of the FITEX algorithm. In the following sections we test its performance to separate BAOs from a complex power spectrum under linear conditions (Sect.\\,\\ref{EXTRACT}), under quasi-non-linear conditions (Sect.\\,\\ref{NONLINEAR}), in redshift space (Sect.\\,\\ref{REDSHIFT}), and from biased samples (Sect.\\,\\ref{BIASING}). Finally, we illustrate the performance of FITEX including all effects and use a cosmological test of $w$ as a benchmark to investigate the quality of the BAO extraction together with the theoretical template (Sect.\\,\\ref{CTEST}). Our results are mainly based on a deep data wedge extracted by the Virgo Consortium from the Hubble Volume Simulation. The paper thus presents for the first time constraints on the application of BAOs for cosmological tests under realistic lightcone survey conditions.  If not mentioned explicitly, we assume a spatially flat Friedmann-Lema\\^{i}tre Robertson-Walker world model with the Hubble constant in units of $h=H_0/(100\\,{\\rm km}\\,{\\rm s}^{-1}\\,{\\rm Mpc^{-1}})$, the present values of the total matter density $\\Omega_{\\rm m}h^2=0.147$ and baryon density $\\Omega_{\\rm b}h^2=0.0196$, the density of relativistic matter (e.g. neutrinos) $\\Omega_\\nu =0$, and the mean CMB temperature $T_{\\rm CMB}=2.728\\,{\\rm K}$ (concordance cosmology). ",
        "conclusions": "\\label{DISCUSS}  We have seen that it is critical for the extraction of BAOs to either fit a phenomenological function to the continuum of the boosted transfer function (FITEX), or to calculate the shape of the boosted transfer function, including all distortions down to the sub-percent level. The latter approach has been carried out recently. Glazebrook \\& Blake (2005), H\\\"utsi (2005), and Springel et al. (2005) used a linear theory power spectrum and divided the `observed' data by this reference. Almost the same approach, but with the addition of linearizing the observational data before the division, to correct for non-linear structure growth was used in Angulo et al. (2005). There are two drawbacks of this method:  First, the ratio of a complex power spectrum and a reference power spectrum does not resemble a wiggle function as defined in the text, but something more complicated (see Eq.\\,\\ref{eq_division}). The simple form of the theoretical wiggle function (Eq.\\,\\ref{eq wiggles}) is strongly deformed by both the chosen transfer function of the reference spectrum and the distortions introduced by redshift space and galaxy biasing. These distortions make the comparison with theoretically expected model spectra unnecessarily complicated. FITEX transforms the observed power spectrum into a boosted transfer function where the subtraction of a non-oscillating phenomenological continuum function is directly related to $T_w(k)$, multiplied by an amplitude factor.  The second drawback is the linear theory power spectrum itself. By computing a correct shape of the power spectrum, one has to rely strongly on assumptions about the nature of the dark matter and various other physical processes. The method is no more the `assumption free' approach that was intended. Furthermore, the calculation of the amplitude factor that includes all the distortions ranging from linear structure growth, non-linear structure growth to redshift space distortions and biasing is a problem. This calculation has to be accurate to the sub-percent level, as the oscillations themselves make up only $\\sim\\,2\\%$ of the transfer function.  A more phenomenological approach was conducted by White (2005). He assumed that redshift space distortions as well as linear and non-linear structure growth could be calculated with reasonable accuracy. Leaving only the ratio between dark matter power spectra and `galaxy' power spectra, he found that one can fit this ratio with about 2\\% accuracy using a cubic polynomial. This adds further evidence that a phenomenological approach can model power spectrum distortions without assuming much about the physics of the distortions themselves.  With the FITEX method, we have shown that one is able to extract the BAOs without any strong priors on the nature of dark matter or the shape of the power spectrum from the complex multi-component transfer function. All the aforementioned effects are modeled by the continuum function and need not be known very accurately, as long as the phases of the oscillations are not distorted or washed out. This extraction is accurate on the sub-percent level within $k$-ranges relevant for the cosmological test. Furthermore, FITDEX is able to disentangle phase information from amplitude information which could in the future be used to test the geometry of the Universe on the one hand and the growth history on the other hand. This type of cosmological test could be used to discriminate Dark Energy from Modified Gravity theories."
    },
    "0609/astro-ph0609024_arXiv.txt": {
        "abstract": "We present $5 - 38\\mu$m mid-infrared spectra at a spectral resolution of $R\\approx 65 - 130$ of a large sample of 22 starburst nuclei taken with the Infrared Spectrograph {\\sl IRS} on board the Spitzer Space Telescope. The spectra show a vast range in starburst SEDs.  The silicate absorption ranges from essentially no absorption to heavily obscured systems with an optical depth of $\\tau_{9.8\\mu m}\\sim 5$.  The spectral slopes can be used to discriminate between starburst and AGN powered sources.  The monochromatic continuum fluxes at $15\\mu$m and $30\\mu$m enable a remarkably accurate estimate of the total infrared luminosity of the starburst. We find that the PAH equivalent width is independent of the total starburst luminosity $L_{IR}$ as both continuum and PAH feature scale proportionally.  However, the luminosity of the $6.2\\mu$m feature scales with $L_{IR}$ and can be used to approximate the total infrared luminosity of the starburst. Although our starburst sample covers about a factor of ten difference in the [Ne\\,III]\\,/\\,[Ne\\,II] ratio, we found no systematic correlation between the radiation field hardness and the PAH equivalent width or the $7.7\\mu$m\\,/\\,$11.3\\mu$m PAH ratio. These results are based on spatially integrated diagnostics over an entire starburst region, and local variations may be ``averaged out''. It is presumably due to this effect that unresolved starburst nuclei with significantly different global properties appear spectrally as rather similar members of one class of objects. ",
        "introduction": "\\label{secintro}  Many nearby galaxies show dramatically increased rates of star formation compared with the Milky Way. In such \"Starburst\" galaxies (e.g., \\citet{wee81}), the primary energy source is driven by high nuclear star formation rates, rapid nuclear gas depletion time-scales and high supernovae rates. Such starburst systems often, though not exclusively, occur in interacting and colliding systems. Since collisions and interactions are believed to be a fundamental part of the evolution of galaxies by hierarchical growth, the full characterization of starburst galaxies is of the great importance in measuring and quantifying the global history of star formation over cosmic time.  Determining the average mid-infrared (mid-IR) spectral properties, and the range of observed behavior within the starburst class at low redshift is vital for interpreting spectra of higher redshift IR sources, providing complementary spectral templates to those parallel Spitzer studies of the more extreme ULIRG systems (e.g., \\citet{arm06b}) and AGN (e.g., \\citet{wee05}).  The term ``starburst galaxy'' is commonly used to describe an apparently well-defined class of objects, although starbursts can be found in the most diverse conditions, ranging from low-pressure dwarf galaxies to high-pressure nuclear starbursts.  Their observed properties are expected to depend on numerous parameters such as the initial stellar mass function (IMF), the duration and epoch of the individual starburst(s), the metallicity of the ISM, the size and distribution of the dust grains, the strength of the magnetic fields, gas pressure and temperature of the ISM, galactic shear, total luminosity, and total mass.  Furthermore, nearby starbursts, for which high resolution imaging is possible, have revealed complex substructures -- in both stellar distributions and ISM -- ranging from ultra-compact H\\,{\\sc ii} regions (UCHIIR) to large complexes of super star clusters (SSC), suggesting small-scale variations of the observables across a starburst region.  We use the low resolution mode of the Infrared Spectrograph\\footnote{The {\\sl 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.} ({\\sl IRS}) \\citep{hou04} on board the {\\sl Spitzer Space Telescope} \\citep{wer04} to observe the central regions of 22 starburst galaxies.  Our objects represent a sample of ``classical'' starbursts for which a wealth of literature exists.  The sample includes both purely starburst and starbursts with weak AGN activity (as determined from X-ray, optical, or radio observations).  The summary in Table~\\ref{tabgenprop} lists the observed targets, their general properties, the classifications we adopt, and the references from which they are derived.  The continuous $5-38\\mu$m {\\sl IRS} spectra include the silicate bands around $10\\mu$m and $18\\mu$m, a large number of PAH emission features, and information on the slope of the spectral continuum.  \\begin{deluxetable*}{l r r l r r c r r r r} \\centering \\tabletypesize{\\footnotesize} \\tablecaption{General Properties} \\tablewidth{0pt} \\tablehead{ \\colhead{Name} & \\colhead{$\\alpha_{\\mbox{J2000}}$\\tablenotemark{a}} & \\colhead{$\\delta_{\\mbox{J2000}}$\\tablenotemark{a}} & \\colhead{Type} & \\colhead{Refs.\\tablenotemark{b}} & \\colhead{D\\tablenotemark{c}} & \\colhead{$\\log(L_{IR})$\\tablenotemark{d}} & \\colhead{$S_{12}$\\tablenotemark{e}} & \\colhead{$S_{25}$\\tablenotemark{e}} & \\colhead{$S_{60}$\\tablenotemark{e}} & \\colhead{$S_{100}$\\tablenotemark{e}}\\\\ & & & & & \\colhead{[Mpc]} & \\colhead{[$L_\\odot$]} & \\colhead{[Jy]} & \\colhead{[Jy]} & \\colhead{[Jy]} & \\colhead{[Jy]} } \\startdata IC   342 &  3 46 48.51  &   +68 05 46.0 & SB & 13,24,25   &  4.6 & 10.17 & 14.92 & 34.48 & 180.80 & 391.66 \\\\ Mrk   52 & 12 25 42.67  &   +00 34 20.4 & SB & 2,9,17     & 30.1 & 10.14 & 0.28 &  1.05 &  4.73 &   5.68 \\\\ Mrk  266 & 13 38 17.69  &   +48 16 33.9 & SB+Sy2 & 14,20  & 115.8 & 11.49 & 0.32 &  1.07 &  7.25 &  10.11 \\\\ NGC  520 &  1 24 35.07  &   +03 47 32.7 & SB & 4,10,12,13 & 30.2 & 10.91 & 0.90 &  3.22 & 31.52 &  47.37 \\\\ NGC  660 &  1 43 02.35  &   +13 38 44.4 & SB+LINER & 10     & 12.3 & 10.49 & 3.05 &  7.30 & 65.52 & 114.74 \\\\ NGC 1097 &  2 46 19.08  & $-$30 16 28.0 & SB+Sy1 & 19,23  & 16.8 & 10.71 & 2.96 &  7.30 & 53.35 & 104.79 \\\\ NGC 1222 &  3 08 56.74  & $-$02 57 18.5 & SB & 1,2        & 32.3 & 10.60 & 0.50 &  2.28 & 13.06 &  15.41 \\\\ NGC 1365 &  3 33 36.37  & $-$36 08 25.5 & SB+Sy2 & 19,26  & 17.9 & 11.00 & 5.12 & 14.28 & 94.31 & 165.67 \\\\ NGC 1614 &  4 33 59.85  & $-$08 34 44.0 & SB & 29,30      & 62.6 & 11.60 & 1.38 & 7.50  & 32.12 & 34.32  \\\\ NGC 2146 &  6 18 37.71  &   +78 21 25.3 & SB & 7,10       & 16.5 & 11.07 & 6.83 & 18.81 & 146.69 & 194.05 \\\\ NGC 2623 &  8 38 24.08  &   +25 45 16.9 & SB & 13,22      & 77.4 & 11.54 & 0.21 &  1.81 & 23.74 &  25.88 \\\\ NGC 3256 & 10 27 51.27  & $-$43 54 13.8 & SB & 9,24,25    & 35.4 & 11.56 & 3.57 & 15.69 & 102.63 & 114.31 \\\\ NGC 3310 & 10 38 45.96  &   +53 30 05.3 & SB & 9,10       & 19.8 & 10.61 & 1.54 &  5.32 & 34.56 &  44.19 \\\\ NGC 3556 & 11 11 30.97  &   +55 40 26.8 & SB & 7,10       & 13.9 & 10.37 & 2.29 &  4.19 & 32.55 &  76.90 \\\\ NGC 3628 & 11 20 17.02  &   +13 35 22.2 & SB+LINER & 10,21  & 10.0 & 10.25 & 3.13 &  4.85 & 54.80 & 105.76 \\\\ NGC 4088 & 12 05 34.19  &   +50 32 20.5 & SB & 3,5,10     & 13.4 & 10.25 & 2.06 &  3.45 & 26.77 &  61.68 \\\\ NGC 4194 & 12 14 09.64  &   +54 31 34.6 & SB & 2,9,17     & 40.3 & 11.06 & 0.99 &  4.51 & 23.20 &  25.16 \\\\ NGC 4676 & 12 46 10.10  &   +30 43 55.0 & SB & 15,16      & 94.0 & 10.88 & 0.11 &  0.33 &  2.67 &   5.18 \\\\ NGC 4818 & 12 56 48.90  & $-$08 31 31.1 & SB & 5,17       &  9.4 & 09.75 & 0.96 &  4.40 & 20.12 &  26.60 \\\\ NGC 4945 & 13 05 27.48  & $-$49 28 05.6 & SB+Sy2 & 7,11,31  & 3.9 & 10.48 & 27.74 & 42.34 & 625.46 & 1329.70 \\\\ NGC 7252 & 22 20 44.77  & $-$24 40 41.8 & SB & 6,15,18    & 66.4 & 10.75 & 0.24 &  0.43 &  3.98 &   7.02 \\\\ NGC 7714 & 23 36 14.10  &   +02 09 18.6 & SB & 8,27,28    & 38.2 & 10.72 & 0.47 &  2.88 & 11.16 &  12.26\\\\ \\enddata \\label{tabgenprop} \\tablecomments{Mrk 52 = NGC 4385, Mrk 266 = NGC 5256} \\tablenotetext{a}{Commanded coordinates of the slit center. $\\alpha$ is given in (h~m~s), $\\delta$ is given in (d~m~s).} \\tablenotetext{b}{References: 1--\\citet{ash95}, 2--\\citet{bal83}, 3--\\citet{ben04}, 4--\\citet{bes03}, 5--\\citet{dev89}, 6--\\citet{fri94}, 7--\\citet{gao04}, 8--\\citet{gon99}, 9--\\citet{hec98}, 10--\\citet{ho97}, 11--\\citet{iwa93}, 12--\\citet{jos85}, 13--\\citet{kee84}, 14--\\citet{lev01}, 15--\\citet{liu95}, 16--\\citet{lon84}, 17--\\citet{may04}, 18--\\citet{mil97}, 19--\\citet{osm74}, 20--\\citet{ost83}, 21--\\citet{rob04}, 22--\\citet{smi98}, 23--\\citet{sto03}, 24--\\citet{tho00}, 25--\\citet{ver03}, 26--\\citet{ver80}, 27--\\citet{wee81}, 28--\\citet{bra04}, 29--\\citet{alo01}, 30--\\citet{ket92}, 31--\\citet{spo00}} \\tablenotetext{c}{Distances adopted from \\citet{san03}, except for Mrk~52, NGC~4676 and NGC~7252, which were derived from measured redshifts via $D = \\frac{cz}{H_{\\mbox{0}}}(1+\\frac{z}{2})$, assuming $H_{\\mbox{0}} = 71$\\,km\\,s$^{-1}$ Mpc$^{-1}$.} \\tablenotetext{d}{Total $8 - 1000\\mu$m infrared luminosity of the entire galaxy, adopted from \\citet{san03}, except for Mrk~52, NGC~4676 and NGC~7252, which were derived from measured IRAS fluxes via $L_{IR} = 312700\\cdot D^2\\cdot 1.8\\cdot (13.48S_{12\\mu m} + 5.16S_{25\\mu m} + 2.58S_{60\\mu m} + S_{100\\mu m})$ where $S_{\\lambda}$ is in Jy.} \\tablenotetext{e}{{\\sl IRAS} flux densities at 12, 25, 60 and $100\\mu$m of the entire galaxy, adopted from \\citet{san03}.} \\end{deluxetable*}   Numerous mid-IR studies of starbursts have been conducted with {\\sl ISO-SWS}, {\\sl ISOCAM}, or {\\sl ISOPHOT}, see for instance, \\citet{rig96}, \\citet{lut98}, \\citet{rig99}, \\citet{dal00}, \\citet{hel00}, \\citet{lau00}, \\citet{stu00}, \\citet{tho00}, \\citet{cha01}, \\citet{for03}, \\citet{lu03}, \\citet{ver03}, \\citet{tac05}, and \\citet{mad06}. An overview of many {\\sl ISO} results is given in \\citet{gen00}.  While the {\\sl ISO} observations provided great new insights in the spectral characteristics of individual starbursts, the sample of continuum spectra remained rather small or was limited to shorter wavelengths or narrow bandwidth scans of the strongest emission lines.  In this paper we will address the question of the mid-IR homogenity of the classical starburst class\\footnote{In this paper the terms `starburst', `starburst galaxy', and `starburst nuclei' all refer to the central, sub-kiloparsec regions of galaxies with significantly enhanced starburst activity}, and attempt to investigate how specific spectral features (especially the mid-IR PAH bands) vary with the total UV continuum flux, UV hardness ratio and dust extinction within the starburst nucleus.  We will also investigate how the shape of the continuum depends on the luminosity source, and if the total luminosity affects the observed spectral shapes, i.e., to what extent starbursts can be scaled up.  We investigate the role of dust, and how well PAH emission correlates with the rate of star formation.  The outline of the paper is as follows: First we give a detailed description of the observations and the data reduction and calibration.  In Section~\\ref{secanaly} we discuss how the relevant spectral features (SED, PAHs, silicate features) have been measured. The main focus will be on the discussion of the numerous results in section~\\ref{secdiscus}, followed by a summary. We note that our sample has also been observed with the {\\sl IRS} high resolution modules, revealing the large, comprehensive zoo of strong and faint fine-structure lines in the $10-38\\mu$m wavelength range. This analysis is complementary to our above science goals and will be presented in a subsequent paper by \\citet{dev06}. ",
        "conclusions": "\\label{secdiscus}  It is important to keep in mind that the starburst spectra presented in this paper represent an entire starburst region, including numerous (super-)star clusters at various ages and evolutionary states, the surrounding PDRs which are internally and externally excited, and the warm and cold dust spread across the entire region as well as localized dust condensations.  While many of the properties of local substructures are expected to vary significantly, the overall significance of these variations may be averaged out in the spatially integrated spectra. Our aim here is to search for global trends between the spectral properties derived in the previous section (silicate absorption features, PAH features, spectral continuum) and the global starburst properties ($L_{IR}$, radiation field).   \\subsection{The Continuum Slope as a SB/AGN Diagnostic} \\label{seccontflux}  The slope of the mid-IR spectral continuum depends on the optical thickness, composition and temperature of graphite dust grains, which are related to the amount of silicate grains \\citep{mat77}.  The dominating species in the {\\sl IRS} spectral range are hot ($\\ge 100$\\,K), large grains heated by ionizing, non-ionizing, and Ly$\\alpha$ photons inside \\2 regions, and small ($\\le 100$\\,\\AA) grains heated by non-ionizing photons outside \\2 region \\citep{mou97}. In Fig.~\\ref{figtauslope} we plot the continuum slope, parametrized by the ratio of the $15\\mu$m and $30\\mu$m flux densities versus the optical depth at $9.8\\mu$m (section~\\ref{secsili}).  The filled symbols in Fig.~\\ref{figtauslope} (and all other figures thereafter) correspond to starbursts with a weak AGN component as identified in Table~\\ref{tabgenprop}.  Although not a tight correlation one can see the general trend that starbursts with stronger silicate absorption tend to have a steeper continuum at longer wavelengths.   \\begin{figure}[ht] \\epsscale{1.2} \\plotone{f7.ps} \\caption{Optical depth at $9.8\\mu$m versus the slope of the continuum as measured by the flux ratio of $15\\mu$m to $30\\mu$m.  The filled diamonds correspond to starbursts with a weak AGN component.  The dotted line is a linear fit to the data points. The highly obscured source NGC~4945 is not included in the fit.. \\label{figtauslope}} \\end{figure}   The slope of the dust continuum depends on the energy distribution and spatial concentration of the heating source(s).  \\citet{dal00} found from {\\sl ISOCAM} data of 61 galaxies a $6.75\\mu$m\\,/\\,$15\\mu$m continuum slope near unity for more quiescent galaxies, whereas that ratio drops (i.e., the slope steepens) for galaxies with increased starburst activity.  Furthermore, the continuum slope can serve as a discriminator between massive stars or an AGN as the underlying power source.  This has already been known since {\\sl IRAS} (e.g., \\citet{wan92}), and was further refined in numerous papers based on {\\sl ISOCAM} observations. (See \\citet{gen00} for a more comprehensive overview). For instance, \\citet{lau00} studied the {\\sl ISOCAM} colors of a large variety of extragalactic objects revealing clear general trends between AGN, photo-dissociation region (PDR), and \\2 region dominated spectra.  However, the characterization of an individual object often turns out to be difficult due to the intrinsically large scatter: the {\\sl IRAS} $12\\mu$m filter includes the silicate absorption band as well as several PAH emission features and strong emission lines.  The wavelength range covered by {\\sl ISOCAM} is limited to shorter wavelengths which are dominated by hot dust and a large variety of emission and absorption features (compare to Fig.~\\ref{figiceabsorb}).  Our narrowband continuum fluxes largely avoid these problems. Figure~\\ref{figcontfluxes}, analogous to a color-color diagram in stellar astronomy, shows the continuum flux ratios at $6\\mu$m/$15\\mu$m versus $15\\mu$m/$30\\mu$m.  The total infrared luminosity $L_{IR}$ is represented by the size of the symbols.  For comparison the figure also contains the classical AGN Cen\\,A (Sy\\,2), I\\,Zw\\,1 (Sy\\,1), Mrk~3 (Sy\\,2), NGC~1275 (Sy\\,2), NGC~4151 (Sy\\,1.5), and NGC~7469 (Sy\\,1.2) from \\citet{wee05}.  We also show the starburst-dominated ULIRGs IRAS~12112+0305, IRAS~22491-1808, and the AGN-dominated Mrk~231, Mrk~463, and Mrk~1014 from \\citet{arm06b}.  NGC~6240 is a peculiar case with an intrinsic fractional AGN contribution to the bolometric luminosity of $~20-24\\%$ \\citep{arm06a}.   \\begin{figure}[ht] \\epsscale{1.2} \\plotone{f8.ps} \\caption{``Color-color plot'' of the $6\\mu$m/$15\\mu$m versus $15\\mu$m/$30\\mu$m flux densities.  The size of the diamonds, representing our starburst sample, increases linearly with $\\log L_{IR}$.  The filled diamonds correspond to starbursts with a weak AGN component and the square refers to the average starburst template (Table~\\ref{tabsbaverage}).  For comparison the figure also contains AGN (pluses) from \\citet{wee05}, as well as ULIRGs (filled circles) from \\citet{arm06a} and \\citet{arm06b}. \\label{figcontfluxes}} \\end{figure}   Several conclusions can be drawn from Fig.~\\ref{figcontfluxes}. First, there is a large scatter along the $y$-axis with no obvious correlation with total starburst luminosity or galaxy type.  Hence, the $6\\mu$m/$15\\mu$m continuum flux ratio does not appear to be a good diagnostic.  Second, using the $15\\mu$m/$30\\mu$m continuum flux ratio, classical (strong) AGN can be clearly separated from the starburst galaxies (including the ones with weak nuclear activity), with the AGN having a significantly shallower mid-IR spectrum.  The grey-shaded bar in Fig.~\\ref{figcontfluxes} indicates the transition region between AGN and starburst-dominated systems, and lies approximately at \\begin{displaymath} 0.17\\le F_{15\\mu m}/F_{30\\mu m} \\le 0.21. \\end{displaymath} Third, the difference in spectral slope also seems to apply to ULIRGs depending on their dominant power source.  Hence, this technique may have a much broader application and should be verified with a much larger sample of different classes of objects.     \\subsection{The Continuum Fluxes as Measures of $L_{IR}$} \\label{secbrandllumi}  The total infrared luminosity is an important parameter to estimate the energetics of a starburst and to characterize the underlying stellar population and rate of star formation. $L_{IR}$ is often derived from the four {\\sl IRAS} filter bands \\citep{san96}. In this subsection we will check how accurately $L_{IR}$ can be derived from the two {\\sl IRS} continuum fluxes at $F_{15\\mu m}$ and $F_{30\\mu m}$ alone.  Numerous attempts to extrapolate $L_{IR}$ from one or two, mainly broad band fluxes can be found in the literature. \\citet{tak05} discuss various estimators of infrared luminosities, and find -- for a very large sample of 1420 galaxies of different type -- correlations of the form $\\log L_{IR} = 1.02 + 0.972 \\log L_{12\\mu m}$, and $\\log L_{IR} = 2.01 + 0.878 \\log L_{25\\mu m}$.  Both $L_{12\\mu m}$ and $L_{25\\mu m}$ allow to predict $L_{IR}$ to an accuracy within a factor of $4-5$ at the 95\\% confidence level over a wide range in luminosities.  These uncertainties are similar to the luminosities derived from the {\\sl MIPS} $24\\mu$m flux alone for a large sample of {\\sl SINGS} galaxies \\citep{dal05}.  \\citet{elb01} found a similar relation fitting $15\\mu$m {\\sl ISO} fluxes: $\\log L_{IR} = (1.05\\pm 0.174) + 0.998\\log L_{15\\mu m}$.  \\citet{for04} found that the monochromatic $15\\mu$m continuum emission is directly proportional to the ionizing photon luminosity, and hence the total infrared luminosity.  As discussed in section~\\ref{seccontflux} the {\\sl IRS} fluxes $F_{15\\mu m}$ and $F_{30\\mu m}$ provide a rather accurate estimate of the ``true'' spectral continuum.  In Figure~\\ref{figbrandlluminosity} we plot a combination of $F_{15\\mu m}$ and $F_{30\\mu m}$ times $D^2$ versus the total infrared luminosity derived from the {\\sl IRAS} bands.  Since the {\\sl IRAS} beam usually covers the entire starburst while the narrower {\\sl IRS} slits can only collect a fraction of the total luminosity for local starbursts (Fig.~\\ref{figslitoverlays}) we have corrected the observed $F_{15\\mu m}$ and $F_{30\\mu m}$ (indicated by the pluses in Fig.~\\ref{figbrandlluminosity}) for the slit losses (see section~\\ref{secstitching} for more details).  A least-square fit to the corrected fluxes yields: \\begin{displaymath} L_{IR}^{est} = D^2\\cdot \\left(4.27\\cdot F_{15\\mu m} + 11\\cdot F_{30\\mu m} \\right) \\end{displaymath} where $D$ is the distance in kiloparsecs, $F_{15\\mu m}$ and $F_{30\\mu m}$ the {\\sl IRS} flux densities in Janskies.   \\begin{figure}[ht] \\epsscale{1.2} \\plotone{f9.ps} \\caption{The estimated infrared luminosity $L_{IR}^{est}$ based on the aperture loss corrected {\\sl IRS} continuum fluxes $F_{15\\mu m}$ and $F_{30\\mu m}$ for a source at distance $D$ [kpc] versus $L_{IR}$ from {\\sl IRAS}.  The dashed line represents a linear fit to the data (in log-log space), and the fitted function is given in the upper left, together with the standard error $\\sigma$.  The uncorrected fluxes are indicated by gray plus signs.  The filled symbols correspond to starbursts with a weak AGN component. \\label{figbrandlluminosity}} \\end{figure}   The correlation in Fig.~\\ref{figbrandlluminosity} is extremely tight, including the weak AGN, with a standard error (mean scatter) of only 0.09 in log-space, i.e., the {\\sl IRS} estimated infrared luminosities agree within 23\\% with $L_{IR}$.  This is much more accurate than the estimates by \\citet{elb01} and \\citet{tak05}.  The excellent correlation suggests that, at least for a homogeneous sample of starburst galaxies, $F_{15\\mu m}$ and $F_{30\\mu m}$ can be used to accurately derive $L_{IR}$.     \\subsection{A Large Variety in Silicate Absorption} \\label{variety}  Figure~\\ref{figoneontop} shows, from top to bottom, a series of starburst spectra, sorted by increasing absorption of the $9.8\\mu$m silicate resonance.  Since the peak of the resonance coincides with a minimum between the $7-9\\mu$m and $11-13\\mu$m PAH emission complexes, the effect of silicate absorption only becomes apparent towards the lower half of the plot.  Nevertheless, the figure strikingly illustrates the strong effect of amorphous silicates on the overall $5-38\\mu$m spectral shape and the large variations even within one class of objects.  Fig.~\\ref{figoneontop} shows that the $10\\mu$m trough can in fact become a dominating feature of the spectral shape of starbursts.   \\begin{figure}[htb] \\epsscale{1.2} \\plotone{f10.ps} \\caption{The {\\sl IRS} spectra of Mkn~52, NGC~7714, NGC~1222, NGC~7252, NGC~2146, NGC~520, NGC~3628, and NGC~4945 arranged to illustrate the gradual effect of increasing silicate absorption from top to bottom.  The flux densities $F_{\\lambda}$ have been arbitrarily scaled for better comparision. \\label{figoneontop}} \\end{figure}   The top two spectra also show a broad emission structure beginning at about $16\\mu$m and extending to $21-25\\mu$m. The shape of this feature is consistent with that of an $18\\mu$m silicate emission feature, which would make these the first detections of this feature in starburst galaxies.  However, the emission structure could also be due to the C--C--C in-plane and out-of-plane bending modes of PAHs \\citep{ker00}.  The identification of this feature will be discussed in a future paper.  The remaining starburst galaxies in Fig.~\\ref{figoneontop} do not show evidence for a similar emission structure. The $13-35\\mu$m spectra, from NGC~1222 (top) to NGC~4945 at the bottom, show an increasingly pronounced depression, peaking at $18.5\\mu$m, signalling increasingly strong silicate absorption. The latter result is in full agreement with the trend found for the $9.8\\mu$m silicate feature.  Among all the galaxies classified as starbursts, NGC~4945 (bottom spectrum in Fig.~\\ref{figoneontop}) is a ``special case'' exhibiting by far the strongest dust obscuration to its nuclear region \\citep{spo00}.  Based on the IRS spectrum, the apparent optical depth in the $9.8\\mu$m silicate feature is at least four and may be higher, depending on the choice of the local continuum.  Apart from strong amorphous silicate absorption, the line of sight also reveals the presence of a $23\\mu$m absorption feature.  Following the analysis of deeply obscured lines of sight towards ULIRG nuclei \\citep{spo06}, we attribute the $23\\mu$m feature to crystalline silicates (forsterite). Their detection in NGC~4945 suggests that crystalline silicates are perhaps a more common component of the ISM and not just limited to ULIRG nuclei.  Figure~\\ref{figsilicabsorb} compares the average starburst template from Fig.~\\ref{figaverage} to NGC~4945, the most extincted source within our sample.  For comparison we also show the heavily embedded ULIRG IRAS~08572+3915 from \\citet{spo06}.  The usually rather shallow $18\\mu$m silicate band reduces the continuum by a large factor compared to the average starburst spectrum.  While NGC~4945 has a similarly strong silicate absorption feature as the ULIRG IRAS~08572+3915, the starburst also shows strong PAH complexes near $16.4\\mu$m and $17.1\\mu$m, which are absent in the ULIRG spectrum. Crystalline silicates, causing the features at 16 and $19\\mu$m in IRAS~08572+3915, appear to be absent in the starburst template, although they are more difficult to discern at our spectral resolution given the presence of PAH features and the [S\\,III] line in that range.   \\begin{figure}[ht] \\epsscale{1.2} \\plotone{f11.ps} \\caption{Comparison of the average starburst template (dotted line) from Fig.~\\ref{figaverage} to the most extincted source within our sample, NGC~4945 (dashed line), and to the ULIRG IRAS~08572+3915 (solid line), chosen as an extreme case of extinction from the sample of \\citet{spo06}.  The spectra have been normalized at $14.5\\mu$m. \\label{figsilicabsorb}} \\end{figure}   While NGC~4945 shows by far the most extreme silicate absorption within our sample it is not an exotic object but defines the endpoint of a sequence of increasing optical depths.  Recently, \\citet{dal06} reported on the {\\sl Spitzer-SINGS} survey of 75 nearby galaxies.  The {\\sl typical} target in their sample shows only modest dust obscuration, consistent with an $A_V\\sim 1$ foreground screen and a lack of dense clumps of highly obscured gas and dust.  Unfortunately, we cannot distinguish between the local and global dust distributions. The differences may be due to evolutionary states, local geometries or other reasons.  However, it is evident that starburst galaxies can have very little or large amounts of extinction, and the presence/absence of a strong dust feature is not a characteristic item.     \\subsection{The Mixture of PAHs and Dust} \\label{secpahmixture}  In this section we will investigate how the derived PAH strengths may depend on extinction within the starburst region.  Table~\\ref{tabpahs} already indicates that the relative fluxes of individual PAH features are not constant for different starbursts.  \\citet{lu03} have found a 25\\% spread in the ratio of the $11.3\\mu$m\\,/\\,$7.7\\mu$m PAH fluxes which they attribute to intrinsic galaxy-to-galaxy variations. However, extinction can affect the relative strength of features at different wavelengths by different amounts (Table~\\ref{tabextlaws}). \\citet{rig99} found that the dominant influence on the (ULIRG) PAH ratio is extinction, and that the $6.2\\mu$m PAH gets suppressed relative to the $7.7\\mu$m PAH for heavily dust extincted systems.  In particular, the $8.6\\mu$m and $11.3\\mu$m PAH features lie at wavelengths that are heavily affected by silicate absorption.  Fig.~\\ref{fig11pahextinction} shows the $11.3\\mu$m\\,/\\,$7.7\\mu$m and $11.3\\mu$m\\,/\\,$6.2\\mu$m PAH flux ratios versus the apparent silicate optical depth $\\tau_{9.8}$.  A trend that starbursts with stronger dust absorption show relatively weaker $11.3\\mu$m PAHs is evident. Fig.~\\ref{fig11pahextinction} suggests that extinction can affect the relative PAH strength in starbursts by up to a factor of about two.   \\begin{figure}[ht] \\epsscale{1.2} \\plotone{f12.ps} \\caption{The PAH flux ratios $11.3\\mu$m/$7.7\\mu$m (diamonds) and $11.3\\mu$m/$6.2\\mu$m (filled circles, divided by a factor 2.05 to match the other ratio) versus the optical depth at $9.8\\mu$m, $\\tau_{9.8}$ from Table~\\ref{tabspecfits}.  The dashed line shows a linear fit to both ratios but excluding NGC~4945. \\label{fig11pahextinction}} \\end{figure}   In the remainder of this subsection we will investigate if the PAH equivalent width is related to the total infrared luminosity of the starburst.  \\citet{rig99} found a ratio of $L_{PAH} / L_{IR}$ which is similar for starburst-dominated ULIRGs and for template starbursts, i.e., no dependency on $L_{IR}$.  In contrast, \\citet{lu03} found a steady decrease of the PAH strength with increasing IR-activity. Figure~\\ref{figpahews} shows the equivalent widths of both the $6.2\\mu$m and $7.7\\mu$m PAH as a function of the total infrared luminosity for our starburst sample.  Within the uncertainties, indicated by the scatter between the $6.2\\mu$m and $7.7\\mu$m PAHs for the same object, the PAH equivalent widths remain constant over a factor of 50 in total luminosity.  This finding is in good agreement with \\citet{pee04} (and references therein) who found that the fraction of the total PAH flux emitted in the $6.2\\mu$m PAH band varies only slightly with an average of $28\\%\\pm 4\\%$.  In other words, the PAH flux and the underlying warm dust continuum scale proportionally, and the PAHs and dust must be well mixed, at least on large scales, to show these correlations.  This rules out a scenario of a luminous, dusty nucleus surrounded by a large PDR with little extinction, in favor of smaller, clumpier structures.   \\begin{figure}[ht] \\epsscale{1.2} \\plotone{f13.ps} \\caption{The $6.2\\mu$m (diamonds) and $7.7\\mu$m (filled circles) PAH equivalent widths versus infrared luminosity $L_{IR}$.  The dashed, horizontal line indicates a zeroth order polynomial fit to all data points, excluding the outlier NGC~1365.  The filled diamonds correspond to starbursts with an AGN component.\\label{figpahews}} \\end{figure}     \\subsection{PAH Luminosity and Star Formation Rate} \\label{pahsfr}  \\citet{ken98} has shown that the $8-1000\\mu$m infrared luminosity $L_{IR}$ of starbursts is a good measure of the star formation rate (SFR) given by \\begin{displaymath} \\mbox{SFR}[M_{\\odot}\\mbox{yr}^{-1}] = 4.5\\times 10^{-44} L_{IR} [\\mbox{erg\\ s}^{-1}]. \\end{displaymath} The SFR determines the number of young, massive stars, which provide the (far-)UV photons to excite both PAH molecules and dust grains.  If both species get excited by the same photons PAHs could potentially be used as quantitative tracers of the SFR (see section~\\ref{secbrandllumi} for the dust luminosity).  Generally, PAHs are considered the most efficient species for photoelectric heating \\citep{bak94} in the PDRs. Molecular gas in these boundary layers, surrounding the \\2 regions, is exposed to far-UV radiation ($6-13.6$\\,eV), which strongly influences its chemical and thermal structure \\citep{tie85}.  PAHs are stochastically heated by these UV photons, predominantly originating from massive stars, and hence expected to be good tracers of star formation.  On large angular scales -- similar to our case -- \\citet{for04} found that the $5-8.5\\mu$m PAH emission constitutes an excellent indicator of the star formation rate in circumnuclear regions and starbursts as quantified by the Lyman continum flux, i.e. in regions where the energy output is dominated by massive star formation.  However, it has been known for a long time that PAHS can also be excited by visible photons (e.g., \\citet{uch98}), and that PAHs can trace other sources besides massive young stars, such as planetary nebulae and reflection nebulae.  If the observed PAH flux is integrated over the whole galaxy it may predominantly trace B stars, which dominate the Galactic stellar energy budget, rather than very recent massive star formation \\citep{pee04}.  This is in agreement with higher angular resolution observations of the $3.3\\mu$m PAH feature at the {\\sl VLT} by \\citet{tac05}, who found a decrease in the PAH/continuum ratio at the sites of the most recent star formation.  In section \\ref{secpahmixture} we have seen that the PAH equivalent width does not dependent on $L_{IR}$.  However, that does not necessarily mean that PAHs aren't good quantitative tracers of star formation if the PAH feature and the underlying continuum scale proportionally.  In Fig.~\\ref{figpahsfr} we compare the flux in the $6.2\\mu$m PAH feature against the total infrared luminosity $L_{IR}$ from {\\sl IRAS}.  To search for a physically meaningful correlation one needs to take the distance of the object into account as well as the fact that the narrow {\\sl IRS} slit misses some of the total flux. Hence we multiply the measured PAH fluxes with the square of the distances and divide by the fractional flux factor FF (Table~\\ref{tabobsprop}).  A remarkably good fit can been achieved with \\begin{displaymath} \\log\\left(L_{IR}^{PAH}\\right) = 1.13\\times \\log\\left( F_{6.2\\mu m PAH} D^2 \\right) \\end{displaymath} where $F_{6.2\\mu m PAH}$ is the $6.2\\mu$m PAH flux in units of $10^{-19}\\mbox{W\\,cm}^{-2}$, $D$ the distance in kiloparsecs, and $L_{IR}$ in units of $L_{\\odot}$.  The standard error is 0.3 in $\\log(L_{IR})$.  Using the above equation, the total infrared luminosity of a starburst galaxy can be derived from the strength of a single PAH emission feature (here: the $6.2\\mu$m PAH) to within a factor of two.  This correlation is less tight than the estimate from the $F_{15\\mu m}$ and $F_{30\\mu m}$ continuum fluxes (section~\\ref{secbrandllumi}), supporting the finding by \\citet{pee04} that PAHs may not (only) trace recent massive star formation.   \\begin{figure}[ht] \\epsscale{1.2} \\plotone{f14.ps} \\caption{The estimated infrared luminosity $L_{IR}^{PAH}$ based on the aperture corrected $6.2\\mu$m PAH flux versus the total infrared luminosity as a measure of SFR.  The gray plus signs refer to the location of the same data points if no aperture correction is applied.  The dashed line represents a linear fit to the data (in log-log space), and the fitted function is given in the upper left, together with the standard error $\\sigma$.  The labeled outliers NGC~1365, NGC~4945 were excluded from the fit.  The filled symbols correspond to starbursts with a weak AGN component. \\label{figpahsfr}} \\end{figure}   Combining the information from Figs.~\\ref{figpahews} and \\ref{figpahsfr} suggests that the continuum and the PAH emission are, to first order, proportional.  Similarly, \\citet{pee04} have found that $L_{FIR}$ is proportional to the PAH luminosity $L_{6.2\\mu m}$.  If our spectra were ``contaminated'' by a significant amount of continuum emission from an AGN or an underlying, older galactic population unrelated to the starburst, the equivalent width would vary with distance (equivalent slit width), which is not observed.  Hence, we conclude that both PAH and continuum emission originate predominantly from the starburst.       \\subsection{PAH Strength and Radiation Field} \\label{secpahvari}  It is well known that the equivalent width of PAHs is much reduced in AGN-dominated environments (e.g., \\citet{stu00}, \\citet{gen00}, \\citet{wee05}).  \\citet{geb89} have discussed the susceptibility of PAHs to destruction by far UV fields, and {\\sl ISO} observations (e.g., \\citet{ces96}, \\citet{tra98}) have shown that more intense far-UV radiation fields may lead to gradual destruction of PAHs around stellar sources.  Recently, \\citet{ywu06} studied a sample of low-metallicity blue compact dwarf galaxies from $1/50 Z_{\\odot}$ to $1/1.5 Z_{\\odot}$ with the {\\sl IRS}.  They find a strong anti-correlation between strength (equivalent width) of the PAH features and the product of the [Ne\\,III]/[Ne\\,II] ratio (as a hardness measure of the radiation field) and the UV luminosity density devided by the metallicity.  A similar trend has been reported by \\citet{mad06} for a small sample of nearby dwarf galaxies. Unfortunately, lower metallicity and harder radiation fields seem to go hand-in-hand in these dwarf galaxies, and one cannot unambiguously distinguish between possibly suppressed PAH formation in low metallicity environments and PAH destruction in harder UV fields. Recently, \\citet{bei06} have investigated the strength of the $11.3\\mu$m PAH feature in the starburst in NGC~5253 for different radial distances, and found that the equivalent width of the PAH feature is inversely proportional to the intensity of the radiation field, suggesting photo-destruction of the aromatic carriers in harsher environments.  Observations of Galactic sources (e.g., \\citet{ver96}, \\citet{ver02}) have also shown that the relative strengths of individual PAH features can depend on the degree of ionization of the molecule: C-C stretching modes at $6.2\\mu$m and $7.7\\mu$m are stronger in ionized PAHs, while the C-H in-plane bending mode at $8.6\\mu$m and the C-H out-of-plane bending mode at $11.3\\mu$m are stronger by more than a factor of two in neutral PAHs.  Comparing the {\\sl ISOCAM} spectra of M\\,82, NGC~253, and NGC~1808, \\citet{for03} found that, while the $5-11\\mu$m spectrum is nearly invariant, the relative PAH intensities exhibit significant variations of $20\\% - 100\\%$, which they attributed to the PAH size distribution, ionization, dehydrogenation, or the incident radiation field.  In Fig.~\\ref{figpahrationeon} we compare the ratio between the $11.3\\mu$m bending mode and the $7.7\\mu$m stretching mode to the hardness of the radiation field, as indicated by the [Ne\\,III]/[Ne\\,II] ratio from \\citet{dev06}, who provide a detailed analysis of the fine structure lines in our starburst sample.  We correct both PAH fluxes for extinction (section~\\ref{secpahmixture}) using the extinction laws by \\citet{dra89} and \\citet{lut99} and the values from Table~\\ref{tabextlaws}.  Within the systematic uncertainties, represented by the scatter of the data points, we find no significant variation of the $7.7\\mu$m\\,/\\,$11.3\\mu$m PAH ratio over more than an order of magnitude in [Ne\\,III]/[Ne\\,II] fluxes. While we cannot exclude variations on scales of individual \\2 regions, our spatially averaged spectra of starburst nuclei do not reveal significant variations between the main PAH features.  We have also looked at the weaker PAH features from Table~\\ref{tabpahs}, but the scatter increases with lower signal-to-noise and more uncertain baseline definition, and does not reveal an obvious trend.  However, more studies, in particular of the PAH complex around $17\\mu$m are planned.   \\begin{figure}[ht] \\epsscale{1.2} \\plotone{f15.ps} \\caption{The ratio of the $7.7\\mu$m\\,/\\,$11.3\\mu$m PAH fluxes versus the [Ne\\,III]/[Ne\\,II] line ratio (from \\citet{dev06}).  The PAH fluxes have been corrected for extinction using the laws from \\citet{dra89} (diamonds) and \\citet{lut99} (squares). The gray pluses indicate the locations of the uncorrected values for reference.  The filled diamonds correspond to starbursts with weak AGN components.\\label{figpahrationeon}} \\end{figure}   Our starburst sample spans a wide range in radiation field hardness. Fig.~\\ref{figpahewsneon} shows the equivalent width of the $7.7\\mu$m PAH feature versus the fine structure line ratio [Ne\\,III]/[Ne\\,II], which has been taken from \\citet{dev06}.  Within the uncertainties the equivalent width of the $7.7\\mu$m PAH feature remains constant over more than an order of magnitude in [Ne\\,III]/[Ne\\,II].  We conclude that, {\\sl on large scales} of starburst nuclei, which typically contain numerous \\2 regions, the PAH~/~continuum ratio does not significantly depend on the average radiation field hardness.   \\begin{figure}[ht] \\epsscale{1.2} \\plotone{f16.ps} \\caption{The equivalent width of the $7.7\\mu$m PAH feature versus the [Ne\\,III]/[Ne\\,II] line ratio (from \\citet{dev06}) as an indicator of the hardness of the radiation field.  The filled diamonds correspond to starbursts with weak AGN components. NGC~1365, which has a significant AGN component, is a clear outlier. \\label{figpahewsneon}} \\end{figure}      We presented and discussed the $5 - 38\\mu$m mid-infrared spectra of a large sample of 22 starburst galaxies taken with the Infrared Spectrograph {\\sl IRS} on board the Spitzer Space Telescope.  The high signal-to-noise spectra contain numerous important diagnostics such as PAH emission features, silicate bands at $9.8\\mu$m and $18\\mu$m, and the shape of the spectral continuum.  The {\\sl IRS} spectral resolution of $R\\approx 65 - 130$ is perfectly matched to study these features.  From our sample we constructed an average starburst spectrum, which can be used as a starburst template.  The availability of continuous mid-infrared spectra of numerous objects within one class over a wide wavelength range enables various important studies.  Remarkably, the spectra show a vast range in starburst SEDs.  We found a trend that more dust extincted starburst systems have a steeper spectral continuum slope longward of $15\\mu$m. The slope can also be used to discriminate between starburst and AGN powered sources, with a transition at $0.17\\le F_{15\\mu m}/F_{30\\mu m} \\le 0.21$.  The monochromatic continuum fluxes, which represent a more accurate estimate of the ``true'' continuum than broad band filters, provide a remarkably accurate estimate of the total infrared luminosity via $L_{IR}^{est} = D^2\\cdot (4.27\\cdot F_{15\\mu m} + 11\\cdot F_{30\\mu m})$ (after correcting for slit losses for nearby, extended systems).  Our starburst spectra cover a wide range of silicate absorption depths, from essentially no absorption to heavily obscured systems with an optical depth of $\\tau_{9.8\\mu m}\\sim 5$.  We present the discovery of crystalline silicates in NGC~4945, which shows many similarities with heavily extincted ULIRGs.  However, unlike the latter, the starbursts in our sample show no signs of water ices or hydrocarbons, suggesting a small amount of self-shielding.  The observed spectra show significant variations in the relative strengths of the individual PAH features at $6.2\\mu$m, $7.7\\mu$m, and $11.3\\mu$m.  However, these variations may be entirely due to extinction and do not necessarily indicate intrinsic variations of the PAH spectrum.  We find that the PAH equivalent width is independent of the total luminosity $L_{IR}$, probably because the PAH strength and the underlying continuum scale proportionally within a ``pure'' starburst.  The luminosity of an individual PAH feature, however, scales with $L_{IR}$.  In particular the $6.2\\mu$m feature can be used to approximate the total infrared luminosity of the starburst (although less accurately than from the $15\\mu$m and $30\\mu$m continuum fluxes).  We investigated possible variations of the PAH spectrum as expected, e.g., from varying degrees of PAH ionization.  The $7.7\\mu$m\\,/\\,$11.3\\mu$m PAH ratios show no significant systematic variation with the hardness of the radiation field.  Although our sample covers about a factor of ten difference in radiation field hardness (as indicated by the [Ne\\,III]\\,/\\,[Ne\\,II] ratio) we found no systematic correlation with the PAH equivalent width.  Furthermore, we found no systematic differences between ``pure'' starbursts and galaxies with a weak, non-dominant AGN component for most of their spectral properties (except for NGC~1365 which shows very weak PAH emission).  We emphasize that these results are based on spatially integrated diagnostics over an entire starburst region.  Local variations of age, IMF, density or geometry on the scales of individual \\2 regions or super star clusters may just average out.  However, it is important to note that, because of this ``averaging out effect'' in unresolved sources, starburst nuclei with significantly different global properties may appear as rather similar members of one class of objects."
    },
    "0609/astro-ph0609815_arXiv.txt": {
        "abstract": "{We apply a matched-filter cluster detection algorithm to the Canada-France-Hawaii Telescope Legacy Survey (CFHTLS) $i-$band data for the Deep-1, Deep-2, Deep-3 and Deep-4 fields covering a total of 4~square degrees.  To test the implemented procedure we carry out simulations for assessing the frequency of noise peaks as well as estimate the recovery efficiency.  We estimate that up to $z\\sim0.7$ the catalogue is essentially complete for clusters of richness class $R\\gtrsim1$. The recovered redshifts are in general overestimated by $\\Delta z=0.1$ with a scatter of $\\sigma_{\\Delta z}\\sim0.1$, except at redshifts $z\\gtrsim1$ where the estimated redshifts are systematically underestimated.  The constructed cluster candidate catalogue contains 162 detections over an effective area of 3.112~square degrees corresponding to a density of $\\sim 52.1$ per square degree. The median estimated redshift of the candidates is $z=0.6$. The estimated noise frequency is $16.9\\pm5.4$ detections per square degree.  From visual inspection we identify systems that show a clear concentration of galaxies with similar colour. These systems have a density of $\\sim20$ per square degree. ",
        "introduction": "Clusters of galaxies, the most massive gravitationally bound systems in the Universe are located at the nodes of the cosmic web and are thus good tracers of the large-scale structures \\citep[e.g.,][]{bahcall88,huchra90,collins00}. Their intimate connection with the power spectrum of the initial density fluctuations as well as to the cosmological parameters make them important targets in observational cosmology \\citep[e.g.  ][]{bahcall97, oukbir97, holder01, schuecker03}. Additionally, they are still accreting matter along filaments and even at $z=0$ a large fraction of them do present signs of substructures in their morphology and in the galaxy velocity distribution. Therefore, a prerequisite for addressing any cosmological issue from a cluster sample is not only to have a large homogeneous sample with a well-understood selection function, but also a good understanding of cluster evolution including constraints on their dynamical state.  The evolution of galaxy clusters has in recent years obtained a lot of attention. This is partly due to the advent of large telescopes and sensitive instruments which allow to carry out detailed studies of high redshift systems. Many studies have concentrated on the colour of the red sequence galaxies. These have shown that most of the known systems exhibit properties similar to those of present day clusters \\citep[e.g., ][]{stanford97, stanford98, vandokkum00,lidman04}, when taking into account the passive evolution of the stellar populations. Most of the studied systems have been X-ray selected, which is likely to trace the most massive systems at any redshift. Therefore, the current studies may be biased towards old, well-evolved systems missing the true progenitors of present day clusters and galaxies \\citep{vandokkum01a}. In order to better understand such biases and selection effects it is important to use different detection methods and wavelength regimes.   A number of different detection algorithms and wavelengths have been used for identifying systems at increasingly higher redshifts.  The largest cluster samples have been constructed using either optical \\citep[e.g.,][]{postman96,olsen99b,gladders01,depropris02,postman02,goto02,kim02,gal03,bahcall03,miller05,tago06} or X-ray \\citep[e.g.,][ and references therein]{bohringer00,bohringer04} data. More recently, the Sunyaev-Zeldovich effect has also been used for detecting clusters of galaxies \\citep{carlstrom02}.  Each method is based on one or at most a few characteristics derived from known clusters and may thus introduce biases against certain types of systems.  Therefore, it is of utmost importance for each method to have a well-defined selection function including a clear understanding of the built-in biases. Furthermore, directly comparing catalogues covering the same area extracted by different methods will yield important insight into the different biases introduced by the different methods.  The Sloan Digital Sky Survey (SDSS) data with its large area and multi-colour coverage, as well as complementing spectroscopic information, has recently provided an important test-bed for a number of optical cluster detection methods at low redshifts \\citep[e.g.,][]{kepner99,kim02,goto02}. Thorough comparisons between the different methods have been carried out by \\cite{kim02,goto02,bahcall03}. These comparisons show that, not surprisingly, there are always differences between the various catalogues, part of which are caused by the different ways the parameters such as for example richness are estimated. The SDSS data are sufficient for detecting clusters to at most intermediate redshifts ($z\\sim0.5$). At higher redshifts only smaller data samples have been available, such as the KPNO/Deeprange survey by \\cite{postman02} and the ESO Imaging Survey \\citep{olsen99a,olsen99b,scodeggio99}. Recently, the Red sequence Cluster Survey by \\citet[RCS,][]{gladders05} covering 100~square degrees has been achieved thus starting to probe large volumes to high redshift. However, a detailed comparison of the efficiency of different methods at high-z has yet to be carried out.   For detailed comparison between different methods, wide, deep and preferentially multi-passband homogeneous surveys are required in order to provide the necessary data for a number of different detection methods. The design of the Wide survey of the Canada-France-Hawaii Telescope Legacy Survey (CFHTLS) provides a data set that is particularly well-suited for carrying out such studies. This survey is currently underway and is planned to cover $\\sim170$~square degrees in 5 passbands spread in 4 patches with limiting magnitudes up to 25$^m$. It will provide the necessary ground for building a large, well-controlled cluster candidate sample at redshifts $z\\lesssim1.3$, derived from a set of different search techniques using the spatial as well as photometric properties in one or more passbands.  Such catalogues will allow us to test accurately the $z>0.5$ component of the cluster distribution. Using automatic search techniques will allow us to build selection functions for each of our independently extracted catalogues.  A careful comparison between the various independently extracted cluster samples will allow us to understand the additional difficulties in detecting clusters at successively higher redshifts as well as improve our knowledge about clusters at these redshifts.  Combining the catalogues from several searches we will create a robust cluster sample well suited for both cosmological and galaxy evolution studies.   This paper is the first in a series describing detection and investigations of primarily the photometric properties of the identified systems. In this first paper we describe our implementation of the matched filter detection method which we apply to the $i-$band data of the Deep Survey of the CFHTLS. The depth presently reached by the Deep CFHTLS corresponds roughly to the final depth of the Wide part of the survey. Furthermore, the small size of the Deep provides a good test-bed for various detection algorithms for investigating the potential of the Wide survey. The motivation for exploring the $i-$band in this first paper is two-folded: first, the survey area is planned to be covered by $i-$band as high priority, later followed by the other passbands; second, most previous optical cluster searches have been carried out in $i-$band, thus focusing on this filter facilitates comparisons with previous surveys minimising the effect of the variation in wavelength. In future papers we will apply the same algorithm to data in other passbands in order to investigate how sensitive this technique is to the choice of passband at different redshifts. To investigate the advantages of using additional colour information we will apply a colour search technique to the same data and cross-compare the detections between the different methods. In parallel we are working on detection methods based on photometric redshift slicing to fully exploit the multi-passband data for identifying structures, as well as on searches based on weak lensing studies \\citep[see][]{gavazzi06}. In future papers we will also characterize the detected systems in terms of total luminosity, concentration, shape parameters and colour properties such as the existence and significance of the red sequence.    The present paper is structured as follows: In Sect.~\\ref{sec:galcats} we describe the galaxy catalogues used in the present work. These catalogues are a modified version of the Terapix CFHTLS galaxy catalogues. Sect.~\\ref{sec:method} describes the matched-filter cluster identification method which is an improved version of that implemented for the ESO Imaging Survey \\citep{olsen99a}. This section also describes a series of simple simulations to account for the rate of false-positives as well as the selection function. In Sect.~\\ref{sec:catalogues} we apply the detection to the $i-$band galaxy catalogue of the four Deep fields of the survey covering a total of $\\sim4~\\mathrm{square\\, degrees}$ and compare the constructed cluster catalogue to those of previous works. In Sect.~\\ref{sec:summary} we summarize our findings.    Throughout the paper we use a cosmological model with $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$ and $H_0=75\\mathrm{km/s/Mpc}$\\footnote{We use $h_{75}=\\frac{H_0}{75\\mathrm{km/s/Mpc}}$}. ",
        "conclusions": "\\label{sec:summary}  In this paper we present the first catalogue of optical cluster candidates extracted from the CFHTLS Deep data. Using an improved implementation of the matched-filter procedure of the EIS project \\citep{olsen99a} we construct cluster catalogues in the $i-$band for the Deep fields of the present survey. Through simple simulations we assess the rate of false detections as well as the recovery rate.  The main properties of the catalogue are the following:  \\begin{itemize}  \\item The catalogue contains 162 clusters covering the redshift range from $z=0.2$ to $1.1$ with a median of $z_{med}=0.6$. The density of candidates is $52.1\\pm7.8$ per square degree; among them, $\\sim20$ per square degree show a concentration of similarly coloured galaxies.  \\item The estimated rate of false detections is $\\sim16.9\\pm5.4$ per square degree. This density is consistent with the fraction of systems not showing any obvious concentration in the visual inspection.  \\item From simulations we find that the catalogue is complete for systems of richness class $\\geq1$ up to $z=0.7$; beyond that the recovery rate decreases to close to zero at $z\\sim1.2$.  \\item We find that the estimated redshifts are in general overestimated by $\\Delta z\\sim0.1$ with a scatter of $\\sigma_{\\Delta z}\\sim0.1$. Correcting for this redshift offset, the recovered richnesses are in good agreement with the input.  \\item The present catalogue extracted from the CFHTLS Deep fields appears to trace poorer systems at higher redshifts than previous matched-filter cluster catalogues.  \\end{itemize}  We have compared our catalogue with that of the XMM-LSS survey \\citep{pierre06}. From this comparison we found that our grade A systems in the region of overlap between the two surveys were detected by XMM-LSS. The remaining X-ray detections that were not included in our catalogue were either in masked regions or appeared optically poor. The grade B and C systems were not included in the known (X-ray) cluster samples.  We conclude that, the CFHTLS imaging survey provides a good basis for constructing large samples of galaxy clusters to study the evolution of cluster proporties as well as for other cosmological applications. Furthermore, the ability to trace the poor cluster population may allow for studying the processes of cluster growth and their impact on the galaxy populations and evolution in more detail than with previous samples. In order to gain a better understanding of the detected systems we will use photometric redshifts to investigate their properties. However, to fully describe their nature we have to carry out a thorough spectroscopic follow-up."
    },
    "0609/astro-ph0609738_arXiv.txt": {
        "abstract": "We select far-infrared (FIR-$60\\;\\mu$m) and far-ultraviolet (FUV-1530~\\AA) samples of nearby galaxies in order to discuss the biases encountered by monochromatic surveys (FIR or FUV). Very different volumes are sampled by each selection and much care is taken to apply volume corrections to all the analyses. The distributions of the bolometric luminosity of young stars are compared for both samples: they are found to be consistent with each other for galaxies of intermediate luminosities but some differences are found for high ($>5 ~ 10^{10} L_\\sun$) luminosities. The shallowness of the {\\sl IRAS} survey prevents us from  securing comparison at low luminosities ($<2 ~ 10^9 L_\\sun$). The ratio of the total infrared (TIR) luminosity to the FUV luminosity is found to increase with the bolometric luminosity in a similar way for both samples up to $5 ~ 10^{10} L_\\sun$. Brighter galaxies are found to have a different behavior according to their selection: the $L_{\\rm TIR}/L_{FUV}$ ratio of the FUV-selected galaxies brighter than $5~10^{10} L_\\sun$ reaches a plateau whereas $L_{\\rm TIR}/L_{FUV}$ continues to increase with the luminosity of bright galaxies selected in FIR. The volume-averaged specific star formation rate (SFR per unit galaxy stellar mass, SSFR) is found to decrease toward massive galaxies within each selection. The SSFR is found to be larger than that measured for optical and NIR-selected sample over the whole mass range for the FIR selection,  and for masses larger than $10^{10}\\; M_\\sun$ for the FUV selection. Luminous and massive galaxies selected in FIR appear as active as galaxies with similar characteristics detected at $z\\sim 0.7$. ",
        "introduction": "Many of the most recent galaxy surveys have attempted to gain a better understanding of the evolution of the star formation rate (SFR) with time and environment.  Because of the spectral redshifting, deep (high redshift) optical surveys in fact sample the far-ultraviolet (FUV) rest-frame emission of the target galaxies.  As a consequence, numerous measurements of the star formation activity of galaxies as a function of redshift ($z$) are based on restframe FUV data obtained from imaging and spectroscopic surveys: at low z \\citep[e.g.,][]{lilly,schim,baldry} or at higher z, using drop-out selection techniques \\citep[e.g.,][]{steidel,bunker,ouchi,giavalisco}. However, the attenuation of the FUV light by interstellar dust is a major issue in the derivation of quantitative SFR from the FUV even at low $z$ \\citep[e.g.,][]{buat05,seibert,cortese}.  Recent infrared surveys \\citep[e.g.,][]{flores,lefloch,perez} have also contributed significantly to the study of the star formation history in the universe \\citep[e.g.,][]{rowan01,takeuchi01a,takeuchi01b,lagache03}: the far-infrared (FIR) emission from the dust heated by hot stars is, by definition, not affected by dust attenuation. However, the FIR emission has its own drawbacks: the calibration of the dust emission into a quantitative SFR usually relies on the strong assumption that most of the dust heating is due to young stars and that all the light from these stars is absorbed by dust and re-emitted in FIR \\citep[e.g.,][]{kennicutt98}.  This is clearly an over-simplification: most galaxies are seen to emit in the FUV-optical and, as a consequence all the light emitted by young stars cannot have been absorbed by dust.  Moreover, substantial dust heating by older stars cannot be ruled out for all galaxies and more complex calibrations have to be undertaken \\citep{buatxu,bell03,hirashita,igle06}.  Another issue is that infrared observations are carried out at one or a few wavelengths (in the mid and/or far-infrared) whereas it is the total infrared (TIR) emission that is required for a star formation rate calculation. Unfortunately, the TIR emission is derived by extrapolation from measured only those few measured fluxes using models, and the correction factors range from 2 to 3, when observations are made in the far infrared, and up to 8-10 when only mid-infrared data are available \\citep[e.g.,][]{ttt3,lefloch}.  The  ultraviolet and  the FIR  emission  are each  strongly linked  to alternate  manifestations  of  the  recent star  formation  rate:  the ''transparent\" one in FUV and the ''hidden\" one in FIR. Obviously a very promising  way to  proceed would  be  to combine  both wavelengths  to perform a  more inclusive and multiwavelength analysis  of the current star formation in the universe.  What do we know today about the FUV and FIR universe?  From a global point of view, the recent observations conducted by {\\sl SPITZER} and {\\sl GALEX} have provided insight into the TIR and FUV luminosity functions and densities from $z=0$ to $z=1$ \\citep[][]{schim,lefloch,perez,arn}. The shapes of the luminosity functions in both wavelength ranges are found to be very different \\citep[][]{ttt2} as previously emphasized by \\citet{buat98}.  The evolution of the luminosity functions and the derived star formation densities were studied in the FUV \\citep[][]{schim} and at infrared wavelengths \\citep[][]{lefloch,perez}: a strong evolution was found at both wavelengths, with a net increase of the luminosity density from $z=0$ to $z=1$. Nevertheless the evolution appears to be stronger in the FIR than in the FUV implying a {\\it global} increase of the dust attenuation from $z=0$ to $z=1$ by $\\sim 1\\;\\mbox{mag}$ \\citep[][]{ttt2}. Such an increase might be explained, at least qualitatively, by the larger fraction of bright galaxies at high redshift, combined with the known positive correlation of the dust attenuation with the absolute luminosity of the host galaxies \\citep[e.g.][]{wang,buat98,hopkins01,martin}. In a recent analysis of {\\sl SWIRE/GALEX} data \\citet{xu06apjs} find no significant differences between the FIR-to-FUV flux ratios of star-forming galaxies at z=0.6 and their local counterparts of similar SFR. They argue that the evolution of the dust attenuation with redshift is primarily due to the SFR evolution and to the strong dependence of the attenuation on SFR itself.  Indeed, most recent studies based on large surveys aim at better understanding of which galaxies are at the origin of the variation of the star formation density with $z$, and especially its decrease from $z\\sim 1$ to 0, which is seen at almost all wavelengths \\citep[e.g.,][and references therein]{hopkins06}. {}From the UV-optical side, almost all recent surveys have found a strong increase in luminosity and/or space density of late-type, blue galaxies although some discrepancies have been noted. Differences in the definition of galaxy types from one a study to another makes a direct comparison of the results problematic \\citep[e.g.,][and references therein]{delapparent,wolf}. {}From the IR side, recent results from the {\\sl SPITZER} Space Telescope seem to attribute the general decrease of the star formation density to a decrease in the SFR in massive spirals \\citep[][]{bell} without strong interactions. Studies, both in the IR and in the optical, suggest a minimal role of strong mergers in the evolution of the star formation density from z=1 to z=0 \\citep[][]{lotz,bell}  {}Connecting from low- to high-$z$ that which is seen in the rest-frame FIR or in FUV, is a new challenge.  Do we see the same galaxy populations in FUV and FIR evolving as a whole and appearing differently in FUV and in FIR, or must we invoke sub-populations of galaxies evolving independently with $z$ and/or being detected only at one wavelength, FUV or FIR? This is a crucial question: if we are observing the same populations in both wavelengths ranges, then, with some justification, we can try to predict the total star formation from single-band surveys (assuming some relevant corrections). But if other populations are present, it would appear to be impossible to reconstruct all of the star formation activity from a single-band survey.  The first step consists of obtaining a reliable reference dataset in the local universe.  Thanks to {\\sl GALEX}, large samples of nearby galaxies observed in the FUV are now available. Making use of the existing IR surveys, we can then build robust reference samples of galaxies selected in the FUV and FIR  which are suitable for comparison.  The aim of this paper is to take advantage of the {\\sl GALEX} shallow survey to build large samples of nearby galaxies selected in the FUV (or in the FIR by {\\sl IRAS}) and to use these samples to analyse the selection biases and the consistency of the FIR and FUV LFs at $z=0$.  We can then build large reference samples of galaxies selected  at $60\\;\\mu$m, say,  in such a way as to allow a very good detection rate at 1530\\AA\\, and vice versa (section 2). In section 3 we emphasize the intrinsic differences existing between a FIR and a FUV selection.  The relative contributions of the TIR and FUV emissions to the luminosity of the young stellar populations in galaxies is assessed in section 4 using a cross-comparison of the luminosity functions in both samples; bolometric luminosity functions are then built to check if we can indeed see all the galaxies at a single wavelength. In section 5, we discuss the star formation activity as a function of the stellar mass in both samples through an analysis of the specific star formation rate (SSFR, SFR per unit galaxy stellar mass). Section 6 is devoted to the conclusions. This works extends the earlier studies of \\citet{martin}, \\citet{igle06} and \\citet{xu06}, which were based on smaller samples.  Throughout this article, we use the cosmological concordance parameters of $H_0 = 72$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M = 0.3$ and $\\Omega_{\\lambda} = 0.7$. All magnitudes are quoted in the AB system. ",
        "conclusions": "\\begin{enumerate} \\item We have built large FIR and FUV-selected galaxy samples fully representative of the nearby universe. The two selections  are found to sample very different volumes, which is a direct consequence of the very different shape of the luminosity functions at the two wavelengths. Therefore one  must apply volume corrections before comparing the mean properties of these samples.  \\item The bolometric luminosity of newly formed stars is estimated by combining the infrared and ultraviolet fluxes and accounting for dust heating by old stars.  The bolometric  LF is calculated and  found to be different in the FUV and the FIR selections. Intrinsically bright galaxies are under-sampled by a FUV selection. No faint galaxy  ($L_{\\rm bol} < 10^9 L_\\sun$) is found  in our FIR sample. We must wait for deeper FIR imaging surveys with a better spatial resolution than {\\sl IRAS} to compare accurately the lowest bins of the  bolometric luminosity functions and investigate the contribution of close pairs and interacting systems.  \\item The ratio of the total IR luminosity to the FUV ($L_{\\rm TIR}/L_{\\rm FUV}$) is found to be strongly related to  the TIR and to the bolometric luminosity for both samples. No universal trend is found with the FUV luminosity, making irrelevant any dust attenuation correction based on a monochromatic FUV luminosity alone. The volume-averaged relation between $L_{\\rm bol}$ and $L_{\\rm TIR}/L_{\\rm FUV}$ is found to be similar for both (FUV and FIR-selected) samples for bolometric luminosities between $10^9 L_\\sun$ to $5~10^{10} L_\\sun$. The monotonic increase of $L_{\\rm TIR}/L_{\\rm FUV}$ with $L_{\\rm bol}$ continues up to $10^{12} L_\\sun$ within the FIR selection. $L_{\\rm TIR}/L_{\\rm FUV}$ saturates for FUV-selected galaxies more luminous than $5 ~ 10^{10} L_\\sun$ at a value corresponding to a dust attenuation of $\\sim 2.5$ mag in FUV.  \\item The specific star formation rate is analysed as a function of the stellar mass. It is found to decrease as the galaxy mass increases at both wavelengths and as soon as volume corrections are applied to the samples. Massive, FUV-selected galaxies and all the FIR-selected ones exhibit a larger specific star formation rate than that deduced from optical-NIR surveys of nearby galaxies with similar stellar masses selected in optical or in the NIR. Luminous FIR-selected galaxies ($L_{\\rm bol}>10^{11}~L_\\sun$)  have a very large specific star formation rates,  similar to those found at $z=0.7$ by \\citet[][]{bell} for FIR-luminous galaxies of similar mass and luminosities. \\end{enumerate}"
    },
    "0609/astro-ph0609212_arXiv.txt": {
        "abstract": "We report subarcsecond-resolution \\textit{VLA} imaging of four nearby spiral galaxies: IC 342, Maffei II, NGC 2903, and NGC 6946. In each galaxy, compact radio continuum sources are identified in the central $\\sim$ $15\\arcsec \\times 15\\arcsec$ region. These compact sources are responsible for 20 - 30\\,\\% of the total emission from the central kpc of the host galaxies at 2 cm, but only $\\sim$ 5 - 10\\,\\% at 6 cm. More than half of the compact sources appear to be \\HII\\ regions. The \\HII\\ regions with rising spectra must be fairly dense ($n_{i} \\sim 10^{4}~cm^{-3}$) and are presumably very young. The largest of these \\HII\\ regions require the excitation of 500 - 800 O stars, within regions of only few parsecs extent. These clusters approach the sizes expected for globular clusters. Thermal free-free emission from compact sources contributes more significantly at 2 cm, while diffuse synchrotron emission dominates at 6 cm. The radio \\HII\\ regions are found near the centers of giant molecular clouds in projection, and do not have obvious visual counterparts. ",
        "introduction": "Young massive star clusters, potential young globular clusters, have been revealed by ground-based imaging \\citep{1985AJ.....90.1163A} and \\textit{Hubble Space Telescope} (\\textit{HST}) observations (see review by \\citealt{2003dhst.symp..153W}). These clusters appear to be a few Myr in age, $10^{3}$ - $10^{7}~M_{\\sun}$ in mass, and 10 pc or smaller in diameter. The youth of these clusters means that, unlike globular clusters, their luminosities are dominated by a population of massive stars. Young, massive star clusters appear to be common \\citep{2002AJ....123.1389B} in the centers of late-type spirals, and may play an important role in galactic nuclear activity. How these large star clusters form and the environments favoring their formation are not understood. The identification of young super star clusters in nearby galaxies gives us the chance to study these clusters in the process of formation.  In nearby starburst galaxies, extremely compact sources with rising spectral indices at centimeter wavelengths have been discovered, which are thought to be ``supernebulae\" surrounding very young super star clusters (SSC) \\citep{1998AJ....116.1212T, 1999IAUS..193..758T, 1999ApJ...527..154K, 2000AJ....120..244B,2000A&A...358...95T, 2002MNRAS.334..912M}. These nebulae have the characteristics of dense, compact \\HII\\ regions in the Galaxy, and by analogy are also expected to be very young, $\\sim$1 Myr or less. However, they are much larger than Galactic compact \\HII\\ regions, requiring the excitation of hundreds to thousands of young O stars \\citep{1998AJ....116.1212T}. These young clusters may be optically obscured by their natal gas and dust clouds \\citep{1989ApJ...340..265W}, particularly in molecular gas-rich galactic centers. Free-free emission from embedded \\HII\\ regions at radio wavelengths is unaffected by extinction from molecular clouds. Radio continuum emission can be used to search for the youngest, potentially embedded young \\HII\\ regions in other galaxies.  The goal of this investigation is to identify young SSC candidates in four nearby and well-studied spiral galaxies, IC 342, Maffei II, NGC 2903, and NGC 6946 using subarcsecond radio continuum imaging. The centers of these galaxies are infrared-bright \\citep{1980ApJ...236..441B, 1980ApJ...235..392T, 1985ApJ...290..108W, 1989ApJ...344..135H} and molecular gas rich \\citep{1977ApJ...218L..51R,1982ApJ...258..467Y,1985ApJS...57..261V}. Previous, lower ($\\sim 1\\arcsec$) resolution \\textit{VLA} maps have shown that the four galaxies have concentrated regions of star formation in their centers \\citep{1983ApJ...268L..79T, 1994ApJ...421..122T, 1985ApJ...290..108W}. However, the low resolution of these maps, $\\sim$ 20 - 70 pc on the galaxy, combines emission from \\HII\\ regions, supernova remnants (SNRs), and extended non-thermal emission within the beam \\citep{1992ARA&A..30..575C}. To isolate nebulae around individual star clusters we need subarcsecond resolution.  Subarcsecond imaging with the \\textit{VLA} in its extended configuration maximizes sensitivity to bright and compact radio nebulae over low brightness disk synchrotron emission. Observing at shorter wavelengths, $\\lambda<$ 2 cm, also minimizes the contribution of synchrotron emission, which falls with frequency. We present maps of the radio continuum emission at 6 and 2 cm in IC 342, Maffei II, NGC 2903, and NGC 6946 with the A configuration of the \\textit{VLA}. The images have resolutions of 0\\farcs3 at 6 cm and 0\\farcs1 at 2 cm, corresponding to size scales of 3 (IC 342 and Maffei~2) - 11 (NGC 6946) pc on the galaxies. ",
        "conclusions": "We have imaged compact radio emission in the centers of four nearby spiral galaxies, IC 342, Maffei II, NGC 2903, and NGC 6946, at subarcsecond resolution at 6 and 2 cm with the VLA. Compact radio sources contribute $\\sim$ 20 - 30\\,\\% of the total 2 cm emission from the central kpc of these galaxies, but only 5 - 10\\,\\% at 6 cm; in NGC 6946, they contribute $\\sim$ 30\\,\\% at both wavelengths. The 2 cm emission within the central 100 - 200~pc of these spirals is dominated by \\HII\\ regions, while diffuse synchrotron emission dominates at 6 cm. Over half of the 38 identified compact sources appear to be \\HII\\ regions, based on their radio spectra; of these, one-third to half appear to have optically thick free-free emission. The high resolution of these maps selects for these compact, high brightness \\HII\\ regions. The brightness temperatures of these optically thick, compact \\HII\\ regions are consistent with this identification, and indicate actual source diameters of about 1-3~pc ($\\sim 0.02\\arcsec$ - $0.05\\arcsec$.) The compact \\HII\\ regions contain ionized gas with densities as high as $n_{i} \\sim 10^{4}~ cm^{-3}$, which suggests that these \\HII\\ regions are relatively young. Few of the compact radio sources have candidate counterparts in the visual wavebands. Comparison with optical $HST$ images show that the young \\HII\\ regions  tend to be found in dust lanes, and probably are subject to high extinction and also dynamical influences from molecular clouds. The predicted infrared emission from these \\HII\\ regions contributes 3 - 5\\,\\% of the total infrared luminosity of the galaxies. The largest \\HII\\ regions require equivalent of $\\sim$ 500 - 1,000 O7 stars to excite them; since they are dense, compact \\HII\\ regions, by analogy with Galactic compact \\HII\\ regions, they are likely to be very youngest of the massive young clusters, a Myr or less in age."
    },
    "0609/astro-ph0609804.txt": {
        "abstract": "\\citeauthor{howell2006} have reported the discovery of SN~Ia SN~2003fg (SNLS-03D3bb) and % from an analysis conclude that SN~2003fg is very likely a super-Chandrasekhar-mass SN~Ia perhaps with a mass $\\sim 2\\,M_{\\odot}$.    % H2006 p. 3 and 13 suggests of order 2 is more what they say than >= 2. Their work is the first strong evidence that has been presented for a super-Chandrasekhar SN~Ia. We have performed an analysis of the SN~2003fg data using the \\citeauthor{yoon2005} binding energy formula for a rotating super-Chandrasekhar-mass white dwarf (also used by \\citeauthor{howell2006}) included in a simple model of SNe~Ia (which we call the SSC~model for Simple Super-Chandrasekhar model for SNe~Ia) which assumes spherically symmetric ejecta and relies on the approximations of an exponential density profile for SN~Ia ejecta and of a sharp boundary of the SN~Ia iron-peak-element core. Our results support the conclusion of \\citeauthor{howell2006}:  SN~2003fg is very probably super-Chandrasekhar and probably has mass $\\sim 2\\,M_{\\odot}$. ",
        "introduction": "}  Recently, \\citet[hereafter H2006]{howell2006}, as part of the Supernova Legacy Survey (SNLS), have reported the discovery of SN~Ia SN~2003fg (supernova SNLS-03D3bb in the SNLS naming scheme). SN~2003fg is a remote supernova at $z=0.2440\\pm0.0003$ and using concordance cosmology ($H_{0}=70\\,{\\rm km\\,s^{-1}\\,Mpc^{-1}}$, $\\Omega_{M}=0.3$, flat universe) % This is sort of a round number concordance model. % See http://cosmos.swin.edu.au/entries/concordancemodel/concordancemodel.html?e=1 is found to be about 2.2 times brighter in absolute $V$ magnitude than the median brightness of SNe~Ia:  it is intrinsically the brightest SN~Ia known with any confidence. SN~2003fg's lightcurves in shape are not out of the range of normal SN~Ia behavior. Thus, SN~2003fg strongly violates the lightcurve width-luminosity relationship \\citep{phillips1993,phillips1999} which is so useful in cosmology. % H2006's analysis of SN~2003fg leads them to conclude that SN~2003fg is super-Chandrasekhar SN~Ia: % i.e., that it has a mass greater than the Chandrasekhar mass $1.38\\,M_{\\odot}$). %% I bring this up below in section 1 in its rightful place.  The one spectrum published by H2006, which comes from 2 days after $B$ maximum (light) (which H2006 assume to also be bolometric maximum light), % (Their discussion is unclear on this, but Figure captions 1, 2, and 4 seem to say % this in an ambiguous fashion????) looks typical of near-$B$-maximum SN~Ia spectra as H2006's Figure~3 shows by a comparison to the very normal SN~Ia SN~1994D (a core normal SN~Ia in the terminology of \\citet{branch2006}). The spectrum line features, however, are distinctly narrow. The narrowness suggests somewhat low ejecta velocities for the region of line formation in comparison to normal SNe~Ia at the same phase. (The width of the lines is determined by the Doppler shift of line opacity in the expanding ejecta.) The line velocity of SN~2003fg's Si~II~$\\lambda6355$ line \\citep[e.g.,][p.~79]{wiese1969} (the line which is the most characteristic line of SNe~Ia) in the spectrum is $8000\\pm500\\,{\\rm km\\,s^{-1}}$. (Line velocity is the Doppler shift velocity corresponding to the wavelength shift of a P~Cygni line absorption trough minimum from the line center wavelength. As supernovae evolve the line velocities generally decrease because the ejecta density is falling and the region of sufficient opacity to form lines recedes into the ejecta in mass fraction and velocity.) Core normal SNe~Ia have photospheric velocities of $\\sim 11000\\,{\\rm km\\,s^{-1}}$ at 2 days after $B$ maximum \\citep{branch2005,branch2006}. Since the photospheric velocity is a lower limit on line velocities with normal line formation, SN~2003fg's photospheric velocity was unusually low for its phase and probably $\\sim 8000\\,{\\rm km\\,s^{-1}}$ which is, in fact, the photospheric velocity given by H2006 in their supplementary information\\footnote{See URL http://www.nature.com/nature/journal/v443/n7109/suppinfo/nature05103.html\\ .}  Note that supernovae after very early times have all mass elements approximately in uniform motion, and thus all ejecta structures just scale up linearly with time since explosion $t$. % (e.g., all mass element radii obey $r\\approx vt$, where $v$ is radial velocity and $t$ is time since explosion): This kind of expansion is called homologous expansion. In homologous expansion, the radii $r$ of all mass elements obey \\begin{equation} r\\approx vt  \\,\\, , \\label{eq-homologous} \\end{equation} where $v$ is the mass element velocity and initial radial position in the progenitor is considered negligible. Because of equation~(\\ref{eq-homologous}), velocity becomes a good comoving frame coordinate for describing supernova ejecta and we conventionally use it as such. Also in homologous expansion, all mass element densities scale as $t^{-3}$.  The photospheric velocity is the characteristic velocity of the layer of continuum emission. In the case of SNe~Ia, the continuum is not a pure blackbody continuum. Because of falling density and thus opacity, the photosphere, like the line formation region, recedes into the ejecta in velocity coordinate with passing time. As the photosphere recedes, the P~Cygni lines gradually transform to emission lines. Eventually, the photosphere vanishes altogether. The photospheric phase of a supernova is when the photosphere is still important to the radiative transfer. After the photospheric phase is the nebular phase. The transition between the two phases is, of course, gradual.  One point to emphasize is that the SN~2003fg spectrum in the optical was dominated by intermediate-mass element (IME) lines and not iron-peak element (IPE) lines. (We consider carbon through calcium as IMEs and scandium through nickel as IPEs.) Optical spectrum dominance in the early post-$B$-maximum phase by IME lines (principally those of Si~II, S~II, and Ca~II) is true for all SNe~Ia, not just normal ones. This shows that the spectrum formation is not primarily in what we call the IPE core of the ejecta, but at higher velocities. Most SNe~Ia show IME-line dominance in the optical spectra in the pre-$B$-maximum and $B$-maximum phases as well.  The IPE core is the interior region of many SN~Ia explosion models that is almost entirely IPEs. At time zero after the explosion the IPE core is in fact mainly radioactive $^{56}$Ni. The IPE core is a standard feature of explosion models that result from the standard SN~Ia model. In the standard SN~Ia model, the progenitor is a carbon-oxygen white dwarf (a CO WD with about equal amounts of C and O) that is very close to the physical Chandrasekhar mass ($1.38\\,M_{\\odot}$ for CO WDs with equal amounts of C and O: see Appendix~\\ref{ap-chan-mass}) with typical central densities of order $3\\times10^{9}\\,{\\rm g\\,cm^{3}}$ and typical central temperatures of order $2.5\\times10^{8}\\,$K \\citep[e.g.,][p.~108]{woosley1994}. With such densities and temperatures unstable carbon burning (i.e., carbon ignition) at or near the WD center will initiate an explosion that totally disrupts the WD and creates a SN~Ia.  The WD is driven to the explosive state by mass accretion from a binary companion. The nature of the companion is not certain at present. Two possibilities exist (which are not exclusive) for the companion: it could be a post-main-sequence star (or very late phase main-sequence star) of some kind  % tornambe2005 gives the late phase case. or it could be another CO WD that is merging with the first WD \\citep[e.g.,][and references therein]{tornambe2005}. The first possibility is the single-degenerate (SD) scenario and the second, the double-degenerate (DD) scenario. In the SD scenario the accreted matter (which would be mostly hydrogen) burns to CO. There are problems with both scenarios. In the SD scenario, the accreted matter may be mostly ejected in some way because of unstable helium burning in the accreted matter and this prevents the WD central region from reaching explosive conditions (e.g., \\citealt[and references therein]{piersanti2003}; \\citealt[and references therein]{tornambe2005}). In the DD scenario, it has been found that an off-center carbon ignition may lead to conversion of the CO WD to an O-Ne-Mg WD \\citep[e.g.,][]{saio2004} and perhaps, depending on conditions, to collapse to a neutron star rather than a SN~Ia explosion \\citep[e.g.,][]{nomoto1991}. Somehow one of the scenarios or both of the scenarios avoid non-explosive fates and allow the CO WD to reach the central conditions that lead to a SN~Ia explosion.  If a WD gets sufficiently close to the physical Chandrasekhar mass (and does not have high rotation:  see below), it should collapse to a neutron star. Electron capture on nuclei brings about collapse before the physical Chandrasekhar mass is reached:  see Appendix~\\ref{ap-chan-mass}. SNe~Ia happen (at least in some cases) because of carbon ignition in the central regions of WDs (as mentioned above) before the point of collapse is reached.  In a SN~Ia explosion, nuclear burning turns the inner part of the ejecta into nearly pure IPEs with the dominant isotope being radioactive $^{56}$Ni:  this inner part is the IPE core. Outside of the IPE core, IPEs are trace elements. Immediately outside of the IPE core, the composition in most explosion models is dominated by explosion-synthesized silicon and sulfur. Above that is a layer dominated by explosion-synthesized and pre-existing oxygen with probably relatively large abundances of explosion-synthesized silicon and magnesium. There may be an outer layer or region of pre-existing CO that survives the burning in some SNe~Ia. Carbon lines have been tentatively identified in a few SNe~Ia \\citep[e.g.,][]{jeffery1992,fisher1999,branch2003}. % However, in normal SNe~Ia it is probable that the unburnt % CO is only a hundredth or much less of the IMEs \\citep{marion2006}, and so a substantial % CO layer is not likely for normal SNe~Ia.  The existence of a somewhat layered composition for SNe~Ia, including in particular an IPE core, seems almost certain. Such compositions are the result of partially successful near-Chandrasekhar-mass, one-dimensional, hydrodynamic explosion models \\citep[e.g.,][]{nomoto1984,thielemann1986,woosley1991,woosley1994,khokhlov1993,hoeflich1996,hoeflich1998} and are indispensable in modeling SNe~Ia with synthetic spectra. A reference SN~Ia explosion model of longstanding exhibiting a layered composition is the one-dimensional deflagration model~W7 \\citep{nomoto1984,thielemann1986}. In many respects model~W7 seems to approximate normal SN~Ia behavior to some degree \\citep[e.g.,][]{baron2006}\\footnote{Plots of model~W7's density profile and composition, and data files that fully specify model~W7 are available at http://www.nhn.ou.edu/{\\tt\\~{}}jeffery/astro/sne/spectra/model/w7/w7.html\\ .}. Model~W7 does, in fact, have an outer layer of unburnt CO (with a trace of Ne) of $0.098\\,M_{\\odot}$ \\citep{thielemann1986}.   % marion2006 reference killed. The unburnt mass fraction is $0.071$. One-dimensional Delayed-detonation models, which have also been partially successful SN~Ia models, typically have much less unburnt mass. For example, the near-Chandrasekhar-mass delayed-detonation models of \\citet{khokhlov1993} have an unburnt mass fraction of order $0.01$ and those of \\citet{hoeflich1998} have an unburnt mass fraction of order $0.001$. Based on this evidence, it is probable that normal SNe~Ia have unburnt mass fraction $\\lesssim 0.07$.  A vital part of the standard SN~Ia model is the synthesis of the radioactive $^{56}$Ni in the explosion. The $^{56}$Ni is the overwhelmingly dominant source of radiation energy for the observable phase of SNe~Ia. The decay chain for  $^{56}$Ni is $^{56}$$\\rm Ni\\to^{56}$$\\rm Co\\to^{56}$$\\rm Fe$ with \\hbox{half-lives} of $6.077\\pm0.012$~days and~$77.27\\pm0.03$~days for the first and second decays, respectively \\citep[e.g.,][]{firestone}.  H2006's analysis of the high luminosity of SN~2003fg led to the conclusion that SN~2003fg had a very high mass of $^{56}$Ni and a (total) mass of $\\sim 2\\,M_{\\odot}$, and thus that SN~2003fg was super-Chandrasekhar (i.e., SNe~Ia with mass greater than the Chandrasekhar mass). H2006 posit that the progenitor WD had high angular momentum, and thus a high centrifugal force. The high centrifugal allowed the progenitor WD to reach a super-Chandrasekhar mass before explosive conditions were reached in the central regions. Both SD and DD scenarios may allow super-Chandrasekhar SNe~Ia. If the WD is spun up by accretion to very fast differential rotation (with mean angular velocity of order a few radians per second on average), then the WD may exceed the physical Chandrasekhar mass by up to some tenths of a solar mass before reaching explosive conditions in the central region \\citep[hereafter YL]{yoon2005}. The simple post-explosion super-Chandrasekhar SN~Ia models discussed in this paper (\\S\\S~\\ref{section-mass} and~\\ref{section-SSC}) are independent of whether the pre-explosion evolution followed the SD scenario or DD scenario, and so we will not discuss their distinctions further. Before SN~2003fg, there was no strong evidence for super-Chandrasekhar SNe~Ia.  In \\S~\\ref{section-mass}, we discuss the SN~2003fg $^{56}$Ni mass and (total) mass and H2006's analysis. Section~\\ref{section-IPE-core-boundary} discusses the SN~Ia IPE-core boundary and IPE-core velocity which are relevant to our analysis of SN~2003fg. Section~\\ref{section-SSC} introduces our SN~Ia model (the SSC~model) for the analysis of the SN~2003fg $^{56}$Ni mass and mass. In \\S~\\ref{section-comparison-sample}, we compare the predictions of the SSC~model to observations for a sample of low-$z$ SNe~Ia from \\citet{benetti2005} and in \\S~\\ref{section-comparison-sn2003fg}, to the observations for SN~2003fg. Conclusions and discussion are given in \\S~\\ref{section-conclusions}. Appendix~\\ref{ap-chan-mass} discusses the meaning and values of the Chandrasekhar mass. Appendix~\\ref{ap-YL-formula} presents the rotating CO WD binding energy formula of YL which H2006 used and which we use in the SSC~model and some related results for this binding energy formula. Appendix~\\ref{ap-SSC-mass} shows how to solve the SSC model for SN~Ia mass when SN~Ia mass is a dependent parameter. Appendix~\\ref{ap-SSC-IPE-core-mass} shows how to solve the SSC model for IPE-core mass when IPE-core mass is a dependent parameter. ",
        "conclusions": "}  Our analysis of the data for SN~2003fg supports the conclusion of H2006 that this SN~Ia is super-Chandrasekhar (\\S~\\ref{section-comparison-sn2003fg}). We find it very improbable, although not absolutely ruled out, that SN~2003fg has only about a Chandrasekhar mass (the physical Chandrasekhar mass being $1.38\\,M_{\\odot}$ Appendix~\\ref{ap-chan-mass}) and probable that its mass is $\\sim 2\\,M_{\\odot}$ as H2006 also concluded. Our conclusion relies on the adequacy of YL's binding energy formula for rotating super-Chandrasekhar-mass WDs (Appendix~\\ref{ap-chan-mass}), the analysis of \\citet{langer2000} that limits WD mass in possible accretion histories to $\\lesssim  2\\,M_{\\odot}$,   % YL-973,978. and our SSC~model of SN~Ia ejecta (\\S~\\ref{section-SSC}). It is likely that improvements to YL's binding energy formula will not change our SSC~model results qualitatively, unless rotating WDs with masses significantly greater than the Chandrasekhar mass are ruled out by improvements in the input physics of YL's theory of differential rotation in mass-accreting WDs.  An ultimate verification that SN~2003fg is super-Chandrasekhar would require that only realistic explosion models of super-Chandrasekhar-mass WDs allow a radiative transfer history (calculated with detailed radiative transfer) consistent with observations. Since we do not yet have completely adequate explosion models for any SN~Ia, it will probably be some time before this ultimate verification is possible.  What is the significance of SN~2003fg? Given that as of 2006 September~28 there are 1347 SNe~Ia or possible SNe~Ia that have been observed according to the \\citet{cbat2006} (with, of course, a wide range of observational coverage from almost nothing to extensive) and SN~2003fg is the first SN~Ia discovered for which there is significant evidence of a significant super-Chandrasekhar mass, we conclude that SN~2003fg-like SNe~Ia must be rather rare. Thus, for cosmological evolution and for cosmological distance determinations, SN~2003fg-like SNe~Ia are probably of small direct importance.  But for the study of SNe~Ia,  SN~2003fg may be quite important. As YL point out, it is somewhat surprising that super-Chandrasekhar SNe~Ia have not been much considered given that WDs rotate, and thus can be super-Chandrasekhar especially if their rotation is accelerated by accretion.   % p. 968. Of course, the lack of observational need to consider super-Chandrasekhar SNe~Ia has been a main factor in the relative neglect of such events.  % super-Chandrasekhar SNe~Ia. Also the relative homogeneity of SNe~Ia pointed to a relative of homogeneity progenitors which pointed to maybe only near-Chandrasekhar mass SNe~Ia. However, it could be (as YL and \\citet{tornambe2005} have discussed at greater length) that     % YL-983. there is a distribution of SN~Ia progenitor masses from slightly below the Chandrasekhar mass to significantly above it and that the distribution of masses may explain part of the dispersion of SN~Ia behavior. This hypothetical distribution is probably fast declining with mass coordinate in the vicinity of the probable mass of SN~2003fg given that SN~2003fg-like SNe~Ia seem so rare. In fact, it is possible that having a small super-Chandrasekhar mass is an important ingredient in normal SNe~Ia explosions and radiative transfer histories. This is a speculation, but one that is hard to ignore in light of SN~2003fg."
    },
    "0609/nucl-th0609040_arXiv.txt": {
        "abstract": "We examine how the properties of inhomogeneous nuclear matter at subnuclear densities depend on the density dependence of the symmetry energy. Using a macroscopic nuclear model we calculate the size and shape of nuclei in neutron star matter at zero temperature in a way dependent on the density dependence of the symmetry energy.  We find that for smaller symmetry energy at subnuclear densities, corresponding to larger density symmetry coefficient $L$, the charge number of nuclei is smaller, and the critical density at which matter with nuclei or bubbles becomes uniform is lower. The decrease in the charge number is associated with the dependence of the surface tension on the nuclear density and the density of a sea of neutrons, while the decrease in the critical density can be generally understood in terms of proton clustering instability in uniform matter. ",
        "introduction": "The inner crust of a neutron star consists of a Coulomb lattice of nuclei embedded in a roughly uniform neutralizing background of electrons and in a sea of neutrons \\cite{PR}.  The equation of state (EOS) of uniform nuclear matter and the surface tension, both at large neutron excess, are essential to understanding of the equilibrium properties of matter in the crust.  However, laboratory data on nuclei only reflect the bulk and surface properties of nearly symmetric nuclear matter \\cite{OI,IO}.  So far, calculations of the equilibrium properties of matter in the crust depend on the way to extrapolate the known bulk and surface properties to large neutron excess, which is different among earlier investigations \\cite{PR}.  In this paper we systematically analyze the question of how the equilibrium properties of inhomogeneous nuclear matter at subnuclear densities depend on the parameter characterizing the density dependence of the symmetry energy.  In doing so, we utilize a macroscopic nuclear model \\cite{OI} in which the equilibrium nucleon distribution depends on the EOS of nuclear matter through minimization of the energy density functional.  One of the most important quantities is the charge number of the equilibrium nuclide.  In a liquid-drop picture \\cite{PR}, this charge number is determined by the size equilibrium condition which controls the ratio between the Coulomb and surface energies.  This condition tells that the charge number squared is proportional to the surface tension and the nuclear volume.  As we shall see, the density dependence of the symmetry energy, which controls the surface tension by affecting the nuclear density and the density of the neutron gas, controls the charge number as well.   We also address the question of how matter with nuclei or bubbles melts into uniform matter with increasing density.  In this melting process, rodlike and slablike nuclei embedded in a gas of neutrons, as well as rodlike and roughly spherical neutron-gas regions (bubbles) surrounded by a nucleon liquid, are expected to occur \\cite{RPW,HSY,LRP,O}.  At a density where roughly spherical nuclei are so closely packed that they occupy about 1/8 of the system volume, the nuclei tend to be elongated and eventually fuse into nuclear rods.  The advantage of this rod formation is a reduction in the total surface area from the roughly spherical case.  However, whether bubbles and nonspherical nuclei actually appear in neutron star crusts depends on the critical density at which proton clustering instability occurs in uniform nuclear matter \\cite{PRL}; they are expected to appear when the density corresponding to the nuclear volume fraction of about $1/8$ is smaller than the critical density for proton clustering.  We find that this critical density is in turn controlled by the symmetry energy at subnuclear densities.  Earlier investigations on such exotic nuclei are more or less based on specific nuclear models \\cite{PR,WS}.  An exception is the work by Watanabe et al.\\ \\cite{WIS} which is systematic in the sense that the liquid-drop model calculations were performed in a way dependent on the proton chemical potential in pure neutron matter, $\\mu_p^{(0)}$, and the surface tension.  It was found that the density at which the system dissolves into uniform matter increases with increasing $\\mu_p^{(0)}$.  However, it remains to be clarified why some nuclear models \\cite{LRP,Chen,SLy} predict the absence of bubbles and nonspherical nuclei.  It is important to note that these models predict relatively high pressure for pure neutron matter (or, equivalently, relatively small symmetry energy) at densities around half the normal nuclear density, while the work by Watanabe et al.\\ \\cite{WIS} uses a parametrization \\cite{BBP} based on the microscopic calculations by Siemens and Pandharipande \\cite{SP} as the EOS of pure neutron matter and fix the density dependence of the symmetry energy. This parametrization is consistent with the recent Green's function Monte Carlo (GFMC) calculations \\cite{CMPR} at neutron densities up to about half the normal nuclear density. We will give a unified picture about the presence of bubbles and nonspherical nuclei by describing the pressure of pure neutron matter in terms of the density dependence of the symmetry energy.    The size and shape of nuclei in the crust bear relevance to the thermal and rotational evolution of neutron stars.  This is because thermal conductivity and neutrino emissivity in the crust are controlled by electron-nucleus scattering \\cite{YP}, while the motion of superfluid neutron vortices in the crust is affected by vortex-nucleus interactions \\cite{EB92}. However, real crustal matter has to be accompanied by defects and impurities, which can play a more important role in the star's evolution \\cite{Jones}.  In a real neutron star, furthermore, the nuclear system is more or less out of equilibrium in the course of mass accretion onto the surface of the star and the star's spin-down \\cite{Sato,HZ,IS}.  Such disordered and nonequilibrium properties are beyond the scope of this paper.   In Sec.\\ \\ref{sec:model}, we construct a model for inhomogeneous nuclear matter at subnuclear densities in a way dependent on the EOS of nearly symmetric nuclear matter near the saturation point.  The equilibrium size and shape of nuclei at given density are then calculated from the model constructed in Sec.\\ \\ref{sec:size}.  Section \\ref{sec:clustering} is devoted to evaluations of the critical density at which uniform matter becomes unstable against proton clustering.  Our conclusions are given in Sec.\\ \\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  We have analyzed the equilibrium properties of inhomogeneous nuclear matter at subnuclear densities in a way dependent on the density symmetry coefficient $L$ by using a macroscopic nuclear model.  We have estimated the upper and lower ends of the density region of the pasta phases from the onset densities of proton clustering in uniform nuclear matter and fission-like instability of spherical nuclei, respectively.  We find that the upper end decreases with $L$, while the lower end is almost flat at 0.05--0.07 fm$^{-3}$.  The former arises from the $L$ dependence of the symmetry energy, while the latter can be understood from the volume fraction $u\\simeq1/8$ at which spherical nuclei become susceptible to fission-inducing deformations.  For a typical EOS model consistent with the GFMC calculations of pure neutron matter, the calculated pasta regime is appreciable.  In fact, the pasta regime is predicted to appear when $L\\lesssim100$ MeV.   The present analysis is the first to attempt a systematic analysis of the pasta region in terms of $L$.  However, much care needs to be taken of the interpretation of the results.  While $L$ is the parameter characterizing the expansion of $w$ with respect to $n$ and $\\alpha$ around $n=n_0$ and $\\alpha=0$, the system of interest here is at large neutron excess and at subnuclear densities.  The relation between the parameter $L$ and neutron star matter depends on the parametrization of $w$ with respect to $n$ and $\\alpha$.  It is thus useful to keep in mind that we confined ourselves to expressions (\\ref{eos1})--(\\ref{vn}) although they are known to be capable of reproducing various existing microscopic calculations of the EOS of uniform nuclear matter.   We have also calculated the charge number $Z$ of spherical nuclei as a function of density for various values of $L$.  Generally, the charge number $Z$ becomes smaller for larger $L$, a feature that could be of relevance to the evolution of neutron stars \\cite{YP}.  In order to make better estimate of $Z$, however, shell and pairing effects should be taken into account."
    },
    "0609/astro-ph0609097_arXiv.txt": {
        "abstract": "{} { XZ~Ceti is the only known anomalous Cepheid in the Galactic field. Being the nearest and brightest such variable star, a detailed study of XZ~Ceti may shed light on the behaviour of anomalous Cepheids whose representatives have been mostly detected in external galaxies. } { CCD photometric and radial velocity observations have been obtained. The actual period and amplitude of pulsation were determined by Fourier analysis. The long time scale behaviour of the pulsation period was studied by the method of the $O-C$ diagram using the archival Harvard photographic plates and published photometric data. } { XZ~Ceti differs from the ordinary classical Cepheids in several respects. Its most peculiar feature is cycle-to-cycle variability of the light curve. The radial velocity phase curve is not stable either. The pulsation period is subjected to strong changes on various time scales including a very short one. The ratio of amplitudes determined from the photometric and radial velocity observations indicates that this Cepheid performs an overtone pulsation, in accord with the other known anomalous Cepheid in our Galaxy, BL~Boo (V19 in the globular cluster NGC~5466). } { Continued observations are necessary to study the deviations from regularity, to determine their time scale, as well as to confirm binarity of XZ~Ceti and to study its role in the observed peculiar behaviour.} {} ",
        "introduction": "\\label{Sect_1}  Variable stars classified as anomalous Cepheids are neglected objects from an observational point of view. Although the first representatives of such variables have been known for more than half a century, the available scanty data are insufficient for the proper explanation of their behaviour. The distinctive features of these variable stars are the short pulsational period (in the period range of RRab stars) and -- strangely enough -- up to 2 magnitudes higher luminosity than that of the RR~Lyrae type stars of the corresponding period. The higher luminosity  relates them to Cepheids, and this is the origin of the term ``anomalous'' Cepheid coined for such variables by Zinn \\& Searle (\\cite{ZS76}). Its first representatives were found in the Sculptor dwarf galaxy (Thackeray \\cite{T50}) and somewhat later in the Ursa Major (van Agt \\cite{vA67}) and Draco dwarf galaxies (Baade \\& Swope \\cite{BSw61}). From the 1970es dozens of anomalous Cepheids have been discovered in various dwarf galaxies of the Local Group. The list of known anomalous Cepheids has been compiled by Nemec et al. (\\cite{NNL94}) and more recently by Pritzl et al. (\\cite{PAJD02}).  Importance of anomalous Cepheids is twofold. On the one hand, comparison with classical Cepheids may shed light on the differences in the internal structure and evolutionary phase of the two types of pulsators. On the other hand, anomalous Cepheids follow a distinct period-luminosity relationship (see Pritzl et al. \\cite{PAJD02} for its most recent form), therefore they serve as standard candles in distance determination for extragalactic systems with no recent star formation.  Existence of pulsating stars in such unusual region in the H-R diagram has been explained in two different ways: (i) by mass transfer (and possibly coalescence) in a close binary system producing a more massive and luminous star than other pulsators on and above the horizontal branch (Zinn \\& Searle \\cite{ZS76}; Renzini et al. \\cite{RMS}) and (ii) assuming that anomalous Cepheids are not old stars but their peculiarity is the extremely low-metallicity (Norris \\& Zinn \\cite{NZ75}; Demarque \\& Hirshfeld \\cite{DH75}). The first reliable model for supporting this latter explanation was calculated by Bono et al. (\\cite{BCSCP}) with confronting evolutionary tracks and pulsational calculations. This nonlinear convective pulsational model was updated by Marconi et al. (\\cite{MFC04}), and these new calculations resulted in a good agreement with the observed behaviour (period, amplitude, colour, luminosity) of the anomalous Cepheids in dwarf spheroidal galaxies. Two recent studies, an observational (Dolphin et al. \\cite{Detal02}) and a theoretical one (Caputo et al. \\cite{CCDFM}) explicitly declare that anomalous Cepheids occur in the region which is the extension of the normal Cepheid sequence toward higher temperature for extreme low metallicity ($Z = 0.0002$ to 0.008). Spectroscopic confirmation of this explanation is, however, very difficult, because the extragalactic anomalous Cepheids are very faint stars for a detailed spectroscopic study.  However, anomalous Cepheids occur in our own Galaxy, in metal-poor environment. Zinn \\& Dahn (\\cite{ZD76}) identified the first such variable: the star V19 in the globular cluster NGC~5466. Curiously, a conventional variable star name, BL~Bootis, has been assigned to this star, though variable stars belonging to globular clusters are catalogued separately. Due to its membership in a globular cluster, luminosity (and luminosity difference with respect to the RR~Lyrae type variables in the same system) of BL~Boo could be reliably determined, and it is this star which advanced to be the prototype of anomalous Cepheids. The BLBOO type variable stars are thus identical with the anomalous Cepheids, and this designation was introduced in the fifth volume of the General Catalog of Variable Stars (Samus \\cite{S95}) listing the extragalactic variables. The first detailed spectroscopic analysis ever made on an anomalous Cepheid was that of BL~Boo (McCarthy \\& Nemec \\cite{McN97}), a 15th magnitude star.  Various pieces of evidence suggest that a much brighter star in the galactic field, XZ~Ceti, also belongs to anomalous Cepheids. In this paper the behaviour of XZ~Ceti is discussed based on new photometric and radial velocity data. Following the overview of the available information on XZ~Ceti (Sect.~\\ref{Sect_2}) our observational data are described (Sect.~\\ref{Sect_3}). Then we discuss the value of the pulsation period and its changes, issues related to the pulsation amplitudes of XZ~Ceti and its possible binarity (Sect.~\\ref{Sect_4}). Our results are briefly summarized in Sect.~\\ref{Sect_5}. ",
        "conclusions": "\\label{Sect_5}  We studied photometric and radial velocity variability of XZ~Ceti, a star classified as an anomalous Cepheid. The term ``anomalous'' is obsolete in its original sense, because the recent pulsation models elaborated by Bono et al. (\\cite{BCSCP}), Marconi et al. (\\cite{MFC04}), and Caputo et al. (\\cite{CCDFM}) give a natural explanation why such variables exist in the given region of the H-R diagram: they are extremely metal-poor classical Cepheids. This new paradigm gives motives for comparing behaviour of XZ~Ceti with classical Cepheids.  Resemblance of the phenomenological parameters of the light and radial velocity curves (period, Fourier parameters, amplitude ratios) of XZ~Ceti and the prototype anomalous Cepheid, BL~Bootis corroborates that XZ~Cet is also an anomalous Cepheid. The large value of the mean radial velocity (about 175~km\\,s$^{-1}$) is a further piece of evidence that XZ~Ceti does not belong to the same subsystem that includes classical Cepheids.  The systematic shift in the radial velocities at given phases of pulsation hints at possible binarity of XZ~Ceti. Presence of a companion is also inferred from the large radial velocity amplitude as compared with the photometric amplitude, even if greatness of this amplitude ratio is partly caused by the overtone mode of pulsation.  XZ~Ceti remains anomalous in the sense that it shows strong period changes on a very short time scale, and possibly slightly unstable pulsation with cycle-to-cycle variations. Summarizing its longer time scale period behaviour, the characteristic periods have been as follows:\\\\ between JD\\,2\\,413\\,000 and JD\\,2\\,434\\,000: 0\\fd8231057;\\\\ between JD\\,2\\,442\\,000 and JD\\,2\\,453\\,400: 0\\fd8231298;\\\\ however, for the interval JD\\,2\\,452\\,000--JD\\,2\\,453\\,400 the most appropriate value is 0\\fd8231561.  Further monitoring of this unique variable star is highly desirable."
    },
    "0609/astro-ph0609742_arXiv.txt": {
        "abstract": "{ {\\it Context.} The nature of the progenitors of type Ia supernovae is still under controversial debate. KPD~1930+2752 is one of the best SN Ia progenitor candidates known today. The object is a double degenerate system consisting of a subluminous B star (sdB) and a massive white dwarf (WD). Maxted et al. (\\cite{maxted}) conclude that the system mass exceeds the Chandrasekhar mass. This conclusion, however, rests on the assumption that the sdB mass is $0.5\\,M_{\\odot}$. However, recent binary population synthesis calculations suggest that the mass of an sdB star may range from $0.3\\,M_{\\odot}$ to more than $0.7\\,M_{\\odot}$.\\\\ {\\it Aims.} It is therefore important to measure the mass of the sdB star simultaneously with that of the white dwarf. Since the rotation of the sdB star is tidally locked to the orbit the inclination of the system can be constrained if the sdB radius and the projected rotational velocity can be measured with high precision. An analysis of the ellipsoidal variations in the light curve allows to tighten the constraints derived from spectroscopy.\\\\ {\\it Methods.} We derive the mass-radius relation for the sdB star from a quantitative spectral analysis of 150 low-resolution spectra obtained with the Calar Alto 2.2 m telescope using metal-rich, line blanketed LTE model atmospheres with and without NLTE line formation. The projected rotational velocity is determined for the first time from 200 high-resolution spectra obtained with the Keck I 10 m, and the ESO-VLT 8.2 m telescopes. In addition a reanalysis of the published light curve is performed.\\\\ {\\it Results.} The atmospheric and orbital parameters are measured with unprecedented accuracy. In particular the projected rotational velocity \\(v_{\\rm rot}\\sin{i} = 92.3 \\pm 1.5~{\\rm km\\,s^{-1}}\\) is determined. Assuming the companion to be a white dwarf, the mass of the sdB is limited between \\(0.45~M_{\\odot}\\) and \\(0.64~M_{\\odot}\\) and the corresponding total mass of the system ranges from \\(1.33~M_{\\odot}\\) to \\(2.04~M_{\\odot}\\). This constrains the inclination to $i>68\\degr$. The photometric analysis allows to constrain the parameters even more. A neutron star companion can be ruled out and the mass of the sdB is limited between \\(0.45~M_{\\odot}\\) and \\(0.52~M_{\\odot}\\). The total mass of the system ranges from \\(1.36~M_{\\odot}\\) to \\(1.48~M_{\\odot}\\) and hence is likely to exceed the Chandrasekhar mass. The inclination angle is $80\\degr$ and the light curve shows weak and shallow signs of eclipses. A high precision light curve is needed in order to accurately measure these eclipses. So KPD~1930+2752 qualifies as an excellent double degenerate supernova Ia progenitor candidate. ",
        "introduction": "Supernovae of type Ia (SN~Ia) play an important role in the study of cosmic evolution. They are regarded as the best standard candles for the determination of the cosmological parameters $H_{\\rm 0}$, $\\Omega$ and $\\Lambda$ (e.g. Riess et al. \\cite{riess}; Leibundgut \\cite{leibundgut}; Perlmutter et al. \\cite{perlmutter}) together with the WMAP probe for anisotropies in the cosmic background radiation (e. g. Bennett et al. \\cite{bennett}). Although some of the most important discoveries in modern cosmology, in particular the accelerated expansion of the universe, were initially derived from distance measurements of SN~Ia with high redshifts, the nature of their progenitors is still under debate (Livio \\cite{livio}). The progenitor population is a crucial information required to back the assumption that distant SN~Ia can be used as standard candles as it was done with the ones in the local universe. There is general consensus that only the thermonuclear explosion of a white dwarf (WD) is compatible with the observed features of SN~Ia. For this a white dwarf has to accrete mass from a close companion to reach the Chandrasekhar limit of $1.4 \\,M_{\\rm \\odot}$ (Hamada \\& Salpeter \\cite{hamada}). Two main scenarios of mass transfer are currently under discussion. In the so-called single degenerate (SD) scenario (Whelan \\& Iben \\cite{whelan}), the mass donating component is a red giant/subgiant, which fills its Roche lobe and is continually transfering mass to the white dwarf. According to the so-called double degenerate (DD) scenario (Iben \\& Tutukov \\cite{iben}) the mass donating companion is a white dwarf, which eventually merges with the primary due to orbital shrinkage caused by gravitational wave radiation.  A progenitor candidate for the DD scenario must have a total mass near or above the Chandrasekhar limit and has to merge in less than a Hubble time. Systematic radial velocity (RV) searches for DDs have been undertaken (e.g. Napiwotzki \\cite{napiwotzki2} and references therein). The largest of these projects was the ESO SN Ia Progenitor Survey (SPY). More than $1\\,000$ WDs were checked for RV-variations (e. g. Napiwotzki et al. \\cite{napiwotzki2}). SPY detected $\\sim 100$ new DDs (only $18$ were known before). One of them may fulfill the criteria for progenitor candidates, even though the error margins are quite high (Napiwotzki et al. \\cite{napiwotzki3}; Karl \\cite{karl}).  KPD~1930+2752 was identified as a subdwarf B (sdB) star in the Kitt Peak -- Downes survey (Downes \\cite{downes}). The parameters which were derived from spectroscopy by model atmosphere fits (Saffer \\& Liebert \\cite{saffer}) are consistent with the theoretical instability strip predicted by Charpinet et al. (\\cite{charpinet}). After the first pulsating sdB stars with short periods (EC 14026 stars) were discovered (Kilkenny et al. \\cite{kilkenny}; Koen et al. \\cite{koen}; Stobie et al. \\cite{stobie}), Bill\\`{e}res et al. (\\cite{billeres}) initiated a survey to search for such objects in the northern hemisphere. KPD~1930+2752 was selected from the list of Saffer \\& Liebert (\\cite{saffer}) for their fast photometry program. They detected multiperiodic variations with short periods and low amplitudes. In addition to 44 $p$-mode pulsations they found a strong variation at a much longer period of about \\(4100 \\, {\\rm s}\\). This variation could be identified as an ellipsoidal deformation of the sdB star most likely caused by a massive companion. Bill\\`{e}res et al. (2000) predicted the period of the binary to be two times the period of this brightness variation (\\(P=8217.8 \\, {\\rm s} = 0.095111 \\, {\\rm d}\\)).  This was proven by Maxted et al. (\\cite{maxted}), who measured a radial velocity curve of KPD~1930+2752 which matched the proper period. The radial velocity amplitude \\(K=349.3 \\pm 2.7 \\, {\\rm km\\,s^{-1}} \\) combined with an assumption of the canonical mass for sdB stars \\(M_{ \\rm sdB}=0.5 \\, M_{\\odot}\\) led to a lower limit for the mass of the system derived from the mass function. This lower limit \\(M \\geq 1.47\\, M_{\\odot} \\) exceeded the Chandrasekhar mass of \\(1.4 \\, M_{\\odot}\\) (Hamada \\& Salpeter \\cite{hamada}). Because there was no sign of a companion in the spectra, it was concluded that the unseen object must be a white dwarf. The system is considered to become double degenerate when the subdwarf eventually evolves to a white dwarf. Orbital shrinkage caused by gravitational wave radiation will lead to a merger of the binary in about $200$ Myr which is much less than a Hubble time. Combining all the evidence, Maxted et al. concluded that KPD~1930+2752 could be a good candidate for the progenitor of a Type Ia supernova.  From the theoretical point of view Ergma et al. (\\cite{ergma}) questioned the double degenerate scenario in the case of KPD~1930+2752. Their simulations based on the derived parameters of Maxted et al. (\\cite{maxted}) suggested that the merger will lead to a single massive ONeMg white dwarf rather than a supernova explosion. The merger of two white dwarfs is a rather complicated process compared to slow accretion of material onto a heavy white dwarf. Detailed merger simulations in 3D are not yet avaible, but urgently needed to learn more about this possible SN~Ia progenitor channel.  The subluminous B stars form the hot end of the horizontal branch, the so called Extreme Horizontal Branch (EHB, Heber \\cite{heber1}). The EHB models imply that they are core helium burning objects with tiny hydrogen dominated envelopes, too thin to sustain hydrogen burning. The core mass is fixed by the onset of the core helium flash at the tip of the red giant branch and depends only slightly on metallicity and helium abundance. Hence the canonical core mass is resticted to a very narrow range of $0.46$ to $0.5\\,M_{\\odot}$. Unlike normal horizontal branch stars, EHB stars do not ascend the asymptotic giant branch (AGB) after core helium burning has terminated, but evolve to higher temperatures until the white dwarf cooling sequence is reached.  Binary evolution scenarios have been proposed (e.g. Han et al., \\cite{han1}, \\cite{han2}) to explain the large fraction of sdB stars in binaries. These scenarios involve strong mass-transfer either through a common-envelope ejection, or through stable Roche lobe overflow. A merger of two helium core white dwarfs is another viable option.  The binary population synthesis models of Han et al. (\\cite{han1}) predict a mass distribution for sdB stars that peaks near the canonical mass, but is much wider, ranging from $0.3\\,M_{\\odot}$ to more than $0.7\\,M_{\\odot}$. The lower limit is based on the fact that no core helium burning can be sustained for stars with masses less than $0.3\\,M_{\\odot}$. The highest masses result from mergers. Hence it might be premature to assume that the mass of an sdB star is half a solar mass, as has been done in previous studies. Therefore we drop this assumption and treat the mass of the sdB component as a free parameter.  Another drawback of all previous investigations was the lack of information on the inclination angle. We aim at deriving constraints on the inclination for the first time by means of accurate measurements of the projected rotational velocity and the surface gravity from appropriate high-quality spectra using sophisticated model atmospheres. Because the rotation of the sdB star is tidally locked to the orbit, we can derive \\(\\sin{i}\\) as a function of the sdB mass from these two quantities. This approach combined with a reanalysis of the published light curve allows to constrain the masses of both components and, therefore, to check whether or not KPD~1930+2752 qualifies as a SN~Ia progenitor candidate through a double degenerate merger.  Section \\ref{sec:obs} describes the newly obtained spectra, from which an improved radial velocity curve has been derived (Sect.~\\ref{sec:rvc}). For the quantitative spectral analysis, existing grids of metal line blanketed LTE models have been used and a new (hybrid) approach is applied that allows to treat both departures from LTE as well as metal line blanketing (Sect.~\\ref{sec:atm}). Combining these results, described in Sect.~\\ref{sec:par}, with the projected rotational velocity (Sect.~\\ref{sec:rot}) we constrain masses and inclination of the KPD~1930+2752 system (Sect. \\ref{sec:mi}). Within the derived parameter space synthetic light curves were modelled and fitted to the published light curve (Sect.~\\ref{sec:phot}) in order to constrain the parameters even further. Our conclusions are given in the final section.  \\begin{figure}[t!] \\centering \\resizebox{\\hsize}{!}{\\includegraphics{KPDRVHIRES.ps}} \\caption{Radial velocity curve. The data points are from the new HIRES spectra, the solid line is the fitted sine curve to these points. The residuals in the lower part show no signs of eccentricity.} \\label{RV1} \\end{figure} ",
        "conclusions": "}  The combination of an extensive set of 350 high- and low-resolution spectra, 2\\,700 radial velocity data points from prior observations and a reanalysis of the published light curve made it possible to determine the parameters of KPD~1930+2752 with unprecedented accuracy. The companion is a heavy white dwarf. KPD~1930+2752 is the first subdwarf whose mass could be constrained in this way. With $0.45 - 0.52\\,M_{\\odot}$ it is remarkebly close to the canonical mass of $0.5\\,M_{\\odot}$. It is worthwhile to note that masses of other sdBV derived by asteroseismology are also close to the canonical value (Brassard et al. \\cite{brassard}; Charpinet et al. \\cite{charpinet2}a,b). Furthermore Orosz \\& Wade (\\cite{orosz}) determine a sdB mass of $0.51\\pm0.05\\,M_{\\odot}$ for the eclipsing sdB+WD binary KPD~0422+5421.\\\\ The total mass and the merging time of the binary indicate that it is an excellent double degenerate candidate for an SN~Ia progenitor. KPD~1930+2752 is one of only two double degenerate systems, which fulfill these requirements (see Fig. \\ref{Progen}; Napiwotzki et al. \\cite{napiwotzki3}). While it fits into the DD scenario, KPD~1930+2752 might also evolve into an SN~Ia as a single degenerate binary system. Since the time for merging due to gravitational wave radiation $t_{\\rm merger} \\approx 2 \\cdot 10^{8}\\,{\\rm yr}$ is of the same order as the EHB life time $t_{\\rm EHB} \\approx 2 \\cdot 10^{8}\\,{\\rm yr}$, Roche lobe overflow could occur well before the sdB becomes a white dwarf due to the shrinkage of the orbit. Recent calculations by Han and Podsiadlowski (priv. comm.) indicate that the mass transfer would be stable.  Follow-up observations should be undertaken to measure an improved light curve and search for clear signs of eclipses. A promising alternative option is provided for asteroseismology. KPD~1930+2752 is the first pulsating subdwarf with a mass estimate from orbital analysis. It could be used to calibrate asteroseismological models (see e.g. Charpinet et al. \\cite{charpinet2}). This is a very difficult task because the frequency spectrum is very complex and modelling is rendered difficult by the ellipsoidal deformation and the rapid rotation of KPD~1930+2752."
    },
    "0609/astro-ph0609268_arXiv.txt": {
        "abstract": "We present optical (\\ubvri\\hspace{-1.0mm}) and near-infrared (\\yjhk\\hspace{-1.0mm})  photometry of the normal Type Ia supernova 2004S.  We also present eight optical spectra and one near-IR spectrum of SN 2004S.  The light curves and spectra are nearly identical to those of SN~2001el.  This is the first time we have seen optical {\\em and} IR light curves of two Type Ia supernovae match so closely. Within the one parameter family of light curves for normal Type Ia supernovae, that two objects should have such similar light curves implies that they had identical intrinsic colors and produced similar amounts of $^{56}$Ni. From the similarities of the light curve shapes we obtain a set of extinctions as a function of wavelength which allows a simultaneous solution for the distance modulus difference of the two objects, the difference of the host galaxy extinctions, and R$_V$. Since SN~2001el had roughly an order of magnitude more host galaxy extinction than SN~2004S, the value of R$_V$ = 2.15$^{+0.24}_{-0.22}$ pertains primarily to dust in the host galaxy of SN~2001el.  We have also shown via Monte Carlo simulations that adding rest frame $J$-band photometry to the complement of \\bvri photometry of Type Ia SNe decreases the uncertainty in the distance modulus by a factor of 2.7.  A combination of rest frame optical and near-IR photometry clearly gives more accurate distances than using rest frame optical photometry alone. ",
        "introduction": "A Type Ia supernova (SN, plural SNe) is widely believed to be a carbon-oxygen white dwarf which explodes when its mass reaches the Chandrasekhar limit owing to mass transfer from a nearby companion \\citep[][and references therein]{Liv00}. The explosion produces several tenths of a solar mass of radioactive $^{56}$Ni, which decays to $^{56}$Co, and finally to stable iron \\citep{Str_etal06}.  Since the progenitors of Type Ia SNe have approximately the same mass (M$_{Ch} \\approx$ 1.4 M$_{\\odot}$), the explosions have approximately the same brightness at maximum light.  \\citet{Phi93} first demonstrated unambiguously that the peak brightness of a Type Ia SN is related to the decline rate of the light curve.  Since the publication of Phillips' classic paper, the common measure of the decline rate has been the number of $B$-band magnitudes that a Type Ia SN fades in the first 15 days after maximum light (designated \\dmm).  A typical value is \\dmm = 1.10 mag.  Values range from $\\sim$0.7 for SN~2001ay (Nugent et al. 2006, in preparation) to 1.93 for SN~1991bg \\citep{Fil_etal92,Lei_etal93,Phi_etal99}.  In the past 13 years supernova researchers have noted various patterns in the light curves of Type Ia SNe.  This has allowed the determination of the peak magnitudes, host galaxy extinction, and distances to these objects even if the data of a particular object are irregularly spaced in time and do not necessarily overlap the light curve maxima.  The three most widely used systems for fitting optical light curves are the \\dmm\\ system itself \\citep{Ham_etal96, Phi_etal99, Ger_etal04, Pri_etal06}, the ``stretch method'' \\citep{Per_etal97, Gol_etal01, Kno_etal03}, and the multi-color light curve shape (MLCS) method \\citep{Rie_etal96a, Rie_etal98, Jha_etal06b}.  There is also the hybrid magnitude-color system of \\citet{Wan_etal03a}.  Host galaxy extinction of Type Ia SNe is measured empirically by first determining reddening-free color loci, determining color excesses, and using empirically determined scale factors to convert color excesses to A$_V$ or extinction in some other band.  \\citet{Lir95} found, from a small sample of Type Ia SNe unreddened in their hosts, that $B-V$ colors are uniform and evolve in the same way from 30 to 90 days after the time of $V$-band maximum.  \\citet{Nob_etal03} investigated the optical colors of Type Ia SNe and found that all the colors except $V-I$ are consistent with zero intrinsic dispersion at 35 days after T($B_{max}$).  \\citet{Wan_etal05} found that the $B-V$ colors at 12 days after T($B_{max}$) provided the best method of determining the $B-V$ color excess.  With MLCS2k2 \\citet{Jha_etal06b} adopt an intrinsic color of $B-V$ = 1.054 $\\pm$ 0.049 at $t$ = 35 days and assume an exponential distribution of host galaxy reddening with scale length E($B-V$) = 0.138 mag. They also adopt a Bayesian prior on the distribution of R$_V$, with a minimum value of 1.8.  Then they marginalize over the optical photometry to obtain the best values of A$_V$ and R$_V$.  The prime problem with using optical photometry to determine A$_V$ along the line of sight to a Type Ia SN can be illustrated as follows.  For normal dust in our Galaxy R$_V$ $\\equiv$ A$_V$ / E($B-V$) $\\approx$ 3.1 \\citep{Car_etal89}.  A small uncertainty in the $B-V$ color excess translates into a much larger uncertainty in A$_V$.  The scale factor R$_V$ is not necessarily equal to 3.1.  The mean value most appropriate to galaxies that host Type Ia SNe is somewhere in the range 2.3 to 2.65 \\citep{Rie_etal96b, Alt_etal04, Rei_etal05, Wan_etal06, Jha_etal06b}. Some very low values have been measured.  The Type Ia SN 1999cl was reddened by dust with R$_V$ = 1.55 $\\pm$ 0.08 \\citep {Kri_etal06}.  The dust in the host of SN~2003cg had R$_V$ = 1.80 $\\pm$ 0.19 \\citep{Eli_etal06}.  Our observations of SN~2006X are consistent with R$_V$ $\\approx$ 1.56.\\footnote[13]{See http://www.nd.edu/~kkrisciu/sn2006X.html.} Perhaps these very low values of R$_V$ associated with highly reddened objects are the result of an interaction of the SN ejecta and the surrounding circumstellar material \\citep{Wan05}.  To account for non-standard reddening laws and to obtain the most robust value of the extinction towards a Type Ia SN, it is best to combine optical and IR photometry.  \\citet{Eli_etal85} first suggested that the $V-K$ colors of Type Ia SNe might be uniform and therefore very useful for determining extinction for these objects. Using the standard reddening law of \\citet{Car_etal89}, A$_V$ $\\approx c_R$ E($V-K$), where $c_R$ = 1.13 for Galactic (R$_V$ = 3.1) dust, and $c_R$ = 1.07 for the extreme case of R$_V$ = 1.55 for SN~1999cl.  Thus, if we can determine a $V-K$ color excess, it needs to be scaled only slightly to give us A$_V$. Using data of 8 Type Ia SNe, we determined zero-reddening $V-J$, $V-H$, and $V-K$ loci of Type Ia SNe with mid-range decline rates \\citep{Kri_etal00}. Subsequently, we determined that the corresponding loci of Type Ia SNe with slow decline rates (\\dmm\\ $\\lesssim$ 1.0) are roughly 0.24 mag bluer \\citep{Kri_etal04b}. We note that $V-J$ colors of Type Ia SNe are the least uniform of the three color indices.  Also, our previous loci are valid for the time range $-$9 $\\leq t \\leq$ +27 days, where $t$ is measured with respect to the time of $B$-band maximum.  Since the Cal\\'{a}n-Tololo SN survey in the early 1990's, extensive efforts have been devoted to discovering and following the light curves of Type Ia SNe, most notably the SN search carried out with the Katzman Automatic Imaging Telescope at Lick Observatory \\citep{Fil_etal01}. Because of the time required to observe the SNe, obtain image subtraction templates for those objects that require templates, and to reduce the data for many objects, one recent trend has been to publish papers on individual objects which are unusual, such as SNe~2000cx \\citep{Li_etal01,Can_etal03} and 2002cx \\citep{Li_etal03}.  The analogy is that pathologists can better understand how normal human organs work by studying organ malfunction and disease.  In this paper we discuss the astronomical case of SN~2004S, a normal Type Ia SN, the light curves of which are essentially clones of another object, SN~2001el \\citep{Kri_etal03}.  SN~2004S was discovered using images taken on 2004 February 3.54 and 4.56 UT by \\citet{Mar04}.  The supernova was located at RA = 06$^h$ 45$^m$ 43.\\hspace{-1.3mm}$^s$5, DEC =$-$31$^{\\rm o}$ 13$^{\\prime}$ 52$\\farcs$5 (J2000), some 47$\\farcs$2 west and 2$\\farcs$5 south of the core of the Sc galaxy MCG -05-16-021 \\citep{Big04}. SN~2004S was confirmed to be a Type~Ia supernova by \\citet{Sun_etal04} from a spectrum taken on 2004 Feb 6.1 UT with the CTIO 1.5-m telescope.  In Fig. \\ref{comp_at_max} we see the \\citet{Sun_etal04} spectrum, along with a spectrum of SN~2001el obtained by \\citet{Wan_etal03b}. Given the similarities of the optical and IR light curves and the spectra of these two objects, we feel fully justified in calling these two objects clones.  According to \\citet{The_etal05}, the heliocentric radial velocity of MCG -05-16-021 is 2808 km s$^{-1}$. In the reference frame of the Cosmic Microwave Background (CMB) radiation the recession velocity is 2957 km s$^{-1}$.  Assuming a Hubble constant of 72 km s$^{-1}$ Mpc$^{-1}$ \\citep{Fre_etal01}, the distance modulus of SN~2004S and its host is $m-M$ = 33.07 $\\pm$ 0.21 mag, where the uncertainty corresponds to an assumed random velocity of 300 km s$^{-1}$.  The Galactic reddening along the line of sight to SN~2004S is E($B-V$) = 0.101 $\\pm$ 0.010 mag \\citep{Sch_etal98}. Assuming R$_V$ = 3.1 for Galactic dust, there must be at least 0.313 $\\pm$ 0.031 mag of $V$-band extinction toward SN~2004S.  Given the offset of the SN from the core of its host (see Fig. \\ref{finder}), we expected the SN to be essentially unreddened by the host.  An analysis of SN~2004S may thus provide useful information on the intrinsic colors of normal Type Ia SNe of mid-range decline rates.  In this paper we present optical and near-IR photometry of SN~2004S. Other optical photometry of the object has been published by \\citet{Mis_etal05}. These authors obtained a light curve solution using their own measurements plus our optical data from the CTIO 1.3-m published here.  Important conclusions can be made about SN~2004S by using IR data as well.  That the optical and IR light curves of SNe 2001el and 2004S are so similar sheds light on the uniformity of explosion models of normal Type Ia SNe and also allows us to investigate the interstellar dust content of the hosts of SNe 2001el and 2004S beyond what is possible from a study of either object on its own.  The paper is organized as follows.  In \\S2 we present optical and infrared photometry of SN~2004S.  In \\S3 we present optical and IR spectra of SN~2004S and compare them to several other Type Ia SNe.  In \\S4 we consider the extinction suffered by SNe 2004S and 2001el as two separate problems and also as a joint exercise.  Using information about the relative uncertainties of the filter by filter photometry from \\S4, in \\S5 we carry out Monte Carlo simulations of extinction and reddening for a general model of the effect of the dust on the uncertainty of distance measurements of Type Ia SNe.  This is relevant for future ground-based and space-based SN surveys.  Our conclusions are presented in \\S6. ",
        "conclusions": "In this paper we provided optical and IR photometry of the normal Type Ia SN~2004S.  Using optical and near-IR data, it was found to have minimal host galaxy extinction, with A$_V$(host) = 0.07 $\\pm$ 0.04 mag.  There are remarkable similarities of the shapes of the optical and IR light curves of SNe 2004S and 2001el.  This is the first time we have seen two objects exhibit such similar $JHK$ light curves at the time of maximum light until 65 days afterwards.  We might assume that the two objects had identical explosion mechanisms, or at least produced the same amount of $^{56}$Ni. Given that Type Ia SNe exhibit an intrinsic dispersion of roughly $\\pm$0.15 mag at any given decline rate, SNe 2001el and 2004S could have produced identically shaped filter-by-filter light curves, yet could have had absolute magnitudes that differ by 0.15 mag or more.  For our purposes it is reasonable to interpret the difference of their extinction-corrected magnitudes in any band (except $U$) as a measure of the difference of their distance moduli.  More generally, we may assume that they had identical extinction-corrected colors.  SNe 2001el and 2004S not only had similar light curves, but very similar optical and near-IR spectra, exhibiting high-velocity absorption due to singly ionized Ca and Fe.  Except for the high velocity features in their ejecta prior to $t \\approx$ +14 d, they can be considered spectroscopically normal \\citep{Bra_etal06}.  The similarities of the light curves of SNe 2001el and 2004S allow us to use the photometry, filter by filter, to derive the {\\em difference} of the host galaxy extinctions without making any a priori assumptions about the unreddened colors of Type Ia SNe, values of R$_V$, or the reddening law between the $K$-band and infinite wavelength.  Since SN~2001el had roughly an order of magnitude more host galaxy extinction compared to SN~2004S, our multi-dimensional $\\chi^2$ minimization primarily tells us something about the dust in the host of SN~2001el. For this object we find R$_V$ = 2.15$^{+0.24}_{-0.22}$ and A$_V$(tot) $\\approx$ 0.59 $\\pm$ 0.05 mag.  Because of the advantages of combining optical and IR photometry to determine extinction, it is important that we know the shapes and zeropoints of the color curves of Type Ia SNe.  Our analysis gives us unreddened solutions for SNe 2001el and 2004S.  The $V-K$ color curves of SNe 2001el and 2004S have the same shape as our template based on the not-so-well sampled light curves of 8 mid-range decliners.  Our $V-J$ and $V-H$ templates do not have quite the same shapes as the color curves of SNe 2001el and 2004S.  In the near future it should be possible to construct more robust color curves from a larger database of objects, which will reflect more realistically the inherent scatter exhibited by Type Ia SNe.  We find that SN~2004S had $VJHK$ absolute magnitudes at maximum light within 1-$\\sigma$ of the mean of other Type Ia SNe of comparable decline rate. Any systematic differences may only be due to the unknown peculiar velocity of the host galaxy.  The $JHK$ absolute magnitudes at maximum of SN~2004S are consistent with the notion that in the near-IR Type Ia SNe (excepting the fastest decliners) are standard candles.  We can derive host galaxy extinction much more accurately using optical and IR photometry than using optical photometry alone, as we have shown graphically and quantitatively.  Our Monte Carlo simulations indicate that rest frame \\bvrij photometry gives distance uncertainties 2.7 times smaller than just using rest frame \\bvri photometry.  Assuming a fixed value of R$_V$ for high-redshift SNe runs the risk of systematic errors if the fixed value is far from the true mean.  Here we showed that adding one additional rest frame band in the near-IR substantially improves the resulting uncertainty of determining the distance moduli of Type Ia SNe.  Rest frame $J$-band corresponds to 3.6 $\\mu$m in the observer frame at $z \\sim$ 2, which is observable with space-based telescopes.  Having rest frame optical and near-IR data for Type Ia SNe will lead to stricter constraints on models of Dark Energy.  \\vspace {1 cm}"
    },
    "0609/astro-ph0609604_arXiv.txt": {
        "abstract": "We model the spectra and spectral energy distribution of V838 Mon which were observed in February, March, and November, 2002. Theoretical spectra are calculated using the classical model atmospheres taking into account absorption of atomic and molecular lines. By fitting the observed spectra we determine the physical parameters of the atmosphere of V838 Mon. These parameters are determined to be Teff = 5330 $\\pm$ 300 K, 5540 $\\pm$ 270 K, 4960 $\\pm$ 190 K, and 2000 $\\pm$ 200 K for February 25, March 2, March 26, and November 6, respectively. ",
        "introduction": "\\subsection{Observational data}  Some spectra obtained on February -- March, 2002 are shown in Fig. \\ref{__obs}.  The first spectrum of V838 Mon taken on Feb. 4 is shown in Fig. \\ref{__obs}. It was obtained by V.Klochkova at the 6 m telescope (SAO, Russia). Details of the data reduction can be found in Kipper et al. (2004).  On February 25 and March 26, 2002 spectra (R $\\sim$ 18000) of V838 Mon were obtained with the Echelle+CCD spectrograph on the 1.82m telescope operated by Osservatorio Astronomico di Padova on Mount Ekar (Asiago).  Spectra  on March 2 were obtained with the echelle fibre-fed spectrograph on the 1.9-m SAAO telescope kindly provided for us by Lisa Crause (see Crause et al. 2003 for details).  The November 6, 2002 spectrum was obtained with the Kast spectograph on the Cassegrain focus of the Shane 3-m telescope at Lick Observatory. Full details of the observing and data-reduction processes are given in Rushton et al. (2005) and are not repeated here.  \\begin{figure*} \\begin{center} \\includegraphics [width=65mm, angle=00]{yp1.eps} \\includegraphics [width=65mm, angle=00]{yp2.eps} \\end{center} \\caption[]{\\label{__obs} Left: a comparison of observed spectra of \\vmon\\ for different dates in Febr. -- March, 2002. Right: comparison of the spectrum of V838 Mon and that of  Arcturus (Hinkle et al. 2000), convolved with macroturbulent profile \\Vm = 50 km s$^{-1}$ See Kaminsky \\& Pavlenko (2002) for more details.} \\end{figure*}   \\subsection{Computation of theoretical spectra}  Through February to March 2002, the spectrum of V838 Mon looks like a normal spectrum of a supergiant of spectral class K -- G. The optical spectra of V838Mon over these dates are governed by absorption by the neutral atoms. A comparison with a synthetic spectrum of Acturus, with lines broadened by macroturbulent velocities of 50 km/s, shows good agreement (Fig. \\ref{__obs}).  To model these spectra we used classical model atmospheres from the Kurucz (1993) grid. Computations of the synthetic spectra were carried out using the program WITA6 Pavlenko (2000) assuming LTE and hydrostatic equilibrium for a one-dimensional model atmosphere without sources and sinks of energy.The equations of ionization-dissociation equilibrium were solved for media consisting of atoms, ions and molecules. Computations were carried out for the VALD (Kupka et al. 1999) line list.  Up until November 2002 both effective temperature and luminosity of V838 Mon drop significantly with time (see also Tylenda 2005). The spectrum of V838 Mon contains strong molecular features. To compute theoretical spectral energy distributions (SEDs) we used a number of model atmospheres, with \\Tef = 2000 --- 2200~K, log g = 0, 0.5 from the NextGen grid of Hauschildt et al. (1999).  In addition to VALD, molecular line data were taken from a variety of sources: \\\\ -- the TiO line lists of Plez (1998). \\\\ -- CN lines from CDROM 18 (Kurucz 1993); \\\\ -- CrH and FeH lines from Burrows (2002) and Dulick et al. (2003) respectively \\\\ -- lines of H$_2^{16}$O from the BT2 database (Barber et al. 2006) \\\\ -- absorption by VO, and by a few molecules of (in the case of \\vmon) lesser importance, was computed in the JOLA approximation (see Pavlenko et al. 2000).  The relative importance of the different opacities contributing to our synthetic spectra is shown in Fig. \\ref{__obs}.  \\begin{figure*} \\includegraphics [width=65mm]{yp3.eps} \\includegraphics [width=65mm]{yp4.eps} \\caption{\\label{__nov}Left: the contribution of different molecules to the formation of the spectrum of \\vmon. Right: our fit to the observed \\vmon spectrum on November 6, 2002 (Pavlenko et al. 2006).}  \\end{figure*} ",
        "conclusions": "From February until November 2002, we do not see any chemical pecularities in the spectrum of V838 Mon. The observed spectrum resembles the spectrum of the photosphere of a normal supergiant of slightly reduced metallicity [Fe/H] = -0.4 or -0.2 (see Kipper et al. 2004, Kaminsky \\& Pavlenko, 2005). It therefore provides indirect observation of the evolution of a massive star or binary system with a massive component.  Then, V838 Mon represents a case of very slow evolution. The spectral sequence of V838 Mon can be fitted by the sets of computations for model atmospheres with hydrostatic equilibrium. Weak or intermediate strong spectral lines form in the atmospheres with decreasing temperature outwards. Only the strongest lines in February-March 2002 show P Cyg profiles. The relaxation time after the stellar flash(es) does not exceed 10 days. All these factors should be accounted for in future theoretical models."
    },
    "0609/astro-ph0609375_arXiv.txt": {
        "abstract": "Sterile neutrinos with masses in the keV range can be the dark matter, and their emission from a supernova can explain the observed velocities of pulsars. The sterile neutrino decays could produce the x-ray radiation in the early universe, which could have an important effect on the formation of the first stars.  X-rays could ionize gas and could catalyze the production of molecular hydrogen during the ``dark ages''.  The increased fraction of molecular hydrogen could facilitate the cooling and collapse of the primordial gas clouds in which the first stars were formed.  \\vspace{1pc} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609696_arXiv.txt": {
        "abstract": "We report the discoveries of the two-image gravitationally lensed quasars, SDSS J0746+4403 and SDSS J1406+6126, selected from the Sloan Digital Sky Survey (SDSS). SDSS J0746+4403, which will be included in our lens sample for statistics and cosmology, has a source redshift of $z_s=2.00$, an estimated lens redshift of $z_l\\sim 0.3$, and an image separation of $1\\farcs08$. SDSS J1406+6126 has a source redshift of $z_s=2.13$, a spectroscopically measured lens redshift of $z_l=0.27$, and an image separation of $1\\farcs98$. We find that the two quasar images of SDSS J1406+6126 have different intervening \\ion{Mg}{2} absorption strengths, which are suggestive of large variations of absorbers on kpc scales. The positions and fluxes of both the lensed quasar systems are easily reproduced by simple mass models with reasonable parameter values. These objects bring to 18 the number of lensed quasars that have been discovered from the SDSS data. ",
        "introduction": "\\label{sec:intro}  Gravitationally lensed quasars are unique astronomical and cosmological tools \\citep[see the review by][]{kochanek06}. They are useful for studying the mass distributions, substructures and interstellar mediums of their lens galaxies, the structure of the quasar host galaxy and accretion disk, and estimating the cosmological model \\citep[e.g.,][]{kochanek91,mao98,chang84,turner90,fukugita90, refsdal64}. Lensed quasars due to foreground galaxy clusters probe the distribution of dark matter in the lens objects \\citep{oguri04a,williams04}, and their statistics can be an useful test of dark matter models \\citep{narayan88,oguri04b}. A continuing problem, however, is the limited number of strongly lensed quasars \\citep[approximately 100,][]{kochanek06} and the still more limited number found in systematic surveys. In particular, the current largest completed survey, the Cosmic Lens All Sky Survey \\citep[CLASS;][]{myers03,browne03}, led to a well-defined statistical sample of only 13 lensed radio sources. While the CLASS survey can be used to obtain constraints on dark energy \\citep{chae02,mitchel2005}, larger well-defined lens samples are needed to be competitive with other approaches to constraining the cosmological model.  The Sloan Digital Sky Survey \\citep[SDSS;][]{york00} is currently the best candidate for building a larger, well-defined lensed quasar sample than the CLASS. The SDSS is expected to eventually provide a catalog of roughly $100,000$ spectroscopically identified quasars \\citep{schneider05,schneider06}, which we have been searching for gravitationally lensed quasars in the SDSS Quasar Lens Search \\citep[SQLS;][]{oguri06,inada06b}. Since the probability of a quasar being multiply imaged is of order $10^{-3}$ \\citep{turner84}, we have a reasonable expectation of obtaining a well-defined sample of roughly 100 lensed quasars. Including the two new lenses reported here, we have now discovered 18 new lensed quasars \\citep{inada03a,inada03b, inada03c,inada05,inada06a,johnston03,morgan03,pindor04,pindor05, oguri04c,oguri05,burles06,morokuma06,kayo06}, and recovered 8 previously known lensed quasars \\citep{walsh79,weymann80,surdej87, bade97,oscoz97,schechter98,morgan01,magain88}, making the SQLS the largest survey for gravitational lensed quasars to date.  In this paper, we present two new doubly-imaged lensed quasars found in the course of the SQLS, SDSS J074653.03+440351.3 (SDSS J0746+4403) and SDSS J140624.82+612640.9 (SDSS J1406+6126). They were confirmed with the observations at the W. M. Keck Observatory's Keck II telescope, the University of Hawaii 2.2-meter (UH88) telescope, and the MDM Observatory's 2.4-m Hiltner (MDM 2.4-m) telescope. The structure of this paper is as follows. We describe our lens candidate selection from the SDSS data in \\S \\ref{sec:sdss}. The results of the imaging and spectroscopic follow-up observations for SDSS J0746+4403 and SDSS J1406+6126 are presented in \\S \\ref{sec:0746} and \\S \\ref{sec:1406}, respectively.  Section \\ref{sec:model} discusses the models of the lensed quasars and we summarize our results in \\S \\ref{sec:conc}. We assume a cosmological model with matter density $\\Omega_M=0.27$, cosmological constant $\\Omega_\\Lambda=0.73$, and Hubble constant $h=H_0/100{\\rm km\\,sec^{-1}Mpc^{-1}}=0.7$ \\citep{spergel03}. ",
        "conclusions": "\\label{sec:conc}  We have discovered two new lensed quasars in the SDSS: SDSS J0746+4403 and SDSS J1406+6126.  SDSS J0746+4403 is a two-image lens ($\\Delta{\\theta}=1\\farcs08$) formed from a $z=2.00$ quasar by a foreground galaxy at $z_l{\\sim}0.3$. The redshift of the lens galaxy was estimated from its colors and $I$-band magnitude. SDSS J1406+6126 is a two-image lens ($\\Delta{\\theta}=1\\farcs98$) formed from a $z=2.13$ quasar by a spectroscopically confirmed lens galaxy at $z_l=0.27$.  Both lenses have properties that are easily reproduced by standard lens models with reasonable parameters.  SDSS J0746+4403 is particularly important because it will  be included in the well-defined statistical sample of the SQLS that we plan to use to constrain the cosmological model. SDSS J1406+6126 is slightly fainter than our magnitude limit for the statistical sample, but the relatively large image separation and the bright lens galaxy should make it a useful lens for studying the structure of early-type galaxies. A particularly interesting feature of SDSS J1406+6126 is that it has two strong \\ion{Mg}{2} absorption line systems that show dramatic changes on the few kpc scales of the path separations, which could provide an useful insight into the size distribution of absorption systems. Spectroscopy to measure the redshifts and velocity dispersions of the lens galaxies, high-resolution imaging to determine the structures of the lens galaxies, and measurements of the time delays are the next steps towards using these lenses to study the structure and evolution of galaxies."
    },
    "0609/astro-ph0609143_arXiv.txt": {
        "abstract": "{} {The aim is to study the polarization of thermal dust emission to see if the alignment of grain by radiative torques could explain the observed relation between the degree of polarization and intensity. Predictions are made for polarimetry observations with the Planck satellite.} {Our results are based on model clouds derived from MHD simulations of magnetized turbulent flows. The continuum radiative transfer problem is solved with Monte Carlo methods in order to estimate the three-dimensional distribution of dust emission and the radiation field strength affecting the grain alignment. The minimum grain size aligned by the radiative torques is calculated, and the Rayleigh polarization reduction factor, $R$, is derived for different grain size distributions. We show maps of polarized thermal dust emission that are predicted by the models. The relation between the intensity and polarization degree is examined in self-gravitating cores. Furthermore, we study the effects of wavelength, resolution, and observational noise.} {We are able to reproduce the P/I-relation with the grain alignment by radiative torques. The decrease in intrinsic polarization and total emission means that sub-mm polarimetry carries only little information about the magnetic fields in dense cores with high visual extinction. The interpretation of the observations will be further complicated by the unknown magnetic field geometry and the fact that what is observed as cores may, in fact, be a superposition of several density enhancements. According to our calculations, Planck will be able to map dust polarization reliably when $A_{\\rm V}$ exceeds $\\sim 2^{\\rm m}$ at spatial resolution of $\\sim 15'$.} {} ",
        "introduction": "Magnetic fields play an important role in the dynamics of molecular clouds and in the star formation process. In protostellar cores the topology of the magnetic fields can be predicted with models of core formation and evolution. However, observations of the magnetic field structure is needed to test such models. The best technique for studying magnetic fields in molecular clouds is to observe the polarization of the light from background stars and the polarization of the thermal dust emission, both arising because dust grains are aligned by the magnetic field.  Although the polarization of starlight by aligned interstellar dust grains was discovered over half a century ago (Hall \\cite{hall}, Hiltner \\cite{hiltner}), the processes responsible for the alignment are still under study. In a recent review by Lazarian (\\cite{lazarian}) the various proposed mechanisms are discussed extensively.  One of the first processes to be suggested was the paramagnetic mechanism (Davis \\& Greenstein \\cite{davisgreen}), which is based on direct interaction of rotating grains with the interstellar magnetic field. While this mechanism requires stronger magnetic fields than have been uncovered by other means, it may contribute at very small grain sizes. Purcell (\\cite{purcell}) introduced several processes to make grains very fast ``suprathermal'' rotators and identified one of them, catalytic sites ejecting ${\\rm H_{2}}$ as the major cause of fast grain rotation. However, it couldn't explain, for example, why the observations indicated that small grains are less aligned than large ones.  Fairly recently, the Purcell mechanism was re-evaluated in a series of papers by Lazarian and colleagues, pointing out processes such as internal grain wobbling (e.g. Jones \\& Spitzer \\cite{josp67}; Lazarian \\cite{la94}; Lazarian \\& Roberge \\cite{laro97}), grain flipping (Lazarian \\& Draine \\cite{ladr99a}) and ``nuclear relaxation'' (Lazarian \\& Draine \\cite{ladr99b}). All these contribute to making the Purcell mechanism inefficient.  The radiative torque mechanism was introduced by Dolginov (\\cite{do72}) and Dolginov \\& Mytrophanov (\\cite{domy76}). Draine \\& Weingartner (\\cite{drwe96}) demonstrated the efficiency of radiative torques using numerical simulations, while Abbas et al. (\\cite{abal04}) demonstrated it in a laboratory setup. The radiative torques use the interaction of radiation with a grain to spin it up. The predictions of the radiative torque mechanism are roughly consistent with observations (e.g. Lazarian et al. \\cite{laal97}; Hildebrand et al. \\cite{hial00}).  In view of this, the radiative torque mechanism is a strong candidate as the primary mechanism for grain alignment inside molecular clouds. Cho \\& Lazarian (\\cite{cholazarian}) showed that even deep inside GMCs ($A_V \\leq 10$) large grains can be aligned by radiative torques, using a simple, radially symmetric core model. Thus, far-infrared and submillimeter polarimetry could reliably reflect the structure of magnetic fields deep inside molecular clouds.  Padoan et al. (\\cite{paal01}) used numerical MHD simulations to model a molecular cloud containing protostellar cores, and computed dust continuum polarization for self-gravitating cores. They made the simplifying assumption that the temperature of the dust is uniform, which allowed them to calculate the polarization based on gas density and magnetic field vectors. They calculated the polarization maps for two cases: one where the grain alignment mechanism efficiency would be independent of the visual extinction and another where it was assumed that the mechanism would stop working at a sharp cut-off of $A_{\\rm V} = 3^{\\rm m}$, and thus dense cores would not contribute to the polarization. The latter case reproduced the observed results of decreasing polarization degree with increasing intensity in the cores, which led them to conclude that the grains at high visual extinction appear not to be aligned. This would mean that submillimeter polarization maps carry little information on the magnetic fields inside protostellar cores at visual extinction larger than $A_{\\rm V} \\approx 3^{\\rm m}$.  Like Padoan et al. (\\cite{paal01}), we base our study on three-dimensional simulations of magnetized gas flow, but we refine the analysis method used to predict the polarized dust emission. Radiative transfer calculations are used to estimate the actual three-dimensional distribution of dust emission. This naturally reduces the weight of high column density regions, because there the emission will be reduced by the attenuation of the heating radiation field. Similarly, radiative transfer modelling allows us to estimate, at every point of the cloud, the radiative torque efficiency of Cho \\& Lazarian (\\cite{cholazarian}). This way we can test the hypothesis, that reduced polarization of dense cores could be due to the non-alignment of smaller grains. We study the general morphology of the resulting dust polarization maps. In particular, we look at the relation between total and polarized dust emission in dense cores. Finally, some predictions are made concerning the dust polarization that will be observable with the Planck satellite. ",
        "conclusions": "\\subsection{Current polarization model}  We have shown dust polarization maps computed for model clouds that are based on three-dimensional MHD simulations. The main improvement with respect to the previous study by Padoan et al. (\\cite{paal01}) is that we have calculated the true distribution of dust emission with the aid of three dimensional radiative transfer calculations. This affects the way different volume elements contribute to the observed signal. Furthermore, instead of resorting to an ad hoc model, the polarization reduction factor has been calculated according to the model of Cho \\& Lazarian (\\cite{cholazarian}). In this model the grain alignment depends on radiative torques and, consequently, radiative transfer calculations have also been used to estimate the strength of the external radiation field at each position of the model clouds. This has been accomplished through the determination of values of $A_{\\rm V,eff}$.  Observations show a decreasing polarization degree toward dense cores. Gon\\c calves et al. (\\cite{goal05}) proposed a geometrical effect, the toroidal magnetic fieldlines bending to resist the gravitational pull, to partly explain this drop. Padoan et al. (\\cite{paal01}) could explain the observed reduction by a qualitative model where the dust alignment was restricted to regions with $A_{\\rm V}$ below 3$^{\\rm m}$. This extinction was calculated in similar fashion to our $A_{\\rm V, eff}$. With our models we find a very similar behaviour in the plots of polarization degree vs. intensity (e.g., Fig.~\\ref{cores}). At low column densities the polarization degree varies between $\\sim$4 and 10\\% while, at the centre of the cores, the degree drops to a few percent or, in some cases, even below one per cent. In other words, our physical model is capable of explaining the general observed behaviour (e.g., Ward-Thompson et al. \\cite{WT2000}; Henning et al. \\cite{Henning2001}). At the same time, the individual fields show considerable variations, and we find also cases with no apparent decrease in the polarization degree (e.g., Core 3 and direction $z$) or even increased polarization around intensity maximum (Core 3 and the filament seen toward direction $x$). This indicates that observations of individual cores may show considerable variation, depending on the actual configuration of the magnetic field. One apparent problem is that from observations it may be difficult to discern real and possibly gravitationally bound cores from projection effects. In particular, a core with large polarization reduction factor may appear very similar to a filament with major axis and field lines parallel to the line of sight.  The individual contributions of local density and $A_{\\rm V, 1D}$ to the reduction of the polarization degree also merit a mention. In Model $C$ the local density varied between $\\sim 1$ and $\\sim 10^5\\;{\\rm cm^{-3}}$ while $A_{\\rm V, 1D}$ had values from 0.005 to 4.3. Just by looking at Eq.~\\ref{eqaalg} it is easy to see that in global scale, the variation in local density dominates over $A_{\\rm V, 1D}$. However, in the cores, where the densities are around the same order of magnitude, the contribution of $A_{\\rm V, 1D}$ comes to its own by emphasizing the difference between the illuminated regions on the outer parts of the core and the shadowed, even denser interior. In cores, both local density and $A_{\\rm V, 1D}$ are of similar importance.  Related to this issue is the isotropic and anisotropic components of the radiation field, which play major part in the efficiency of grain alignment by radiative torques. In this work, we neglected detailed simulation of the anisotropy of the radiation field needed to derive the actual values of the anisotropy factor, $\\gamma$, and instead used the fitting formula (Eq.~\\ref{eqaalg}) by Cho and Lazarian (\\cite{cholazarian}) where $\\gamma = 0.7$. Draine and Weingartner (\\cite{drwe96}) give $\\gamma = 0.1$ for diffuse clouds, due to the anisotropy of the star distribution in the Galaxy. We can qualitatively estimate the effect that the change of $\\gamma$ would have as it would be similar to the increase of local density by the corresponding inverse factor. For example, the difference between $10^2$ and $10^3\\;{\\rm cm^{-3}}$ at low $A_{\\rm V, 1D}$ (see Fig.~\\ref{Rfactors}) is a drop from 0.95 to 0.8 in $R$. Thus, we would expect a slightly lower polarization in those regions, but not by much. If anything, it would help to reduce the 'contamination' of the diffuse background and make the cores stand out more. Also, because we would expect to see a more isotropic field at the centre of the core than at its exterior, the inclusion of $\\gamma$ would make the core centres drop in polarization even more sharply. In other words, our model is likely to overestimate the polarization at the core centres, rather than the opposite.  In the modelling we assumed dust properties to remain constant throughout the model clouds, while only the alignment of the grains was dependent on the local density and the visual extinction. However, there is ample evidence that optical dust properties do change already above a few magnitudes of $A_{\\rm V}$. If this were caused primarily by a change in grain size distribution, this could counteract the reduction obtained via the $R$ factor. As an example, let us consider an optically thick core with density $n\\sim 10^4$\\,cm$^{-3}$ and $A_{\\rm V, 1D}\\sim 6^{\\rm m}$. According to Fig.~\\ref{Rfactors} the $R$-factor could increase by a factor of several if the L\\&D grain distribution were replaced with a distribution with much larger $a_{\\rm max}$. On the other hand, Fig.~\\ref{scatterR} shows that, in the actual model, the effect of an increasing $a_{\\rm max}$ was rather limited. In that particular case, when L\\&D dust model is replaced with MRN type distribution with $a_{\\rm max}\\sim 1\\mu$m, the polarization degree increases in the core centre only from $\\sim$2\\% to $\\sim$3\\%. It is very unlikely that in such a core $a_{\\rm max}$ could grow close to one micron. Another explanation for the relative insensitivity to $a_{\\rm amax}$ is that the observed intensity is weighted toward low density material. There are two reasons for this. First, the polarization reduction factor decreases rapidly, whatever the value of $a_{\\rm max}$. Second, in a dense core the dust is colder and, therefore, the emission is lower for a given column density. In a core with $A_{\\rm V,1D}\\sim 7^{\\rm m}$ to the centre, the dust temperature may already be reduced to $\\sim 14$\\,K or below (Stamatellos \\& Whitworth \\cite{stwh03}). This can be compared with $\\sim$18\\,K in diffuse material. In the Rayleigh-Jeans tail the effect is small, and even at 353\\,GHz only a factor of 1.5. However, the effect becomes important when one approaches the peak of dust emission. At 100$\\mu$m the intensity produced by an 18\\,K grain is already ten times as large as the intensity produced by a 14\\,K grain. From Fig.~\\ref{Rfactors} one can see that at low densities the slope of the $R$ vs. $A_{\\rm V, 1D}$ relation is very similar for all values of $a_{\\rm max}$, and the percentile difference in the absolute value of $R$ are similarly small.  The above picture may change if dust grains undergo more profound changes. The sub-mm emissivity of dust has been observed to increase significantly already in clouds with visual extinction only a few magnitudes (Bernard et al. \\cite{bernard1999}, Stepnik et al. \\cite{stepnik2003}). This has been interpreted as evidence of not only grain growth but the appearance of dust aggregates and fluffy grains of low packing density. At long wavelengths the high emissivity might again increase the contribution that the densest regions have on the total intensity observed towards cores. On the other hand, the formation of dust aggregates can also affect the intrinsic polarization efficiency of the grains and change the efficiency of the alignment mechanisms (Lazarian \\& Efroimsky \\cite{Lazarian1999}; Gupta et al. \\cite{Gupta2006}). Such effects are not considered in this paper.  \\subsection{Comparison with other simulations}  On larger scales, Prunet et al. (\\cite{Prunet1998}) presented simulations of polarized dust emission at high Galactic latitudes. The calculations were based on three-dimensional cloud structure that was estimated based on HI data. The HI line velocity provided information on the line-of-sight density distribution. The simulations assumed constant dust properties with intrinsic polarization efficiency $\\sim$30\\%. Furthermore, three field geometries were considered. A uniform field direction resulted in high polarization degree with very narrow distribution, and was dismissed as being incompatible with observations. In our simulations this corresponds closely to model $A$, where the field lines have remained almost parallel because of the strength of the magnetic field. Of course, strong polarization also requires that the cloud is not observed parallel to the field lines. In our model $A$, when the sightlines were perpendicular to the field lines, the polarization degree peaked around 10\\%. For a random orientation the mean polarization would be lower by half. In the other two cases studied by Prunet et al. the field was assumed to be either parallel with the major axis of the cloud structures, or to have a random orientation perpendicular to it. Both resulted in a wider distribution of polarization degrees that, in the case of $R=1$, peaked around 2-3\\%. These are again similar to our models $B$ and $C$, where, irrespective of the viewing angle, we have an approximately log-normal distribution of polarization degrees that peaks around the value of $\\sim$3\\%. Therefore, although the results were obtained by entirely different methods, the basic properties of the dust polarization seem very similar. This is reassuring for attempts to build statistically meaningful all-sky templates for the polarized dust emission that will be observed, e.g., by the Planck satellite.  Another important statistic is the power spectrum of the polarized signal. Figure~\\ref{fig:powerspectra} shows the power spectra computed for the model $C$. Spectra are shown separately for the I, Q, and U components, as well as the polarized intensity ${\\rm I_{pol}}$, from direction $y$. These can be compared with the power spectra presented by Prunet et al. (\\cite{Prunet1998}). The slope of intensity, in interval $2 \\leq k \\leq 20$, is about $k^{-3}$, in agreement with Prunet et al. (\\cite{Prunet1998}). The Stokes parameters Q and U have a slightly more shallow slope, but not as flat as by Prunet et al. (\\cite{Prunet1998}). It can also be seen that in small scales, the power in the polarized components drops slower than the total intensity, making the polarized power an appreciable fraction of the total intensity power. The reason for this would appear to be the magnetic field morphology causing abrupt, pixel to pixel variations or larger structures, to about 10-pixel level, while the background intensity stays about the same.   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{5838_18.eps}} \\caption{ The power spectra of the polarized intensity ${\\rm I_{pol}}$ and the Stokes parameters ${\\rm I_{tot}}$, Q and U are plotted for model $C$ from direction $y$. Column density $N$ follows the intensity power spectrum closely, and is thus not shown. While model $A$ showed a distinct shift in the level of the Q and U power spectrum in direction $z$, it - and model $B$ - were not different enough in shape to justify another figure. } \\label{fig:powerspectra} \\end{figure}  \\subsection{Comparison with observations}  Direct observations of polarized dust emission are still very scarce for all but the very densest regions. The first detection of diffuse polarized dust emission was reported by Beno\\^it et al. (\\cite{Benoit2004}). The observations were done with the balloon-borne Archeops experiment, with instruments that closely resemble those of the Planck satellite. Polarized signal could be detected at 353\\,GHz with the resolution of 13 arcmin, the average polarization degree being of the order of 4-5\\%. However, a large variation was seen between different regions. In the Cygnus region the polarization was low, as can be expected if the magnetic field is preferentially parallel to the spiral arms. On the other hand, in some regions of the Galactic plane the polarization reached values of $\\sim$15\\%. In our models, polarization degrees this high could be reached only in model $A$ and only when the polarization reduction degree $R$ was equal to one (see Fig.~\\ref{ld_poldeg_353}). However, our maximum polarization degree is somewhat arbitrarily set by our choice of $\\alpha$, or to be more exact, our choice of $C_{\\rm pol}/C_{\\rm ran} = 0.15$. This ratio scales our maximum polarization degree to about 15.3\\%, which explains why we do not reach $P = 15\\%$ in all our models, since that would demand a total alignment and no depolarization on the line of sight. A slightly larger, yet still a moderate deviation from sphericity, can double our polarization degrees: at axial ratio of 1.2, $C_{\\rm pol}/C_{\\rm ran}= 0.3$. Thus one needs to be very careful not to overstate the case. Certainly a high polarization degree hints to an ordered magnetic field and high polarization efficiency throughout the cloud, yet even small changes in our assumptions about the grain shapes or lowering the density of our very dense model clouds would change the maximum polarization degree so that the other two models could duplicate even the high polarization degree of 15\\%.  Ponthieu et al. (\\cite{poal05}) presented the measurements of temperature and angular power spectra of the diffuse emission of Galactic dust at 353\\,GHz as seen by Archeops on 20\\% of the sky. They concluded that if the fraction of the sky observed by Archeops is representative, dust polarized radiation will be a major foreground for determining the polarization power spectra of the CMB at high frequencies above 100\\,GHz. However, their power spectra cannot be directly compared with our results, because they cover much larger area, where the turbulence is no longer the dominant factor affecting the structure of the density and the magnetic fields.   \\subsection{Future}  In Fig.~\\ref{fig:planck_fwhm3} it is shown that when the spatial resolution of the Planck satellite is reduced to 15 arcmin, it will be able to provide a relatively clean map of the polarization of a diffuse cloud at 353\\,GHz. The general morphology of the magnetic field is more or less correct above $A_{\\rm V} \\sim 0.5^{\\rm m}$, even if the polarization degree P itself has a relative error of $\\sim$60\\%, as seen in Fig.~\\ref{fig:planck_noise}. The situation improves with increasing $A_{\\rm V}$, until at $A_{\\rm V} \\sim 2^{\\rm m}$, the errors have dropped to the relatively low values of $\\sim$25\\% for P and $\\sim$15 degrees for the polarization angle $\\chi$. Thus, we conclude that Planck will be able to trace the general morphology of the magnetic fields at $A_{\\rm V} \\sim 1^{\\rm m}$ and to map dust polarization reliably when $A_{\\rm V}$ exceeds $\\sim 2^{\\rm m}$.  In this study we have used simulations with relatively low spatial resolution. The results remain reliable at large scales, and the discussion of core properties is appropriate for general density enhancements and the initial phases in the formation of gravitationally bound cores. In a rotating and collapsing core the magnetic field lines become wound so that both the field strength and the depolarization caused by the field geometry increase. If the system were not perfectly symmetric one might, at high resolution, detect sharp variation in both the strength and direction of the polarization vectors.  However, in order to be able to follow the development of magnetic fields inside such collapsing objects, simulations with considerably higher spatial resolution are needed. Spatially variable $\\gamma$ and the inclusion of other grain alignment mechanisms would allow a more realistic polarization model. The work presented in this paper will be continued with such more detailed studies."
    },
    "0609/astro-ph0609469_arXiv.txt": {
        "abstract": "Recent works showed that the absorbing material in broad absorption line (BAL) quasars is optically thick to major resonant absorption lines. This material may contribute significantly to the polarization in the absorption lines. In this paper, we present a detailed study of the resonant line scattering process using Monte-Carlo method to constrain the optical depth, the geometry and the kinematics of BAL Region (BALR). By comparing our results with observed polarized spectra of BAL quasars, we find: (1) Resonant scattering can produce polarization up to 9\\% at the absorption trough for doublet transitions and up to 20\\% for singlet transitions in radially accelerated flows. To explain the large polarization degree in the CIV, NV absorption line troughs detected in a small fraction of BAL QSOs, a nonmonotonic velocity distribution along the line of sight or/and additional contribution from the electron scattering region is required. (2) The rotation of the flow can lead to the rotation of the polarization position angle (PA) in the line trough. Large extending angle of BALR is required to produce the observed large PA rotation in a few BAL QSOs. (3) A large extending angle of BALR is required to explain a sub-trough in the polarized flux that was observed in a number of BAL QSOs. (4) The resonant-scattering can contribute a significant part of NV emission line in some QSOs, and may give rise to anomalous strong NV lines in these quasars. (5) The polarized flux and PA rotation produced by the resonant scattering in non-BAL is uniquely asymmetric, which may be used to test the presence of BALR in non-BAL QSOs. ",
        "introduction": "Blue-shifted, broad absorption lines (BAL) are observed in the ultraviolet spectra of about 10-20\\% luminous quasars. These lines are formed in partially ionized outflows with velocities up to 0.1 c. The outflow is likely driven by intensive radiation of the quasar probably along the equatorial directions to the extension at least larger than the broad emission line region, and is likely several 10 parsecs. Disk wind and material evaporating from the putative dust torus are two plausible scenarios for the origin of the gas. It is usually believed that BAL region exists in every quasar, but only covers a small fraction of quasar sky (Weymann et al. 1991; Reichard et al. 2003; Green et al.  2001; Hamann, Korista \\& Morris 1993, hereafter HKM93). The outflow may carry a significant fraction of power released by the accretion process and momentum into the host galaxy of the quasar, so that it will influence the subsequent evolution of the galaxy. However, in order to establish its role, we need to understand many properties of the outflow such as the global covering factor of BAL region, the column density and velocity field as a factor of the distance to the continuum sources.  The absorption line profiles are usually quite stable over time scales of several to ten years, suggesting of smooth outflow and/or saturation of the UV emission lines. Similar line strength from ions of very different abundance and strong absorption detected in soft X-rays supports latter interpretations (Brinkmann et al. 1999; Wang et al. 1999; Gallagher et al. 2002). The broad absorption lines have now probably been detected also in X-rays with much large column densities. If confirmed, this will suggest that the very high velocity outflow is already there at very close to the continuum emission region. Efficient acceleration at small scale is required.  However, census has yet to be reach on a number of critical issues: (1) The covering factor of BALR is likely a function of fundamental parameters, such as the black hole mass and the accretion rate, which may lead to some difference in the statistical properties of the BAL QSOs and non-BAL QSOs (Boroson 2002), but we still need to find concrete evidence for this and their potential relation. (2) Whether the outflow is equatorial or polar is still a matter of controversy. Recent VLBI observations of a small number of radio loud BAL quasars with equal number of steep and flat radio spectra reveal only compact structure in most case (Jiang D.R. et al. in preparation). While based on the radio variability, Zhou et al. (2006) proposed polar outflows in six radio loud quasars. It is still unclear that whether radio loud BAL quasars are special. Hydrodynamic simulation of accretion disk wind models, however, results in an equatorial outflow. We note including poloid magnetic field in the accretion disk may change the simulation results as for the radio jet model (Blandford \\& Payne 1982); (3) Whether the outflow carry significant angular momentum, i.e., whether the massive disk wind serves as the driver of the accretion process. (4) There is big concern whether certain ultraviolet emission lines such as NV, CIV will be significantly affected by resonant scattering process, such that we need to revise our metallicity determination in some objects.  Our current knowledge about the BALR is almost exclusively from either absorption lines or X-ray absorption edge. Unfortunately, both types of absorption carries only information of the absorbing gas along the line of sight to the continuum source, and we have to rely on the statistics of a large sample of BAL quasars to obtain the average information of the global properties. Note such information (on the global properties) is contained in the scattered light, i.e., the polarized flux.  Broad Absorption Line QSOs (BAL QSOs) are the only highly polarized population among radio quiet QSOs (e.g., Stockman, Moore \\& Angel 1984). Its optical/UV continuum shows polarization (e.g., Stockman et al. 1984; Schmidt \\& Hines 1999; Hutsem$\\grave{e}$kers \\& Lamy 2001), much larger than that of non-BAL QSOs. The high continuum polarization is believed due to the electron scattering probably in the BAL region (BALR, Stockman et al. 1984; Ogle 1997; Wang, Wang, Wang 2005, hereafter Paper I). In Paper I we also show that if the BALR exists in all QSOs, and covers around 20\\% of the solid angle, the electron scattering in the BALR can successfully explain the observed continuum polarization for both BAL QSOs and non-BAL QSOs. %  Observations show that the polarization is even stronger in the BAL trough than in the continuum. Ogle et al. (1999, hereafter O99) presented a spectropolarimetric survey of 36 BAL QSOs, and found: \\begin{itemize} \\item The BAL troughs are usually more polarized than the continuum, whereas the broad emission lines are less polarized (also see Cohen et al. 1995; Hines et al. 1995). Deeper BAL troughs tend to have higher polarization degrees. The polarization in the trough can be as high as $\\sim$20\\%. \\item Position angle (PA) of the polarization in the troughs are quite common, and smaller rotations across the corresponding emission lines were found in some objects. \\item In the spectra of the polarized flux, the absorption line troughs are usually evident but appear shallower and show various characteristics: the troughs in polarized flux are more blueshifted than that in the total flux spectrum in some objects; a sub-trough emerges to the red side of the CIV, SiIV and/or NV absorption trough in the polarized flux in several objects, e.g., 0105-0265, 0226-1024, 1413+1143 $\\&$ 1333+2840; occasionally, a boxy absorption trough, similar to that in the total flux, was also detected (e.g., in 2225-0534 and 1232+1325). \\item Excess polarized flux across the corresponding emission lines is observed in several objects (see also Lamy \\& Hutsem$\\grave{e}$kers 2004). \\end{itemize}  Two processes were proposed to explain the origin of the polarization in the absorption line troughs. If the BALR does not cover or only-partially covers the electron scattering region, the leaked scattered photons will fill the troughs, thus produce high polarization (Hines \\& Wills 1995; Goodrich \\& Miller 1995; Ogle 1997; O99; Lamy \\& Hutsem$\\grave{e}$kers 2004). On the other hand, Lee \\& Blandford (1997, hereafter LB97) showed that resonant scattering can produce polarization degree as large as 15\\% in the troughs for doublet transitions. Note LB97 does not calculate the PA rotation for resonant scattering light. Since then, our knowledge about the column density, and the geometry of BALR has changed considerably from the X-ray observations as well as high resolution UV spectroscopy. In particularly, as showed in Paper I, the electron scattering in the BALR can produce the continuum polarization. By taking these new information into consideration, we will explore in detail the polarization properties of resonantly scattered light, including the polarization degree, polarized flux and position angle of the polarization for different models using Monte Carlo method.  This paper is organized as follows. The geometries and dynamics of the outflow model and the Monte-Carlo method will be described in \\S 2. The results of Monte-Carlo simulation are given separately for singlet and doublet transitions in \\S 3 and \\S 4, respectively. In \\S 5, we will compare our results with the observed polarized spectra of a sample of BAL QSOs (O99) to put constraints on the geometries and kinematics of the flow. Resonantly scattered lines are discussed in \\S 6. ",
        "conclusions": "Polarization provides rich information on the structure and kinematics of BALR in QSOs, which is complementary to those derived from absorption lines, which only depend on the physical condition and kinematics of gas on the LOS. To extract this information, we carry out extensive Monte-Carlo simulations of electron and resonance scattering process in the BALRs. Both singlet and doublet transitions are considered for radial outflows of two different geometries: equatorial outflows and hollow-conical outflows with and without rotational velocities.  In an axisymmetric scatterer model, PA rotation in the absorption trough can be produced only when the outflow carries angular momentum. The PA rotation increases with the angular velocity as well as the subtending angle of the flow. In order to produce a PA rotation $>$ 10\\deg observed in a few BAL QSOs, subtending angle of the outflow should be larger than 25\\deg. Similar requirement is imposed to explain the red sub-trough in the polarized flux observed in some BAL QSOs.  Resonant scattering of doublet transition produces much lower polarization at large optical depths (about 6\\% for $\\tau_0=5$, $\\beta=0.5$) for an accelerating outflow. Large ($>$10\\%) polarization in the absorption trough detected in a few BAL QSOs may indicate that the optical depth to the resonant scattering is at most moderate, otherwise resonant scattering of $T_{1/2}$ transition will erase all polarization. A slowly-accelerating model can produce larger polarization. If the observed high-velocity outflow in X-ray is real, the flow is likely decelerating, which will produce much larger polarization.  A large electron scattering region can also produce larger polarization. The jump which appears at the starting velocity of the absorption trough in PA rotation and polarized flux in models with a SESR does not appear in the model with a LESR. This characteristic may be an important indicator to distinguish different electron scattering models. We find that LESR models can explain that the absorption troughs are more blueshifted in polarized flux, relative to that in the total flux, and PA swings in the troughs, $\\sim 90^{\\rm o}$ relative to continuum, in some objects.  We show that the resonantly scattered light will contribute a significant part of NV in some QSOs and can give rise to anomalous strong NV lines in these QSOs. A correlation between the EW of NV$\\lambda$1240 with the continuum polarization is expected. We propose that the polarized flux and the PA rotation of the scattering emission can be used to test the presence of BALR in non-BAL QSOs. Further observations with large telescopes should allow us to extract the important information about the flow geometry and kinematics."
    },
    "0609/astro-ph0609233_arXiv.txt": {
        "abstract": "We present mid-IR N-band $(\\lambda_{eff} = 10.2\\micron)$ photometry of a carefully selected sample of T Tauri stars thought to be single from the Taurus-Auriga molecular cloud. Infrared excesses in these stars are generally attributed to circumstellar dust-disks. Combining observations at 2.16$\\micron$ (K$_{s}$-band) and 10.2$\\micron$ (N-band) we probe a region in the circumstellar dust-disk from a few stellar radii through the terrestrial planet zone (0.02-1.0AU). By analyzing the distribution of the $(K_{s}-N)$ color index with respect to previously measured photometric rotation periods we investigate what role circumstellar disks play in the rotational evolution of the central star. The resulting positive correlation between these two variables is consistent with the notion that a star-disk interaction facilitates the regulation of angular momentum during the T Tauri stage. We also demonstrate, how including non-single stars in such an analysis will \\textit{weaken} any correlation in the relation between $(K_{s}-N)$ color and period. To further understand disk properties we also present SEDs for a few objects with new ground based M-band $(\\lambda_{eff} = 4.8\\micron)$ and Q-band $(\\lambda_{eff} = 20\\micron)$ data and compare them to a geometrically thin, optically-thick disk model. ",
        "introduction": "As a rotating molecular cloud core collapses to form a star it must shed vast amounts of angular momentum to explain the observed differences between the specific angular momentum of molecular cloud cores and stars on the main-sequence \\citep{bod95}. Various physical processes, which serve as solutions to this ``angular momentum problem'' are thought to occur as an object transitions from a cloud core to a Zero-Age Main Sequence (ZAMS) star. For low-mass stars in the pre-main sequence stage (T Tauri stars) the regulation of stellar rotational angular momentum is thought to be facilitated by the formation of and subsequent interaction with a circumstellar disk. Circumstellar disks are thought to be common by-products of the star formation process \\citep{BS96}. They are easily identified by peculiar features like infrared and ultraviolet excesses, and in some rare circumstances can also be detected through direct imaging \\citep{Odell93,MO96}. \\citet{KH95} estimate $> 50\\% $ of stars in the \\object{Taurus-Auriga} star forming region have circumstellar disks, while \\citet{hillenbrand98} estimate that $> 80\\%$ of stars in the Trapezium cluster possess circumstellar accretion disks \\citep[see also][]{haisch01}. The IR excesses from these disks are the result of heating circumstellar gas and dust by radiation from the central star and through viscous accretion. The disk temperature distribution is often modeled as a power-law. If the disk is optically thick it will radiate as a blackbody with the shortest wavelengths dominated by higher temperature material in the inner disk and the longer wavelengths dominated by cooler material in the outer regions \\citep{beckwith99}. Thus, emission at different wavelengths can be used to probe different regions in the disk.  The pioneering study by \\citet{edwards93} for a set of stars from the Taurus, Auriga, Chameleon, Orion and Lupus star--forming regions, compared stellar rotation periods with the $(H-K)$ color index. This color index probes a region in the disk that lies within a few stellar radii $(\\sim 10R_{\\sun})$. They found stars whose $(H-K)$ colors indicated the presence of an accretion disk were rotating more slowly than stars whose $(H-K)$ colors indicated the absence of an accretion disk. A similar study by \\citet{bouvier93} suggested a bimodal distribution of slowly rotating classical T Tauri stars (cTTS, showing evidence of active accretion from a circumstellar disk) and fast rotating weak T Tauri stars (wTTS, weak-lined $H\\alpha$ emission-line stars with no signs of accretion). Studies of the \\object{Orion Nebula Cluster} (ONC) have yielded mixed results. \\citet{CH96} report a bimodal distribution in photometric periods of T Tauri stars in Orion as evidence for the disk-locking phenomena \\citep[see also][]{herbst02}. Theoretical models have always relied heavily on some form of direct or indirect angular momentum transfer between stars and their disks to explain the redistribution of rotational angular momentum during the transition from the Pre-Main Sequence (PMS) stage to the ZAMS stage. However, an extensive study conducted by \\citet{stassun99,stassun01} report no bi-modality in period distribution and no dichotomy of disked slow rotators and disk-less fast rotators. A recent review of this literature can be found in \\citet{mathieu04}. In this study we use the $(K_{s}-N)$ color index to probe a region in T Tauri dust-disks that extends from a few stellar radii through the terrestrial planet zone (0.02-1.0AU). Comparing the $(K_{s}-N)$ color index of carefully selected single stars with previously measured rotation periods allows us to investigate any effects the disk might have on the angular momentum of the star.  This paper outlines an observational study that expands on the work presented in \\citet{MB97}, and hopes to further our understanding of the role of disks in the rotational evolution of T Tauri stars. In section 2 we present a description of our sample, the observations and the data reduction procedure. In section 3 we present the relationship between the $(K_{s}-N)$ color index and rotation period of single PMS stars in Taurus-Auriga, and the spectral energy distributions (SEDs) of a few stars with new M and Q-band data. In section 4 we discuss the results in light of current theories for the rotational evolution of T Tauri star and disk systems. We also demonstrate that including binaries in a study of $(K_{s}-N)$ vs. rotation period diminishes any observed correlation. We conclude with a summary in section 5. ",
        "conclusions": "In the following sections we discuss the significance of our results in the context of current theoretical frameworks for rotational evolution. We also compare our observational study to others conducted in Tau-Aur and the ONC. Finally we emphasize the importance of sample selection criteria for studying rotational evolution of low mass PMS stars.  \\subsection{Disk-regulation} The positive correlation, at the 99$\\%$ confidence level, between the $(K_{s}-N)$ color index and rotation period supports the idea that disks play a role in regulating the angular momentum of T Tauri stars. Most models of PMS angular momentum evolution rely on some form of disk-assisted angular momentum regulation. In the absence of a disk, we would expect young stars to conserve total angular momentum as they contract towards the ZAMS, spinning-up as their radii decrease over time. Our results suggest that the disk is countering the spin-up effect caused by ordinary stellar contraction and to some extent the spin-up torque caused by the accretion of high angular momentum material from the disk boundary \\citep{GL79a,GL79b}. Evidence for accretion is seen in the U-band excess (Figure 3) of even the fast rotator \\object{GK Tau}. \\citet{camenzind90}, \\citet{konigl91} and \\citet{CC93} explore the idea of magnetospheric star-disk coupling. In this model the stellar magnetic field is thought to couple with the disk, extending to regions where the disk is spinning slower than the stellar surface. As a result the star is braked and eventually slows its rotation significantly. \\citet{shu94} and \\citet{OS95} suggest it is unlikely for magnetic coupling to occur over extended regions of the disk. Their models allow for the disk to be penetrated by the stellar magnetic field only near the co-rotation radius, hence making the transfer of angular momentum via rotational coupling insignificant. Instead, angular momentum is lost from the stellar surface through the expulsion of X-winds (eXtraordinary winds) near the star-disk boundary. The magnetized X-winds are driven by accretion from the disk. Similarly \\citet{vRB04} and \\citet{MP05} explore a scenario where angular momentum is lost via accretion-powered stellar winds. In all theoretical frameworks (disk-locking, X-winds or stellar winds) the presence of an extended disk (0.02-1.0AU) is vital to angular momentum regulation. The magnetospheric disk locking models require an extended magnetized disk beyond the co-rotation radius to provide a spin-down torque. The wind models require the presence of accretion from the inner disk, which in turn depends on the extended disk to persist.  T Tauri stars are thought to dissipate their disks over a range of times during their evolution \\citep{strom89,S90,haisch01}. We conducted a two-sided Kolmogorov-Smirnov(K-S) test to see if the age distributions for disked and disk-less stars (as defined by the $(K_{s}-N)$ index) in our sample are similar. We derived a D-statistic of 0.32 indicating that the age distributions are consistent with having been drawn from the same parent population. If disk-assisted regulation is a dominant force the angular momentum evolution of a T Tauri star, then the fact that stars dissipate their disks at different times would explain the wide dispersion of rotational rates observed for stars on the ZAMS \\citep{BFA97,K97}. In other words stars that dissipate their disks early are believed to spin-up and become rapid rotators, while stars that dissipate their disks late, end up as slow rotators \\citep{bouvier94,rebull04,HM05}.  \\subsection{Sample Selection} Most rotational evolution models explore only the case of a single star's interaction with its disk. Comparing the $(K_{s}-N)$ index with the rotation period of single stars enables us to compare our results to the theoretical frameworks which use single star + disk scenarios. Furthermore, there is some evidence that the presence of a companion affects the evolution of a disk around a PMS object \\citep{meyer97b,ghez94}. \\citet{jensen96} and \\citet{JM97} present convincing evidence that disks from 50-100AU are affected by the presence of binary companions over this range of separations.  In Figure 4 we combine data points from Figure 1 with $(K_{s}-N)$ and period data available for unresolved binaries (listed in Table 4). Performing a linear correlation test similar to the one described in section 3, for $(K_{s}-N)$ vs. rotation period, resulted in a correlation coefficient of 0.35, for 38 data points (not including upper-limits). This means the probability that linear correlation is random is $< 0.05$, suggesting that the presence of unresolved binaries in a sample of young stars could \\emph{weaken} any expected correlation. We also conducted a two sided K-S test on the objects in Figure 4, separating them into disked and disk-less stars, as described in section 3. We derived a D-statistic of 0.45, which indicates that the probability the distributions are similar is $7.0\\%$. This is a $1.0\\sigma$ result. So we cannot confidently assert that the period distributions were drawn from different parent populations. \\citet{ghez93} estimate that 2/3 of the stars in the Tau-Aur region are multiple systems. Yet, the linear correlation test only includes 13 binary systems compared to the 25 single systems (excluding upperlimits). This is because there is a lack of photometric period data for approx 85$\\%$ of the stars surveyed to be multiple systems, compared to 50$\\%$ for those survyed to be single systems \\citep{ghez93,leinert93,simon95,bouvier93,bouvier95,osterloh96}. This is probably due to the difficulty in interpreting data from binary systems. Most photometric period surveys that include binaries do not resolve companions and hence it is hard to know for certain a) which component's period was determined (if the stars are of comparable brightness) or b) if the light fluctuations were due to oribital motions (eclipsing) or actual rotational motion (spots). In addition, components in a close binary system are hard to resolve in the infrared. The source of N-band flux from an unresolved multiple system cannot be easily determined. Contributions from each stellar component, the dust disk around each object, or a circumbinary disk \\citep{ghez94} could add to the N-band flux. In investigating star-disk interactions in the PMS, it is important to restrict the sample to well-studied single stars, or resolved binaries.  \\subsection{Resolving Apparent Discrepancies} The results from \\citet{edwards93} for the $(H-K)$ color index shows evidence for disk-less fast rotators and disked slow rotators in Tau-Aur.  However, the $(H-K)$ color index probes a region in the dust-disk very close to the stellar surface (approx $<0.02$ AU). \\citet{muzerolle03} estimate that the typical temperature at which dust sublimates around T Tauri stars is approximately 1400K, which roughly corresponds to the temperatures of the hottest dust emitting in K$_{s} (\\lambda_{eff} = 2.16\\micron)$ and L-band $(\\lambda_{eff} = 3.54\\micron)$. So emission from the dust-disk is limited to the region with $T_{disk} < 1400K$. If we consider a case where the disk truncation radius lies far beyond the sublimation zone (where $T_{inner-disk} << 1400K)$ at $r_{inner} >> 0.02 AU$, the disk might still be effectively disk-locked but will not display any excess in H and K-bands. There are 8 stars from our single star set which are also studied by \\citet{edwards93}. Interestingly, 1 of the 8 stars (\\object{DN Tau}) lacks a $(H-K)$ excess but \\textit{does display a} $(K_{s}-N)$ \\textit{excess} (P$_{rot}$ = 6.0 days). This could mean that its disk is truncated outside of the sublimation radius. The other 7 stars either lack both $(H-K)$ and $(K_{s}-N)$ excesses or display both $(H-K)$ and $(K_{s}-N)$ excesses. There are \\textit{no cases} where an object lacks a $(K_{s}-N)$ excess but shows an $(H-K)$ excess.  In a complementary study, \\citet{stassun99} surveyed a large sample of stars in the ONC at near-IR wavelengths, with periods $< 8.0$ days, and discovered no bimodality in the distribution of rotation periods and no correlation of (I-K) color excess with rotation periods. However, subdividing the sample by mass to select only stars $> 0.25M_{\\sun}$ demonstrated that the higher-mass PMS stars display the bimodal distribution of rotation periods \\citep{herbst00}. The \\citet{stassun99} sample is dominated by very low-mass stars which may possess different magnetic star-disk coupling mechanisms. \\citet{CB00} argue that the ONC population differs significantly from the \\object{Taurus} population in that PMS stars in Taurus are generally slow rotators. \\citet{hartmann02} suggests that this could be because the typical disk-braking time-scale of T Tauri stars is comparable to the ages of T Tauri stars in the ONC. Hence, the mechanism might not have had time to brake most disked stars in the ONC. \\citet{stassun01} obtained mid-IR data for a sample of stars from the ONC and Tau-Aur selected not to have near-IR excesses. These data indicate no trend of slow rotators with mid-IR excesses in their sample for both clusters. \\citet{stassun01}, like most previous studies, include single, binary and stars that have not been surveyed for companions in the analysis that led to a null result. Interpretation of period and infrared fluxes of binary stars is not straightforward as discussed in the previous subsection. In fact, for the single stars in their sample (also included in our sample) their results are entirely consistent with ours.  Some of our results echo the observations made by \\citet{stassun01} that there is no strict dichotomy of disked slow rotators and disk-less fast rotators. A handful of stars qualify as abnormal cases. \\object{GK Tau} is a fast rotator with a period of 4.65 days. Yet it displays a strong infrared excess, $E(K_{s}-N) = 3.20$. The well-known PMS star \\object{SU Aur} appears to be a disked $(E(K_{s}-N) = 3.33)$ fast rotator (P = 3.50d) and \\object{IP Tau} is definitely a disked $((K_{s}-N) = 2.30)$ fast rotator (P = 3.25d). \\object{LkCa 21} is a disk-less slow rotator and \\object{V819 Tau} is disk-less ``moderate-to-fast'' rotator, classified as a ``transition object'' by \\citet{S90}. We queried 2MASS H and K$_{s}$-band magnitudes and derived the $E(H-K_{s})_{o}$ index for these 5 objects. $4/5$ objects have near-IR properties similar to their mid-IR $E(K_{s}-N)_{o}$ properties. GK Tau, SU Aur and IP Tau have significant $E(H-K_{s})_{o}$ excesses while LkCa 21 has no $E(H-K_{s})_{o}$ excesses. Perhaps LkCa 21 has very recently lost its disk and has not yet had time to spin up (cf. DI Tau; \\citet{meyer97b}).  Of these ``special case'' objects, only V819 Tau is included in the study by \\citet{edwards93}. It has no near--IR excess, a marginal (K-N) excess, and a period of 5.6 days.  Including all these objects in our analysis, we see a dispersion in the $E(K_{s}-N)_{o}$ vs. rotation period plot comparable to that seen in the $E(H-K)_o$ analysis of \\citet{edwards93}. It would be interesting to further investigate the physical properties (M$_{\\star}$,$\\dot{M}$, R$_{\\star}$,B) of these objects that are expected to determine the co-rotation point in their disks. For example, if $\\dot{M}$ is large enough in the protostellar phase, the inner disk might crush the magnetosphere resulting in a population of rapidly rotating objects \\citep{najita95,covey06}. Note that the disked fast rotators here have moderate to low accretion rates \\citep{GHBC98}. A sample of stars with a fixed M$_{\\star}$,R$_{\\star}$, and B should exhibit a correlation between P$_{rot}$ and $\\dot{M}$ in the x-wind model. Such studies might also elucidate the \\emph{dispersion} in the $(K_{s}-N)$ vs. period graph.  From a theoretical perspective, we don't necessarily expect a linear correlation between $(K_{s}-N)$ and P$_{rot}$. However, our fit demonstrates the low probability that these variables are uncorrelated. Our two-sided K-S test of the periods of disked vs. disk-less stars also hints at a correlation. \\citet{rebull04} point out that correlations between P$_{rot}$ and dust-disk or accretion indicators might be difficult to establish given that the spin-up time is of order the age of typical pre-main sequence samples \\citep{hartmann02}. The fact that we do see a correlation and that our disked and non-disked stars appear to be of similar age (1-3 Myr) suggests that these weak-emission T Tauri stars in Taurus lost their disks very early in their pre-main sequence evolution. Deep far-infrared or millimeter observations might yet reveal whether they retain any disk signatures at all (e.g. \\citet{AW05}). New surveys for disks at all radii enabled with the Spitzer Space Telescope combined with photometric rotation studies will shed considerable light on these remaining questions."
    },
    "0609/astro-ph0609719_arXiv.txt": {
        "abstract": "We study the 36 brightest radio sources in the Subaru/\\textit{XMM-Newton} Deep Field (SXDF). Using MIPS 24 $\\mu \\rm m$ data from Spitzer we expect to trace accretion activity, even if it is hidden at optical wavelengths, unless the obscuring column is extreme. Our results suggest that in the decade or so below the break in the radio luminosity function that at least half, and potentially nearly all, radio sources are associated with accreting quasar-like objects. This is not true at lower radio luminosities where the quasar-like fraction approaches zero once compact sources are excluded. ",
        "introduction": "In unified AGN models, powerful radio sources are believed to have a central source emitting in the optical, UV and X-rays. The dusty torus that surrounds the nuclear region absorbs this light and re-emits it in the infrared. This mechanism makes it difficult to observe the objects viewed through the torus directly in the optical, UV and soft X-rays. Willott et al. (2000) define the quasar fraction $f_{q}$ to be the ratio of unobscured to obscured quasars. In \\cite{eleniv06} we expressed the need to study radio sources in the mid-infrared in order to investigate the existence of any hidden  accretion activity in radio sources. Spitzer observations at 24 $\\mu \\rm m$ are ideal for this task, since they trace hot dust emission which can be obscured only by extreme columns (\\cite{ogle06}).  Our sample is the 36 brightest radio sources from the VLA survey of the SXDF (\\cite{simpson06}) with flux densities greater than 2 mJy at 1.4 GHz. Most of our objects are not as yet spectroscopically confirmed, so we use photometric redshifts for our analysis. We use a radio spectral index $\\alpha$ = 0.8 ($S_{\\nu} \\propto \\nu^{-\\alpha}$). We assume throughout a low-density, $\\Lambda$-dominated Universe in which $H_{0}=70~ {\\rm km~s^{-1}Mpc^{-1}}$, $\\Omega_{\\rm M}=0.3$ and $\\Omega_{\\Lambda}=0.7$.  More details of this work will appear in Vardoulaki et al. (in prep.). ",
        "conclusions": "We have constructed spectral energy distributions (SEDs) for our SXDF sample and fitted Bruzual-Charlot galaxy templates using the HyperZ photometric code (\\cite{hyperz}) to estimate photometric redshifts. An example is shown on the left side of Fig. \\ref{fig1}. This object is spectroscopically confirmed at redshift $z$ = 1.095. From its SED we can see excess emission at 24 $\\mu \\rm m$, indicating that this object has hidden accretion activity. This object is not a spectroscopically confirmed quasar but has high-excitation narrow lines in its optical spectrum; its radio map shows an unresolved source.  We define objects to be `quasar-like' when their rest-frame luminosity at $\\lambda =$ 24 $\\mu \\rm m$ is $\\lambda L_{24 \\mu \\rm m}>10^{37.3}$ ${\\rm W}$, since this corresponds to  $\\lambda L_{24 \\mu \\rm m} > 10^{-1.8} L_{Edd}$ for a quasar with $M_{\\rm BH}$ $\\gtsim$ $10^{8}$ $\\rm M_{\\odot}$ (see \\cite{mclure04}; \\cite{mclureetal04}). The luminosity at 24 $\\mu \\rm m$ is calculated using the spectral index measured from 24 $\\mu \\rm m$ to the nearest lower-$\\lambda$ detection in the observed frame SED. The quasar-like fraction $q_{\\rm l}$ in the region between the RLF break and the FRI/II break is high, $q_{\\rm l} \\sim 0.5-0.9$, depending on how many of the 24 $\\mu \\rm m$ upper limits turn out to be close to detections. [The most radio luminous of the two non-quasar-like objects, is an FRI at $z \\sim 0.7$]. Recall that $f_{q} \\sim 0.2$ in this region (\\cite{willott00}), i.e. the 24 $\\mu \\rm m$ data have revealed at least some hidden accretion activity.  The quasar-like fraction drops significantly below the FRI/II break, and if one excludes compact ($D <$ 100 kpc) sources as potentially part of a separate (beamed) population, we see that there are no definite quasar-like objects.  We conclude that Spitzer 24-$\\mu \\rm m$ data on the SXDF radio sources are consistent with the hypothesis that quasar-like objects are almost always associated with powerful (above the FRI/II break) radio sources, but are rarely connected to less luminous sources, unless they are part of the compact population.   \\begin{figure} \\plottwo{vardoulaki_e_fig1.ps}{vardoulaki_e_fig2.ps} \\caption{{\\itshape Left:\\/} Observed frame SED of sxds\\_0033 with a redshift estimated using HyperZ and data from Spitzer, UKIDSS and SXDF surveys. {\\itshape Right:\\/} $\\log_{10}( L_{1.4{\\rm GHz}})$ versus largest projected linear size $D$: possible quasar-like objects have upper limits on $\\lambda L_{24 \\mu \\rm m}$ consistent with $\\lambda L_{24 \\mu \\rm m}>10^{37.3}$ ${\\rm W}$.} \\label{fig1} \\end{figure}"
    },
    "0609/astro-ph0609005_arXiv.txt": {
        "abstract": "{The~VIMOS~VLT~Deep~Survey~(VVDS)~was~undertaken to map the evolution of galaxies, large scale structures, and active galaxy nuclei from the redshift spectroscopic measurements of $\\sim$10$^{5}$~objects down to an apparent magnitude $I_{AB} = 24$, in combination with a multi-wavelength acquisition for~radio,~infrared,~optical,~ultraviolet, and~X-rays~data.} {We~present~the~evolution~of~the~comoving~star~formation~rate (SFR)~density~in the redshift range $0 < z < 5$ using the first epoch data release of the VVDS, that is 11564 spectra over 2200~arcmin$^2$ in two fields of view, the VVDS-0226-04 and the VVDS-CDFS-0332-27, and the cosmological parameters ($\\Omega_\\mathrm{M}$, $\\Omega_{\\Lambda}$, $h$)~=~(0.3, 0.7, 0.7).} {We study the multi-wavelength  non dust-corrected luminosity densities at~$0 < z < 2$ from the rest-frame far ultraviolet to the optical passbands, and the rest-frame 1500~\\AA\\ luminosity~functions~and~densities~at~$2.7 < z < 5$.} {They~evolve from $z=1.2$ to $z=0.05$ according to $(1+z)^{x}$ with $x = 2.05, 1.94, 1.92, 1.14, 0.73, 0.42$, and 0.30 in the FUV-1500, NUV-2800, U-3600, B-4400, V-5500, R-6500, and I-7900 passbands, respectively. From~$z=1.2$~to~$z=0.2$ the $B$-band density for the irregular-like galaxies decreases markedly by a~factor~$3.5$ while it increases by a~factor~$1.7$ for the elliptical-like galaxies. We identify several SFR periods; from $z=5$~to~$3.4$ the FUV-band density increases by at most 0.5~dex, from $z=3.4$~to~$1.2$ it decreases by 0.08~dex, from $z=1.2$ to $z=0.05$ it declines steadily by 0.6~dex.  For the most luminous $M_{AB}(1500~\\AA)<-21$ galaxies the FUV-band density drops by 2~dex from $z=3.9$ to $z=1.2$, and for the intermediate $-21<M_{AB}(1500~\\AA)<-20$ galaxies it drops by 2~dex from $z=0.2$ to $z=0$.  Comparing with dust corrected surveys, at $0.4 \\la z \\la 2$ the FUV seems obscured by a constant factor of $\\sim1.8-2$ mag, while at $z<0.5$ it seems progressively less obscured by up to $\\sim0.9-1$ mag when the dust-deficient early-type population is increasingly dominating the $B$-band density.} {The~VVDS~results~agree~with~a~downsizing~picture~where~the~most~luminous~sources~cease~to~efficiently~produce~new stars 12 Gyrs ago (at $z\\simeq 4$), while intermediate luminosity sources keep producing stars until 2.5 Gyrs ago (at $z\\simeq 0.2$). A modest contribution of dry mergers and morphologies evolving towards early-type galaxies might contribute to increase the number density of the bright early types at $z<1.5$.  Our observed SFR density is not in agreement with a continuous smooth decrease since $z\\sim4$.} ",
        "introduction": "\\subsection{The global star-formation history one decade ago} The history of the comoving space density of the star formation rate (SFR) is a key study which has undergone a spectacular explosion of publications since the pioneer work of \\citet{madau96} in which galaxy surveys were used for the first time.  \\citet{madau96} derived the global star formation as a function of redshift combining emissivities from three distinct surveys: the H$\\alpha$-selected UCM survey \\citep{gallego95} at $z < 0.05$, the largest spectroscopic sample at that time, the $I$-band selected Canada-France Redshift Survey at $0.2 < z < 1$ \\citep[CFRS;][]{1996ApJ...460L...1L}, and their own galaxy sample at $2 < z < 4.5$ using the Lyman continuum break colour-selection technique \\citep{steidel96} applied on the deep optical imaging, the Hubble Deep Field (HDF) survey.  Nevertheless high-$z$ sources were lower limits, and incompleteness corrections and uniform limiting magnitude were required.  The SFR density is usually derived from the mean luminosity density, or emissivity, that is $\\cal{L} = \\int^{\\infty}_\\mathrm{0} \\phi(\\mathrm{L}) \\mathrm{L} \\mathrm{dL}$. At first sight the SFR density appears a simple and useful tool to trace back the evolution of star formation and to link it with the evolution of stellar mass, but a large spread between different measurements have led to controversy.  Uncertainties in conversion factors from luminosity to SFR, and in the amount of dust obscured SFR, coupled with the different selection criteria of each survey and with the uncertainty about the shape of the lumino\\-si\\-ty function, meant that the SFR history of the Universe was poorly determined, and thus hotly debated.  In the mid-90's the main questions that arose were: is there a peak of the cosmic star-formation history at $1.3<z<2.7$? By how much does interstellar dust attenuate the ionizing flux? Is the evolution so rapid below $z = 1$? Is the high-redshift dropout population representative? Does the red galaxy population evolve passively or not? \\subsection{A wealth of multi-wavelength studies} Measuring the cosmic star-formation history has advanced rapidly in the last decade meanwhile one had to translate results using the standard cold dark matter (CDM) into the $\\Lambda$CDM cosmology.  It adds uncertainties since the cosmological cons\\-tant introduces a redshift dependence to the luminosity function unless one can take back all the data needed to compute the luminosity function. Still the inferred larger distances and vo\\-lumes result in decreasing luminosities and densities, and thus in a shallower SFR evolution.  On the one hand, a wealth of multi-wavelength emissivities (far- and near-ultraviolet, far- and near-infrared, radio, H$\\alpha$, etc.) has been obtained, all of them being more or less directly proportional to the ionizing ultraviolet stellar spectra at $\\lambda < 912$~\\AA\\ mainly produced by massive, young, short-lived stars (OB stars, $t <$ few x $10^6$ yr). On the other hand, supernovae events and their by-products (neutrinos and gamma rays) are proportional to the star-death rate and are also used to probe the SFR.  Nevertheless these results depend most strongly on stars more massive than 3 solar mass, and require extrapolation of the initial mass function (IMF) to lower masses to obtain the SFR for all masses.  Futhermore, the ultimate desired value is the amount of the interstellar gas mass transformed into stars, and the correlated gas mass return into the interstellar medium as stars deplete their initial fuel sources; it requires assumptions and models for stellar atmospheres and stellar evolution tracks.  An example of an extensive compilation drawn from the literature of SFR density measurements at $0 < z < 6$ was done by \\citet{hopkins04}, and \\citet{hopkins06}, and it led to a cosmic SFR history constrained to within factors of about $\\sim3$. There is now growing evidence that the evolution has no peak at $z\\sim1.5$, however it is still unclear whether at $z>3$ the evolution flattens, or declines or continues to increase.  Those remaining uncertainties show the necessity to further investigate galaxy redshift surveys to constrain galaxy assemblies. In particular, one needs to explore multi-wavelength datasets over the same field of view, and to obtain homogeneous datasets over a large redshift range.  A detailed picture of the SFR history is emerging with the advent of large, deep redshift surveys, coupled with multi-wavelength ground and space observations over the {\\it same} sky area.  These new multi-wavelength surveys are providing a better understanding of the nature and the evolution of the galaxy population \\cite[i.e., e.g.][and references within]{bell04}. Indeed on one hand, one can probe the same galaxy population at different wavelengths which gives insights about stellar masses (near-infrared) and star formation (far-ultraviolet, far-infrared), possibly coupled with spectroscopic indexes. On the other hand, one can measure reliable comoving volumes with accurate redshifts and intrinsic luminosities of objects which both are key measurements to estimate luminosity functions. Futhermore observing in various windows of the electro\\-ma\\-gne\\-tic spectrum enable to select the largest galaxy sample which includes sources with specific energy distributions, and detectable at only some wavelengths.  The accuracy in the measurement of the shape of the luminosity function, coupled with the stellar mass-to-light ratio of galaxies, is crucial to estimate the amount of baryons in stars at a given epoch of the Universe. \\subsection{The detailed study of the VVDS} The VIsible Multi-Object Spectrograph (VIMOS) installed on the European Southern Observatory (ESO) Very Large Telescope (VLT) was built to produce systematic large redshift surveys thanks to its high multiplex capabilities, e.g. $\\sim$550 $R\\simeq230$ spectra of sources observed simultaneously over 218 arcmin$^2$ \\citep{lefevre03}.  We conducted the VIMOS VLT Deep Survey (VVDS), a major multi-wavelength spectroscopic survey, to investigate the evolution of galaxies, Active Galaxy Nuclei (AGN), and large-scale structures\\footnote{http://www.oamp.fr/virmos/}.  Here we detail our work in using the $I$-selected VVDS first epoch data described in \\citet{lefevre05a} to study the evolution of the luminosity density within the redshift range $0<z<5$.  It is part of a series of papers which analyse different aspects related to the luminosity function evolution with this data set.  \\citet{ilbert05} describe the global optical luminosity functions over $0<z<2$. \\citet{zucca06} explore the color-type luminosity functions over $0<z<1.5$.  \\citet{ilbert06a} investigate the contribution of different morphological types to the luminosity functions. \\citet{ilbert06b} analyse the luminosity functions in different environments over $0<z<1.5$.  \\citet{paltani07} analyse in detail the 1700~\\AA\\ luminosity function at $3<z<4$.  \\citet{arnouts05} des\\-cribe the 1500~\\AA\\ luminosity functions over $0<z<1.3$ using GALEX-VVDS data. The infrared luminosity functions with K-band data and with SWIRE-VVDS data, and the luminosity functions for different spectroscopic-based classes are in preparation.  This paper is organized as follows. In Section~2 we present the data. In Section~3 we detail the methods used to estimate the comoving luminosity densities and associated uncertainties. In Section~4 we present the multi-wavelength global luminosity densities at $0<z<2$ derived in the UBVRI and in the near-UV, far-UV passbands.  In Section~5 we compare our results to other surveys at $z<2$.  In Section~6 we investigate the luminosity densities for different galaxy types at $z<2$.  In Section~7 we measure the rest-frame 1500~\\AA\\ luminosity functions and densities at $2.7<z<5$ and we compare our results to other surveys.  In Section~8 we detail the evolution of the global far-UV luminosity density all the way from $z=5$ to $z=0$, and we analyse its dependency to the luminosity.  In Section~9 we derive the history of the star formation rate density since $z=5$ and discuss the issue of dust obscuration.  Finally in Section~10 we recap our conclusions about the evolution of the rest-frame luminosity densities in a well controlled and homogeneous $I$-selected population over the large redshift range $0<z<5$ as observed by the VVDS. Throughout this paper we use the AB flux normalization \\citep{oke74}.  We adopt the set ($\\Omega_\\mathrm{M}$, $\\Omega_{\\Lambda}$, $h$)~=~(0.3, 0.7, 0.7) for the cosmological parameters. ",
        "conclusions": "We studied the first epoch VVDS data purely $I$-selected at [17.5-24]~mag in (AB).  The sample is unique in the sense that it goes deeper than previous $I$-selected spectroscopic samples, it has a well-defined selection function and it has enough data to study sub-samples. Within a single survey, we trace the evolution of the galaxy population dominating the total light at different redshifts all the way from $z=5$ to $z=0$.  The main results of our comprehensive study are summarised below. \\begin{itemize} \\item[i)]To study the luminosity density evolution, observing deeper in flux is superior to assembling large samples as far as the galaxy sample is not dependent on cosmic variance and has a well-defined selection function.  Nevertheless, the survey field must be large enough to identify the rare bright objects required to constrain the bright end of the LF. \\item[ii)]The luminosity density evolution is substantially wavelength-dependent since the rest-frame passband luminosity is related more or less directly to very different stellar populations. We find that the non dust-corrected $\\cal{L}$ evolve with time over $0.05\\le z \\le1.2$, as $\\cal{L} \\propto \\rm (1+z)^{x}$ with $x=2.05, 1.94, 1.92, 1.14, 0.73, 0.42$, and $0.30$ in the FUV-1500, NUV-2800, U-3600, B-4400, V-5500, R-6500, and I-7900 passbands, respectively. Although error bars are still large, we note that most luminosity densities exhibit a transition at $z\\simeq1.1$ in the evolutionnary tendency. Over the last 8.5~Gyrs, the SFR-related FUV-band luminosity density drops by a factor 4. Recent merger events should produce little star formation. \\item[iii)]The $B$-band luminosity density evolution is strongly type-dependent. From $z=1.2$ to $z=0.2$ the irregular-like, and Scd-like type galaxies decrease markedly by factors 3.5, and 1.7, respectively while the elliptical-like, and Sab-like types increase by factors 1.7, and 1.6, respectively. The late-type galaxy population undergoes a downsizing scenario where most star formation is shifting to intermediate- and low-mass galaxies at $z<1.2$.  The early-type galaxy population suggests a little contribution from merging or simple rest-frame color evolution phenomena, but it evoles passively in terms of luminosity, which suggests a very early formation stage ($z\\gg4$) for the majority of this red population. <\\item[iv)]The SFR density as seen in the rest-frame FUV-1500 luminosity density without dust correction undergoes the several following up-and-down phases as redshift decreases, and is strongly luminosity dependent.  From $z=5$ to $z=3.4$, it increases by at most a factor $\\sim3.5$, and it corresponds to the end of the mass assembly of the most luminous galaxies, $M_{AB}(1500~\\AA)<-21$ (L$>$L$^*_{\\rm LBG}$), which present a peak of SFR at $z\\simeq3.5$. From $z=3.4$ to $z=1.2$ it decreases by a factor $1.2$. Notwithstanding the large error bars this small decrease might be the result of two opposite processes; a decrease by a factor $\\sim1.4$ from $z=3.4$ to $z=2$ due to the fading of the giant galaxy population by a factor $40$, and an increase by a factor $\\sim1.3$ from $z=1.2$ to $z=1.0$ due to the shift of the star formation activity towards the less luminous ($-21<M_{AB}(1500~\\AA)<-19$) galaxy population which presents a peak of SFR at $z\\sim1$.  From $z=1.2$ to $z=0.05$ it declines steadily by a factor $4$, since this phase undergoes the fading of both the giant and the intermediate galaxy populations. Nevertheless it hides a strong SFR drop by a factor $100$ at $z<0.2$ of the intermediate galaxy population. This phenomena might be combined to the local underdensity seen in optical redshift surveys (see \\citet{zucca97}, \\citet{tresse98}), and thus the local Universe might not be an excellent reference local point to trace back the evolution of the global galaxy population.  Our observed global evolution does not seem to be in agreement with a continous smooth decrease from $z\\sim2$ to $z\\simeq0$ as predicted by the simulations. \\item[v)]Comparing our SFR FUV-derived densities with mid-infrared or H$\\alpha$ SFR-derived densities, we find that at $0.4 \\la z \\la 2$ the FUV is obscured by a constant factor of $\\sim1.8-2$ mag, and at $z<0.5$ it is progressively less obscured down to $\\sim0.9-1$ mag. In parallel, we find that from $z=0.4$ to $z=0.05$ the $B$-band luminosity density is more and more dominated by the early-type (E/Sab-like) galaxy population which is known to be dust deficient. Futher analysis combining VVDS-GALEX-SWIRE data will refine this result. \\item[vi)]We conclude that the old, most luminous and massive galaxies have exhausted their cold gas reservoir during their early intense star formation which has occured in the early Universe at $z\\gg4$, and since $z\\simeq3.5$, i.e. $\\sim$ 12~Gyrs, they undergo passive evolution as star formation cease.  This creates excellent dry merger candidates, that is, in which gas has been sufficiently depleted that it prevents new star formation.  The small level of contribution of dry-mergers remains to be quantified. Younger, less massive and less luminous $L^{*}$ ga\\-la\\-xies continue to efficiently form stars with a large reservoir of cold gas up to $z=0.2$.  And at $z<0.2$ these galaxies appear to have also exhausted their gas supply as suggested by the $FUV$-band luminosity density noticeable drop, and start to evolve passively, and might be combined with a local underdensity.  This picture is consistent with the downsizing scenario for the star formation \\citep{cowie96}. We recall that the luminosity and SFR densities are mainly dominated by the $L^{*}\\phi^{*}$ galaxies rather than the dwarf population which is usually undergoing density evolution \\citep{ilbert05}. At $z<3.5$ dry mergers and morphologies evolving towards early-type galaxies might contribute to increase the number density of the bright early-types in maintaining a passive luminosity evolution as observed by \\citet{zucca06} at $0<z<1.5$, and also to increase the emissivity in over-dense regions as observed by \\citet{ilbert06b}. \\end{itemize}"
    },
    "0609/hep-ph0609310_arXiv.txt": {
        "abstract": "Quantum electrodynamics predicts that photons undergo one-loop scattering. The combined effect of this on the behaviour of a photon gas for temperatures above $\\sim10^{10}$~K results in a softening of the equation of state. We calculate the effect this has on the effective equation of state in the early universe, taking into account all the species of the Standard Model. The change to the dynamics of the early universe is discussed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609249_arXiv.txt": {
        "abstract": "The SFI++ consists of $\\sim$5000 spiral galaxies which have measurements suitable for the application of the I-band Tully-Fisher (TF) relation. This sample builds on the SCI and SFI samples published in the 1990s but includes significant amounts of new data as well as improved methods for parameter determination. We derive a new I-band TF relation from a subset of this sample which consists of 807 galaxies in the fields of 31 nearby clusters and groups. This sample constitutes the largest ever available for the calibration of the TF template and extends the range of line-widths over which the template is reliably measured. Careful accounting is made of observational and sample biases such as incompleteness, finite cluster size, galaxy morphology and environment. We find evidence for a type-dependent TF slope which is shallower for early type than for late type spirals. The line-of-sight cluster peculiar velocity dispersion is measured for the sample of 31 clusters. This value is directly related to the spectrum of initial density fluctuations and thus provides an independent verification of the best fit WMAP cosmology and an estimate of $\\Omega^{0.6} \\sigma_8 = 0.52\\pm0.06$. We also provide an independent measure of the TF zeropoint using 17 galaxies in the SFI++ sample for which Cepheid distances are available. In combination with the ``basket of clusters'' template relation these calibrator galaxies provide a measure of $H_0 = 74\\pm2$ (random) $\\pm6$ (systematic) \\kmsMpc. ",
        "introduction": "The Tully-Fisher relation \\citep{TF77} is an empirical relationship between the rotational velocity and absolute magnitude (or luminosity) of spiral galaxies. The application of this relation to measure distances to spiral galaxies has had an enormous impact on extragalactic astronomy and cosmology since its discovery. The Tully-Fisher (TF) relation for spirals (along with the $D_n-\\sigma$ relation for ellipticals) dramatically increased the volume in which redshift-independent distances could be measured. The TF relation and its scatter can also be used to constrain models of disk galaxy formation, and by studying the relations variation with environment, galaxy type and its evolution over the age of the universe many clues into how the disks of galaxies are assembled have been found. A reliable template relation is an essential starting point to all this, without which no trust can be placed in other conclusions. Like all empirically determined scaling relations, the selection effects and make-up of a TF sample can introduce non-trivial biases into the derived template which if not taken care of could introduce spurious results. With improved statistics and parameter measurement both in galaxy surveys and simulations the need for an accurate, unbiased template grows even stronger.  The TF relation has historically played an important role in cosmological parameter determination, in particular by allowing measurements of $H_0$ to be extended out of the very local universe. In fact a review of TF measurements of $H_0$ over time shows that the technique has done surprisingly well. The value of $H_0$ from TF studies played a significant role in the debate between high and low values of $H_0$, with cluster peculiar velocities, the Malmquist bias and even dust extinction intervening to provide a range of measures of $H_0$ from TF. The Hubble Key Project was initiated to resolve this debate and found $H_0=71\\pm2\\pm7$ \\kmsMpc ~from a combination of various TF surveys and Cepheid data \\citep{S00}. Similarly \\citet{G97a} found $H_0=69\\pm5$ \\kmsMpc ~from a combination of the I-band TF relations in 24 clusters and a small set of Cepheid calibrators; and \\citet{TP00} found $H_0=77\\pm8$ \\kmsMpc ~from their sample of 5 clusters. In Section 7 below we derive $H_0 = 74\\pm2\\pm6$ \\kmsMpc ~from the sample discussed in this paper. Even in todays era of ``precision cosmology'' the TF constraints on $H_0$ have similar random error to the best available measurements combining WMAP and large scale structure data (\\eg ~$H_0 = 74\\pm2$ \\kmsMpc ~from a combination of WMAP and 2dFRGS; \\citealt{Sa06}), and better than WMAP alone which measures $H_0 = 73\\pm3$ \\kmsMpc ~(with the assumption that the universe is flat; \\citealt{Sp06}).  \\citet{TF77} first calibrated the TF relationship in the B band using a sample consisting of 10 galaxies with $D<8$ Mpc and eight galaxies in the Virgo cluster. They measured $M - 5 \\log h = -19.13 - 6.25 (\\log W - 2.5)$, where $M$ is the absolute magnitude, $W = 2 v_{\\rm max}$ is the width of the Doppler broadened spectral line and $h = H_0/100$ \\kmsMpc. This is equivalent to  $L \\propto v_{\\rm max}^{2.5}$. In the B band, extinction (both internal to the galaxies and from our Galaxy) plays an important role in adding to the scatter of the TF relation. \\citet{A79} argued that the relationship in the NIR would have a smaller scatter. They observed that the slope steepened in the H band, more nearly approximating the expected $L \\propto v_{\\rm max}^4$ from simple physical arguments. Much of the most successful early work using TF used the \\citet{A80} calibration based on the distance to M31 and M33, and the assumption that $L \\propto v_{\\rm max}^4$. The uncertainty in these (and other) early TF templates from small samples of nearby calibrator galaxies was obviously dominated by small number statistics and since the samples did not cover a wide dynamic range of galaxy masses they also provided little leverage on the measurement of the TF slope. Several immediate attempts were made to extend the calibration, both by looking at cluster samples (where all galaxies could be assumed to be at roughly the same distance) and obtaining more local calibrators. In the late 1980s it was argued that the I-band TF relation could improve the scatter in the relation further \\citep{BM87}. This band, intermediate between B and H, retained most of the advantages gained by moving into the NIR but had the additional advantage of allowing CCD photometry.  It should be obvious that using a strictly magnitude-limited sample to calibrate the TF relation will result in an underestimate of the slope, as at the low-width end only galaxies scattered above the line will be included. An early solution to this potentially strong bias was the adoption of so-called {\\it inverse} fits in which the magnitude was treated as the independent variable. Such a fit would not suffer from incompleteness bias in the simple case where selection depended only on absolute magnitude. Unfortunately, realistic samples do not have such simple selection criteria, and in particular internal extinction corrections, which have been shown to depend on galaxy luminosity, and therefore implicitly on the rotation width \\citep{G94,M03} make bias corrections more complicated. Arguments also exist in the literature that cluster samples constitute {\\it complete} or {\\it volume-limited} samples and therefore the TF relation derived from such a sample will not suffer from incompleteness bias. As first discussed in \\citet{S95} this is not the case. An explicit or implicit magnitude limit in a cluster sample will have a the same impact on the measured TF relation as a sample with a full spread of distances.  A biased TF template slope can introduce several subtle, {\\it qualitative} biases into many of the conclusions which might be drawn from its application. For example Malmquist bias-like effects mean that the more distant objects in any given sample are more likely to be at the large-width end of the relation while the most nearby objects are more likely to be at the low-width end of the sample. If the template is biased shallow, this means that distant objects will preferentially have spurious negative magnitude offsets (\\ie ~they appear to be artificially brighter than the biased template), while nearby objects will preferentially have spurious positive offsets. If these offsets were interpreted as being due to the peculiar velocities of the galaxies, this would produce a spurious infall region in the local universe and excess expansion in the more distant parts of the sample.  Determinations of the morphological dependence of the TF relation can be impacted by bias on the slope too. Early type spirals (Sas) are preferentially brighter and have larger rotation widths than later type spirals. A TF template biased shallow might then lead to the conclusion that the zeropoint for earlier type spirals was brighter than that for late type spirals. In a high redshift sample which will only have galaxies at the large width end of the local relation, comparison with such a biased local relation could easily lead to a spurious determination that the zeropoint of the TF relation brightens with redshift.  In fact most high redshift studies of the TF relation use the relatively small samples of \\citet{PT92} or \\citet{MFB} as a $z=0$ comparison. The \\citet{PT92} B, R and I-band calibration is based on the TF relation of only 6 local spirals for which Cepheid distances were available, in combination with a small sample ($\\sim 30$ galaxies) from the Ursa Major cluster. \\citet{MFB} use a sample of only 14 galaxies in the Fornax cluster to provide an I-band calibration. No attempt is made to account for sample selection biases in either sample. As discussed in \\citet{MK06}, larger and better calibrated local samples already exist. The \\citet{TP00} calibration uses a significantly larger sample (155 galaxies in 5 clusters at I-band; 91 galaxies at B and R and 65 at B-band), however the work still makes no attempt to account for selection biases relying on the erroneous idea that cluster samples are complete. \\citet{KF02} use a sample of 196 nearby field galaxies but just fit the {\\it inverse} U, B and R-band relations in an attempt to account for bias, and make no attempt to account for the effect of peculiar velocities which also have the potential to bias the template slope. No fully bias corrected local template exists in optical wavelengths.  The first comprehensive discussion of the effect of incompleteness bias on the TF template relation came from \\citet{W94}. Here it was explicitly demonstrated that the {\\it inverse} TF relation is not bias free. A method was introduced which used Monte Carlo simulations to calculate the bias based on the assumption of a global linear TF relation, a fixed scatter and a single magnitude limit. Giovanelli \\etal~ (1997b; hereafter G79b) extended this technique to account for a TF scatter which varied with line width, and to provide realistic accounting of sample selection criteria. They used this to calibrate I-band TF relation with a sample of 555 spiral galaxies in the fields of 24 clusters. They found $M - 5 \\log h = -21.01 - 7.68 (\\log W - 2.5)$, or $L \\propto v^{3.1}$. They observed that the scatter in the Tully-Fisher relation varies with the rotational velocity of the galaxy, from about 0.4 mag at the low width end, to less than 0.3 mag at the high width end, resulting in distance errors for individual galaxies of 15--20\\%.   Much of this paper is devoted to the derivation of a new template TF relation in the I-band. We follow a method similar to that used in G97b and use a sample of galaxies which builds on that cluster sample, but has a significant amount of new data (including the addition of several new clusters) making a total sample of 807 galaxies in the fields of 31 clusters. This sample constitutes the largest ever available for the calibration of the TF relation, and extends the range of the calibration between $v_{\\rm max}=$ 50--360 \\kms ~($\\log W =$ 2.0--2.9). In combination with the much larger field sample (Springob \\etal ~2007) this work will finally allow for studies of the extent, non-linearity and morphological dependence of the TF relation at high statistical significance.  This new I-band template is warranted both because of the addition of significant amounts of new data to the sample and because of changes in corrections to the raw data. In particular, new corrections for rotation velocities derived from observations of both the 21cm and H$\\alpha$ lines of neutral hydrogen (HI and H$\\alpha$) are expected to have an impact on the derived template. New simulations have been performed to study the instrumental effects on measuring HI widths \\citep{SH05}, and much greater care has been taken in combining HI and H$\\alpha$ rotation measures (Catinella, Haynes \\& Giovanelli 2006). The addition of significant amount of new data also alters the completeness and morphological distribution of the sample, which must be taken into account in rederiving the template.  In Section 2 we discuss the observables needed for TF. Section 3 covers the sample selection and bias corrections. In Section 4 we discuss individual cluster TF and the impact of environment. Section 5 combines the individual cluster templates into a global template. In the process we measure the peculiar velocity of each cluster, from which a cluster peculiar velocity dispersion is calculated. This quantity depends relatively simply on the power spectrum of initial density fluctuations and is used to estimate $\\Omega^{0.6} \\sigma_8 = 0.52\\pm0.06$. In Section 6 we describe the characteristics of the scatter of the TF template. Section 7 discuss an alternate zero-point calibration which uses Cepheid distances to 17 galaxies in the SFI++. In combination with the ``basket of clusters'' zeropoint this gives a measure of $H_0 = 74\\pm2\\pm6$ \\kmsMpc. The final section of the paper provides a summary of the conclusions. In future work including Masters \\etal ~(2007, in prep.) and Springob \\etal ~(2007; also see \\citealt{KLM}; \\citealt{CS}) this new template will be applied to a larger sample of spiral galaxies in the field, from which the peculiar velocity field, and therefore the mass distribution, of the local universe will be studied. ",
        "conclusions": "In this paper we present the SFI++ sample of spiral galaxies. Galaxies in this sample have both I-band photometry and velocity widths (either from HI global profiles or optical rotation curves) suitable for use in the Tully-Fisher relation. From a subset of this sample consisting of 807 galaxies in the vicinity of 31 nearby clusters we rederive the I-band TF relation. This sample constitute by far the largest ever available to calibrate the TF relation. We argue that it provides evidence for a type-dependent TF slope that is steeper for later type spirals, in effect fitting for two TF relations -- one for the 465 Sc galaxies in the sample, the other for the 342 Sa/Sb galaxies. We measure the cluster peculiar velocity dispersion for the 31 cluster samples. This quantity is directly related to the shape of the initial power spectrum of density fluctuations and is used to estimate a value of $\\Omega^{0.6}\\sigma_8=0.52\\pm0.06$. By comparing the TF template with the Cepheid distances to 17 galaxies we measure $H_0=74\\pm2\\pm6$ \\kmsMpc; an accuracy comparable to the best precision cosmology. A summary of the numeric results is presented in Table \\ref{conclusions}. A more detailed summary of the conclusions follows in this section.  \\begin{deluxetable*}{ll} \\tablecolumns{1} \\tablewidth{0pc} \\tablecaption{Summary of Numeric Results\\label{conclusions}} \\tablehead{\\colhead{Parameter} & \\colhead{Values}} \\startdata Template relation (bivariate fit to \\inp ~sample) & $M - 5 \\log h = -20.85 \\pm 0.03 - (7.85 \\pm 0.1)(\\log W - 2.5)$ mag \\\\ &  $L = 3.7 \\times 10^{10} (v_{\\rm max}/200 ~{\\rm km~s}^{-1})^{3.1} L_\\sun/h^2$\\\\ Template scatter  & $\\epsilon_{\\rm total} = 0.41 - 0.44 (\\log W - 2.5)$ mag\\\\ & $\\epsilon_{\\rm intrinsic} = 0.35 - 0.37 (\\log W - 2.5)$ mag\\\\ Cepheid calibrated zeropoint & $M = -20.49 \\pm 0.05 \\pm 0.06 - (7.85\\pm0.1)(\\log W - 2.5)$ mag\\\\ Hubble's constant & $H_0=74\\pm2\\pm6$ \\kmsMpc \\\\ Cluster peculiar velocity dispersion & $\\sigma = 298\\pm34$ \\kms \\\\ &  $\\Omega^{0.6} \\sigma_8 =0.52\\pm0.06$ \\enddata \\end{deluxetable*}  In Section 3.1 we introduced a subset of the SFI++ sample which consists of 807 galaxies in the vicinity of 31 nearby clusters (within $v_{\\rm CMB}=$10,000 \\kms). This sample is divided into an {\\bf in} sample consisting of 483 {\\it bona fide} cluster galaxies and an {\\bf in+} sample which also includes galaxies considered to be peripheral members of the cluster. Cluster membership assignments are discussed further in Springob \\etal ~(2007).  In Section \\ref{bias}.2, we discuss various biases which modify the observed TF relation from the intrinsic relation within a cluster sample. We derive corrections for these biases which include morphological type bias and incompleteness bias due to the implicit magnitude limit of the sample. The magnitude of the incompleteness bias depends most strongly on the assumed scatter in the TF relation.  We find evidence to suggest that the slope of the TF relation gets shallower for early type spirals and construct a width-dependent morphological correction to reconstruct the relation for Sc galaxies. This is equivalent to fitting two separate TF relations to the sample - one to the major subset of 465 Sc galaxies, a second to the 342 galaxies in the sample with types earlier than Sc. The slope difference can be explained either by differing completeness characteristics for the different morphological types in the sample, or a real physical difference in the galaxies. We argue for the latter as the difference is shown to be largest at the high width end of the relation where the bias corrections are smallest. In the numerical models of \\citet{M00} and \\citet{NS00} this appears consistent with the idea that earlier type spirals have a smaller fraction of their mass in the disk at a given rotational velocity, and implies that as well as being dimmer than later types at a given rotational velocity they have less concentrated halos.  Any TF relation fit to a sample with an explicit or implicit magnitude limit will suffer from biases relating to incompleteness. The largest such bias is caused by it being possible to observe galaxies which are scattered brighter than the TF relation from near the magnitude limit, while those scattered dimmer will not make it into the sample. We calculate and correct for this bias, as well as those relating to the finite depth of the clusters. All bias corrections applied are summarized in Section 3.6.  Individual TF relations are fit to each of the 31 clusters samples and presented in Section \\ref{indivTF}. We search for environmental dependence of the TF relation by looking for correlations in global cluster properties with the slope, zeropoint and scatter. No correlations are found. There is also no dependence of TF residuals on projected cluster center difference for galaxies in the sample. We therefore argue that there is no evidence for environmental dependence of the TF relation. It is noted however that because of the morphology--density relation we are already implicitly correcting for an environmental dependence by applying morphological type corrections.  The individual cluster samples are combined in Section \\ref{globalfit} to create a global TF template. In order to do this, peculiar velocities are derived for the clusters and a cluster velocity dispersion of $\\sigma = 298\\pm 34$ \\kms ~is measured. This measurement provides information on cosmological parameters as it depends on the initial spectrum of density fluctuations. The value is completely independent, yet within $2\\sigma$ of that predicted from the best fit cosmological model to the third year WMAP data release which implies $\\sigma = 239\\pm23$ \\kms \\citep{Sp06}. With minimal assumptions about cosmology we use the measurement to estimate $\\Omega^{0.6} \\sigma_8 = 0.52\\pm0.06$ from TF data alone.  Direct, inverse and bivariate linear fits to the combined {\\bf in} and {\\bf in+} samples (and subsets of them) are presented. Quadratic fits do not reduce the observed scatter and so are not reported. We favor the bivariate fit to the {\\bf in+} sample, which gives a template TF relation of $M - 5 \\log h = -20.85 - 7.85 (\\log W - 2.5)$, or $L \\propto v^{3.1}$. The statistical error on the zeropoint is $\\sim 0.02$ mag, an additional 0.03 mag is estimated to account for the possible net motion of the cluster reference frame in the CMB. The statistical error on the slope is $\\sim 0.10$ mag$^{-1}$. The TF relation obtained here will be applied to the rest of the galaxies in the SFI++ sample in future papers including Masters \\etal ~(2007, in prep.) and Springob \\etal ~(2007) to study the local peculiar velocity field.  The scatter in the TF relation is as important as the relation itself, both in models of galaxy formation and evolution and in deriving reliable distance errors.  We measure the total incompleteness bias corrected scatter from the template relation to be $\\epsilon = 0.41 - 0.44 (\\log W - 2.5)$. Once measurement errors have been accounted for this results in an intrinsic TF scatter of  $\\epsilon = 0.35 - 0.37 (\\log W - 2.5)$. This scatter corresponds to distance errors of 10\\% for the largest width galaxies increasing to 26\\% for the smallest widths.  As an independent measure of the zeropoint, that does not rely on the assumption that a subset of the most distant clusters are at rest with the CMB, we use 17 galaxies in the SFI++ that also have published distances from Cepheid variables. We find a negligible mean offset in the zeropoint of the TF relation for these galaxies relative to the ``basket of clusters\" bivariate fit to the \\inp~ sample when $H_0=74\\pm2$ \\kmsMpc ~is assumed, providing a reassuring check that the ``basket of clusters\" method provides a reliable measure of the TF zeropoint. Turning this method around to estimate $H_0$ from a combination of the template and Cepheid distances we find $H_0=74\\pm2\\pm6$ \\kmsMpc. The dominant source of error on this measurement comes not from TF directly, but rather from the uncertainty on the distance to the LMC (via the Cepheid calibration).  The TF relation still has a place in the era of ``precision cosmology'' to provide independent checks of the best fit cosmological model. It has the potential to measure some parameters (notably the cluster peculiar velocity dispersion and $H_0$) to an accuracy comparable to the best available. In addition the TF relation and its scatter provides valuable information about the formation of disk galaxies. However, the reliability of any conclusion drawn from TF depends critically on the availability of an unbiased template relation, such as the one provided here."
    },
    "0609/astro-ph0609763_arXiv.txt": {
        "abstract": "Nearby dwarf galaxies exhibit tight correlations between their global stellar and dynamical properties, such as circular velocity, mass-to-light ratio, stellar mass, surface brightness, and metallicity.  Such correlations have often been attributed to gas or metal-rich outflows driven by supernova energy feedback to the interstellar medium.  We use high-resolution cosmological simulations of high-redshift galaxies with and without energy feedback, as well as analytic modeling, to investigate whether the observed correlations can arise without supernova-driven outflows.  We find that the simulated dwarf galaxies exhibit correlations similar to those observed as early as $z\\approx 10$, regardless of whether supernova feedback is included.  We also show that the correlations can be well reproduced by our analytic model that accounts for realistic gas inflow but assumes no outflows, and star formation rate obeying the Kennicutt-Schmidt law with a critical density threshold. We argue that correlations in simulated galaxies arise due to the increasingly inefficient conversion of gas into stars in low-mass dwarf galaxies rather than supernova-driven outflows. We also show that the decrease of the observed effective yield in low-mass objects, often used as an indicator of gas and metal outflows, can be reasonably reproduced in our simulations without outflows.  We show that this trend can arise if a significant fraction of metals in small galaxies is spread to the outer regions of the halo outside the stellar extent via mixing. In this case the effective yield can be significantly underestimated if only metals within the stellar radius are taken into account.  Measurements of gas metallicity in the outskirts of gaseous disks of dwarfs would thus provide a key test of such explanation. ",
        "introduction": "\\label{sec:intro}  Observational studies of galaxies have revealed the existence of tight correlations between their global stellar and dynamical properties, such as circular velocity, mass-to-light ratio, stellar mass, surface brightness, and metallicity. For example, scalings of the mean metallicity of a galaxy with other global properties such as its luminosity, stellar mass, or total mass, have been established for galaxies of a large range of masses and morphologies (e.g., \\citealp{Leq79, Gar87, Zar94, Garn02, PB02, DW03, Trem04, Pil04, simon_etal06, lee06}). Theoretical modeling of these correlations can help us to identify key physical processes shaping global properties of galaxies.  Supernova energy feedback to the interstellar medium (ISM) with putative associated gas outflows from galaxies has long been a favored mechanism to explain these correlations and other properties of low-mass galaxies \\citep[e.g.,][]{Lar74,DS86,arimoto_yoshii87}.  The strong effect of the SN feedback on properties of low-mass galaxies is also a standard assumption of the semi-analytic models of galaxy formation \\citep[e.g.,][]{lacey_etal93,KWG93,cole_etal94,somerville_primack99, benson_etal03,croton_etal06}. \\citet{DW03} argue, for example, that the correlations between circular velocity, stellar mass, metallicity, and stellar surface densities they find for the nearby dwarf galaxies can be reproduced with a simple semi-analytic model incorporating the effect of supernova feedback on the gas component of dwarf galaxies.   The existence of galactic winds has indeed been observationally established in star-forming galaxies at low and high redshifts \\citep[e.g.,][see \\citeauthor{veilleux_etal05} \\citeyear{veilleux_etal05} for a recent review]{Hetal00,Pett01,Metal02,strickland_etal04,martin05,ott_etal05}. An indirect evidence for metal loss in winds is the high-metallicity of the diffuse intergalactic medium in groups and clusters \\citep[e.g.,][]{renzini_etal93}. However, the extent to which the outflows affect the global properties of galaxies and whether the outflow gas escapes the gravitational potential of its host halo or rains back down to the disk remains uncertain. \\citet{maclow_ferrara99} and \\citet{dercole_brighenti99} used numerical simulations to show that global gas blowout is inefficient in all but the smallest mass halos \\citep[$M\\lesssim 10^7{\\rm\\ M_{\\odot}}$; see also][]{ferrara_tolstoy00,marcolini_etal06}, although metal-enriched SN ejecta may be removed efficiently. It is also uncertain whether correlations of metallicity with galaxy stellar mass and mass-to-light ratios can be attributed solely to such winds.  In the context of cosmological simulations, the SN feedback is usually ineffective, unless \\emph{ad hoc} phenomenological recipes enhancing SN feedback are employed \\citep[e.g.,][]{navarro_white93,navarro_steinmetz97, thacker_couchman01,marri_white03,springel_hernquist03,scannapieco_etal06}. It is likely that this inefficiency is at least partially due to the inability of current cosmological simulations to resolve the scales and processes relevant to stellar feedback.  However, it is also possible that the actual effects of feedback on the global properties of galaxies are fairly small in reality.  Cosmological simulations at this point cannot make \\emph{ab initio} predictions about the importance or inefficiency of stellar feedback.  Observations, on the other hand, although showing evidence for the presence of large-scale winds in actively starforming galaxies \\citep[e.g., see][for review]{veilleux_etal05}, are rather uncertain in their estimates of wind mass loss to provide reliable direct observationally-motivated feedback recipes for implementation in simulations. Nevertheless, we can gauge the importance of supernova energy injection by comparing galaxies formed in simulations performed with different assumptions about feedback to observations.  In this respect, the smallest dwarf galaxies present the ideal test case because they can be expected to be the most susceptible to the effects of SN feedback due to their shallow potential wells \\citep{Lar74,DS86}.  The role of feedback in the formation of dwarf galaxies has been investigated in several recent studies.  \\cite{Tetal03} used Eulerian AMR simulations of galaxy formation with star formation and SN feedback and found that correlations between mass-to-light ratio and metallicity and stellar mass and metallicity, similar to those observed for the nearby dwarfs, are exhibited by dwarf galaxies in their simulations at $z \\gtrsim 3$.  The simulations of \\citet{Tetal03} used a rather extreme amount of energy per supernovae to maximize the effects of feedback. They did not test, however, whether feedback or some other mechanism is in fact the dominant factor in shaping the correlations.  More recently, \\cite[][]{Kob06} found a tight relation between stellar metallicity and stellar mass at all redshifts in their Smooth Particle Hydrodynamics (SPH) cosmological galaxy formation simulations, which they attributed to the mass-dependent galactic winds. The winds in the model of \\citet[][]{Kob06} do affect the gas and metal content of galaxies significantly, with the effect increasing towards lower mass systems (see their Fig.~16), which plays an important role in shaping the resulting correlation of metallicity and stellar mass.  \\citet{derossi_etal06} have recently studied the origin of luminosity- and stellar mass-metallicity relations in cosmological simulations. They find that correlations similar to those observed arise already at high redshifts, although supernova feedback in their simulations is not efficient. This is consistent with our results presented below. \\cite{simon_etal06} seeked observational correlations between metallicity, surface brightness and mass-to-light ratio within a single galaxy (M33) and found that indeed such correlations exist. They used SPH simulations and analytical calculations to show that feedback is not responsible for these correlations. They found that the presence of feedback decreases the scatter in such relations, although this may at least in part be due to their numerical implementation of feedback. Most recently, \\citet{brooks_etal06} presented a study of the origin and evolution of the mass-metallicity relation in cosmological simulations. They also concluded that the relation arises not due to the mass loss in winds, but due to increasing inefficiency of star formation in smaller mass galaxies. However, they argue that low effective yields of dwarf galaxies can only be explained by the loss of metals in winds. We argue below that there is an alternative explanation for the low effective yields of dwarfs.  In this paper we investigate the origin of the observed mass-metallicity and other correlations of dwarf galaxies with the specific goal of testing the role of SN feedback. In particular, we discuss results of several high-resolution cosmological simulations of galaxy formation started from the same initial conditions and run with the same prescriptions for star formation and metal enrichment of the ISM, but with different models for gas cooling, UV heating (optically thin approximation vs. self-consistent radiative transfer of the UV radiation), and supernova feedback (see \\S~\\ref{sims} for details). We follow the simulations until $z\\sim 3$.  Although our models simulate galaxies at high redshifts, we study the evolution of the correlations with redshift from $z=9$ to $z\\sim 3$ and show that the trends with metallicity and stellar mass are approximately preserved as the galaxy population evolves allowing extrapolation of the results to the present epoch. Therefore, our results and conclusions give us insight to the processes responsible for the properties of observed low-redshift dwarfs. There is indeed observational indication that the stellar mass-metallicity correlation, for example, is already established by $z\\sim 1-2$ \\citep{kobulnicky_kewley04,savaglio_etal05,Erb06}, although the metallicities of galaxies of a given stellar mass appear to be lower by $\\approx 0.3$~dex at $z\\sim 2$ compared to $z=0$ \\citep{savaglio_etal05,Erb06}. In addition, for some of the Local Group dwarf galaxies that exhibit these correlations, a significant fraction of their stellar mass was built at $z\\gtrsim 3$ \\citep[e.g.,][]{dolphin_etal05}.  The correlations for these dwarfs should thus be largely set at high redshifts, which we probe with our simulations.  Our main result is that correlations similar to those observed can be reproduced in all runs, which suggests that supernova energy feedback is not required to explain the observed properties of dwarfs.  To interpret the simulation results we use a simple open-box analytic model for the evolution of the baryonic component of dwarf galaxies. We show that a model with gas accretion but no outflows and the Kennicutt law for star formation with the critical density threshold for star formation reproduces the results of our simulations.  This paper is organized as follows. In \\S \\ref{obs} we summarize the observationally established correlations between global properties of dwarf galaxies that we wish to interpret. Our analytic model and chemical evolution calculations are described in \\S \\ref{ce}. In \\S \\ref{sims} we describe in detail our cosmological simulations, and in \\S \\ref{res} we discuss how global properties of the simulated galaxies are correlated with each other, and we compare these correlations to the observed ones, and to the analytic model of \\S \\ref{ce}.  A discussion of observational results on the effective yield and corresponding constraints on the existence of outflows is given in \\S \\ref{syeff}.  We discuss our findings in \\S \\ref{disc}, and we summarize our conclusions in \\S \\ref{sum}. ",
        "conclusions": "\\label{sum}  In this paper, we use a suite of high-resolution cosmological simulations to examine the role of different physical processes in establishing the observed scaling relations of dwarf galaxies. All simulations included a recipe for star formation and metal enrichment of the ISM. Three simulations (FEC, FNEC-RT, F2NEC-RT) additionally included energy feedback from supernovae to the ISM. Simulations FNEC-RT and F2NEC-RT followed detailed 3-D radiative transfer from individual UV sources and the chemistry network of ionic species of hydrogen, helium, and molecular hydrogen to determine the radiative heating and cooling of the baryonic gas, while the other two simulations used tabulated heating and cooling rates and a uniform ionizing background.  In order to interpret our simulation results, we develop and use an analytic model, which follows the evolution of the dynamical mass, gas, and stellar components as well as the metallicity of galaxies. We assumed an open box model with no outflows but with accretion of mass at a rate proportional to the galaxy mass, and star formation following the Kennicutt law with a critical density threshold for star formation.  Our main results and conclusions can be summarized as follows. \\begin{itemize}  \\item In all simulations, correlations between global quantities, such as the stellar mass-metallicity relation, very similar to those observed for the nearby dwarf galaxies arise in simulated galaxies with stellar masses $10^5 \\msol \\lesssim \\ms \\lesssim 10^8 \\msol$ for $z \\gtrsim 3.3$.  \\item In galaxies of higher masses these correlations exhibit a break and the dependence of global properties on stellar mass flattens off. A similar flattening of the correlations is also found in observational data, although the location of the break may have evolved to higher masses by the present epoch (e.g., \\citealp{Trem04}).  \\item The key finding of this study is that neither the inclusion of supernova energy feedback, nor the inclusion of 3-D radiative transfer significantly affects any of these correlations, a strong indication that thermal and radiative processes (as well as associated galactic outflows and winds) are not required for the observed correlations to arise.  Using our analytic model, we show that these correlations in the simulated galaxies arise due the increasing inefficiency of star formation with decreasing mass of the object. The inefficiency of conversion of gas into stars and its dependence on mass are due to the critical gas column density threshold of the star formation law.  \\item We thus argue that trends similar to those exhibited by observed dwarf galaxies can arise without loss of gas and metals in winds, although winds may affect the exact form of the relation (normalization and slope) and its evolution.  \\item The observed trends of a decreasing effective yield with decreasing galaxy mass can be reproduced in our simulated galaxies reasonably well, and without winds, if the effective yield is calculated taking into account only the gas, stars, and metals that would be accessible to observations. We show that this trend is due to efficient mixing in the ISM, driven by disk instabilities, mergers, and possibly cold accretion flows. The efficient mixing redistributes the metals far beyond the stellar extent in smaller mass galaxies. Thus a significant fraction of metals in small galaxies is outside the stellar extent leading to a decrease of the estimated yield if only metals within the stellar radius are taken into account. This explanation can potentially be tested by measurements of the gas metallicity in the outskirts of the gaseous disks of dwarf galaxies. \\end{itemize}"
    },
    "0609/astro-ph0609280_arXiv.txt": {
        "abstract": "% {% Recent work suggested that the traditional picture of the corona above the quiet Sun being rooted in the magnetic concentrations of the chromospheric network alone is strongly questionable. } {% \\NEW{Building on that previous study we} explore the impact of magnetic configurations in the photosphere and the low corona on the magnetic connectivity from the network to the corona. \\NEW{Observational studies of this connectivity are often utilizing magnetic field extrapolations. However, it is open to which extent such extrapolations really represent the connectivity found on the Sun, as observations are not able to resolve all fine scale magnetic structures. The present numerical experiments aim at contributing to this question.} } {% We investigated random salt-and-pepper-type distributions of kilo-Gauss internetwork flux elements carrying some $10^{15}$ to $10^{17}$\\,Mx, which are hardly distinguishable by current observational techniques. \\NEW{These photospheric distributions are then extrapolated into the corona using different sets of boundary conditions at the bottom \\emph{and} the top. This allows us to investigate} the fraction of network flux which is connected to the corona, as well as the locations of those coronal regions which are connected to the network patches. } {% \\NEW{We find} that with current instrumentation one cannot really determine from observations, which regions on the quiet Sun surface, i.e.\\ in the network and internetwork, are connected to which parts of the corona through extrapolation techniques. \\NEW{Future spectro-polarimetric instruments, such as with Solar\\,B or {\\sc{Gregor}}, will provide a higher sensitivity, and studies like the present one could help to estimate to which extent one can then pinpoint the connection from the chromosphere to the corona.} } {% } ",
        "introduction": "The appearance of the large-scale magnetic field on the quiet Sun photosphere is dominated by the so called network structure, which is found at the boundaries of the super-granular cells. This magnetic network with a typical scale of 20\\,Mm is made up by concentrations each carrying a magnetic flux in the range of $10^{18}$ to $10^{19}$\\,Mx with field strengths of the order of kilo-Gauss \\citep{Schrijver+al:1997sp}. The region in-between these network structures, the internetwork, is not field-free \\citep{Livingston+Harvey:1975}, but shows a small average flux density appearing to be of the order of a few to 50\\,Mx/cm$^2$. Depending on the spatial resolution, polarimetric sensitivity and diagnostic technique (Hanle and Zeeman effect), authors find a weak volume filling component \\citep{Faurobert+al:2001}, or small concentrations with kilo-Gauss flux tubes \\citep[e.g.][]{Dominguez-Cerdena+al:2003,Sanchez-Almeida+al:2003}.   As there are flux tubes with high, kilo-Gauss field strengths to be found in the network as well as the internetwork, a better concept to distinguish between the network and internetwork would be the amount of flux carried by each individual flux elements. With some $10^{16}$ to $10^{17}$\\,Mx each internetwork (flux tube) element carries  about one to two orders of magnitude less magnetic flux than the network patches \\citep{Lin+Rimmele:1999,Socas-Navarro+Sanchez-Almeida:2002}.   Even though each internetwork patch carries a very small amount of flux as compared to the stronger network patches, the weak internetwork patches can contribute a significant amount to the total (unsigned) flux on the solar surface, because they are so abundant. In contrast to the old canopy concept \\citep[e.g.][]{Giovanelli:1980}, where the internetwork patches have been neglected, the abundant low flux but high field strength concentrations in the internetwork do interact with the strong network patches: a sizable fraction of the field lines originating from the internetwork patches do not close back to another internetwork patch very close by, but can connect either to the stronger network patches or to the corona.  This was first pointed out by \\cite{Schrijver+Title:2003} who constructed a network patch surrounded by many smaller internetwork patches by placing magnetic charges at a plane, performed a potential field extrapolation and then studied the connections from their simulated photosphere, i.e.\\ the distribution of point charges, to the corona. They found that, depending on the average absolute flux density, only part of the network flux is actually connected to the corona. At an internetwork average absolute flux density of about 20\\,Mx/cm$^2$, which is compatible with solar observations, they found that only half of the network flux connects into the corona, the other half connects to the surrounding internetwork. Consequently, half of the magnetic flux in the corona originates from the internetwork, which is in strong contrast to the traditional canopy picture.   In this paper we will go one step further than \\cite{Schrijver+Title:2003} and not only demonstrate that the coronal magnetic field to a large part originates in the internetwork, but also investigate the role of the (spatial) distribution of the magnetic field in the low corona and the photosphere, i.e.\\ at the upper and lower boundary, for the link between the photosphere and the corona. This investigation will show the limitations for the knowledge we might gain from magnetic field extrapolations on the connections from the network and internetwork to the corona. As it will turn out, this is basically through our imperfect ability to determine the photospheric magnetic field at small enough scales with sufficient polarimetric sensitivity.   The organization of this paper is as follows. In order to also employ a boundary condition at the top of the computational domain we briefly describe a Fourier transform technique for the  potential field extrapolation, the boundary conditions representing the photosphere and the low corona as well as how to derive the fraction of network flux reaching the corona (\\sect{method}). Based on these results for the extrapolations with different boundary conditions we first compare our results to previous work in \\sect{comparison} before we discuss our new findings on the impact of the low corona (\\sect{impact_ubc}) and the distribution of photospheric flux (\\sect{impact_ubc}) on the magnetic connection from the photosphere to the corona. We will conclude the paper with a discussion of the consequences of this study for further investigations of the chromosphere--corona coupling in \\sect{conclusions}. ",
        "conclusions": "\\label{conclusions}  In the present paper we investigated the role of the magnetic field distribution in the photosphere and the low corona for the mapping from the chromospheric magnetic network into the corona utilzing magnetic field extrapolations. The classic picture was that the field fans out from strong network concentrations  in the chromosphere creating an often called canopy \\citep{Gabriel:1976}, which could be seen as a more or less stable foundation of the coronal magnetism. Extending the work of \\cite{Schrijver+Title:2003}, who are challenging this scenario, we manipulate the boundary conditions of a potential field extrapolation at the top \\emph{and} bottom of a computational box stretching from the photosphere to the low corona at some 15\\,Mm height. This allows us to investigate the mapping from the photosphere into the corona in great detail.   The good news are that the photosphere--corona connectivity is very stable against changes of the top boundary, i.e.\\ the magnetic field in the lower boundary at heights of around 15 Mm. This is very reassuring when studying different regions on the Sun, such as coronal holes, active region patches or quiet Sun, where one might expect different magnetic field structures in the low corona: when investigating the connection e.g.\\ of coronal Doppler shifts and photospheric fields to study the driving of coronal dynamics, a magnetic field extrapolation from an observed high-quality magnetogram seems to be sufficient to identify the respective areas in the photosphere and the corona which are magnetically connected (cf.\\ \\sect{impact_ubc}).   However, the reliability of the magnetic connection derived from the field extrapolation hinges on the quality of the information about the weak small-scale magnetic field in the photosphere. Changes in the distribution of the weak internetwork magnetic flux patches, which are hardly observable with current instrumentation, can lead to dramatic differences in the connectivity from the photosphere into the chromosphere, qualitatively as well as quantitatively. Especially the mapping from the photosphere into the corona, i.e.\\ which photospheric regions are magnetically connected to which coronal areas, is quite sensitive to the distribution of weak internetwork flux. With current instrumentation only beginning to resolve the structure of the internetwork magnetic field, one has to be careful when drawing conclusions on magnetic connections from network patches to the corona, as one cannot really tell about the role of the internetwork (cf.\\ \\sect{impact_lbc}).   It will be exciting to investigate the polarimetric studies with new high-resolution instruments from space as well as from the ground, e.g.\\ the soon to be started Solar\\,B mission or the 1.5\\,m solar telescope {\\sc{Gregor}} with their sensitive full Stokes polarimeters. Such studies will be of great importance to better understand the driving of the corona through the underlying photospheric convection.          \\def\\apj       {ApJ}% \\def\\apjl      {ApJ}% \\def\\apjs      {ApJS}% \\def\\apss      {Astrophys. Space Sci.}% \\def\\jgr       {J. Geophys. Res.}% \\def\\nat       {Nature}% \\def\\solphys   {Solar Phys.}% \\def\\philtrans {Phil.\\ Trans.\\ Roy.\\ Soc.\\ Lond.}% \\def\\aap       {A\\&A}% \\def\\aaps      {A\\&A Suppl.}% \\def\\mnras     {Mon.\\ Not.\\ Roy.\\ Astron.\\ Soc.} \\def\\araa      {Ann.\\ Rev.\\ Astron.\\ Astrophys.} \\def\\naturwiss {Naturwiss.} \\def\\za        {Zeitschrift f.\\ Astrophys.} \\def\\an        {Astronomische Nachrichten} \\def\\pasp      {PASP}  {\\small"
    },
    "0609/astro-ph0609555_arXiv.txt": {
        "abstract": "Dusty debris disks around main-sequence stars are signposts for the existence of planetesimals and exoplanets.  From cross-correlating \\hip stars with the \\iras catalogs, we identify 146 stars within 120 pc of Earth that show excess emission at 60\\,$\\micron$.  This search took special precautions to avoid false positives. Our sample is reasonably well distributed from late B to early K-type stars, but it contains very few later type stars. Even though \\iras flew more than 20 years ago and many astronomers have cross-correlated its catalogs with stellar catalogs, we were still able to newly identify debris disks at as many as 33 main-sequence stars; of these, 32 are within 100~pc of Earth.  The power of an all-sky survey satellite like \\iras is evident when comparing our 33 new debris disks with the total of only 22 dusty debris disk stars detected first with the more sensitive, but pointed, satellite \\iso.  Our investigation focuses on the mass, dimensions, and evolution of dusty debris disks. ",
        "introduction": "Dusty debris disks that surround nearby main-sequence stars were first detected by the \\textit{Infrared Astronomical Satellite} (\\iras) in 1983.  These circumstellar disks were inferred from an infrared excess flux between 25 and 100\\,$\\micron$ many times brighter than expected from the stellar photosphere.  The IR excess was modeled by disk distributions that would absorb optical and ultraviolet flux from the host star and then isotropically radiate this energy at infrared wavelengths.  The first dusty debris disk was discovered around the bright main-sequence star Vega \\citep{aum84}; consequently a dusty disk around a main-sequence star is commonly referred to as the Vega phenomenon.  Numerous studies of T Tauri stars dating back many years indicate the characteristic timescale for the dispersal of a surrounding dusty, gaseous disk is a few million years.  Following dissipation of the gaseous component, the remaining dust can further dissipate during the following few million years via coagulation into large objects, Poynting-Robertson and stellar wind drag, radiation pressure, and collisional destruction (e.g., \\citealt{bac93,lag00,dom03,pla05}). Vega-like stars are, however, generally much older than 10 Myr, thus the observed dust should be of secondary origin, most likely replenished via collision and fragmentation of planetesimals.  Furthermore, the Vega-phenomenon overlaps with the important planetary system formation epochs in our solar system: giant gas planet formation within $\\sim$10 Myr, terrestrial planet formation within $\\sim$30 Myr, and the era of heavy bombardment in the inner solar system within $\\sim$600 Myr.  Therefore, studies of IR-excess stars can provide crucial information on extrasolar planetary formation and evolution.  During the past two decades, about two dozen papers have been published that describe \\iras, Infrared Space Observatory (\\iso), and \\spit \\textit{Space Telescope} searches for stars with excess IR emission ($\\S$~3; \\citealt{lag00,zuc01,dec03}).   These searches employed different techniques for cross-correlating IR and stellar sources with no consistent definition of an IR excess.  To date, several hundred main-sequence IR-excess stars have been reported in the literature including those that have an IR excess at 25\\,$\\micron$.  A major goal of debris disk research has been to characterize the temporal evolution of the quantity of dust present in the disks. Notwithstanding almost two decades of debris disk research using data from three IR satellites, a convincing assessment of this temporal evolution remains incomplete.  Such an assessment requires a large sample of stars and a reliable estimate of the dust mass and the age for each debris disk system.  False positive IR-excess stars due to the large beam size of \\iras and improper search or calibration techniques have contaminated some previous studies.  Such contamination of the debris disk population has not only plagued many follow-up observations from ground and/or space observatories but also precludes a global assessment of the distribution and evolution of the dust population.  If the IR excess is from a bona fide dust disk, then the best estimator of dust mass comes from submillimeter flux.  Unfortunately, submillimeter flux measurements are difficult, time-intensive observations.  A more readily accessible observable is $\\tau$, the ratio of excess infrared luminosity due to dust divided by the total energy output from a star. We compute values of $\\tau$ for each of the IR-excess stars presented in this paper.  In $\\S$~5.1 we discuss the relationship between submillimeter flux and $\\tau$ for those Vega-like stars for which both are known, and we derive our own relationship, which is used to predict a dust mass if both $\\tau$ and the dust disk radius are known.  Furthermore, estimation of stellar age is often troublesome since most nearby IR-excess stars are isolated field stars.  In order to obviate the shortcoming of stellar age estimation, several groups are using \\spit to search for IR-excess stars in nearby young stellar groups with well-determined ages (e.g., \\citealt{sta05}).  However, because the distance to all rich clusters is substantial (except for the Hyades), it remains difficult to obtain statistically significant results even with \\spit \\citep{sta05}.  Thus, the large, clean sample of relatively nearby field stars we discuss in the present paper can contribute in a statistically meaningful way to our understanding of the Vega phenomenon and its evolution with time. ",
        "conclusions": "The 1983 all-sky \\iras far-IR survey yielded a wealth of information about the properties of cool dust in orbit around main sequence stars.  However, notwithstanding decades of ground- and space-based follow-up projects including \\iso, as of 2004 when we began the research reported here, in our opinion, a consistent, convincing evolutionary picture of these dusty stars had not been published.  In particular, while various researchers had cross-correlated various stellar catalogs against the \\iras catalog, none had used the \\hip catalog.  Stellar distances and proper motions provided by the \\hip and Tycho catalogs yield information useful for establishing ages of dusty stars; reliable ages are essential if correct evolutionary sequences are to be deduced.  Also, as a consequence of the rather large \\iras beam-size and inadequate attention to elimination of background confusion, some previous stellar studies with \\iras have suffered from the inclusion of false positive far-IR-excess stars.  In the research reported here we have taken special effort to deduce stellar ages and to eliminate false positives.  Just as it is possible to deduce many properties of stellar clusters and associations even though some stars are mistakenly included as members, we trust that our Table~2 \\iras sample is clean enough that our conclusions will stand the test of time.  Nonetheless, because ages of nearby field stars are notoriously difficult to estimate accurately and because of limitations with the \\iras database, we recognize that some entries in the tables and figures presented in this paper will be in error.  \\iras was most effective for the study of luminous B- and A-type main-sequence stars.  In agreement with some earlier studies, we find that \\iras detected excess emission at 60\\,$\\micron$ from about 20\\% of nearby A-type stars. This percentage will certainly rise as the A stars are examined with far-IR photometers more sensitive than those aboard \\iras.  In particular, we find that about 10\\% of stars of various spectral classes are revealed to display far-IR dust emission at brightness levels between \\iras and \\spit sensitivities.  Although this 10\\% subsumes stellar age, spectral types, and distance from Earth and thus is potentially subject to selection effects, it is consistent with the well-defined TW Hydra association sample of \\citet{low05}.  Using heterogeneous samples, \\citet{smi06} and \\citet{bry06} also reported about 10\\% of stars show dust excess in the MIPS 70\\,$\\micron$ band, but below \\iras sensitivity.  From their analysis of \\iso data sets, especially the volume-limited sample of \\citet{hab01}, \\citet{dec03} deduced that the percentage of stars with detectable 60\\,$\\micron$ emission diminishes with age. However, the small data set of Habing et al. (2001) and difficulties with estimating stellar age, precluded a meaningful quantitative result in our opinion.  With our larger and more robust database we can derive that the probability of 60\\,$\\micron$ excess emission detectable with the sensitivity of \\iras is about 35 times larger for stars with ages in the range 10-50 Myr compared to such stars with ages $>$ 500 Myr within the volume within 120 pc of Earth.  While it is generally agreed that measurements at submillimeter wavelengths are best for the derivation of dust masses, by means of a simple model that relates submillimeter and far-IR fluxes, we are able to derive dust masses for numerous stars that lack submillimeter data. These masses lie in the range between 0.0005 and 0.5 \\textit{M}$_{\\earth}$. For stars with ages between 30 and 1000 Myr, these dust masses appear to depend little, if at all, on age.  Based on Figure 5, and as described in $\\S$~5.1, our model indicates that far-IR data can be used, quite reliably, to predict a submillimeter flux and, thus, a disk dust mass.  As a consequence, disk dust masses can generally be derived based solely on \\spit data provided that excess flux is measured at two or more well-separated wavelengths with MIPS and/or the Infrared Spectrograph (IRS).  Four Algol binary stars appear to display excess emission at 60\\,$\\micron$ wavelength, although the existence of the excess is perhaps not compelling in all cases.  We considered models in which the emission is generated by free-free and bound-free emission in orbiting ionized gas or by orbiting dust particles, dust perhaps associated with a tertiary (third)  stellar component.  Future studies will be required to clarify the dominant physical mechanism(s) involved.  Additional results of our study include:  (1) Peak $\\tau$ ($\\sim10^{-3}$) does not vary much at all ages later than $\\sim$10\\,Myr; this might be because occasional catastrophic dust-generating events can occur at any age. (2) The spread of measured $\\tau$ at ages $\\sim$10 Myr is about a factor of 10, increasing to about 100 at later ages; given the measured peak $\\tau$ (item 1) and \\iras threshold ($\\sim10^{-5}$), the measured spread of $\\tau$ cannot be greater than 100. (3) At any given age late-type stars tend to have the largest $\\tau$.  (4) For stars with ages between 10 and 1000 Myr, the mean $\\tau$ of stars with \\iras detectable far-IR excess emission declines in proportion to (age)$^{0.7}$.   (5) For early-type stars between ages of $\\sim$10 and 100 Myr, the typical radius of a dusty debris disk appears to be smaller than for stars with ages between 100 Myr and 1 Gyr. (6) The very largest taus ($>$ $10^{-3}$) are associated only with disks that have relatively small radii.  (7) \\iras detected excess 60\\,$\\micron$ emission from $\\sim$20\\,\\% of nearby A-type stars.  (8) Four Algol-type eclipsing binaries, including Algol A itself, display 60\\,$\\micron$ emission, generated by free-free and bound-free transitions in ionized gas, by dust grains, or by both.  (9) Gl~803 (AU~Mic, 12 Myr old) is the only M-type, non-T Tauri, \\hip dwarf star to display 60\\,$\\micron$ excess emission in the \\iras Catalogs."
    },
    "0609/astro-ph0609139_arXiv.txt": {
        "abstract": "The bulge of our Galaxy is illuminated by the 0.511 MeV gamma-ray line flux from annihilations of nonrelativistic positrons. The emission is strongly concentrated at the Galactic Center (GC), in contrast to gamma-ray maps tracing nucleosynthesis (e.g., the 1.809 MeV line from decaying $^{26}$Al) or cosmic ray processes (e.g., the 1--30 MeV continuum), which reveal a bright disk with a much less prominent central region. If positrons are generated at relativistic energies, higher-energy gamma rays will also be produced from inflight annihilation of positrons on ambient electrons.  The comparison of the gamma-ray spectrum from inflight annihilation to the observed diffuse Galactic gamma-ray data constrains the injection energies of Galactic positrons to be less than $3$ MeV. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609413_arXiv.txt": {
        "abstract": "A commonly used perturbative method for computing large-scale clustering of tracers of mass density, like galaxies, is to model the tracer density field as a Taylor series in the local smoothed mass density fluctuations, possibly adding a stochastic component.  I suggest a set of parameter redefinitions, eliminating problematic perturbative correction terms, that should represent a modest improvement, at least, to this method.  As presented here, my method can be used to compute the power spectrum and bispectrum to 4th order in initial density perturbations, and higher order extensions should be straightforward. While the model is technically unchanged at this order, just reparameterized, the renormalized model is more elegant, and should have better convergence behavior, for three reasons:  First, in the usual approach the effects of beyond-linear-order bias parameters can be seen at asymptotically large scales, while after renormalization the linear model is preserved in the large-scale limit, i.e., the effects of higher order bias parameters are restricted to relatively high $k$.  Second, while the standard approach includes smoothing to suppress large perturbative correction terms, resulting in dependence on the arbitrary cutoff scale, no cutoff-sensitive terms appear explicitly after my redefinitions (and, relatedly, my correction terms are less sensitive to high-$k$, non-linear, power).  Third, the 3rd order bias parameter disappears entirely, so my model has one fewer free parameter than usual (this parameter was redundant at the order considered).  This model predicts a small modification of the baryonic acoustic oscillation (BAO) signal, in real space, supporting the robustness of BAO as a probe of dark energy, and providing a complete perturbative description over the relevant range of scales. ",
        "introduction": "In the past year, significant progress has been made in using perturbation theory (PT) to calculate the large-scale clustering of collisionless mass in the Universe, to the point where essentially perfect calculations of the quasi-linear clustering may soon be available \\cite{2006PhRvD..73f3520C,2006PhRvD..73f3519C, 2006astro.ph..6028M}. This motivates a fresh look at the bias models that are needed to couple these calculations to most observable tracers of large-scale mass density, e.g., galaxies \\citep{2004ApJ...606..702T,2006MNRAS.366..189S}, the \\lyaf\\ \\citep{2005ApJ...635..761M,2006ApJS..163...80M,2006MNRAS.365..231V}, galaxy cluster/Sunyaev-Zel'dovich effect (SZ) measurements \\citep{2005astro.ph.11060D}, and possibly future 21cm surveys \\citep{2005MNRAS.364..743N}.  Baryonic acoustic oscillation surveys aimed at probing dark energy \\citep{1998ApJ...504L..57E, 1998ApJ...496..605E,2001ApJ...557L...7C,2003astro.ph..1623E, 2003ApJ...594..665B,2003PhRvD..68h3504L,2003ApJ...598..720S, 2004ApJ...615..573M,2005ApJ...631....1G,2005astro.ph..7457G, 2005MNRAS.357..429A,2005MNRAS.363.1329B,2006MNRAS.365..255B, 2006MNRAS.366..884D,2006astro.ph..7122M,2006astro.ph..4075J}, in particular, fall in the range of scales where perturbation theory may be most useful.  Probably the most straightforward, commonly used, model for the bias when coupled to perturbation theory is to write the tracer density, $\\delta_g$, as a Taylor series in the mass density fluctuations, $\\delta$ \\citep{1993ApJ...413..447F,1998MNRAS.301..797H}, i.e., $\\delta_g = \\sum_{i} b_i \\delta^i/i!$. Note that I will often refer to the tracer as galaxies (this explains the subscript $g$), but essentially everything in the paper could be applied to any tracer. The terms in this series generally will not decrease in size, so $\\delta$ is usually taken to be a smoothed version of the density field, to reduce the size of its fluctuations.  When $\\delta$ is computed using perturbation theory for gravitational clustering \\citep{1980lssu.book.....P,1981MNRAS.197..931J,1983MNRAS.203..345V, 1984ApJ...279..499F,1986ApJ...311....6G,1994ApJ...431..495J, 1996ApJ...473..620S,2002PhR...367....1B}, clustering of the tracers is in principle fully described. \\citep{1998MNRAS.301..797H} showed that applying this approach to compute the power spectrum up to 4th order in the initial density perturbations (the lowest order that gives a correction to linear theory) leads to interesting, but not completely satisfactory results.  They found that the 2nd order bias term (i.e., $b_2$) produces a white ($k$-independent) contribution to the power on large scales.  Furthermore, both the 2nd and 3rd order bias parameters contribute terms to the effective large-scale bias, defined as the ratio of the galaxy to mass power.  These terms all depend on the scale of the smoothing applied to the density field.  Note that the most straightforward expectation is that the smoothing scale should be quite large, as one can see by considering the measured value of rms density fluctuations in $8\\hmpc$ radius spheres, $\\sigma_8\\simeq0.85$ \\citep{2006JCAP...10..014S}.  This level of smoothing dramatically affects the power on relatively large scales, e.g., an $8\\hmpc$ radius top-hat smoothing suppresses the power at $k=0.2\\ihmpc$ by almost 50\\%.  It may be intrinsically interesting, when studying galaxy formation, to determine the coefficients $b_i$ in the series $\\delta_g = \\sum_{i} b_i \\delta^i/i!$, and to see generally how well this model works, as a function of smoothing scale. However, I will suggest a cleaner repackaging of the model, for the specific purpose of large-scale structure phenomenology, which has the desirable properties of preserving the form of the linear theory model at very large scales (i.e., the power at small $k$ is described by a single linear theory bias parameter and a single shot-noise parameter), and being free of any explicit smoothing scales. The approach I take is more or less precisely the same as the one used in classic renormalization of quantum field theories.  When perturbation theory is pursued naively, infinities appear, but only in ways that can be reinterpreted as corrections to the values of parameters of the model (e.g., particle masses, analogous to our linear model parameters).  The infinities are removed by simply redefining the free parameter in the model to be the most directly observable quantity (e.g., the mass you would measure by weighing the particle, or the ratio of large-scale galaxy to mass power), which would be given by a series of terms if one continued to use the original parameterization \\citep{Peskin:1995ev}.  Section \\ref{seccalc} of this paper gives my primary calculations, followed by some discussion in section \\ref{secdiscuss}. Before proceeding I reiterate the main goals:  (1) To formulate the bias model in a way that does not require arbitrary smoothing scales, while simultaneously being insensitive to the small-scale, highly non-linear regime.  (2) To formulate the model in a way that preserves the linear theory model on very large scales (low-$k$), i.e., higher order bias terms only affect the relatively high-$k$ power. To be honest, I should also note here that, while I refer throughout the paper to the approach of \\citep{1998MNRAS.301..797H} as the ``usual'' or ``standard'' approach, this perturbative approach to the galaxy power spectrum (including beyond-linear corrections) has not to my knowledge actually been used to interpret real data.  However, between improvements in perturbation theory and the need to interpret increasingly precise observations, the time for this kind of approach may have arrived \\citep{2006astro.ph..4075J}. ",
        "conclusions": "}  The standard method for including biasing of tracers of mass density in perturbation theory calculations of large-scale clustering \\citep{1993ApJ...413..447F,1998MNRAS.301..797H} uses five parameters to model the power spectrum and bispectrum up to 4th order in density perturbations, and requires at least one arbitrary cutoff parameter (smoothing scale). In the standard approach the effects of the beyond-linear order bias parameters and the cutoff propagate to asymptotically large scales. (Admittedly, this ``standard'' method has not yet been used much for interpreting data, but this may change soon.) After my redefinitions we have only four parameters, with much more cleanly separated effects: \\begin{eqnarray*} {\\rm standard~ parameters} &\\longrightarrow & \\rho_0,~\\rho_0^\\prime,~\\rho_0^{\\prime\\prime},~\\rho_0^{\\prime\\prime\\prime},~ N_0,~{\\rm and~ messy~ cutoffs} \\nonumber \\\\ {\\rm new~ parameters} &\\longrightarrow & \\bar{\\rho}_g,~b_1,~\\tilde{b}_2,~{\\rm and}~ N~. \\end{eqnarray*} \\cite{1998MNRAS.301..797H} provided the core calculations for this paper. The innovation here is the suggestion that their effective linear bias and shot-noise-like term should not be treated as output predictions of the model but instead should be absorbed into renormalized linear bias and shot-noise parameters. Significantly, these new free parameters of the model contain most of the high-$k$ (small, non-linear scale) sensitivity in the bias calculation. Rather than having a large-scale effective bias (ratio of tracer power to mass power) which is a function of all three bias parameters in the model, and the smoothing scale, I have only a single parameter for the large-scale bias. Rather than shot-noise that is a function of the original stochasticity, the 2nd order bias, and the smoothing scale, I have a single shot-noise parameter that describes all asymptotically large-scale deviation from the linear bias model.  Quasi-linear deviation from the linear bias plus shot-noise model is described by a single parameter, with no need for smoothing.  One might ask if it is desirable to bury time and cosmology dependence, in the form of the values of $\\sigma^2$ and $\\int dq~ q^2 P(q)^2$, in $b_1$ and $N$. Do we not lose information this way? The answer is: not really, because $c_1$, $c_2$, $c_3$, and $N_0$ are also mixed into $b_1$ and $N$, and, even for fixed smoothing scale, generally none of these parameters will be independent of the cosmological model or of time.  Furthermore, $\\sigma^2$ and $\\int dq~ q^2 P(q)^2$ are relatively sensitive to high-$k$ power which is poorly described by perturbation theory.  Note that simulations or other microscopic models of galaxy formation can still be used to calibrate the renormalized large-scale structure model.  The way this works exemplifies the usefulness of the revised model.  The most straightforward method for calibrating the unrenormalized model is to pick some smoothing scale and measure directly the parameters of the original Taylor series, Eq. (\\ref{eqtaylor}).  If this is done very literally (e.g., by making a scatter plot of galaxy density vs. smoothed mass density and reading off derivatives at $\\delta=0$), one is not even guaranteed that the calibrated model will reproduce the mean density of galaxies in the simulations used to calibrate it (without the trivial renormalization of the first term in the Taylor series, discussed above), let alone the large-scale bias, the shot-noise, or really anything else about large-scale structure.  (I emphasize here that I am not saying this method will surely fail, only that it is not guaranteed to work.) To calibrate the renormalized model one computes the mean density and power spectrum from the simulation, and if desired the bispectrum, and then fits for the parameters of the model. This guarantees that the large-scale structure in the simulation will be reproduced by the model as well as possible.  Of course, one can be less literal and calibrate the unrenormalized model by fitting the simulated statistics, but this still includes the extra steps of choosing a smoothing scale and computing the simple observables (e.g., mean density, large scale bias) as a function of several parameters.  Of course, as I discussed above, there is an additional problem with the literal use Eq. (\\ref{eqtaylor}) in that any smoothing aggressive enough to render the density fluctuations small will also corrupt the power on interesting, perturbative scales.  Within the perturbative approach there is a simulation-independent means of estimating the validity of the results.  One can go through the full process of extracting the cosmological parameters of most fundamental interest from the data, marginalizing over the bias parameters, using the perturbative model calculated to two different orders, and compare the results.  As presented here, the results from linear theory could be compared to the model with the first non-linear corrections, although computing to another order might be useful in the future (of course, higher order is necessary if one wants to use the bispectrum to two different non-zero orders). The level of agreement would surely depend on the maximum $k$ used in the fit.  This maximum $k$ could be chosen small enough to ensure agreement, i.e., consistency of the perturbative model.  There are several obvious next steps for this line of work: Extension to redshift space \\cite{2003ApJ...585...34M,2004PhRvD..70h3007S} should be straightforward, following the calculations of \\cite{1998MNRAS.301..797H}. A more general time and space-local model could be implemented, e.g., once one goes beyond linear theory the Taylor series for $\\rho_g$ should probably include velocity divergence terms as well as density terms. As discussed by \\cite{2006astro.ph..6028M}, it may be possible to make non-trivial predictions about the time evolution of bias by starting the perturbative calculation from a model for the local formation (and merging, etc.) rate of the tracer \\citep{1998ApJ...500L..79T,2005A&A...430..827S}, rather than the instantaneous density of the tracer (this will probably require a renormalization group approach as in \\cite{2006astro.ph..6028M}, rather than simple parameter redefinitions as in this paper). Extension to multiple luminosities or types of galaxies (or more divers tracers \\cite{2000ApJ...537...37T, 2002ApJ...580...42M,2006MNRAS.369...68S}) can be done by writing a Taylor series like Eq. (\\ref{eqtaylor}) for each.  This will lead to extra parameters, but cross-correlation will provide extra observables to constrain them.  Finally, I observe that, beyond just removing badly behaved perturbative terms at the present order of calculation, by treating the linear bias and shot-noise as free parameters we are in effect already accounting for similar terms that appear at any order in perturbation theory. This is probably the most compelling reason to believe that the approach of this paper can be expected work better than the standard approach as presented (the problem in the standard approach of smoothing corrupting the power on interesting scales is also significant). It may be possible to show that all higher order corrections to the tracer power as $k\\rightarrow 0$ can be included in these two terms, i.e., that the renormalized linear bias and shot-noise give a fundamentally complete descriptions of very large-scale clustering \\cite{1998ApJ...504..607S}."
    },
    "0609/astro-ph0609226.txt": {
        "abstract": "\\abstract{} {We study the EUV/soft X-ray emission generated by charge transfer between solar wind heavy ions and interstellar neutral atoms and variations of the X-ray intensities and spectra with the line of sight direction, the observer location, the solar cycle phase and the solar wind anisotropies, and a temporary enhancement of the solar wind similar to the event observed by Snowden et al. (2004) during the XMM-Hubble Deep Field North exposure.} {Using recent observations of the neutral atoms combined with updated cross-sections and cascading photon spectra we have computed self-consistent distributions of interstellar hydrogen, helium and highly charged solar wind ions for a stationary solar wind and we have constructed monochromatic emission maps and spectra. We have evaluated separately the contribution of the heliosheath and heliotail, and included X-ray emission of the excited solar wind ions produced in sequential collisions to the signal.} {In most practicable observations, the low and medium latitude X-ray emission is significantly higher at minimum activity than at maximum, especially around December. This occurs due to a strong depletion of neutrals during the high activity phase, which is not compensated by an increase of the solar wind flux. For high latitudes the emission depends on the ion species in a complex way. Intensity maps are in general significantly different for observations separated by six-month intervals. Secondary ions are found to make a negligible contribution to the X-ray line of sight intensities, because their density becomes significant only at large distances. The contribution of the heliosheath-heliotail is always smaller than 5\\%.  We can reproduce both the intensity range and the temporal variation of the XMM-HDFN emission lines in the 0.52-0.75 keV interval, using a simple enhanced solar wind spiral stream. This suggests a dominant heliospheric origin for these lines, before, during and also after the event.}{} % ",
        "introduction": "That objects as cold as comets emit copious X-rays was a surprising discovery (Lisse et al, 1996). The emission mechanism, first proposed by Cravens (1997), is now convincingly demonstrated as a major source of cometary and planetary X-rays. Charge exchange (CX) collisions of highly charged ions of the solar wind (SW) with the neutral atoms and molecules from the coma produce cascades of photons in the extreme ultraviolet (EUV) and soft X-ray regions (H\\\"{a}berli et al. 1997, Krasnopolsky, 1997, Kharchenko and Dalgarno, 2000, Schwadron and Cravens, 2002).  Cox (1998) pointed out that neutral interstellar (hereafter IS) and geocoronal atoms would also undergo CX with the SW ions, and generate soft X-rays throughout the heliosphere. The heliospheric  emission has been subsequently modelled by Cravens et al (2001) and Robertson et al (2001), who explained satisfactorily the global correlation observed by Freyberg (1994, 1998) between  the so-called Long Term Enhancements (LTE's) measured by the ROSAT satellite  and strong solar wind events. The geocoronal emission has been investigated by Robertson and Cravens (2003) and measured during Chandra observations of the dark moon (Wargelin et al, 2004). Earlier, Gruntman (1998) modelled the EUV emission produced by the recombination of alpha particles of the SW colliding with interstellar atoms, and showed the substantial influence of the solar wind characteristics on the emission.  The intensity of the soft X-rays in the heliosphere was first estimated by Cravens (2000), who found it to be of the same order as the soft X-ray emission from the so-called Local Bubble (hereafter LB), the 100 parsecs wide cavity surrounding the Sun, filled with tenuous hot (10$^6$ K) gas (Snowden et al, 1994, 1999). More generally, for faint and diffuse astronomical X-ray sources the CX heliospheric emission may contaminate significantly the X-ray spectra below 1.5 keV as has been demonstrated by a long duration XMM-Newton exposure towards the Hubble Deep Field-North (Snowden et al, 2004, hereafter SCK04).  The cosmic diffuse X- ray emission is the sum of extragalactic, halo and LB contributions (Kuntz and Snowden, 2000). It is difficult to separate out the CX emissions from those of the Local Bubble (LB). Shadowing is of limited use because of the absence of clouds of sufficient density and emission lines dominate the continuum in the CX spectrum and the LB spectrum.  As seen from the Sun, heliospheric CX emission maxima will be aligned along the interstellar wind axis (Cox, 1998). The ROSAT 6-month survey necessary to reconstruct full sky soft X-ray maps did not show this alignment so that low upper limits to the local emission were derived. However, it has been demonstrated that parallax effects connected with the ROSAT all-sky survey geometry destroy the axial symmetry and cause a much smoother emission pattern, so that substantial contamination by CX is not excluded (Lallement, 2004).  Robertson and Cravens (2003a,b) constructed sky maps of the heliospheric and geocoronal CX intensity, using global efficiency factors representing all ions at once. Pepino et al. (2004) carried out similar calculations of the spectra emitted by the individual ions and of the corresponding power densities taking separate account of collisions of hydrogen and helium and distinguishing between the fast and slow solar winds.  We extend the study to the calculation of line of sight (LOS) spectral emission maps. We explore potential sources of variability of the received signals and quantify the resulting intensities. We take into account the observer location, the solar cycle phase, the LOS and we calculate the intensities of the strongest emission lines and the contributions of the outer heliosphere. We investigate the contribution of secondary ions created sequentially by CX. We examine the effects of temporary solar wind enhancements and solar rotation (Cravens 2000) using a realistic model of the expanding solar wind. Observations of the soft diffuse X-ray background with the XMM - Newton telescope are analysed.  In section 2 we describe the model developed for the calculation of the X-ray emission. In section 3, we present the heliospheric X-ray emission maps derived from the model and discuss the effects of the solar cycle, observing location on the emission pattern. In section \\ref{secCE_section}\\ we evaluate the contribution of the secondary ions and in \\ref{outerHelio_section} the contributions of the heliosheath and heliotail. In section \\ref{helio_spectra}, we present X-ray and EUV spectra for lines of sight at different helio-ecliptic latitudes, corresponding to slow and fast solar winds, and for neutral gas distributions that depend on solar activity conditions. In section \\ref{Snow_mod} we simulate the temporal variations of the heliospheric X-ray emission generated by a solar impulsive event for solar wind conditions and a geometry corresponding to the XMM-Newton Hubble Deep Field-North (HDFN) exposure of June 1, 2001. The model results are compared with the XMM-Newton observations of the diffuse background spectra by SCK04. Finally, in the last section \\ref{conclusions}, we summarize and discuss the results. ",
        "conclusions": "\\label{conclusions} We have performed a parametric study of the heliospheric X-ray background emission and an analysis of the multiple factors which can influence the emission level. We have demonstrated the complex response to these factors, which include the date of observation (through the difference between the earth longitude and the interstellar wind axis longitude), the solar cycle phase (through the H and He densities which vary significantly with the phase), the solar wind type (through the high ion absolute and relative abundances), and finally the solar wind history (through the potentially very strong solar wind enhancements and abundance variations). The contribution to the signal generated in the inner heliosphere by secondary collisions has been shown to be negligible. The contribution to the signal of the external heliosphere (heliosheath and tail) is estimated and found to be small, except in the tail direction for some particular ions. This type of model results should allow an estimate of the expected range for the emission for a given direction of sight and date.  We have modelled in a very simplified way the temporal variation of the signal in the case of a sudden and temporary solar wind increase, and made a data/model comparison in the case of the XMM-HDFN.  It is possible to reproduce the intensity level and the temporal evolution of the line intensities, although no solution has been found here matching perfectly all the data simultaneously. It is likely however from this study that a large fraction of the signal is of heliospheric origin, including in the post-event (LOW) period of time. This is in agreement with the results of Cravens et al (2001), based on correlations with the solar wind, and of Lallement (2004) based on the Local Bubble emission pattern.  It is clear from figure \\ref{Snowgeo} that each specific geometry will produce a different temporal variation. A LOS at 180 degrees from the HDFN direction would have had a totally different light curve, for the same solar wind history. Such complex temporal variations preclude a good correlation on small time scales between the X-ray emission and the locally measured solar wind. On the other hand, this diversity can be useful because the combination of numerous examples such as the above one may allow constraints on the actual local and cosmic contributions.  To reach this goal a number of improvements in the modelling have to be worked out. We have been using here a single fluid model for the IS H density, but we know that H atoms are separated in the outer heliosphere in two populations with different temperature, velocity and arrival direction (e.g. Lallement et al., 2005). Two population models should be developed. More important, in our simulation the eruptive region is stationary in the Sun's frame, like a garden hose, and apart from one enhanced stream we have only used simplified stationary states of solar cycle minimum and maximum conditions.  Realistic models should include 3D solar wind hourly data (ideally from different vantage points) as well as precise measurements of the high ion relative abundances. The exact 3D observation geometry should be, as well, taken into account in the model, which requires the knowledge of the active region geometry. In parallel, a more precise time-dependent model of the neutral distributions, consistent with the solar wind flux and velocity variations is needed.\\\\\\\\ \\small{\\textit{Acknowledgements.} We wish to thank John Raymond for useful discussions and valuable suggestions which led to the correction of a mistake.\\\\ R.L. and D.K., V.I. acknowledge funding by CNRS and RFBR respectively under PICS contract 3205.\\\\ A.D. and V.K. acknowledge NASA for support through grant NNG04GD57G.}\\small"
    },
    "0609/astro-ph0609631_arXiv.txt": {
        "abstract": "Type Ia supernovae (SNe Ia) are one of the major tools to determine the cosmological parameters. Utilizing them as distance indicators, it is possible to geometrically survey the universe. To this end, the intrinsic scatter in the luminosities of these events needs to be calibrated using empirical relations between observables. A theoretical explanation for these relations is still lacking and can only be provided by a sound understanding of the mechanism of SNe Ia. Recently there has been significant progress in modeling SNe Ia. We report on numerical simulations of the explosion process, compare the results with observations of nearby SNe Ia, and discuss current uncertainties of the models. The presented simulations shed some light on the origin of the diversity of SNe Ia. Such simulations will pave the way towards an understanding of SN Ia diversities and correlations of their properties and ultimately provide a tool to validate the cosmological implications of SN Ia distance measurements. ",
        "introduction": "Almost 80 years after Einstein's attempts to derive a cosmological model from his theory of General Relativity, subsequently introducing and discarding the cosmological constant, two cosmological surveys \\citep{riess1998a, perlmutter1999a} of the universe based on distance measurements of type Ia supernovae (SNe Ia) revolutionized the cosmological paradigm. They found that the Universe is currently in a phase of \\emph{accelerated} expansion. This revived the discussion of the cosmological constant. Independent determinations of cosmological parameters based on anisotropies in the Cosmic Microwave Background radiations and on large-scale galaxy surveys provided independent confirmation of the SN Ia measurements.  Yet the question of the applicability of SNe Ia as distance indicators is still not satisfactory answered. SNe Ia are remarkably uniform events by astrophysical standards,but evidently no standard candles. Only the calibration of the distance measurements according to empirical correlations between observables provides the necessary accuracy for the determination of cosmological parameters. A firm theoretical reasoning of such correlations is, however, still lacking.  It is still an open question whether the accelerated expansion of the Universe is caused by a simple cosmological constant or some more complicated constituent. Therefore one of the most fascinating problems of cosmology is to determine the equation of state of this ``Dark Energy'' component. SNe Ia may be a suitable tool for this task, but the accuracy of measurements required here can only be reached if systematical errors that may arise from calibration of the distance measurements are under control. A sound understanding of the physical mechanism of SNe Ia is inevitable in this project.  This poses a great challenge to SNe Ia theory since it requires self-consistent parameter-free modeling so that the predictive power of the models is sufficient to tackle questions of the origin of diversities and correlations.  Consequently, as a first step, an astrophysical model has to be built and implemented into a numerical simulation code. As will be explained below, the nature of relevant physical effects requires three-dimensional simulations. This approach is notably different from earlier one-dimensional SN Ia models which parametrized physical effects and therefore could be tuned to fit the observational data. Although this approach provided valuable hints to underlying physical effects, it cannot be used to constrain the correlations. Three-dimensional models have been implemented over the past years by different groups on the basis of different assumptions \\citep[e.g.][]{hillebrandt2000b,reinecke2002d,gamezo2003a, calder2004a}. Most of them led to viable explosions and the task is now to compare the results of these simulations with observations of nearby SNe Ia. Once a model matching the observational requirements has been established, it can be used to explore the diversity arising from variations of physical parameters as suggested by progenitor evolution and environmental effects. The results of such a survey should then be analyzed to identify correlations in the characteristics. This may eventually provide a theoretical basis for distance measurement calibrations and a handle on possible systematical errors in those. ",
        "conclusions": "We discussed the mechanism of SN Ia explosions and reviewed the common astrophysical scenarios applied to explain these events. The implementation of such models into numerical simulations was described focusing on the pure deflagration model for SNe Ia. Subsequently, recent progress in numerical simulations was reported. On the basis of comparison of these results with observations, the conclusion is drawn, that pure deflagration models are at least capable of reproducing the gross features of observed SNe Ia.  Such models therefore can be applied to study the origin of the diversity of SNe Ia. In a first systematic survey of initial parameters of three-dimensional models, the effects of the WD's carbon-to-oxygen ratio, its central density and the progenitor's metallicity on the explosion process were analyzed. Although this was only possible on the basis of simplified and therefore weakly exploding models, clear trends could be identified.  From this set of models, synthetic light curves were derived. The relation between the peak-luminosities and the decline rates of the light curves were found to be consistent with the slope predicted by \\citet{phillips1999a}. Due to the simplicity of the applied models, the absolute magnitude of effects could not be predicted.  Forthcoming studies will focus on the questions how a realistic stellar evolution of the progenitor interrelates the effects. A second important question is whether the range of scatter predicted by the observations can be reproduced with more elaborate models. Although it captures main aspects of observations, it may be possible that the applied SN Ia model is still incomplete and additional physical effects (like delayed detonations) need to be taken into account."
    },
    "0609/astro-ph0609588_arXiv.txt": {
        "abstract": "We describe an automated method for assigning the most probable physical parameters to the components of an eclipsing binary, using only its photometric light curve and combined colors. With traditional methods, one attempts to optimize a multi-parameter model over many iterations, so as to minimize the chi-squared value. We suggest an alternative method, where one selects pairs of coeval stars from a set of theoretical stellar models, and compares their simulated light curves and combined colors with the observations. This approach greatly reduces the parameter space over which one needs to search, and allows one to estimate the components' masses, radii and absolute magnitudes, without spectroscopic data. We have implemented this method in an automated program using published theoretical isochrones and limb-darkening coefficients. Since it is easy to automate, this method lends itself to systematic analyses of datasets consisting of photometric time series of large numbers of stars, such as those produced by OGLE, MACHO, TrES, HAT, and many others surveys. ",
        "introduction": "Eclipsing double-lined spectroscopic binaries provide the only method by which both the masses and radii of stars can be estimated without having to resolve spatially the binary or rely on astrophysical assumptions. Despite the large variety of models and parameter-fitting implementations [e.g. WD \\citep{Wilson71} and EBOP \\citep{Etzel81, Popper81}], their underlying methodology is essentially the same. Photometric data provide the light curve of the eclipsing binary (EB), and spectroscopic data provide the radial velocities of its components. The depth and shape of the light curve eclipses constrain the components' brightness and fractional radii, while the radial velocity sets the length scale of the system. In order to characterize fully the components of the binary, one needs to combine all of this information. Only a small fraction of all binaries eclipse, and spectra with sufficient resolution and signal-to-noise can be gathered only for bright stars. The intersection of these two groups leaves a small number of stars.  Over the past decade, the number of stars with high-quality, multi-epoch, photometric data has grown dramatically due to the growing interest in finding gravitational lensing events \\citep{Wambsganss06} and eclipsing extrasolar planets \\citep{Charbonneau06}. In addition, major technical improvements in both CCD detectors and implementations of image-difference analysis techniques \\citep{Crotts92, Alard98, Alard00} enable simultaneous photometric measurements of tens of thousands of stars in a single exposure. Today, there are many millions of light curves available from a variety of surveys, such as OGLE \\citep{Udalski94}, MACHO \\citep{Alcock98}, TrES \\citep{Alonso04}, HAT \\citep{Bakos04}, and XO \\citep{McCullough06}. Despite the increase in photometric data, there has not been a corresponding growth in the quantity of spectroscopic data, nor is this growth likely to occur in the near future. Thus, the number of fully-characterized EBs has not grown at a rate commensurate with the available photometric datasets.  In recent years, there has been a growing effort to mine the wealth of available photometric data, by employing automated pipelines which use simplified EB models in the absence of spectroscopic observations and hence without a fixed physical length scale and absolute luminosity \\citep{Wyithe01, Wyithe02, Devor04, Devor05a}. In this paper, we present a method that utilizes theoretical isochrones and multi-epoch photometric observations of the binary system to estimate the physical parameters of the component stars, while still not requiring spectroscopic observations.  Our Method for Eclipsing Component Identification\\footnote{The source code and running examples of MECI, as well as a suite of utilities, can be downloaded from: http://cfa-www.harvard.edu/$\\sim$jdevor/MECI.html} (MECI), finds the most probable masses, radii, and absolute magnitudes of the stars. The input for MECI is an EB's photometric light curve and out-of-eclipse colors (we note that in the absence of color information, the accuracy in the estimation of the stellar parameters is significantly reduced; \\S\\ref{subsecSimulate}). This approach can be used to characterize quickly large numbers of eclipsing binaries; however it is not sufficient to improve stellar models, since underlying isochrones must be assumed.  In a previous paper \\citep{Devor05b}, we outlined the ideas behind both MECI and a closely related, ``quick and dirty'' alternative, which we called MECI-express. Though MECI-express is much faster and easier to implement, it is also far less accurate. For this reason we will not discuss it further, and instead concentrate exclusively on MECI. We discuss its applications (\\S\\ref{secMotivation}), aspects of its implementation (\\S\\ref{secMethod}), tests of its accuracy (\\S\\ref{secTests}), and finally summarize our findings (\\S\\ref{secConclusions}). ",
        "conclusions": "\\label{secConclusions}  We have described a method for identifying an EB's components using only its photometric light curve and combined colors. By utilizing theoretical isochrones and limb-darkening coefficients, this method greatly reduces the EB parameter space over which one needs to search. Using this approach, we can quickly estimate the masses, radii and absolute magnitudes of the components, without spectroscopic data. We described an implementation of this method, which enables the systematic analyses of datasets consisting of photometric time series of large numbers of stars, such as those produced by OGLE, MACHO, TrES, HAT, and many others.  Such techniques are expected to grow in importance with the next generation surveys, such as Pan-STARRS \\citep{Kaiser02} and LSST \\citep{Tyson02}.  In a future publication, we shall describe a specific application of these codes, namely to search for low-mass eclipsing binaries in the TrES dataset."
    },
    "0609/astro-ph0609294_arXiv.txt": {
        "abstract": "We compute the ultra-high energy cosmic ray horizon, i.e. the distance up to which cosmic ray sources may significantly contribute to the fluxes above a certain threshold on the observed energies. We obtain results both for proton and heavy nuclei sources. ",
        "introduction": "Soon after the discovery of the cosmic microwave background (CMB)  radiation it was realized by Greisen, Zatsepin and Kuzmin \\cite{gzk}  that the fluxes of cosmic ray (CR) protons with energies of order $10^{20}$~eV and above would be strongly attenuated over distances of few tens of Mpc. This is due to the energy losses caused by the photo-pion production processes in the interactions of the protons with the CMB photons. Similarly, if CR sources accelerate heavy nuclei, these can photodisintegrate into lighter ones as they interact with CMB and infrared (IR) photons on their journey to us. In this way the fragments may arrive to the Earth with significantly smaller energies than the parent nuclei produced at the sources. Moreover, both protons and heavy nuclei can further loose energy by pair production processes, although due to the small inelasticities involved the typical attenuation length associated to $e^+e^-$ production at ultra-high energies is large ($\\sim 1$~Gpc for protons).  Many works have studied these processes (see e.g. refs. \\cite{st68,hi85,be88,be90,ah90,yo93,ra93,ah94,wa95,be05} for the case of protons, and \\cite{psb,ep98,st99,kh05,ho06,al05,al06} for nuclei), and our purpose here is to make a detailed  analysis to present the results in a way which can be useful for the study of correlations between the arrival directions of the highest energy cosmic rays (we will focus here on energies above 50~EeV, where 1~${\\rm EeV}\\equiv 10^{18}$~eV) and candidate astronomical sources. Assuming that the sources are distributed uniformly and have similar absolute cosmic ray luminosities, we  will obtain the distances within which the major part  of the observed events should be produced (i.e. the horizon for the potential sources) as a function of the threshold adopted for the energies of the events. ",
        "conclusions": "As a summary, we have shown that the sources of the vast majority of the CR events observed with energies above 50~EeV should be at relatively close distances (under the plausible assumption that CRs consist of ordinary nucleons or nuclei), as is illustrated in a quantitative way in figs.~2, 4 and 5\\footnote{Some particular values of the results in fig.~2 can be read off from plots in previous works \\cite{wa95,wa96,cu05}. To our knowledge the results for nuclei are not available in the literature.}. These results should be relevant in order to restrict  the distances to the candidate astrophysical objects  that could act as sources for the observed ultra-high energy cosmic rays."
    },
    "0609/astro-ph0609777_arXiv.txt": {
        "abstract": "We present a complete particle physics model that explains three major problems of modern cosmology: inflation, dark matter and dark energy, and also gives a mechanism for leptogenesis. The model has a new gauge group $SU(2)_Z$ that grows strong at a scale $\\Lambda\\sim 10^{-3}$ eV. We focus on the inflationary aspects of the model. Inflation occurs with a Coleman-Weinberg potential at a low scale, down to $\\sim 6\\times 10^5\\gev$, being compatible with observational data. ",
        "introduction": "The recent three-year Wilkinson Microwave Anisotropy Probe (WMAP3) results \\cite{wmap3} have put quite a severe constraint on inflationary models. In particular, new results on the value of the spectral index $n_{\\rm s}=0.95 \\pm 0.02$ are sufficiently ``precise'' as to rule out many models with an exact Harrison-Zel'dovich-Peebles scale-invariant spectrum with $n_{\\rm s}=1$ and for which the tensor-to-scalar ratio $r \\ll 1$. Any model purported to describe the early inflationary era will have to take into account this constraint. However, by itself, it is not sufficient to narrow down the various candidate models of inflation. In particular, ``low-scale'' inflationary models are by no means ruled out by the new data. By ``low-scale'' we refer to models in which the scale that characterizes inflation is several orders of magnitude smaller than a typical Grand Unified Theory (GUT) scale $\\sim 10^{15}- 10^{16}\\gev$. It is in this context that we wish to present a  model of low-scale inflation which could also describe the dark energy and dark matter \\cite{hung}, \\cite{hung1}.  The model of dark energy and dark matter described in \\cite{hung1} involves a new gauge group $SU(2)_Z$ which grows strong at a scale $\\Lambda_Z \\sim 3 \\times 10^{-3}$ eV starting with the value of the gauge coupling at $\\sim 10^{16}\\gev$ which is not too different from the Standard Model (SM) couplings at a similar scale. (This is nicely seen when we embed $SU(2)_Z$ and the SM in the unified gauge group $E_6$ \\cite{hung2}.) The model of \\cite{hung1} contains, in addition to the usual SM content, particles which are SM singlets but $SU(2)_Z$ triplets, $\\psi^{(Z)}_{(L,R),i}$ with $i=1,2$, particles which carry quantum numbers of both gauge groups, $\\tilde{\\varphi}_{1,2}^{(Z)}$, which are the so-called messenger fields with the decay of $\\tilde{\\varphi}_{1}^{(Z)}$ being the source of SM leptogenesis \\cite{hung3}, and a singlet complex scalar field, \\be \\label{phiz} \\phi_{Z} = (\\sigma_Z+v_{Z})\\,\\exp(ia_Z/v_{Z})\\,, \\ee whose angular part $a_Z$ is the axion-like scalar. We have defined the radial part of $\\phi_Z$ as the sum of a field $\\sigma_Z$ and a vacuum-expectation-value (v.e.v.) $v_Z$. The $SU(2)_Z$ instanton-induced potential for $a_Z$ (with two degenerate vacua) along with a soft-breaking term whose dynamical origin is discussed in \\cite{hung4}, is one that is proposed in \\cite{hung1} as a model for dark energy. In that scenario, the present universe is assumed to be trapped in a false vacuum of the $a_Z$ potential with an energy density $\\sim \\Lambda_Z^4$. The exit time to the true vacuum was estimated in \\cite{hung1} and was found to be enormous, meaning that our universe will eventually enter a late inflationary stage.  What might be interesting is the possibility that the real part of $\\phi_{Z}$, namely $\\sigma_Z$, could play the role of the inflaton while the imaginary part, $a_Z$, plays the role of the ``acceleron'' as we have mentioned above. This unified description is attractive for the simple reason that {\\em one complex} field describes both phenomena: Early and Late inflation. (This scenario has been exploited earlier \\cite{Masso:2006yk}, \\cite{others} in the context of GUT scale inflation.) Although the structure of the potential describing the accelerating universe is determined, in the model of \\cite{hung1}, by instanton dynamics of the $SU(2)_Z$ gauge interactions \\cite{hung1}, the potential for $\\sigma_Z$, which would describe the early inflationary universe, is arbitrary as with scalar potentials in general. In this case, the only constraint comes from the requirement that this potential should be of the type that gives the desired spectral index and the right amount of inflation corresponding to the characteristic scale of the model.  We now briefly describe the model of \\cite{hung1}. The key ingredient of that model is the postulate of a new, unbroken gauge group $SU(2)_Z$ which grows strong at a scale $\\Lambda_Z \\sim 3 \\times 10^{-3}$ eV. The model also contains a global symmetry $U(1)_{A}^{(Z)}$ which is spontaneously broken by the v.e.v. of $\\phi_{Z}$, namely $\\langle \\phi_{Z} \\rangle = v_Z$, and is also {\\em explicitly} broken at a scale $\\Lambda_Z \\ll v_Z$ by the $SU(2)_Z$ gauge anomaly. Because of this, the pseudo-Nambu-Goldstone boson (PNGB) $a_Z$ acquires a tiny mass as discussed in \\cite{hung1}. Its $SU(2)_Z$ instanton-induced potential used in the false vacuum scenario for the dark energy is given by \\be \\label{totpot} V_{tot}(a_Z,T) = \\Lambda_Z^4\\left[1-\\kappa(T)\\,\\cos\\frac{a_Z}{v_Z}\\right] + \\kappa (T)\\Lambda_Z^4\\,\\frac{a_Z}{2\\pi\\,v_Z} \\,, \\ee where $\\kappa (T) = 1$ at $T < \\Lambda_Z$. ($SU(2)_Z$ instanton effects become important when $\\alpha_Z=g_Z^2/4\\,\\pi \\sim 1$ at $\\Lambda_Z \\sim 3 \\times 10^{-3}\\,eV$.) The universe is assumed to be presently trapped in the false vacuum at $a_Z = 2 \\pi v_Z$ with an energy density $\\sim (3 \\times 10^{-3}\\,eV)^4$. As such, this model mimicks the $\\Lambda$CDM scenario with $w(a_Z)= \\frac{\\frac{1}{2}\\dot{a}_Z^2 - V(a_Z)} {\\frac{1}{2}\\dot{a}_Z^2 + V(a_Z)} \\approx -1$, at present and for a long time from now on, but not in the distant past \\cite{hung1}.  What could be the form of the potential for the {\\em real part}, namely $\\sigma_Z$, of $\\phi_{Z}$? As with any scalar field, the form of the potential is rather arbitrary, with the general constraints being gauge invariance and renormalizability. In this paper, we would like to propose a form of potential for $\\sigma_Z$ which is particularly suited to the discussion of the ``low-scale'' inflationary scenario: a Coleman-Weinberg (CW) type of potential \\cite{Coleman:1973jx}. (The CW-type of potential has been recently used \\cite{Shafi:2006cs} to describe a GUT-scale inflation using the WMAP3 data.) There are three types of contributions to the potential. The sources of these three types are the following terms in the lagrangian:  \\bi  \\item[a)] $\\phi_{Z}-\\psi^{(Z)}_{(L,R),i}$ coupling  \\be \\label{coupling1} \\sum_{i} K_{i} \\, \\bar{\\psi}^{(Z)}_{L,i}\\,\\psi^{(Z)}_{R,i}\\,\\phi_{Z} + h.c. \\ee Let us recall from \\cite{hung1} that (\\ref{coupling1}) is invariant under the following global $U(1)_{A}^{(Z)}$ symmetry transformations: $\\psi^{(Z)}_{L,i} \\rightarrow e^{-i\\alpha}\\,\\psi^{(Z)}_{L,i}$, $\\psi^{(Z)}_{R,i} \\rightarrow e^{i\\alpha}\\,\\psi^{(Z)}_{R,i}$, and $\\phi_{Z} \\rightarrow e^{-2i\\alpha}\\,\\phi_{Z}$.  \\item[b)] $\\phi_Z - \\tilde\\varphi_1^Z$ mixing (we ignore the $\\phi_Z - \\tilde\\varphi_2^Z$ coupling since it is assumed to have a mass of the order of a typical GUT scale) \\be\\label{coupling2} \\tilde \\lambda_{1Z}(\\phi_Z^{\\dag}\\,\\phi_Z)( \\tilde\\varphi_1^{Z,\\dag}\\,\\tilde\\varphi_1^Z)\\ee    \\item[c)] $\\sigma_Z$ self-interaction \\be\\label{coupling3} \\frac{\\lambda}{4!} \\sigma_Z^4\\ee Both terms, (\\ref{coupling2}) and (\\ref{coupling3}), arise from the general potential for all fields.   \\ei Let us look into constraints on these couplings coming from issues discussed in \\cite{hung1}: dark matter and leptogenesis.  Since the coupling (\\ref{coupling1}) will, in principle, contribute to the CW potential for $\\sigma_Z$, it is crucial to have an estimate on the magnitude of the Yukawa couplings $K_{i}$. In \\cite{hung1}, an argument was made as to why it might be possible that $\\psi^{(Z)}_{i}$ could be Cold Dark Matter (CDM) provided \\be \\label{cdm} m_{\\psi^{(Z)}_{i}}= |K_i| v_Z \\leq O(200\\,GeV)\\,, \\ee or \\be \\label{Ki} |K_i| \\leq O(200\\,GeV/v_Z) \\ee Roughly speaking, in order for $\\Omega_{CDM} \\sim O(1)$, the annihilation cross sections for $\\psi^{(Z)}_{i}$ are required to be of the order of weak cross sections. In this case, they are approximately $\\sigma_{\\rm annihilation} \\sim \\alpha_Z^2( m_{\\psi^{(Z)}_{i}})/m_{\\psi^{(Z)}_{i}}^2$ ($\\alpha_Z^2( m_{\\psi^{(Z)}_{i}})$ is the coupling evaluated at $E = m_{\\psi^{(Z)}_{i}}$) and have the desired magnitude when $m_{\\psi^{(Z)}_{i}} \\sim O(200\\,GeV)$ with $\\alpha_Z^2(m_{\\psi^{(Z)}_{i}}) \\sim \\alpha_{SU(2)_L}^2 (m_{\\psi^{(Z)}_{i}})$, a characteristic feature of the model of \\cite{hung1}.   A second requirement comes from a new mechanism for leptogenesis as briefly mentioned in \\cite{hung1} and described in detail in \\cite{hung3}. This new scenario of leptogenesis involves the decay of a messenger scalar field, $\\tilde{\\varphi}_{1}^{(Z)}$, into $\\psi^{(Z)}_{i}$ and a SM lepton. In order to give the correct estimate for the net lepton number, a bound on the mass of $\\tilde{\\varphi}_{1}^{(Z)}$ was derived. In \\cite{hung3}, it was found that \\be \\label{lepto} m_{\\tilde{\\varphi}_{1}^{(Z)}} \\alt 1\\,TeV \\,. \\ee This came about when one calculates the interference between the tree-level and one-loop contributions to the decays \\be \\label{decay1} \\tilde{\\varphi}_{1}^{(Z)} \\rightarrow \\bar{\\psi}^{(Z)}_{1,2}+ l \\ee \\be \\label{decay2} \\tilde{\\varphi}_{1}^{(Z),*} \\rightarrow \\psi^{(Z)}_{1,2}+ \\bar{l} \\ee where $l$ represents a SM lepton. By requiring that the asymmetry coming from this scenario to be $\\epsilon^{\\tilde{\\varphi}_{1}}_{l} \\sim -10^{-7}$ in order to obtain the right amount of baryon number asymmetry through the electroweak sphaleron process, \\cite{hung3} came up with the constraint (\\ref{lepto}) which could be interesting for searches of non-SM scalars at the Large Hadron Collider (LHC). On the other hand, as discussed in \\cite{hung1}, the mixing between $\\tilde{\\varphi}_{1}^{(Z)}$ and $\\phi_{Z}$ results in an additional term in the mass squared formula for $\\tilde{\\varphi}_{1}^{(Z)}$, namely $2 \\tilde{\\lambda}_{1Z}v_{Z}^2$. Taking into account the leptogenesis bound (\\ref{lepto}), one can write \\be \\label{cross} \\tilde{\\lambda}_{1Z} \\alt (1\\,TeV/v_Z)^2 \\,. \\ee  All these constraints will be used to estimate the contributions to the effective potential. The $\\psi^{(Z)}_{i}$ fermion loop contribution to the $\\sigma_Z$ CW potential is given by $-|K_i|^4/16\\pi^2$, and therefore will be bound by (\\ref{Ki}). The $\\tilde{\\varphi}_{1}^{(Z)}$ loop contribution to the potential is given by $\\tilde{\\lambda}_{1Z}^2/16\\pi^2$ and is constrained by (\\ref{cross}). The third contribution c) coming from the $\\sigma_Z$ loop is given by $\\lambda^2/16\\pi^2$. There are no constraints on it coming from dark matter or leptogenesis arguments, as we have in the other cases a) and b).  Below, we will constrain both $v_Z$ and the coefficient of the CW potential (which includes contributions from various loops) using the latest WMAP3 data. Next, we use these results to further constrain $|K_i|$ and $\\tilde{\\lambda}_{1Z}$. We will finally comment on the implications of these constraints. ",
        "conclusions": "In this paper, we show that the model presented in \\cite{hung1}, which explains dark matter and dark energy, also provides a mechanism for inflation in the early universe. We find that it is conceivable to have a low-scale inflation.  The complete model contains a new gauge group $SU(2)_Z$ which grows strong at a scale $\\Lambda\\sim 3\\times 10^{-3}$ eV, with the gauge coupling at GUT scale comparable to the SM couplings at the same scale. In addition to the SM particles, the model contains new particles: $\\psi_{(L,R),i}^{(Z)} (i=1,2)$ which are $SU(2)_Z$ triplets and SM singlets; $\\tilde\\varphi_{1,2}^{(Z)}$ which are the so-called messenger fields and carry charges of both $SU(2)_Z$ and SM groups; and $\\phi_Z$, which is a singlet complex scalar field. The model also contains a new global symmetry $U(1)_A^{(Z)}$, which is spontaneously broken by the v.e.v. of the scalar field, $\\langle\\phi_Z\\rangle=v_Z$, and also explicitly broken at the scale $\\Lambda_Z\\ll v_Z$ by the $SU(2)_Z$ gauge anomaly. The real part of the complex scalar field, namely $\\sigma_Z$, is identified with the inflaton field. We considered a CW-type of potential for $\\sigma_Z$ and obtained the constraints on the parameters of the model in order to have a right description of inflation. The angular part of the complex scalar field, namely $a_Z$, acquires a small mass due to the explicit breaking of $SU(2)_Z$ and is trapped in a false vacuum, being responsible for the dark energy of the universe. The new particles $\\psi_{(L,R),i}^{(Z)}$, with masses of order $200\\gev$, explain the dark matter. They are produced at reheating, by the decay of the inflaton. For values of the $SU(2)_Z$ breaking scale $v_Z\\sim 3\\times 10^9\\gev$, we obtain a low-scale model of inflation, namely a scale of $\\sim 6\\times 10^5\\gev$. Notice that, in order to have a realistic reheating mechanism, this ``low-scale'' is also bounded from above by $\\sim 2\\times 10^7\\gev$ as we have discussed in the last section. Because of this fact, our model is a bona-fide ``low-scale'' inflationary scenario. It is an exciting possibility because the model might be indirectly probed at future LHC experiments."
    },
    "0609/astro-ph0609011_arXiv.txt": {
        "abstract": "We derive the bulk viscous damping timescale of hybrid stars, neutron stars with quark matter core. The r-mode instability windows of the stars show that the theoretical results are consistent with the rapid rotation pulsar data, which may give an indication for the existence of quark matter  in the interior of neutron stars. Hybrid stars instead of neutron or strange stars may  lead to submillisecond pulsars. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609820_arXiv.txt": {
        "abstract": "{ The structure of the high galactic latitude Cometary Globule 12 (CG~12) has been investigated by means of radio molecular line observations.  Detailed, high signal to noise ratio maps in C$^{18}$O~(1--0), C$^{18}$O~(2--1) and molecules tracing high density gas, CS (3--2), \\dcopl (2--1) and \\htcopl (1--0), are presented. The C$^{18}$O line emission is distributed in a $10\\arcmin$ long North-South elongated lane with two strong maxima, \\cgnp(orth) and \\cgsp(outh).  In \\cgs the high density tracers delineate a compact core, \\dcopl core, which is offset by 15\\arcsec \\ from the \\ceo maximum. The observed strong \\ceo emission traces the surface of the \\dcopl core or a separate, adjacent cloud component.  The driving source of the collimated molecular outflow detected by \\cite{white1993} is located in the \\dcopl core. The \\ceo lines in \\cgs have low intensity wings possibly caused by the outflow.The emission in high density tracers is weak in \\cgn and especially the \\htcoplp, \\dcopl and \\nthp lines are +0.5 \\kmps \\  offset in velocity with respect to the \\ceo lines. Evidence is presented that the molecular gas is highly depleted. The observed strong \\ceo emission towards \\cgn originates in the envelope of this depleted cloud component  or in a separate entity seen in the same line of sight. The \\ceo lines in CG 12 were analyzed using Positive Matrix Factorization, PMF.  The shape and the spatial distribution of the individual PMF factors fitted separately to the \\ceo (1--0) and (2--1) transitions were consistent with each other.  The results indicate a complex velocity and line excitation structure in the cloud.  Besides separate cloud velocity components the \\ceo line shapes and intensities are influenced by excitation temperature variations caused by e.g, the molecular outflow or by molecular depletion.  Assuming a distance of 630 pc the size of the CG 12 compact head, 1.1 pc by 1.8 pc, and the \\ceo mass larger than 100 M$_{\\sun}$ are comparable to those of other nearby low/intermediate mass star formation regions. ",
        "introduction": "Herschel~(1847) noted that a 10$^{\\rm th}$ mag. star, now known as \\object{h4636} or \\object{CoD~--39\\degr~8581}, is a binary with a separation of 3\\farcs7. In the optical, h4636 illuminates the bright reflection nebula \\object{NGC~5367}. \\cite{HawardenBrand1976} showed that it lies in the head of an impressive \\object{Cometary Globule 12}, \\object{CG~12}, with a tail stretching about one degree to the SE. With a galactic latitude of 21\\degr \\ and at the distance of 630~pc estimated by \\cite{williams1977}, CG~12 lies more than 200~pc above the plane.  CG 12 has an associated low/intermediate mass stellar cluster which has at least 9 members (Williams et al. 1977). The clouds  cometary structure could be due to the passage of a supernova blast wave. Curiously, the cometary tail stretches towards the Galactic plane which would place the putative supernova even farther away from the Galactic plane than the globule.  According to \\cite{maheswar2004} the head of CG~12 is pointing towards the centre of an HI shell. Such a shell is, however, not readily evident in the whole sky HI survey~(\\cite{kalberlaetal2005}) which merges the northern Leiden/Dwingeloo Survey~(\\cite{hartmanburton1997}) and the southern Instituto Argentino de Radioastronomia Survey~(\\cite{arnaletal2000}).  The CG~12 cloud cometary shape is also seen in the IRAS surface emission~(\\cite{odenwald1988}).    \\begin{figure*} \\centering \\includegraphics[bb=50 10 480 700, width=13.0cm,angle=-90, clip]{haikala_fig_1.eps} \\caption{ Maps of integrated emission of C$^{18}$O~(1--0) and C$^{18}$O~(2--1) $\\int$\\tastar d{\\it v} from $-8$ \\kmps \\ to $-4$ \\kmps \\ in CG12.  The lowest contour and the contour increment are 0.3 $\\kkms$.  The SEST half power beam sizes are shown in the upper corners of the panels.  Offsets from the map centre position, $13\\hour 57\\minute 43\\fs 1, -39\\degr 58\\arcmin 43\\farcs 3 $~(J2000), are shown on the axes of the right panel.  The scale tick in the lower right corner assumes a distance of 630 pc to the cloud.} \\label{figure:CG12_CO_maps}  \\end{figure*}      White~(1993) mapped the region around and south of h4636 in $^{12}$CO~(2--1) and C$^{18}$O~(2--1) lines with a spatial resolution of 22\\arcsec.  He found a small size C$^{18}$O core near the binary, and a molecular outflow with a centre close to the binary system. Further large scale \\twco, \\thco and \\ceo observations with 2\\farcm7 resolution are presented in \\cite{yonekuraetal1999a}.  CG 12 contains two compact 1.2mm continuum sources, one in the direction of the \\ceo core detected by White~(1993) (Reipurth et al. 1996) and another one two arcminutes north of it (Haikala 2006, in preparation). The centres of  the continuum sources were observed in  \\ceo (3--2) by \\cite{haikalaetal2006}.   Near infrared~(J, H and K) images and photometry of stars in CG12 are available in  the 2MASS survey and \\cite{santos1998}. A deeper J, H and Ks imaging study is presented in Haikala (2006, in preparation). Far infrared emission in CG 12 is dominated by a strong point source, \\object{IRAS 13547-3944}, near the binary h4636.   CG 12/NGC 5367 is an intriguing object. It has the appearance of a cometary globule, such as usually found in the outskirts of HII regions. The linear size of cometary globules, like those in the Gum nebula, are however much smaller (e.g., Reipurth 1983).  The linear extent of CG12, 10 pc, is four times larger than, e.g., the archetype object CG 1. A cluster of low/intermediate mass stellar cluster is associated with CG 12. Perhaps the most curious feature of CG 12 is its location over 200 pc above the Galactic plane with no sign of other nearby dark clouds or star formation.  CG 12 has not attracted much interest since the identification of the stellar cluster (Williams et al.1977) and the \\cite{HawardenBrand1976} cometary globule paper. The subsequent papers have either concentrated on the Herbig~AeBe binary, h4636 or the cloud has been included in various surveys. The \\twco (2--1) and \\ceo (2--1) observations by White~(1993), though detailed, cover only the very centre of the cloud. What has been missing is a detailed but at the same time extended study of the dense CG 12 molecular cloud in molecular transitions sensitive to the large scale structure (cloud envelope) and to the detailed structure of the high density material (cores).    In this paper we report the mapping of the head of CG12 in the C$^{18}$O~(1--0) and (2--1) and in \\thco (1--0) lines. The  \\ceo emission maxima  were further mapped in \\dcoptop, \\htcopozp \\ and CS~(2--1) and (3--2).  Further pointed observations with long integration time in CO~(1--0) and (2--1) (and isotopologues), CS~(2--1) and (3--2), \\ctfs~(2--1), \\htcopl (1--0), \\dcopto and \\nthp (1--0) were made towards selected positions in the cloud.  Observations, data reduction and calibration procedures are described in Sect.~\\ref{sec:observations} and the observational results in Sect~\\ref{sec:results}. The new results are compared with the optical and NIR images  in Sect.~\\ref{sec:compare}. In Sect.~\\ref{sec:finescale} and Appendix \\ref{sec:PMF_app} Positive Matrix Factorization is used to analyze the \\ceo small scale structure in the globule. In the discussion part in Sect.~\\ref{sec:discussion} the \\ceo column densities and the cloud mass are derived and the observational and calculated results are summarized. The conclusions are drawn in  Sect.~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have performed a detailed, high signal-to-noise ratio, mm line study of CG 12 in various molecular transitions, principally of C$^{18}$O~(1--0) and C$^{18}$O~(2--1), as well as in molecular lines probing dense material, and have obtained the following results:  1. The C$^{18}$O line emission is distributed in a 10$\\arcmin$ North-South elongated lane with two strong, compact maxima, \\cgn and \\cgsp, and a weaker maximum, \\cgswp.  2. High density tracers CS (2--1), (3-2), \\dcopl and \\htcopl are detected in both strong \\ceo maxima. Emission from these molecules is weak in \\cgn but in \\cgs it defines a compact core (referred to as \\dcopl core) which is spatially offset from the \\ceo maximum.  3. The emission from the high density tracers in \\cgn takes place at a velocity where emission from \\ceo is weak. The molecules associated with the cloud component  detected in high density tracers are likely to be heavily depleted.  4. Positive Matrix Factorization was applied to study the cloud \\ceo fine scale structure. The observed strong \\ceo emission in \\cgn (PMF factor 2{\\it n}) originates in the envelope of the depleted cloud component or in a separate entity seen in the same line of sight.  In \\cgs the most intense PMF factor 3{\\it s} traces warm gas on the surface of the \\dcopl core or a separate adjacent cloud component.  5. The driving source of the collimated molecular outflow detected by \\cite{white1993} lies in the \\dcopl core.   6. The average \\ceo LTE mass is $\\sim$80 M$_{\\sun}$ for \\cgnp, $\\sim$70 M$_{\\sun}$ for \\cgs and $\\sim$20 M$_{\\sun}$ for \\cgswp. These numbers can only be considered as a zeroth order estimate because of the uncertainty in defining the \\ceo excitation temperature and possible molecular depletion in the \\ceo maxima.  7. If the distance to CG 12, 630 pc, is correct the linear size and the mass of this cometary globule approaches that of a typical low mass star forming region like eg. Chamaeleon I."
    },
    "0609/astro-ph0609227_arXiv.txt": {
        "abstract": "We present a new unbiased search and analysis of all B stars in the solar neighbourhood (within a volume of 400 pc diameter) using the {\\sf Arivel} data base to track down the remains of the OB associations, which hosted the supernovae responsible for the Local Bubble in the interstellar gas. We find after careful dereddening and by comparison with theoretical isochrones, that besides the Upper Scorpius the Upper Centaurus Lupus and Lower Centaurus Crux subgroups are the youngest stellar associations in the solar neighbourhood with ages of 20 to 30 Myr, in agreement with previous work. In search for the ``smoking gun'' of the origin of the Local Bubble, we have traced the paths of the associations back into the past and found that they entered the present bubble region 10 to 15 Myr ago. We argue that the Local Bubble began to form then and estimate that 14 to 20 supernovae have gone off since. It is shown that the implied energy input is sufficient to excavate a bubble of the presently observed size. ",
        "introduction": "The Local Bubble (LB), a low-density X--ray emitting cavity deficient of H{\\sc i}, is our Galactic habitat. Yet, until recently, its origin remained mysterious. It was conjectured to be the result of one or several supernova explosions (e.g.\\ Cox \\& Anderson 1982, Innes \\& Hartquist 1984, Smith \\& Cox 2001), but firm evidence was lacking, as no OB association was found within its boundaries, extending about 200 pc in the Galactic plane, and 600 pc perpendicular to it, but inclined by about $20^\\circ$ with respect to the axis of Galactic rotation, similar to Gould's Belt (cf.\\ Lallement et al.~2003). Further problems arose, since the X--ray and EUV spectra measured in the Wisconsin Survey, by {\\sf ROSAT PSPC, DXS, XQC} and {\\sf EUVE} were severely at odds with a thermal hot plasma in collisional ionization equilibrium (CIE) as was pointed out by Jelinsky et al.~(1995), Sanders et al.~(2001), McCammon et al.~(2002). Most recently Hurwitz et al.~(2005) analyzed {\\sf CHIPS} data and found an extremely low emissivity of EUV iron lines. The underabundance of soft X--ray lines can be naturally explained if the plasma is in a state of delayed recombination (Breitschwerdt \\& Schmutzler 1994, Breitschwerdt 2001), but a high resolution numerical hydrodynamical evolution model is needed to better constrain non-equilibrium models. Spectral discrepancies between models and observations can be alleviated if there is a substantial contribution from very local sources, such as the Earth's exosphere (Freyberg 1998) or charge exchange reactions between solar wind ions (SWCE) and heliospheric gas (Lallement 2004). At present it is unclear what fraction can be attributed to these very nearby sources, although there is fairly robust evidence that even in the extreme case of all of the X--ray emission being due to SWCE in a certain direction, a substantial LB fraction remains, especially perpendicular to the disc. For further details on LB properties we refer to the review of Breitschwerdt (2001) and the conference proceedings \\emph{The Local Bubble and Beyond} (Breitschwerdt et al.~1998).  All these shortcomings have led several authors to speculate that if the LB is not a classical superbubble, but rather an appendix of the neighbouring Loop~I superbubble, which was expanding into an interarm region  between the Sagittarius and the Perseus spiral arms of the Galaxy (Bochkarev 1987, Frisch 1995). However, the existence of a ``wall'' between the two bubbles, showing up in absorption of soft X--rays in {\\sf ROSAT PSPC} images (Egger \\& Aschenbach 1995) renders this scenario not very plausible.  The search for the ``smoking gun'' of the origin of the LB proved partially successful by discovering that moving groups of young stars in the solar neighbourhood could provide an adequate number of SN explosions while crossing the path of the LB. Bergh\\\"ofer \\& Breitschwerdt (2002, henceforth BB02) calculated the trajectory of the Pleiades subgroup B1 backwards in time, and found that 19 SNe could have exploded between 10--20 Myr ago in the region that is occupied by the LB. The remaining stars of B1 are now part of the Scorpius Centaurus OB association. It could be shown that this is in good agreement with the size of the LB and the present soft X--ray emissivity. A similar analysis was carried out by Ma\\'{i}z-Apell\\'{a}niz (2001), who calculated backwards in time the trajectories of Sco Cen subgroups and claimed that about 6 SNe that went off in the Lower Centaurus Crux subgroup of the Sco OB2 association 7 to 9 Myr ago formed the LB.  While these analyses represent a major step towards the understanding of the origin of the LB they are not free from bias, in particular the assumption that \\emph{certain stellar groups should} be responsible for the sought SN explosions. The purpose of this paper is to scrutinize \\emph{all stars} that are within a volume of about 400 pc in diameter centered around the Sun, and to perform a selection according to spectral and kinematical properties. The latter is based on three dimensional space velocities of the stars. Thus our approach is complementary to studies like by de Zeeuw et al. (1999) which are based on proper motions alone. Sartori et al. (2003) do include radial velocities when analyzing the subgroups of the Sco OB2 association, but work from a list of stars preselected by de Zeeuw et al. (1999). From their position in the HR diagram and the turnoff point from the ZAMS, we can reliably determine the age of the stars and estimate the number of SNe within a defined region, such as the LB.  The paper is organized as follows. In the next section we describe our search strategy for the remnants of the OB association responsible for the origin of the LB. In section 3 we discuss the consistency of our findings with the properties of the LB as observed today, and present a high--resolution 3D hydrodynamical simulation of the formation of the LB in the local interstellar gas. In the final section we summarize our conclusions.     \\section[]{Search for nearby OB associations}  Even though the Sco\\,OB2 association has been claimed with good reason to be responsible for the origin of the LB (Ma\\'{i}z-Apell\\'{a}niz 2001, BB02), a fresh, unprejudiced search for the OB association, that might have triggered the formation of the LB, will improve and harmonize previous studies. As starting point we used the {\\sf Hipparcos} catalogue (ESA 1997) from which we selected all stars bluer than $(B-V) < -0.05$ with parallaxes larger than 5 milliarcseconds, because otherwise the distances would be too inaccurate. Drimmel et al.~(2000) find by a comparison with the Tycho catalogue that the {\\sf Hipparcos} catalogue is about 97 percent complete down to $V=7.5$ which corresponds at a distance of $1/(5mas)=200$ pc to an absolute magnitude of $M_{\\rm V}=1$. Allowing for an extinction of $A_{\\rm V} \\leq 0.2$ we have $M_{\\rm V}<1.2$, and we have chosen the colour cut in $(B-V)$ accordingly (cf. Figs. 4 and 5). We have obtained this way an unbiased complete sample of 762 B stars within a distance of 200 pc from the Sun. One of the bluest stars is the B0.5 star $\\alpha$\\,Cru. The stars of the sample show at this stage already a well defined main sequence in the colour--magnitude diagram (cf.~Fig.~4). There are a few stars lying several magnitudes below the main sequence, which are probably subdwarfs. We have omitted these, $M_{\\rm V}>3.0$ mag, because we are interested in young stars. All stars in our list appear in the {\\sf Arivel} data base, which combines {\\sf Hipparcos} parallaxes and proper motions with radial velocity data collected from the literature (Wielen et al., in preparation). We find that for 610 stars radial velocities are available. These are accurate enough to be used for a kinematical analysis (cf.~Fig.~3). We have tested if the availability of radial velocities is correlated with the colours of the stars and find that the 152 out of 762 stars for which no radial velocity is known lie above $(B-V) > -0.1$ . We discuss the implications of this effect on the determination of the expected number of supernovae below. For each star spatial $X$, $Y$, and $Z$ coordinates and the corresponding $U$, $V$, and $W$ velocity components have been calculated. The coordinates are centered on the Sun with $X$ pointing towards $l=0^\\circ$, $b=0^\\circ$, $Y$ towards $l= 90^\\circ$, $b = 0^\\circ$, and $Z$ towards $b = 90^\\circ$, respectively. \\begin{figure} \\centering \\caption[]{Positions of 610 stars drawn from the {\\sf Hipparcos} catalogue. The selected stars have colours $B-V < -0.05$ and for each star its radial velocity is known. The $X$--axis points towards the Galactic centre, $Y$ into the direction of Galactic rotation, and $Z$ towards the Galactic north-pole, respectively.} \\label{fig:1}       % \\end{figure} Since OB associations disperse slowly on time scales of the order of several $10^7$ years (Blaauw 1964), we have searched for kinematically coherent structures in our sample. For this purpose we have traced the positions of the stars backward in time over 3$\\cdot$10$^7$ years. Stellar orbits have been calculated using the epicyclic equations of motion for the stars in the sample (Lindblad 1959, Wielen 1982) \\begin{eqnarray} X(t) &=& X(0)-\\frac{V(0)}{-2B}(1-\\cos{(\\kappa t)})+\\frac{U(0)}{\\kappa} \\sin{(\\kappa t)}, \\nonumber \\\\ U(t) &=& U(0)\\cos{(\\kappa t)}-\\frac{\\kappa}{-2B}V(0)\\sin{(\\kappa t)}, \\nonumber \\\\Y(t)& =& Y(0) + 2 A \\left( X(0)-\\frac{V(0)}{-2B} \\right) t \\\\ & + & \\frac{\\Omega_0}{-B \\kappa}V(0)\\sin{(\\kappa t)}  + \\frac{2 \\Omega_0}{\\kappa^2}U(0) ( 1 -\\cos{(\\kappa t)}), \\nonumber \\\\ V(t)& = &\\frac{-2 B}{\\kappa}U(0)\\sin{(\\kappa t)} + V(0)\\cos{(\\kappa t)}, \\nonumber \\\\ Z(t) &=& \\frac{W(0)}{\\nu}\\sin{(\\nu t)} + Z(0)\\cos{(\\nu t)}, \\nonumber \\\\ W(t)&=& W(0)\\cos{(\\nu t)} -Z(0)\\nu \\sin{(\\nu t)}\\,. \\nonumber \\end{eqnarray} In Eqns.~(1) $\\kappa$ denotes the epicyclic frequency, $\\kappa = \\sqrt{-4\\Omega _0B}$. $A$ and $B$ are the Oort constants, and $\\Omega _0$ is the angular frequency of the rotation of the local standard of rest around the Galactic centre, $ \\Omega_0=V_{\\rm LSR}/R_\\odot$. $\\nu$ denotes the vertical oscillation frequency which is related to the local density $\\rho_0$ by the Poisson equation  as $\\nu= \\sqrt{4\\pi G\\rho_0}$, where $G$ is the constant of gravitation. For the angular velocity of the local standard of rest we have adopted a value of $\\Omega_0$ = 220 km/s/8kpc. The choice of the Oort constants was guided by the consideration that they describe in eqns.~(1) the smooth Galactic gravitational potential. The latter is consistent with an essentially flat shape of the local Galactic rotation curve, $A$ = $-B$ = $\\Omega_0/2$ (Feast \\& Whitelock 1997). This must not be confused with determinations of $A$ and $B$ using OB stars as, for instance, in the studies of Torra et al.~(2000) or Elias et al.~(2006). These reflect peculiarities of the orbits of the OB stars in Gould's belt related to the velocities with which they were born, but not the characteristic smooth shape of the Galactic potential. For the local density we adopt a value of $\\rho_0$ = 0.1 $\\mathcal{M}_\\odot$/pc$^3$ (Holmberg \\& Flynn 2004). These parameter values imply $\\kappa$ = 0.039 km/s/pc = 4$\\cdot$10$^{-8}$yr$^{-1}$ and $\\nu$ = 0.074 km/s/pc = 7.5$\\cdot$10$^{-8}$yr$^{-1}$. In Fig.~2 we show the positions of the stars today and 3$\\cdot$10$^7$ years ago. Apparently most stars came from directions $-90^\\circ < l < 90^\\circ$ and stayed close to the Galactic midplane. Most of the 610 stars do not belong to the OB association, which hosted the SNe responsible for the origin of the LB, and have space velocities different from the velocity of the association. Thus they are dispersed away into a wide cloud. However, the overdense regions in Fig.~2 indicate that there is a considerable number of stars which stayed together. The larger size of the overdense regions in the back projected sample compared to its size today is obviously due to the observational errors. The typical accuracy of {\\sf Hipparcos} proper motions is about 1 mas/yr which corresponds at a distance of 100 pc to a velocity of 0.5 km/s, whereas the accuracy of the radial velocities is several km/s. Taken together with an expansion velocity of the order of 10 km/s (Blaauw 1964), this implies a spreading of the overdensity, which represents the kinematically homogenous group of stars, to a size of roughly 500 $\\times$ 500 pc in X and Y. As can be seen from Fig.~2 there is an outer shroud of stars which lies at greater distances from the core of the overdensity. These must be stars with genuinely different space velocities from the kinematically homogenous group of stars. We identify this kinematically homogeneous group of stars as an OB association and select 302 stars lying in the windows indicated as dashed lines in Fig.~2. As expected these stars are more or less closely related to the Sco\\,OB2 association. \\begin{figure} \\centering \\caption[]{Positions of the originally selected stars today (pink) and 3$\\cdot$10$^7$ years ago (blue). The Sun is at rest in the diagrams. Stars lying in the windows indicated by dashed lines are identified as putative members of the searched for OB association.} \\label{fig:2}       % \\end{figure}  In Fig.~3 we show the present day velocity distribution of the 302 selected stars. Since the velocity dispersion of an OB association is of the order of 10 km/s (Blaauw 1964) or even less (Kamaya 2004), we make a second selection indicated by windows drawn as dashed lines in Fig.~3. This leaves a sample of 236 stars which we analyze in the following. {\\sf Hipparcos} numbers of these stars are listed in the appendix. \\begin{figure} \\centering \\caption[]{Present day velocity distribution of the 302 selected stars. A second selection is made of the stars lying in the windows indicated by dashed lines. } \\label{fig:3}       % \\end{figure}  The final sample is shown as a colour--magnitude diagram in Fig.~4. For this purpose we have cross--identified the sample stars in the {\\sf Geneva photometry} data base (Mermilliod, Hauck \\& Mermilliod 2000) and replaced the $(B-V)_T$ colours given in the {\\sf Hipparcos} catalogue by $(B-V)_J$ colours, because they can be then directly compared with theoretical isochrones available in the literature. In the colour range, which we consider here, $B-V$ given in the {\\sf Tycho} system can not be transformed directly to the Johnson system (ESA 1997). The absolute magnitudes have been determined from the visual magnitudes given in the {\\sf Hipparcos} catalogue in the Johnson system.  We have compared our sample with the extensive membership list of the Sco\\,OB2 association compiled by de Zeeuw et al.~(1999) who applied a combination of a modified convergent point method and the so called spaghetti method (Hoogerwerf \\& Aguilar 1999) to {\\sf Hipparcos} data. Of particular interest are the membership lists of the subgroups Upper Scorpius (US), Upper Centaurus Lupus (UCL), and Lower Centaurus Crux (LCC). With only very few exceptions all stars in the membership lists, which fulfill our colour selection criterion, appear also in our sample, which gives confidence in our selection procedure. A few stars from our final sample could be identified additionally in the membership list of de Geus et al.~(1989) as members of the subgroups. The 79 stars common to both lists are colour coded in Fig.~4 and listed separately in the appendix. \\begin{figure} \\centering \\caption[]{Colour--magnitude diagram of the final sample (236 stars). Members of the Upper Centaurus Lupus subgroup of Sco\\,OB2 are highlighted in orange, Lower Centaurus Crux in yellow, and Upper Scorpius in grey, respectively.} \\label{fig:4}       % \\end{figure} \\begin{figure} \\centering \\caption[]{Dereddened colour--magnitude diagram of the members of the Upper Scorpius (grey), Upper Centaurus Lupus (orange) and Lower Centaurus Crux (yellow) subgroups. The solid lines are theoretical isochrones colour coded according to their ages.} \\label{fig:5}       % \\end{figure}  \\section[]{Results and discussion} \\subsection[]{The search for the ``smoking gun''}   The colour--magnitude diagram presented in Fig.~4 shows a clearly discernible main sequence, which is particularly well delineated by the members of the UCL group. The turn-off point at the tip is defined by both the members of the UCL and the LCC subgroups. Apparently these are together with the US subgroup indeed the youngest OB associations in the solar neighbourhood (de Geus et al. 1989, Sartori et al. 2003). In order to determine their age we have compared the colour--magnitude diagram with theoretical isochrones calculated by Schaller et al.~(1992) for solar metallicities. Fortunately de Bruijne (1999) and Sartori et al.~(2003) have determined individually for most members of the US, UCL, and LCC subgroups, respectively, the extinction and colour excess by comparing the observed $(V-I)_{\\rm C}$ colours with the intrinsic colours of stars of the same spectral type and luminosity class. Dereddened data of the 79 stars are shown together with isochrones in Fig.~5. We conclude from Fig.~5 that the ages of the UCL and LCC subgroups lie in the range of 20 to 30 Myr, whereas we cannot date the age of the US subgroup on the basis of our data. We note that this estimate of the ages of the subgroups is nearly twice of that of de Geus et al.~(1989), who determined an age of 11 - 12 Myr of the LCC subgroup and 14 - 15 Myr of the UCL subgroup, respectively. These age estimates were revised by Sartori et al. (2003) to 16 - 20 Myr on the basis of the Padova isochrones (Bertelli et al. 1994) instead of the Maeder (1981a, b, c) isochrones, which were used by de Geus et al. (1989). The Schaller et al. (1992) isochrones, which we used, are an upgrade of Maeder's isochrones by the Geneva group. Moreover we note that Sartori et al.~(2003) have adopted for the majority of their stars the spectral types given in the {\\sf Hipparcos} catalogue, which might not be as reliable as the Geneva photometric data which we used. Given these uncertainties we conclude that our age datings of the LCC and UCL subgroups are consistent with Sartori et al.'s result. This agrees also well with the age of Pleiades subgroup B1, which was suggested to be responsible for the origin of the LB by BB02, but is significantly larger than assumed by Ma\\'{i}z-Apell\\'{a}niz (2001), especially for the LCC subgroup. Moreover, we have examined with the help of the {\\sf Simbad} data base each star of the subgroups lying not on the main sequence and found that practically all these stars are either binaries or peculiar in the sense that they are variable, emission line stars etc.~(cf.~the notes to the tables), so that their position off the main sequence in the colour--magnitude diagram shown in Fig.~5 can be explained in our interpretation by such effects.  In Fig.~6 we trace back the positions of the UCL and LCC subgroup members over the last 30 Myr using again the epicycle equations (1). However, we have not used the individual space velocities of the stars, but adopted for each star the mass--weighted mean velocity of the combined subgroups. This avoids any unphysical spread of the spatial distribution of the stars at earlier times due to the errors of the space velocity components of the stars. The stellar masses have been determined with the mass--to--magnitude relation \\begin{equation} \\mathcal{M}_\\ast/\\mathcal{M}_\\odot = 3.857-1.453 M_{\\rm V} + 0.183M_{\\rm V}^2 + 0.069 M_{\\rm v}^3 \\,, \\end{equation} which we have derived from a fit to Schaller et al.'s (1992) isochrone data. To the mean velocity of the stars we have added the solar motion $(U,V,W)_\\odot = (10,5.3,7.2)$ km/s (Dehnen \\& Binney 1998), so that the orbits are calculated in the reference frame of the local standard of rest. Since interstellar gas has usually only small peculiar motions, the local interstellar gas, and with it the LB which is indicated in Fig.~6 by the contour line taken from Lallement et al.~(2003), will basically corotate with the local standard of rest around the Galactic centre. This means that the LB is at rest in the reference frame of Fig.~6. As can been seen from Fig.~6 the path of the association has aligned itself 15 Myr ago nearly parallel to the tangential $Y$--direction, the direction of Galactic rotation. Remnants of supernova explosions occurring during this period will have experienced very little shear due to the differential rotation of the Galaxy. The shear effect is described quantitatively by the term linearly proportional to time in the epicycle equation for $Y(t)$ (cf.~Eqs.~1). The coefficient $X(0)-\\frac{V(0)}{-2B}$ is the mean guiding centre radius of an orbit. If the spread of these radii is small, as was the case in the last 15 Myr, the shear effect of the Galactic differential rotation is minimized. In our view this might well explain why all supernovae occurring during that time have combined together to form the LB, while supernova remnants formed at earlier times have drifted away (cf.~ Figs.~6 and 7). \\begin{figure} \\centering \\caption[]{Path of the Upper Centaurus Lupus and Lower Centaurus Crux associations over the last 30 Myr projected onto the Galactic plane. The look--back time is colour coded. The orbits are calculated backwards in the reference frame of the local standard of rest assuming for each star the same mass--weighted mean velocity of the stars. The position of the Local Bubble is indicated by the dash--dotted contour line and is at rest in this reference frame.} \\label{fig:6}       % \\end{figure}  Next we illustrate in Fig.~7 the position of the Upper Centaurus Lupus and Lower Centaurus Crux associations relative to the LB today and at earlier times and reproduce the present day LB contours in meridional sections through the bubble. From Fig.~6 we estimate the Galactic longitude in which direction we expect the associations to move. Choosing then the appropriate meridional section through the bubble from the paper by Lallement et al.~(2003), we can determine immediately the positions of the stars in that longitude range relative to the LB. As can be seen from the upper panels of Fig.~7 the associations are today just about to exit the bubble. Five and 10 Myr ago they were inside. The bottom right panel of Fig.~7 indicates that they entered 15 Myr ago the region occupied by the LB today. In this scenario the LB was starting to form about 15 Myr ago, which is consistent with the estimates of the age of the LB by Ma\\'{i}z-Apell\\'{a}niz (2001) and BB02. In this context it should be kept in mind that although the contours determined by Lallement et al.~(2003) are the presently best available, they are derived from Na{\\sc i} absorption line measurements, which allow to trace the  H{\\sc i} distribution under certain conditions, such as low temperatures ($< 10^4$ K), line saturation for high column densities etc.~(see Sfeir et al.~1999 for a discussion). In particular, the extension of the LB is \\emph{unknown} in directions where there are no background stars. \\begin{figure} \\centering \\caption[]{Meridional sections of the contours delineating the outer boundary of the Local Bubble together with the positions of the stars in the Upper Centaurus Lupus and Lower Centaurus Crux associations. The horizontal axis in the upper left panel points into the direction $l=300^\\circ$, in the upper right towards $l=315^\\circ$ and so on. The vertical direction is always perpendicular to the Galactic midplane. The ages of the associations are colour coded as in Fig.~6.} \\label{fig:7}       % \\end{figure} Regardless of the uncertainties of the outer boundary of the LB one might wonder, moreover, how realistic a scenario is, in which the SNe explode rather close to the edge of the present day bubble. As our high resolution simulation discussed in the next section shows, the location of the star cluster with respect to the centre of the bubble is not crucial. The bubble expands always fastest in the direction of the lowest ambient density and pressure. Since in the direction of the Galactic centre the Loop~I superbubble was formed almost at the same time as the LB by SNe exploding within the Sco--Cen association, the pressure in this direction is very high. Hence the LB was forced to expand rather towards the anticentre direction and perpendicular to the plane, in agreement with the observations.  \\subsection[]{Unravelling the supernovae of the Local Bubble} Next we shall derive the number and masses of stars that have exploded during the journey of the OB association through the LB. Following BB02, the number of supernovae, which created the LB can be estimated with the aid of the initial mass function (IMF). We have fitted an IMF of the form (Massey et al.~1995) \\begin{equation} \\frac{d N }{d \\mathcal{M}}= \\frac{d N }{d \\mathcal{M}}\\Big|_0 \\mathcal{M}^{\\Gamma -1} \\end{equation} with an index $\\Gamma = -1.1\\pm 0.1$ to the data. Masses are given in units of solar masses. The lower end of the main sequence at $M_{\\rm V} = 1$ mag corresponds to A0 stars with masses of $\\mathcal{M}_{\\rm l} = 2.6\\,\\mathcal{M}_\\odot$ and the upper tip at $M_{\\rm V} = -3.7$ mag to B0 stars with masses of $\\mathcal{M}_{\\rm u} = 8.2\\,\\mathcal{M}_\\odot$ (Schaller et al.~1992), respectively. The total number of stars in the UCL and LCC associations, respectively, allow the determination of the normalization constants, \\begin{equation} N=\\int_{2.6}^{8.2} \\frac{d N }{d \\mathcal{M}}\\Big|_0 \\mathcal{M}^{-2.1} d \\mathcal{M}= \\, 0.228\\,\\frac{d N }{d \\mathcal{M}}\\Big|_0 \\,, \\end{equation} implying $\\frac{d N }{d \\mathcal{M}}\\big|_0$ = 184 for the 42 UCL and $\\frac{d N }{d \\mathcal{M}}\\big|_0$ = 118 for the 27 LCC stars, respectively. As we have shown in the previous section, OB stars entered the LB region $10 - 15$ Myr ago, setting the clock for its origin to $t=0$. From a further fit to Schaller et al.'s isochrone data we estimate that the main sequence life time of such bright stars scales with mass as \\begin{equation} \\tau=\\tau_0\\, {\\mathcal{M}}^{-\\alpha} \\quad (2\\,{\\mathcal{M}}_\\odot \\leq {\\mathcal{M}}\\leq 67\\,{\\mathcal{M}}_\\odot ) \\label{fuchs_formula} \\end{equation} with $\\tau _0 = 1.6 \\cdot 10^{8 }$yr and $\\alpha = 0.932$. This means that the masses of the most massive stars $\\mathcal{M}_{\\Delta \\tau}$ in the associations at a lookback time of $\\Delta \\tau $ years ago are given by \\begin{equation} \\mathcal{M}_{\\Delta \\tau}=\\left( \\mathcal{M}_{\\rm u}^{-\\alpha}-\\frac{\\Delta \\tau}{\\tau_0} \\right)^{-\\frac{1}{\\alpha}}\\,, \\end{equation} implying $\\mathcal{M}_{10}=15.4\\,\\mathcal{M}_\\odot $ if $\\Delta \\tau= 10$ Myr or $\\mathcal{M}_{15}=26.6\\,\\mathcal{M}_\\odot $ if $\\Delta \\tau= 15$ Myr depending on the entry time of the associations into the volume occupied by the LB today. The expected number of supernovae, i.e. the number of 'missing' stars, is then calculated by \\begin{equation} N_{\\rm SN}=\\int_{8.2}^{\\mathcal{M}_{\\Delta \\tau}} \\frac{d N }{d \\mathcal{M}}\\Big|_0 \\mathcal{M}^{-2.1} d \\mathcal{M}\\,. \\end{equation} We thus obtain estimates of $N_{\\rm SN} = 8 - 12$ from the Upper Centaurus Lupus and $N_{\\rm SN} = 6 - 8$ from the Lower Centaurus Crux associations, respectively. The estimate of 14 to 20 supernovae, which created the LB, is in good agreement with the values determined by Ma\\'{i}z-Apell\\'{a}niz (2001) and BB02. Extrapolating the IMF to masses beyond ${\\mathcal{M}_{\\Delta \\tau}}$ we estimate that 12 to 5 SNe exploded before the associations entered the present LB volume.  We have noted above that our original sample is complete in radial velocities for stars with colours $(B-V) < -0.1$ which corresponds to $M_{\\rm V}=0.7$, if an extinction of $A_{\\rm V}=0.1$ is assumed. According to Eq.~(2) such stars  have a mass of 2.95 ${\\mathcal{M}}_\\odot$. If we remove 15 stars with colours $(B-V) > -0.1$ from the final sample (79 stars) and modify Eq.~(4) for the cut-off at 2.95  at the low mass end, we find $\\frac{d N } {d \\mathcal{M}}\\big|_{\\rm 0, UCL+LCC}$ = 291 instead of 302. Thus the incompleteness of the original sample has not introduced any significant bias in our sample.  In order to assess the question, whether the estimated number of supernovae would suffice to excavate the LB, we consider the energy input by the supernova explosions into the interstellar gas. According to the initial mass function (3) there are \\begin{equation} d N = \\frac{d N }{d \\mathcal{M}}\\Big|_0 \\mathcal{M}^{\\Gamma -1} d\\mathcal{M} \\label{starnumbers1} \\end{equation} stars in the mass range $(\\mathcal{M}, \\mathcal{M}+d\\mathcal{M})$ with main sequence life times $(\\tau ,\\tau -d\\tau )$. Thus \\begin{equation} d N = \\frac{d N }{d \\mathcal{M}}\\Big|_0 \\mathcal{M}^{\\Gamma -1} \\left(-\\frac{d\\mathcal{M}}{d\\tau}\\right)d\\tau \\,, \\label{starnumbers2} \\end{equation} and the energy input rate is given by (cf.\\ BB02) \\begin{eqnarray} && \\dot{\\mathcal{E}}_{\\rm SN} = \\frac{d}{dt}\\mathcal{E}_{\\rm SN} N_{\\rm SN} = \\mathcal{E}_{\\rm SN}\\frac{dN_{\\rm SN}}{dt} \\\\ && =\\mathcal{E}_{\\rm SN}\\frac{d N }{d \\mathcal{M}}\\Big|_0 \\mathcal{M}^{\\Gamma -1}(-1)\\frac{d}{d\\tau} \\left(\\frac{\\tau}{\\tau_0}\\right)^{-\\frac{1}{\\alpha}}\\frac{d\\tau}{dt}\\,, \\nonumber \\label{en-input-rate1} \\end{eqnarray} where $\\mathcal{E}_{\\rm SN}$ denotes the energy released by a single supernova, $\\mathcal{E}_{\\rm SN} = 10^{51}$ ergs. According to the way we have set up Eq.~(9) $\\frac{d\\tau }{dt }$ is equal to 1. Equation (10) describes the trade--off of the increasing number of supernova progenitors and their increasing main--sequence life times with decreasing mass. Inserting the age--to--mass relation (5) into Eq.~(10) leads then to \\begin{equation} \\dot{\\mathcal{E}}_{\\rm SN} = \\dot{\\mathcal{E}}_{\\rm SN0}\\,t_7^{-\\frac{\\Gamma +\\alpha}{\\alpha}} \\label{en-input-rate2} \\end{equation} with $t_7$ defined as  $t_7 = t/10^7$yr. For the constant $\\dot{\\mathcal{E}}_{\\rm SN0}$ we find \\begin{eqnarray} && \\dot{\\mathcal{E}}_{\\rm SN0} = \\frac{\\mathcal{E}_{\\rm SN}}{\\alpha \\tau_0} \\frac{d N }{d \\mathcal{M}}\\Big|_0 \\left(\\frac{10^7 yr}{\\tau_0}\\right)^{- 0.1803} \\\\ \\nonumber && = 3.5\\cdot 10^{35}\\frac{d N }{d \\mathcal{M}}\\Big|_0 {\\rm erg/s}\\,. \\nonumber \\label{en-input-rate3} \\end{eqnarray} Equation~(\\ref{en-input-rate2}) shows a rather weak decline of the supernova energy input rate into the LB as a result of partial compensation between the increasing number of stars with decreasing mass and a corresponding increase in main sequence life time. It is quite remarkable -- although probably fortuitous -- that the distribution of stellar masses during the star formation process is nearly anti-correlated with the main sequence life time of stars.  BB02 have derived a bubble wind equation which describes the growth of the size of the bubble with time. The radius of the bubble is given by \\begin{equation} R_{\\rm b}(t) = R_{\\rm b0}\\,t^{\\frac{2 \\alpha-\\Gamma}{5 \\alpha}} \\label{simsol} \\end{equation} with the constant \\begin{equation} R_{\\rm b0}= \\Big[  \\frac{475 \\alpha}{(4\\alpha-7\\Gamma)(3\\alpha-4\\Gamma)}\\Big]^\\frac{1}{5} \\times \\Big[  \\frac{\\alpha\\dot{\\mathcal{E}}_{\\rm SN0} \\, \\tau_0^{1+\\Gamma/\\alpha}}{2 \\pi (2-\\Gamma)\\rho_0} \\Big]^\\frac{1}{5}\\,. \\label{simsol_const} \\end{equation} We note in passing that Eqs.~(13) and (14) are consistent with Eqs.~(14) - (16) of BB02, except for a different value of $\\alpha$ used here, and a normalization error in BB02, where $L_0$ and $\\rho_0$ should be replaced by $\\tilde L_0 = L_0/t_0^\\delta$ with $\\delta=-(1+\\Gamma/\\alpha)$ and $\\tilde \\rho_0 = \\rho_0/R_0^\\beta$. In Eq.~(13) a constant density $\\rho_0$ of the ambient interstellar gas is assumed for which we adopt a value of $\\rho_0 = 2 \\cdot 10^{-24}$ g/cm$^3$. The index in Eq.~(13), ($2\\alpha  - \\Gamma )/5\\alpha = 0.564$, lies between the index of 0.4 of the Sedov equation, describing supernova remnants, and the index of 0.6 of the stellar wind/superbubble expansion law. For a LB age of 10 to 15 Myr Eq.~(13) predicts a bubble radius of 78~pc to 100~pc, respectively. This in good agreement with the observed size of the LB in the Galactic disk, as determined by Lallement et al.~(2003; cf. also Fig.~7). For the determination of the expected LB size we have used the expected numbers of supernovae both from the LCC and UCL subgroups. Ma\\'{i}z-Apell\\'{a}niz (2001) has argued that the LB owes its existence only to the 6 SNe stemming from the LCC subgroup, because stars from this subgroup came closest to the Sun in the past. We find the same when tracing the orbits of the stars backwards in time. However, the members of the UCL subgroup did enter the region occupied by the LB today and SNe stemming from the UCL subgroup have to be taken into account, in our view, in the energy considerations as well. The energy input of 6 SNe would excavate a bubble with radius of only 65 pc, which is more difficult to reconcile with the fact that the walls of the LB have been blown out above and below the Galactic plane so that the LB has become effectively a chimney. In general, however, similarity solutions as applied here can only give a rough estimate of the LB age and size due to several severe restrictions. Firstly, the ambient medium has to be assumed to be either homogeneously distributed or to follow a power law distribution in density, and its pressure has to be small compared to the bubble pressure. Secondly, turbulent mixing and mass loading, which occur in real bubbles, are hard to incorporate without further assumptions (cf.\\ Dyson et al.~2002). Therefore the most realistic approach to model \\emph{existing bubbles} is to perform 3D high--resolution numerical simulations of their formation. A first simulation of this kind was carried out by Breitschwerdt \\& Avillez (2006) which was based on the older and less detailed LB formation scenario of BB02. In the next section we present an upgrade of that simulation which is now based on the better understood supernova rate and the calculated paths of their progenitors through the LB as derived in this paper.  \\subsection{High Resolution Simulations of the Local Bubble Evolution}  We have simulated the effects of the explosions of the stars formerly belonging to the UCL and LCC subgroups as their trajectories have crossed the LB volume towards their present positions. The crucial physical boundary conditions we have to apply to our simulations are the locations, the masses, and derived from this, the explosion times of the supernovae responsible for the origin of the LB. The latter can be inferred from Eqs.~(\\ref{starnumbers1}) and (\\ref{en-input-rate1}) to be \\begin{equation} \\frac{d N}{d\\tau} = \\frac{1}{\\alpha \\tau_0}\\frac{d N }{d \\mathcal{M}}\\Big|_0 \\, \\left(\\frac{\\tau}{\\tau_0}\\right)^{-\\frac{\\Gamma + \\alpha}{\\alpha}} \\,, \\label{snrate} \\end{equation} which can be integrated to \\begin{equation} N(\\tau) = \\frac{1}{\\Gamma} \\frac{d N }{d \\mathcal{M}}\\Big|_0 \\, \\left[ \\mathcal{M}^{\\Gamma}_{\\Delta \\tau} - \\left(\\frac{\\tau}{\\tau_0}\\right)^{-\\frac{\\Gamma}{\\alpha}}\\right] \\,, \\label{snnumber} \\end{equation} taking the mass of the most massive star as an upper boundary. We then proceed, somewhat arbitrarily, to bin the number of exploded stars between $\\mathcal{M}_u$ and $\\mathcal{M}_{\\Delta \\tau}$ modulo integer solar masses, and derive their main sequence and hence explosion times from Eq.~(\\ref{fuchs_formula}). Next, the explosion locations are fixed by assuming that the presently ``missing stars'' were following the centres of mass of their respective subgroups.  The 3D high--resolution simulations are based on a hydrodynamical Godunov scheme (cf.\\ Godunov \\& Ryabenki 1964) supplemented by adaptive mesh refinement (AMR) along the lines described by Avillez \\& Breitschwerdt (2004), Breitschwerdt \\& Avillez (2006). This entails a detailed treatment of the evolution of the interstellar gas in a volume of the Galaxy with a square area of 1 kpc$^{2}$ and a vertical extent of 10 kpc on either side of the Galactic midplane based on the 3D SN--driven ISM model of Avillez (2000) and Avillez \\& Breitschwerdt (2004). In these calculations the ISM is disturbed by background supernova explosions at the Galactic rate. Initial conditions for the \\emph{ambient medium} were chosen from a data cube of a previous hydrodynamical run where the highest adaptive mesh refinement resolution was 1.25 pc (Avillez \\& Breitschwerdt 2004, Breitschwerdt \\&  Avillez 2006). As a specific boundary condition we have to include the simultaneous evolution of the Loop~I superbubble, which has been observed to interact with the Local Bubble according to ROSAT PSPC observations (Egger and Aschenbach 1995). We therefore selected a site with enough mass to form all the high mass stars which are expected to explode as supernovae. Using the same IMF for Galactic OB associations we derived in total 81 stars with masses ${\\mathcal{M}}$ between 7 and 31 ${\\mathcal{M}}_\\odot$ which in our simulations compose the Sco Cen cluster; 39 massive stars with $14 \\leq {\\mathcal{M}} \\leq 31 \\,{\\mathcal{M}}_\\odot $ have already gone off, generating the Loop~I cavity (see Egger 1998, see also Avillez \\& Breitschwerdt 2005a). Presently the Sco Cen cluster, which is located at $(375,400)$ pc in the top panel of Figure~\\ref{lb1}, hosts 42 stars to explode within the next 13 Myr. Periodic boundary conditions are applied along the four vertical boundary faces of our computational volume, while outflow boundary conditions are imposed at the top ($z=10$ kpc) and bottom ($z=-10$ kpc) boundaries. The simulation time of this run was 30 Myr.  \\begin{figure} \\centering \\caption[]{ Temperature (top panel) and pressure (bottom panel) distributions in the Galactic midplane 13.4 Myr after the first explosion in UCL occurred. The pressure is given in units of cm$^{-3}\\,K$, i.e.~divided by Boltzmann's constant $k$. The dimensions and morphology of the Local Bubble are similar to the present observations. Loop~I, to the right of the LB, is bounded by an X--ray illuminated shell (top panel). \\label{lb1}} \\end{figure}   Figure~\\ref{lb1} shows the temperature (top) and pressure (bottom) distributions in the Galactic midplane 13.4 Myr after the explosion of the first SN, only a few thousand years after the last UCL and LCC supernovae with masses of 8.2 ${\\mathcal{M}}_\\odot$ have exploded. This can be seen as a red spot at ($x, y$) = (200, 300) pc. The LB is located in the region between $100 \\leq x\\leq 300$ pc and $250 \\leq y\\leq 550$ pc, its centre being located at $(x,y) = (200,400)$ pc. The shock waves of the last two SNe occurring within the LB are most noticeable in the $P/k$ distribution by the high pressure peak shown in the bottom panel. To the right of the LB the shell of Loop~I can be seen, which due to its high temperature will emit in soft X--rays (top panel), consistent with {\\sf ROSAT PSPC} observations.  Another striking feature in Fig.~\\ref{lb1} (bottom panel) are the coherent bubble structures within a highly disturbed background medium with a pressure in the range $2 \\leq \\log(P/k) \\leq 4 $  which are due to the locally enhanced SN rates in the vicinity of the Sun and in the Loop~I region. The successive explosions close to the Sun heat and pressurize the LB, which at first looks smooth, but develops internal temperature and density structures at later stages. About 13.4 Myr after the first explosion the LB cavity, which is bounded by an outer shell will start to fragment due to Rayleigh--Taylor instabilities, in agreement with a linear stability analysis carried out by Breitschwerdt et al. (2000). It then fills a volume roughly corresponding to the present day LB size.  A more detailed analysis of these results and their observational consequences will be the subject of forthcoming papers. ",
        "conclusions": "In contrast to previous analyses of the origin of the local bubble we have not merely selected presently known stellar subgroups and traced their kinematics back in time. Instead we have scrutinized \\emph{ab initio} a large sample volume of stars for stellar groups by analyzing their spatial and kinematical properties. From such an unbiased search among nearby B stars we confirm the rather robust result that besides the Upper Scorpius (US) subgroup the Upper Centaurus Lupus (UCL), and Lower Centaurus Crux (LCC) subgroups with ages of 20 - 30 Myr are the youngest stellar associations in the solar neighbourhood. Our search volume is presently limited to a diameter of 400 pc, because the {\\sf Hipparcos} parallaxes are not accurate enough at distances larger than 200 pc. Hence the analysis of a larger volume has to await the launch of {\\sf GAIA}.  Our search strategy relied mainly on kinematical criteria, and we found many other B stars with the same kinematics as the subgroups.  We have followed the paths of the associations into the past and find that they entered the region of the present LB 10 to 15 Myr ago. Deriving O{\\sc vi} column densities from a numerical simulation of the general ISM (Avillez  \\& Breitschwerdt 2005b) as well as of LB and Loop~I evolution (Breitschwerdt \\& Avillez 2006) in a realistic background medium, excellent agreement was found with O{\\sc vi} absorption line data obtained with {\\sf FUSE} (Oegerle et al.~2005, Savage \\& Lehner 2006). According to numerical LB evolution simulations by Breitschwerdt \\& Avillez (2006), who used supernovae from the subgroup B1 of the Pleiades to power the LB, the O{\\sc vi} data can be fitted with a LB age of $14.4\\pm^{0.7}_{0.4}$ Myr. The age of 13.9 - 14.1 Myr estimated from the present simulation is thus consistent with the age estimated from the slightly different simulation by Breitschwerdt \\& Avillez (2006). We therefore conclude that the LB must have been excavated during this time. We find that about 14 to 20 SNe originated from the associations LCC and UCL. The implied energy input into the ambient interstellar gas explains quantitatively the present size of the LB.  The LB serves as an ideal test laboratory for superbubble models due to the wealth of observations against which they can be tested. Apart from the important O{\\sc vi} test, we will also compare EUV and soft X--ray emission data with our models in order to derive the excitation history of ions in the LB and a possible deviation from collisional ionisation equilibrium."
    },
    "0609/astro-ph0609157_arXiv.txt": {
        "abstract": "{} {We search for Polycyclic Aromatic Hydrocarbon (PAH) features towards young low-mass (T Tauri) stars and compare them with surveys of intermediate mass (Herbig Ae/Be) stars. The presence and strength of the PAH features are interpreted with disk radiative transfer models exploring the PAH feature dependence on the incident UV radiation, PAH abundance and disk parameters.} {Spitzer Space Telescope 5--35 $\\mu$m spectra of 54 pre-main sequence stars with disks were obtained, consisting of 38 T Tauri, 7 Herbig Ae/Be and 9 stars with unknown spectral type.} {Compact PAH emission is detected towards at least 8 sources of which 5 are Herbig Ae/Be stars.  The 11.2 $\\mu$m PAH feature is detected in all of these sources, as is the 6.2 $\\mu$m PAH feature for the 4 sources for which short wavelength data are available. However, the 7.7 and 8.6 $\\mu$m features appear strongly in only 1 of these 4 sources. Based on the 11.2 $\\mu$m feature, PAH emission is observed towards at least 3 T Tauri stars, with 14 tentative detections, resulting in a lower limit to the PAH detection rate of 8\\,\\%.  The lowest mass source with PAH emission in our sample is \\object{T Cha} with a spectral type G8. All 4 sources in our sample with evidence for dust holes in their inner disk show PAH emission, increasing the feature/continuum ratio. Typical 11.2 $\\mu$m line intensities are an order of magnitude lower than those observed for the more massive Herbig Ae/Be stars. Measured line fluxes indicate PAH abundances that are factors of 10--100 lower than standard interstellar values. Conversely, PAH features from disks exposed to stars with $T_{\\rm eff}\\leq 4200$~K without enhanced UV are predicted to be below the current detection limit, even for high PAH abundances.  Disk modeling shows that the 6.2 and 11.2 $\\mu$m features are the best PAH tracers for T Tauri stars, whereas the 7.7 and 8.6 $\\mu$m bands have low feature over continuum ratios due to the strongly rising silicate emission.} {} ",
        "introduction": "Polycyclic Aromatic Hydrocarbons (PAHs) have been observed in a wide variety of sources in our own and external galaxies. Within the Milky Way PAHs are observed in the diffuse medium, dense molecular clouds, circumstellar envelopes, and (proto-)planetary nebulae (see \\citealt{pee04} for a summary). A common characteristic of all of these sources is that they are exposed to copious ultraviolet (UV) photons. The UV radiation drives the molecules into excited electronic states, which subsequently decay to lower electronic states through a non-radiative process called internal conversion, followed by vibrational emission in the available C-H and C-C stretching and bending vibrational modes at $3.3$, $6.2$, $7.7$, $8.6$, $11.2$, $12.8$ and 16.4 $\\mu$m.  Thus, PAH molecules form an important diagnostic of UV radiation.  In recent years, PAH emission has also been detected from disks around young stars in ground-based and {\\it Infrared Space Observatory} (ISO) spectra \\citep{ker00,hon01,pee02,boe03,prz03,ack04}. The PAH emission is thought to originate from the surface layer of a (flaring) disk exposed to radiation from the central star \\citep[cf. models by][ hereafter H04]{man99,hab04}. Indeed, ground-based spatially resolved observations show that the features come from regions with sizes typical of that of a circumstellar disk (radius $<$12 AU at 3.3 $\\mu$m, $<$100 AU at 11.2 $\\mu$m) \\citep{gee04,boe03,hab05}.  Searches for PAHs in disks are important because in addition to being a tracer of the strength of the stellar radiation field and disk geometry, the PAHs also affect the disk structure and chemistry. For example, the high opacity of PAHs at FUV wavelengths \\citep{mat05a} could significantly reduce the stellar UV radiation in the inner disk while photoionization of PAHs produces energetic electrons which are a major heating mechanism for the gas in the upper layers of the disk where the gas and dust temperatures are not well coupled \\citep{jon04,kam04b}.  Detections of PAH features also provide diagnostics of the presence of small grains in the surface layers of disks and the dust evolution through grain growth and dust settling. Evidence for grain growth has been found from modeling of the silicate features from Herbig Ae/Be (hereafter HAeBe) \\citep{boe03} and T Tauri disks (\\citealp{prz03}; \\citealp{kes05}, 2006) and more indirectly from the modeling of the observed H$_2$ emission features from T Tauri disks \\citep{ber04}. PAHs are considered to be on the small end of the size distribution of grains. An important question is whether PAHs have a different timescale for settling and/or growth compared to that of larger silicate/carbon dust grains.  So far, most data obtained on PAHs refer to relatively bright features in the spectra of intermediate mass HAeBe stars. A recent ISO spectroscopic survey has detected PAH features, in particular the most frequently observed 6.2\\,$\\mu$m feature, towards 57\\,\\% of a sample of 46 HAeBe stars \\citep{ack04}. \\citet{mee01} classified the ISO observed spectral energy distributions (SEDs) of these intermediate mass young stars into two groups (group I and II). \\citet{dom03} interpreted these SEDs in the context of a passive disk model with a puffed-up inner rim \\citep{dul01} and proposed that group I sources with larger mid-infrared excess have flaring disks while group II sources are consistent with small and/or self-shadowed disks. \\citet{ack04} showed that the group I sources with strong mid-infrared relative to near-infrared excess display significantly more PAH emission than the group II sources with weaker mid-infrared excesses, consistent with the idea that the PAH emission originates mostly from the disk surface.  The arrival of the Spitzer Space Telescope \\citep{wer04} with the InfraRed Spectrograph (IRS) \\citep{hou04} provides the opportunity to extend these studies to disks around fainter low mass T Tauri stars. For sources of spectral type G and later, the stellar UV field is orders of magnitude weaker than for HAeBe stars, which will directly affect the PAH excitation and emission. On the other hand, enhanced UV radiation has been detected for some T Tauri stars (e.g., Costa et al.\\ 2000, Bergin et al.\\ 2003). Such an enhanced UV field should be directly reflected in the intensity of the PAH features if these molecules are present in normal abundances. Furthermore, ionized PAHs can be excited by less energetic, optical photons from these cooler stars \\citep{mat05a}.  In this paper we present detections of PAH features toward T Tauri stars from our initial set of Spitzer IRS 5-38 $\\mu$m spectroscopic observations that were taken as part of the Spitzer Legacy program ``From Molecular Cores to Planet-Forming Disks'' \\citep{eva03} (hereafter c2d). The c2d targets consist of a large number of sources with infrared excess in five of the nearest large star-forming regions: Chamaeleon, Lupus, Ophiuchus, Perseus and Serpens. In Paper I, the silicate 10 and 20 $\\mu$m features from pre-main sequence stars with disks are presented and analyzed (Kessler-Silacci et al.\\ 2006). The data are used here to search for PAH features and address a number of outstanding questions. For how many low mass stars can PAH features be found and how does this compare to HAeBe stars? Can limits be put on the abundance of PAHs and thus indirectly on that of the smallest grains?  Which factors influence the appearance of PAH features in disks?  Can we quantify any additional UV or optical radiation from the strength of the PAH features?  Section \\ref{sec:obs} describes the sample selection, observations and reduction method. The results for the observed PAH features are presented in Sect.\\,\\ref{sec:results} and discussed in Sects.\\,\\ref{sec:pahfeatures} and \\ref{sec:pahindisks}.  Conclusions are presented in Sect.\\,\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} Spitzer-IRS spectra were obtained for a set of 54 pre-main sequence stars with disks, including 38 T Tauri stars and 7 Herbig Ae stars. The observations are an order of magnitude more sensitive than those used in previous surveys of PAHs in disks. We detect PAH features in at least 3 T Tauri stars, with an additional 14 tentative detections to be confirmed, resulting in a lower limit to the PAH detection rate in T Tauri stars of 8\\,\\%.  Spitzer SL observations are needed to confirm the presence of the PAH features for sources where we currently only see the $11.2$\\,$\\mu$m feature in the SH spectrum. All 4 sources that show hints for inner holes in their dust disk also a show clear 11.2\\,$\\mu$m PAH feature.  The lowest mass source with PAH emission in our sample is \\object{T Cha} with spectral type G8.  The derived 11.2 $\\mu$m line intensities are between a few $\\times 10^{-15}$ and $3.3 \\times 10^{-16}$ W m$^{-2}$, which is typically an order of magnitude lower than what was observed for HAeBe stars with ISO.  Radiative transfer modeling of disks coupled with PAH emission models shows that for stars of late spectral type, the $11.2$ $\\mu$m feature is expected to be the best tracer of the presence of PAHs. The models also show that the 7.7 and 8.6 $\\mu$m PAH features are most affected by veiling by the continuum due to strongly rising silicate emission, resulting in low feature over continuum ratios.  For the small number of T Tauri sources detected as well as for the Herbig stars, the measured PAH line fluxes and feature/continuum ratios are lower than those found from our template disk model which maximizes the PAH emission. Variations of the model parameters indicate that the most likely explanation is lower PAH abundances by factors of 10--100, with geometry affecting these conclusions at the level of a factor of a few. The high feature/continuum ratios for the detected T Tauri stars are due to their abnormally low continuum at 11 $\\mu$m caused by the dust holes in their inner disks.  The template model predictions indicate that the 11.2 $\\mu$m feature strength becomes undetectable at our sensitivity limit when $T_{\\rm eff}<4200$~K (K6). This likely explains the absence of PAH features for the majority of sources in our sample which have spectral types later than K0 ($\\sim$\\,70\\%). If a large fraction of these sources would have excess UV radiation at the level detected for some K7 stars, then the absence of PAH features implies a lower PAH abundance by at least an order of magnitude. The same conclusion holds for the 11 out of 38 sources with spectral types earlier than K6 which are not detected. Geometry affects these conclusions by a factor of a few.  Thus, the lack of PAH detections does not mean that PAHs are absent in these disks; even at lower abundances they do need to be considered since they can affect disk structure, chemistry and gas heating.  In a follow-up paper, the combination of 5-35\\,$\\mu$m spectra for the entire sample with IRAC and MIPS photometry will be used to improve the uncertainty in the PAH detection rate and to compare the PAH sources in terms of their disk SEDs and interpret this with the disk modeling.  Ground-based spectra using instruments like VLT-ISAAC and VLT-VISIR on 8-m class telescopes will be able to both search for presence of the $3.3$\\,$\\mu$m PAH feature as well as better characterize the shape of the 8.6 and 11.2 $\\mu$m PAH feature through higher spectral resolution. Through higher spatial resolution, they will also permit us to put a stronger constraint on the spatial extent of the PAH features from T Tauri disks."
    },
    "0609/astro-ph0609682_arXiv.txt": {
        "abstract": "{} {In large gradual solar energetic particle (SEP) events, especially the ground-level enhancement (GLE) events, where and how energetic particles are accelerated is still a problem.} {By using imaging data from TRACE, Yohkoh/HXT, SOHO/MDI and SOHO/EIT, along with the data from the GOES, Apatity NM, and SOHO/LASCO CME catalog, the evolution of the X5.7 two-ribbon flare and the associated SEP event on 14 July 2000 are studied.} {It is found that the magnetic reconnection in this event consists of two parts, and the induced electric field $E_{rec}$ is temporally correlated with the evolution of hard X-ray and $\\gamma$-ray emission. In particular, the first hard X-ray and $\\gamma$-ray emission peak occurred at 10:22 UT, corresponding to the magnetic reconnection in the western part of the flare ribbons and the maximum $E_{rec}$ of $\\sim$ 9.5 V/cm; the second emission peak at 10:27 UT, corresponding to the eastern part and the maximum $E_{rec}$ of $\\sim$ 13.0 V/cm. We also analyze the SEP injection profiles as functions of time and CME-height, and find two-component injection which may result from different acceleration mechanisms.} {A reasonable conclusion is that reconnection electric field makes a crucial contribution to the acceleration of relativistic particles and the impulsive component of the large gradual SEP event, while CME-driven shocks play a dominant role in the gradual component.} ",
        "introduction": "Solar energetic particles are one manifestation of violent energy releases on the Sun. It is widely accepted that SEP events can be divided into two classes: impulsive and gradual (Cane et al. \\cite{Cane86}; Kallenrode et al. \\cite{Kallenrode92}). The impulsive SEPs are accelerated in flares, while the gradual ones are accelerated by CME-driven shocks (Reames \\cite{Reames90}; Reames \\cite{Reames99}; Kahler \\cite{Kahler01}; Kallenrode \\cite{Kallenrode03}). ``Mixed\" events consisting of both particles accelerated in flares and at the shocks were considered to be nonexistent (Reames \\cite{Reames02}). In particular, Kahler (\\cite{Kahler94}) examined the SEP injection profiles of three ground-level enhancement (GLE) events as the function of the height of the associated coronal mass ejections (CMEs). They suggested that in these events SEP injection appears to only result from a single CME-driven shock and not from the flare impulsive phase.  In 1992, Kallenrode et al. found the most common energetic proton events are those exhibiting both an impulsive phase and shock emission, the ``pure gradual\" events are relatively rare. They suggest that between these two classes there should be ``mixed\" or ``hybrid\" events with both impulsive and shock acceleration. Cliver et al. (\\cite{Cliver95}) then expanded the two-type classification to include the ``mixed\" events, in which the SEP event contains a mixture of flare-accelerated and CME/shock-accelerated particles. Recent observations and studies provide more and more evidence of the existence of mixed SEP events. Cohen et al. (\\cite{Cohen99}) found that the charge states of 12 elements with energies of 12 -- 60 MeV/nucl and source temperatures of (3 -- 6)$\\times 10^{6}$ K, are significantly higher than at lower energies. It rules out acceleration out of the ambient material (Kallenrode \\cite{Kallenrode03}). In addition, Cane et al. (\\cite{Cane03}) find that in some intense SEP events, the time-intensity profiles exhibit two peaks: the earlier one with a high Fe/O ratio, the later one a low Fe/O ratio. A much simpler interpretation is that the earlier one is flare-related and the later one is CME-driven shock related. Dai et al. (\\cite{Dai05}) find that some large gradual SEP events have an initial impulsive component, which may stem from flare acceleration. Shocks could also add substantial protons and contribute to the gradual component (Cane et al. \\cite{Cane86}). And a recent simulation of mixed particle acceleration by Li \\& Zank (\\cite{Li05}) indicated that the SEP time-intensity profile has an initial rapid increase followed by a plateau similar to a pure shock case.  An outstanding solar event occurred on 14 July 2000, ``Bastille Day'', comprising an X-class flare, a fast halo CME, and a large particle intensity enhancement in interplanetary space. It offers an excellent opportunity to extend our knowledge about the acceleration and transport process of solar energetic particles. The GOES X5.7/3B two-ribbon flare, which occurred in the NOAA region 9077 (N22W07) at 10:03 UT, has been thoroughly studied by Aschwanden \\& Alexander (\\cite{Aschwanden01}), Fletcher \\& Hudson (\\cite{Fletcher01}), and Masuda et al. (\\cite{Masuda01}). Their studies mainly focus on the generation of EUV and hard X-ray sources and on the corresponding magnetic structure. One of their important results is that flare ribbons map the footprints of magnetic field lines newly linked by reconnection, along which nonthermal particles bombard the lower atmosphere and lose their energy.  In this paper, we go on to use high-cadence observations from the Transition Region and Coronal Explorer (TRACE), along with the longitudinal magnetic field from the Michelson Doppler Imager (MDI) on board SOHO, to evaluate the magnetic reconnection electric field. This enables us to estimate a charged particle's energy gain in the active region. Because this event was associated with a fast halo CME, we also analyzed the SEP injection profiles as the functions of time and the CME-height. We then discuss the roles of flare and CME-driven shock in producing energetic particles. This paper is organized as follows: In Sect. 2, we analyze the reconnection electric field, comparing it with hard X-ray emission, and investigate the SEP injection profiles. Section 3 is dedicated to discussions, and conclusions are drawn in Sect. 4. ",
        "conclusions": "In this paper, we have evaluated the reconnection electric field inferred from the 14 July 2000 X5.7 two-ribbon flare, and analyzed the associated SEP event. Based on the above analysis, our conclusions can be briefly summarized as follows:  \\begin{enumerate}  \\item The magnetic reconnection in this event consists of two parts, corresponding to the evolution of the two parts of the flare ribbons, respectively. The induced electric field is temporally correlated with the evolution of hard X-ray and $\\gamma$-ray emission, which indicates that the reconnection electric field plays an important role in accelerating nonthermal particles. Given the maximum $E_{rec} \\sim 13.0$ V/cm, acceleration length $l_{acc} \\sim 7.0\\times10^{7}$ cm is needed to accelerate protons to GeV energy in the RCS by DC electric field.  \\item The injection of SEPs is composed of two components. The impulsive component occurs very close to the Sun, and the solar protons' release time is consistent with the hard X-ray, $\\gamma$-ray emission peak and the inferred maximum reconnection electric field.  \\item The reconnection electric field probably makes a crucial contribution to the acceleration of relativistic particles and to the impulsive component of the large gradual SEP events, while CME-driven shocks play a dominant role in the gradual component. \\end{enumerate}"
    },
    "0609/astro-ph0609361_arXiv.txt": {
        "abstract": " ",
        "introduction": "The current standard model of structure formation, in which the Universe is dominated by cold dark matter, predicts that the radial density profile of dark matter halos is best described by the functional form proposed by \\citet{Navarro1997} (hereafter NFW):  \\begin{equation} \\label{NFWspherical} \\rho(r) = \\frac{\\rho_c}{\\left( r/r_s \\right)^\\alpha(1+r/r_s)^{3-\\alpha}} \\ , \\end{equation}  \\noindent with the two fitting parameters $\\rho_c$ and $r_s$, and $\\alpha$ was found to be equal to unity, i.e. $\\alpha=1$.  During the last decade great strides have been undertaken to verify and better understand the origin of this density profile of dark matter halos forming in cosmological $n$-body simulations. The functional form proposed by NFW is independent of cosmology and the two fitting parameters $\\rho_c$ and $r_s$ are actually related.  Equation~\\ref{NFWspherical} therefore describes a one-parameter family of functions. While \\cite{Navarro1997} still relied on systems containing of the order $10^4$ particles (and hence resolving the profile down to about 10\\% of the virial radius, according to the convergence criteria of \\citet{Power2003}) studies containing about two orders of magnitude more particles within the virial radius were only able to confirm their findings (e.g. \\citealt*{Fukushige1997}; \\citealt*{Fukushige2001}; \\citealt*{Moore1998}; \\citealt*{Jing2000}; \\citealt*{JingSuto2000}; \\citealt*{Jing2002}). However, there has been a debate about the exact value of the logarithmic inner slope ranging from $\\alpha \\sim 1$~to~$1.5$ (e.g., \\citealt{Navarro1997}, \\citealt{Moore1999}). Recently a great deal of effort has gone into even higher resolution simulations which have revealed density profiles of cold dark matter halos down to scales well below one percent of the virial radius (\\citealt*{Fukushige2004}; \\citealt*{Tasitsiomi2004}; \\citealt*{Navarro2004}; \\citealt*{Reed2005}; \\citealt*{Diemand2004}). The highest resolved simulation to-date reached  an effective mass resolution of about 130 million particles in a cluster sized dark matter halo still supporting the evidence for a central cusp with a logarithmic inner slope of about $\\alpha \\sim 1.2$ (\\citealt*{Diemand2005}). However, these steep central density cusps are not supported by observations. By measuring the rotation curve of a galaxy, one can in principle infer the density profile of its dark matter halo. High resolution observations of low surface brightness galaxies are though best fit by halos with a core of constant density (Simon~et al. 2005; Swaters~et al. 2003; de Block \\& Bosma 2002). This discrepancy between the prediction of theory and actual observations has caught the attention of cosmologists for nearly a decade now and is far from being resolved.  Aside from addressing the mass distribution of dark matter halos from a simulator's point of view, a fair amount of labour went into a theoretical understanding of the origin of the NFW profile equation~(\\ref{NFWspherical}). But only when making strong assumptions about a connection between the density and the velocity dispersion can one can make analytical predictions. For example, under the assumption that the phase-space density is a power law in radius as suggested by \\cite{Taylor2001} one can solve the spherical Jeans equation to obtain the density profile (\\citealt*{Dehnen2005}, \\citealt*{Austin2005}, \\citealt*{Hansen2004}). These studies, guided by the results of the numerical simulations, confirm the inner logarithmic slope to be in the range $\\alpha \\sim 1$ to $2$. But see also \\cite{Hansen2006} for recent claims that equilibrated dark matter structures have a central slope of $\\alpha \\sim 0.8$.   It needs to mentioned that (nearly) all these studies are based upon the assumption of spherical symmetry utilizing radially distributed shells. But \\cite{Jing2002} already noted that halo density profiles as found in cosmological simulations are better fit by applying an elliptical or triaxial modeling of dark matter halos. Detailed investigations of the distributions of shapes of halos forming in cosmological $n$-body simulations now demonstrate that the majority of objects has intrinsic shapes best characterized by axis ratios of about 1:0.74:0.64 (e.g., \\citealt{Kasun2005}; \\citealt{Bailin2005}, \\citealt{Allgood2006}). \\cite{Oguri2005}, for instance, showed that such prolate mass distributions can have a significant influence on the interpretation of lensing signals observed in galaxy clusters, i.e. halo triaxiality may act to bias the profile constraints derived under the assumption of a spherically symmetric mass distribution. It therefore appears mandatory to systematically study the error introduced by the supposition that halo profiles are spherical while their actual distribution is triaxial. Can the debate about the central cusp be connected to this error? Answering this question is the major motiviation for the current study. ",
        "conclusions": "Motivated by the findings of cosmological simulations that dark matter halos have triaxial mass distributions (e.g., \\citealt{Kasun2005}, \\citealt{Bailin2005}, \\citealt{Allgood2006}) we investigated the errors in the derivation of the density profile under the assumption of spherical symmetry. This supposition is rather common and nearly all results presenting evidence for the \"cuspy\" behaviour of dark matter halos rely upon this premise. While others already showed that a spherical assumption will lead to discrepancies in the interpretation of observational data (e.g., \\citealt{Jing2002}; \\citealt{Oguri2005}) we presented an estimator to quantify this error.  We set out to shed light on the question whether the discrepancy and debate about the exact value of the logarithmic central slope (e.g., \\citealt{Navarro1997}, \\citealt{Moore1999}) can be related to the error introduced by the simplifying assumption of spherical symmetry. We now conclude that the inner parts are only marginally affected and the error will be too small to explain variations in the slope as large as $\\pm 0.5$. A change in $\\alpha$ from, for instance, $1.5$ down to $1$ entails a difference in the central density of a factor of about 3 (at approx. 10\\% of the scale radius) rather than the low percentage level indicated by equation~(\\ref{error}) and seen in figure~\\ref{fig:error}. However, there are claims that dark matter halos tend to be more elliptical in the central regions (\\citealt{Dubinski1991}, \\citealt{Jing2002}, \\citealt{Kazantzidis2004}, \\citealt{Bailin2005b}, \\citealt{Allgood2006}, \\citealt{Power2006}). But those shape changes are only minor and will hence not be able to lift the error to the required level.   We further provided a fitting function that allows to estimate the error introduced by spherically averaging a triaxial mass distribution as a function of radius for (almost) arbitrary axis ratios $a/c$ and $b/c$."
    },
    "0609/hep-th0609193_arXiv.txt": {
        "abstract": "I review recent progress in defining a probability measure in the inflationary multiverse. General requirements for a satisfactory measure are formulated and recent proposals for the measure are clarified and discussed. ",
        "introduction": "String theory appears to have a multitude of solutions describing vacua with different values of the low-energy constants. The number of vacua in this vast ``landscape'' of possibilities can be as large as $10^{500}$ \\cite{BP,Susskind,Douglas}. In the cosmological context, high-energy vacua drive exponential inflationary expansion of the universe. Transitions between different vacua occur through tunneling and quantum diffusion, with regions of different vacua nucleating and expanding in the never-ending process of eternal inflation \\cite{AV83,Linde86}. As a result, the entire landscape of vacua is explored.  If indeed this kind of picture describes our universe, we will never be able to calculate all constants of Nature from first principles. At best we may only be able to make statistical predictions. The key problem is then to calculate the probability distribution for the constants. It is often referred to as {\\it the measure problem}.  The probability $P_j$ of observing vacuum $j$ can be expressed as a product \\beq P_j=P_j^{(prior)} f_j \\label{Pj} \\eeq where the prior probability $P_j^{(prior)}$ is determined by the geography of the landscape and by the dynamics of eternal inflation, and the selection factor $f_j$ characterizes the chances for an observer to evolve in vacuum $j$. The distribution (\\ref{Pj}) gives the probability for a randomly picked observer to be in a given vacuum.  It seems natural to identify the prior probability with the fraction of volume $P_j^{(V)}$ occupied by a given vacuum and the selection factor with the number of observers $n_j^{(obs)}$ per unit volume\\footnote{The product in (\\ref{Pj}) should of course be properly normalized.} \\cite{AV95}, \\beq P_j^{(prior)}\\propto P_j^{(V)}, \\label{old1} \\eeq \\beq f_j \\propto n_j^{(obs)}. \\label{old2} \\eeq This approach, however, encounters a severe difficulty: the result sensitively depends on the choice of a spacelike hypersurface (a constant-time surface) on which the distribution is to be evaluated. This problem was uncovered by Andrei Linde and his collaborators when they first attempted to calculate volume distributions \\cite{LLM,LM,GBL}. It eluded resolution for more than a decade, but recently there have been some promising developments, and I believe we are getting close to completely solving the problem. Here, I will briefly discuss the nature of the dificulty and then review the new proposals for $P_j$. Most of this discussion is based on my work with Jaume Garriga, Delia Schwartz-Perlov, Vitaly Vanchurin, and Serge Winitzki \\cite{GSPVW,Vitaly} (see also \\cite{landscape1}). ",
        "conclusions": "The above definition of the measure is just a proposal. We have not derived it from first principles. In fact, there is no guarantee that there is some unique measure that can be used for making predictions in the multiverse. How, then, can we ever know that we made the right choice out of all possible options?  What I find encouraging is that even a single definition of measure that satisfies some basic requirements proved very difficult to find. It is also reassuring that alternative prescriptions suggested in \\cite{ELM} and \\cite{Vitaly} turned out to be equivalent to the pocket-based measure of \\cite{GSPVW}.  Here, we required that the measure should not depend on any arbitrary choices, such as the choice of gauge or of a spacelike hypersurface, and that it should be independent of the initial conditions at the onset of inflation. These conditions, however, do not specify the measure uniquely. For example, a flat measure, $p_j={\\rm const}$ for all $j$, clearly satisfies the conditions. It would be interesting to formulate a set of requirements which selects a unique definition of the measure.  Bubbles of different types can generally collide, with domain walls forming to separate the different vacua. Our prescription for the measure needs to be generalized to include these processes. Another necessary extension is to the case where transitions betwen vacua can occur through quantum diffusion. (Some steps in this direction have been made in \\cite{GSPVW}.)  The ultimate test of any proposed measure will be a comparison of its predictions with observations. The first attempts in this direction have already produced some intriguing results \\cite{Aguirre,Freivogel,Hall1,QLambda,Tegmark2,Hall2,SPV,Hawking}.   ~~~~~~~~~~~~~~~~~~~  I am grateful to Raphael Bousso, Jaume Garriga, Alan Guth, Delia Schwartz-Perlov, Leonard Susskind, Vitaly Vanchurin and Serge Winitzki for discussions and useful comments.  This work was supported in part by the Foundational Questions Institute and by the National Science Foundation."
    },
    "0609/astro-ph0609271_arXiv.txt": {
        "abstract": "The concept of expanding space has come under fire recently as being inadequate and even misleading in describing the motion of test particles in the universe. Previous investigations have suffered from a number of shortcomings, which we seek to correct. We study the motion of test particles in the universe in detail, solving the geodesic equations of General Relativity for a number of cosmological models. In particular, we use analytic methods to examine whether particles removed from the Hubble flow asymptotically rejoin the Hubble flow, a topic that has caused confusion because of differing definitions and invalid reasoning. We conclude that particles in eternally expanding but otherwise arbitrary universes do not in general rejoin the Hubble flow. ",
        "introduction": "The paradigm of modern cosmology is the Friedmann-Lema\\^\\i tre-Robertson-Walker (FLRW) model \\citep{friedmann,1931MNRAS..91..483L,1935ApJ....82..284R,walker}. In recent years there has been debate over paradoxical features of the model, and the physical interpretation of its dynamics. Attention has been drawn in particular to superluminal recession velocities \\citep{2001AIPC..555..348D,2001astro.ph..4349D,2006astro.ph..1171C,2006astro.ph..2102S} and to the motion of test particles in expanding space \\citep{2001astro.ph..4349D,2004Obs...124..174W,2006astro.ph..3162G,Peacockweb}. It is to the second of these issues that we turn our attention.  The motion of test particles in the FLRW model is a fascinating illustration of the interaction between physical concepts and quantitative theories. One of the defining characteristics of physics is the mathematical precision of its predictions. Yet there is more to applying physical laws than simply solving equations. In order to make physical laws more transparent and accessible, we use physical concepts that develop an intuition or a mental picture of the scenario. A successful physical concept allows us to shortcut the mathematics, qualitatively understanding a scenario without having to solve the equations. As an example, consider the north pole of a bar magnet approaching a loop of wire. Looking from behind the magnet, we know that an anti-clockwise current will be set up in the wire, but we need not come to this conclusion by solving Maxwell's equations---Lenz's law will give us the answer without the mathematics.  Attached to the equations of the FLRW model is the physical concept that ``space is expanding''. Galaxies, we are taught, are receding not because they are moving through space but because space itself is being stretched between us and the galaxy. On the face of it, this concept gives us a good intuitive understanding of many cosmological phenomena---it helps us understand why the velocity-distance law is linear and why light is redshifted as it moves through the universe. However, it has been attacked recently, most notably by \\citet{2004Obs...124..174W} and \\citet{Peacockweb}, as being inadequate in describing local dynamics. Whilst \\citet{Peacockweb} will allow a global form of the expanding space concept---the total volume of a closed universe increases with time---he contends that: \\begin{quote} there is no local effect on particle dynamics from the global expansion of the universe \\ldots `Expanding space' is in general a dangerously flawed way of thinking about an expanding universe. \\end{quote}  Previous attempts at resolving this debate have suffered from a number of shortcomings. The most common is the overwhelming desire to approximate General Relativity (GR) by something else---Newtonian gravity, Special Relativity (SR) or a weak field limit of GR. This is probably wrong and certainly unnecessary---why approximate when you can use the exact geodesic equations? Another problem is the small range of cosmological models considered---\\citet{2006astro.ph..3162G}, for example, do a marvellous job of solving the geodesic equations and then apply the solution to just two models, both of which are observationally disfavoured. There is also a dangerous reliance on numerical calculations to determine asymptotic ($t \\to \\infty$) behaviour. \\citet{2004Obs...124..174W} makes this point about \\citet{2001astro.ph..4349D}, but doesn't say why analytic solutions to a Newtonian approximation are better than numerical solutions of the exact GR equations.  In this paper we address these problems by treating the asymptotic behaviour of the full GR equations analytically, for general (up to non-oscillating, bounded equations of state) cosmological models. Section \\ref{GRequations} will give the full equations of motion for particles in the FLRW model. In Section \\ref{testparticles} we will analyse test particle motion and correct some errors made in previous work. In Section \\ref{joininghubble} we consider the asymptotic behaviour of particles in the universe, and in particular the notion of joining the Hubble flow. Finally, in Section \\ref{space} we briefly outline the relevance of our results for expanding space. Our aim is not to be the final word in this debate but instead to clarify the relevant features of Robertson-Walker spacetime. ",
        "conclusions": ""
    },
    "0609/astro-ph0609792_arXiv.txt": {
        "abstract": "{\\it Hubble Space Telescope} (\\hst) Wide Field Planetary Camera 2 ({\\it WFPC2}) images of the Virgo Cluster dwarf elliptical galaxy VCC 128 reveal an apparently double nucleus.  The two components, which are separated by $\\sim 32$ pc in projection, have the same magnitude and color.  We present a spectrum of this double nucleus and show that it is inconsistent with one or both components being emission-line background objects or foreground stars.  The most likely interpretation is that, as suggested by \\citet{lauer_etal_96} for the double nucleus in NGC 4486B, we are seeing a nuclear disk surrounding a supermassive black hole.  This is only the second time an early-type dwarf (dE/dSph) galaxy has been suggested to host a SMBH. ",
        "introduction": "\\label{s:intro}  The centers of galaxies often host supermassive black holes (SMBHs) with masses \\Mbh\\ ranging from $\\sim 10^6$ to $\\sim 10^9$ \\Msun\\ \\citep{kor_ric_95}.  SMBH masses exhibit a variety of scaling relations including the \\Msig\\ relation \\citep{geb_etal_00, fer_mer_00} and an $\\Mbh-M_{DM}$ relation, where $M_{DM}$ is the mass of the host dark halo \\citep{ferrar_02, bae_etal_03}.  These correlations are evidence for a coupling between the formation of SMBHs and galaxies.  One key to unraveling this coupling is to identify the seeds of SMBHs in dwarf galaxies.  Ground-based stellar kinematic data of dwarf elliptical galaxies (dEs) \\citep{geh_etal_02} rule out $\\Mbh \\ga 10^7$ \\Msun.  In the M31 globular cluster G1, thought to be the stripped nucleus of an accreted dE, \\citet{geb_etal_02} report finding an intermediate mass black hole (IMBH) with $\\Mbh \\sim 10^4~\\Msun$, consistent with the extrapolation of the \\Msig\\ relation.  Nuclear activity in low-luminosity dwarf galaxies has turned up a number of IMBHs with estimated masses from $10^4$ to $10^6~ \\Msun$ \\citep{fil_ho_03, bar_etal_04, gre_ho_04}. There is instead no evidence of a SMBH in either M33 \\citep{mer_etal_01, geb_etal_01} or NGC 205 \\citep{val_etal_05}. The paucity of SMBHs in dwarf galaxies led \\citet{ferrar_02} to suggest that low mass galaxies are inefficient at forming SMBHs. \\citet{cot_etal_acsvcs8_06}, \\citet{fer_etal_06} and \\citet{weh_har_06} find instead a continuity in scaling relations, with \\Mbh\\ replaced by $M_{\\rm CMO}$, the mass of a central massive object, which is either a SMBH or a compact nucleus.  \\hst\\ does not permit many dynamical \\Mbh\\ measurements below $\\sim 10^6~ \\Msun$.  An alternative approach to finding SMBHs in low surface brightness dwarf galaxies is a morphological one.  In a few galaxies, \\hst\\ has found double nuclei; the clearest examples are in M31 and in NGC 4486B \\citep{lauer_etal_93, lauer_etal_96, lauer_etal_05}.  The two nuclei in M31, dubbed P1 and P2, are separated by $0\\farcs 49$ or 3.6 pc and have unequal surface brightness, with P2 having the lower surface brightness but sitting very close to the global photocenter. In contrast, in NGC 4486B the two nuclei are equal in magnitude, color and displacement from the photocenter, with a total separation of 12 pc. The generally accepted model of the nucleus in M31 identifies P2 as the center of the galaxy surrounding its SMBH and P1 as the bright off-center apoapsis of an eccentric Keplerian disk \\citep{tremai_95}. Self-consistent models and simulations of such nuclei can be constructed \\citep{sal_sta_01, jac_sel_01, sam_sri_02}. \\citet{lauer_etal_96} proposed a similar model for the double nucleus in NGC 4486B with the difference that the SMBH is located on the global photocenter between the two peaks.  They suggested that binary nuclei are morphological indicators of SMBHs.  Here we present another example of a galaxy with a double nucleus, the Virgo Cluster dE VCC 128.  Its nucleus is very similar to that of NGC 4486B.  Section \\ref{sec:hst} presents the archival \\hst\\ data, Section \\ref{sec:sed} presents new spectra and their modeling and in Section \\ref{sec:discussion} we propose that the double nucleus is best explained by a disk surrounding a SMBH at the center of VCC 128 and explore its likely position in the \\Msig\\ plane. ",
        "conclusions": "\\label{sec:discussion}  Since the nucleus is old, the present configuration could not have resulted from recent gas infall.  It is unlikely, though not impossible, that the two nuclei are both globular clusters 32 pc apart in the last stages of merging because of the short lifetime of such a configuration ($\\sim 5$ Myr using the standard \\citet{chandr_43} formula with $\\ln \\Lambda = 4$ and $v_c = 20 \\kms$).  Two well-separated globular clusters projected so close to each other, so close to the center of the galaxy, and so similar in their properties is even more unlikely. The double nucleus in VCC 128 is similar in many ways to that in NGC 4486B.  As in NGC 4486B, a nuclear disk surrounding a SMBH provides the best explanation: this would be a stable, long-lived configuration, and would account for the similar colors and magnitudes of the two components and for their location near the center of the galaxy.  If the double nucleus is a disk orbiting a SMBH, we can estimate a lower limit to \\Mbh\\ by assuming that it is larger than the nuclear disk mass, ${\\rm M}_d = 10^6 \\Msun$. Alternatively we can use the ratio of ${\\rm M}_d$, to \\Mbh\\ in M31 \\citep[${\\rm M}_d/\\Mbh =0.16$][]{tremai_95} and in NGC 4486B \\citep[${\\rm M}_d/\\Mbh = 0.019$][]{kor_etal_97, lauer_etal_96} to estimate $\\Mbh \\sim 6 \\times 10^6 - 5 \\times 10^7~ \\Msun$ in VCC 128. We estimate a velocity dispersion, $\\sigma$, for VCC 128 from the Faber-Jackson relation of dE's \\citep{der_etal_05}.  We obtain $\\sigma \\sim 35-65 \\kms$.  Thus within large uncertainties the postulated SMBH in VCC 128 could satisfy the \\Msig\\ relation, as we show in Figure \\ref{fig:starmsig}.  \\citet{ferrar_02} proposed that galaxies with a dark matter virial velocity $v_{vir} \\la 200 \\kms$ would not be able to form SMBH's. Using her scaling relations between $\\sigma$ and $v_{vir}$, we find that $v_{vir} \\la 65 \\kms$ assuming the $\\sigma$ estimated above.  It is unlikely that there is enough scatter in the scaling relations as to accommodate $v_{vir} \\simeq 200 \\kms$ in VCC 128.  Thus its SMBH probably violates the proposed limit unless VCC 128 has been heavily stripped in the cluster environment.  Curiously however, VCC 128 would still satisfy the relation between $M_{\\rm CMO}$ and $B-$band magnitude $M_B$ \\citep{weh_har_06, fer_etal_06}.  We have shown that the nucleus of VCC 128 is double, with the two components of equal magnitude and color.  We proposed that this can be explained by the presence of a disk surrounding a SMBH.  Simple estimates of its mass all put it in the regime $\\Mbh \\sim 10^6~\\Msun$ or larger.  If VCC 128 is on the fundamental plane, then these estimates put the SMBH at, or above, the \\Msig\\ relation.  However the halo virial velocity would be smaller than the proposed limit for galaxies to be able to form SMBHs.  Amongst the dwarf elliptical/spheroidal population of galaxies, this is only the second example of a system with evidence for a SMBH \\citep{maccarone_etal_05}.  Given the very interesting nature of this object, we suggest that further high-resolution imaging and spectroscopy of VCC128 would be well worthwhile."
    },
    "0609/astro-ph0609337_arXiv.txt": {
        "abstract": "The major obstacle to the direct detection of companions to nearby stars is the overwhelming brightness of the host star. Current instruments employing the combination of adaptive optics (AO) and coronagraphy can typically detect objects within 2$^{\\prime \\prime}$ of the star that are $\\sim 10^{4-5}$ times fainter. Correlated speckle noise is one of the biggest obstacles limiting such high-contrast imaging. We have obtained a series of 284 8 s, AO-corrected, coronagraphically occulted $H$-band images of the star Vega at the 3.63 m AEOS telescope located on Haleakala, Hawaii. This dataset is unique for studying the temporal behavior of speckle noise, and represents the first time such a study on highly corrected {\\it coronagraphic} AO images has been carried out in a quantitative way. We find the speckle pattern to be highly stable in both position and time in our data. This is due to the fact that the AO system corrects disturbances to the stellar wave front at the level where the instrumental wave front errors dominate. Because of this, we find that our detection limit is not significantly improved simply with increased exposure time alone. However, we are able to improve our dynamic range by 1.5-2 magnitudes through subtraction of static/quasi-static speckles in two rotating frames: the telescope pupil frame and the deformable mirror frame. The highly stable nature of speckles will exist for any program using a combination of coronagraphy and high-order AO, and underscores the importance of calibration of non-common path errors between the wave front sensor and the image plane. Such calibration is critical for high-contrast AO systems and we demonstrate this using empirical data. From our data, we are able to constrain the mass of any purported companion to Vega to be less than $\\sim 45 M_J $ at $8$ AU and less than $\\sim 30 M_J$ at 16 AU, radii not previously probed at these sensitivities. ",
        "introduction": "While the indirect detection of more than 200 planets over the past decade \\citep{mpv05,mbf05} has been a significant accomplishment for exoplanetary science, a current major thrust for the field is the direct detection of these objects through high-contrast imaging. Very large telescope apertures, although extremely effective for gathering photons \\citep{l02}, are not a necessity for direct imaging \\citep{osm03}. The key requirement is the suppression of the host star's overwhelming flux. A ``hot jupiter'' will typically be $10^7$-$10^9$ times fainter than its host \\citep{bcb03,b05} and often situated within a fraction of an arcsecond of the star. A promising method for direct imaging of stellar companions involves two techniques working in conjunction: high-order adaptive optics (AO), providing control and manipulation of the starlight \\citep{a94}, and an optimized Lyot coronagraph \\citep{skm01,odn04} to suppress this light. Together, these two techniques can obtain contrast levels of $10^5$ or better within $2^{\\prime\\prime}$. After this level of contrast has been achieved, the major source of noise in the images is correlated speckle noise. Without a coronagraph, \\citet{rwn99} has demonstrated that speckle noise will dominate over photon noise by a factor $\\sim10^4$. With a coronagraph, the contribution from the ``pinned'' speckle noise can be significantly reduced \\citep{slh02,psm03,as04}. Such speckle noise is largely due to non-common path errors (those not measured by the wave front sensor) e.g. small aberrations in the coronagraphic optics. Other sources include AO correction errors such as fitting error, aliasing, lag \\citep{jvc06} and Fresnel wavefront propagation effects.  The Lyot Project (Oppenheimer et al.\\ 2003; Oppenheimer et al.\\ 2004; Oppenheimer et al.\\ 2006a, in preparation) employs an optimized, diffraction-limited Lyot coronagraph \\citep{l39,m96,skm01}. It is deployed at the $D=3.63$ m Advanced Electro-Optical System (AEOS) telescope in Maui, with an AO system equipped with a 941 actuator deformable mirror (DM). The telescope is an altitude-azimuth design, with a beam traveling under the dome floor to a coud\\'e room containing the AO system and the DM  \\citep{rn02}. During an observation, the DM frame of reference is always aligned with the frame of our infrared camera. The frame of the telescope pupil containing the secondary mirror support struts (the ``spiders''), however, rotates with respect to the camera during an observation. For the purposes of this study, we mention that these frames contain two identifiable sources of speckle noise in our observations.  In March 2004, the project began a survey of approximately 100 nearby stars in coronagraphic mode in the $J$,$H$, and $K_s$ bands with the goal of detecting faint companions and disks orbiting the stars. On 2005 May 14 (UTC), we obtained a sequence of 284 AO-corrected, coronagraphically occulted images of the bright star Vega (HD172167, A0V, $V=0^m$). We estimate our uncorrected wavefront errors at $\\sim 150$nm, and we observed with adequate seeing ($r_0=14.1 \\pm 2.5$cm). With {\\it atmospheric} phase errors well under control by the AO, any speckle noise should come largely from the AO system or the coronagraph itself. Such noise largely arises from non-common path errors between the wave front sensor and the image plane, and can be mitigated through careful calibration \\citep{wgs04, hbf03,bfh03}.  Although quasi-static speckle noise has been shown to be the main limitation to earlier instruments \\citep{mdn05,bcl03,brb04} this dataset is unique for studying the temporal characteristics of correlated speckle noise in such a highly corrected system coupled to an optimized Lyot coronagraph.  After describing our observations (\\S2), we will describe our method for quantifying our sensitivity in the following section. We will show that our sensitivity does not increase substantially when many images are simply coadded together, but demonstrate  that the dynamic range can be improved by subtracting those speckles that are static in the telescope pupil frame and also those that are fixed with respect to the DM. ",
        "conclusions": "The relative floor in sensitivity shown in Figure~\\ref{vsimages} is due to the highly correlated and persistent nature of the speckle noise in the images. In contrast to the way the random, uncorrelated noise is suppressed through averaging, the data are subject to a static and quasi-static speckle pattern, placing a limit on the sensitivity of an observing sequence. Among other factors, such an influx of speckle noise is likely due to a known malfunction in a subset of the AEOS deformable mirror actuators. In addition, high-order AO systems in general can suffer from a variety of problems (misalignment, mirror figure errors, etc.), each of which can be a significant source of speckle noise. We anticipate fixed speckles like those discussed here to be a major obstacle for other AO+Coronagraphy programs with similar objectives. But more importantly, this study underscores the primary importance of non-common path errors between the wave front sensor and the science image plane as other authors have discussed \\citep{hbf03, bfh03}. Calibration of these errors is critical for AO imaging to exceed the $\\sim 70\\%$ Strehl level. We have demonstrated this with empirical data.  This paper demonstrates the extent to which fixed speckles can dominate the overall noise budget of high-contrast AO observations, far exceeding the contributions from residual atmospheric effects. As we have shown, an increase in sensitivity (1.5-2  magnitudes) can be obtained by subtracting off individual contributions to the speckle pattern. Another solution may rely on the wavelength-dependent nature of the speckle noise. Although speckles remain fixed in space for a given filter, they will reside in different locations in other passbands, especially if they result from phase errors in or near pupil planes in the optical train. This key feature will allow a faint companion---with a fixed position in all filters---to be selected out from the speckle noise. Thus, the addition of multi-wavelength observations will enhance the efficiency with which companions may be detected \\citep{mdr00,sf02,bcl04,mdn05}. Moreover, an integral field unit \\citep{lqk03} will allow the retrieval of spectra across an image, further disentangling companions from the speckle noise."
    },
    "0609/hep-ph0609214_arXiv.txt": {
        "abstract": "\\noindent We emphasize the inelasticity distribution of events detected at the IceCube neutrino telescope as an important tool for revealing new physics. This is possible because the unique energy resolution at this facility allows to separately assign the energy fractions for emergent muons and taus in neutrino interactions. As a particular example, we explore the possibility of probing second and third generation leptoquark parameter space (coupling and mass).  We show that production of leptoquarks with masses $\\agt 250$~GeV and diagonal generation couplings of ${\\cal O}(1)$ can be directly tested if the cosmic neutrino flux is at the Waxman-Bahcall level. ",
        "introduction": "Leptoquarks are $SU(3)$-colored particles which simultaneously carry non-zero baryon and lepton quantum numbers. They are predicted in several models (such as $SU(5)$~\\cite{Georgi:1974sy} or Pati-Salam $SU(4)$~\\cite{Pati:1974yy}) addressing the unification of the lepton and quark sectors of the standard model (SM). In such models, the masses of the leptoquarks are generically superheavy, on the order of the GUT scale, which puts them out of reach of direct experimental access. Nevertheless, since leptoquarks with electroweak scale masses are not disallowed for any fundamental reason, it is of interest to conduct experimental searches to delimit their properties~\\cite{Dorsner:2005fq}. It is, of course, important to note that in order to avoid rapid baryon decay, the simultaneous trilinear coupling of the leptoquark to a purely hadronic channel needs to be excluded~\\cite{Hinchliffe:1992ad}.  In general, the couplings of the leptoquark need not be generation-diagonal, and the problem of extracting limits on couplings and masses is complicated by the presence of a large-dimensional parameter space.  Experiments at HERA have placed lower limits of ${\\cal O} (300~{\\rm GeV})$ on first generation leptoquark masses, for trilinear couplings of electroweak gauge strength~\\cite{Adloff:2003jm}. Similar bounds, under the same assumptions, have been found at LEP from their search for anomalous 4-fermion vertices~\\cite{Abbiendi:1998ea}. For first generation leptoquark trilinear  couplings which are much smaller than gauge strength (as is the case for the Yukawas in the SM), the mass bounds are greatly weakened.  At the Tevatron, the leptoquarks could be produced in pairs, with identification made through decay topologies. In this way, the bounds are not dependent on the trilinear couplings, except for decay branching fractions. In the case of first~\\cite{Abbott:1997mq} and second~\\cite{Abazov:2006vc} generation leptoquarks, the final state topology consists of 2 hadronic jets + 2 charged leptons, and the resulting lower limits on the leptoquark mass are around 250 GeV.  In the case of the third generation, a lower limit of 219 GeV has been recently reported by the D\\O~Collaboration~\\cite{D0}, by tagging on 2 $b$ jets + missing energy.  For this value of the leptoquark mass, the decay into $t \\tau$ is largely suppressed compared to the $b \\nu_\\tau$ channel, so that the mass bound is nearly independent of even the branching fraction. As the explored mass region becomes larger, the $t \\tau$ channel becomes more available and thus the mass limit obtained is pushed a bit lower (to $\\ge$ 213 GeV).  In this work, we explore the possibility of probing second and third generation leptoquark parameter space (coupling and mass) with the IceCube neutrino detection facility~\\cite{Ahrens:2002dv}. This experiment, located below the surface of the Antarctic ice sheet at the geographic South pole, is required to be sensitive to the best estimates of potential cosmic ray neutrino fluxes.  When completed, the telescope will consist of 80 kilometer-length strings, each instrumented with 60 10-inch photomultipliers spaced by 17~m.  The deepest module is 2.4~km below the surface. The strings are arranged at the apexes of equilateral triangles 125\\,m on a side. The instrumented (not effective!) detector volume is a cubic kilometer. IceTop, a surface array of {C}erenkov detectors deployed over 1~km$^{2}$ above IceCube, augments the deep-ice component by providing a tool for calibration, background rejection and air-shower physics.  The expected energy resolution is $\\pm 0.1$ on a log$_{10}$ scale.  Construction of the detector started in the Austral summer of 2004/2005 and will continue for 6 years, possibly less. At present, data collection by the first 9 strings has begun.  The event signatures are grouped as tracks, showers, or a combination of the two. Tracks include muons resulting from both cosmic ray showers and from charged current (CC) interaction of muon neutrinos. Tracks can also be produced by $\\tau$ leptons arising in ultra-high energy $\\nu_\\tau$ CC interactions. Showers are generated by neutrino collisions $(\\nu_e$ or low energy $\\nu_\\tau$ CC interactions, and all neutral current interactions) inside or near the detector, and by muon bremsstrahlung radiation near the detector.  The experimental situation is greatly simplified for neutrino energy $E_\\nu \\agt 10^6~{\\rm GeV}.$ A cut at this energy is sufficient to reduce the great majority of background from muon bremsstrahlung and tracks arising from muons produced in cosmic ray showers. Moreover, the flux of atmospheric neutrinos is low above this energy~\\cite{Lipari:1993hd}, so this cut generates a very pure sample of extraterrestrial neutrinos. Of particular interest here, for $E_\\nu > 10^{6}~{\\rm GeV}$, there is sufficient energy resolution ($\\pm 0.2$ on a log$_{10}$ scale) to separately assign the energy fractions in the muon track and the hadronic shower, allowing the determination of the inelasticity distribution. Similarly, in the energy decade $10^{6.5} < E_\\nu/{\\rm GeV} < 10^{7.5}$ one can expect good resolution (less than 5\\%) in ``double bang'' events generated by incoming $\\nu_\\tau$'s. Again, this will allow a reasonable measurement of the inelasticity distribution.   In this study, we emphasize the inelasticity $(y)$ distribution of events as an important tool for detection of new physics. In particular, we will find that the $y$ distribution of events generated through resonant leptoquark production differs substantially from the SM prediction. If the event rate for the new physics turns out to be comparable to SM expectations, then the $y$ profile of the measured data can be used to probe the coupling-mass leptoquark parameter space. The outline of the paper is as follows: In Sec.~\\ref{2} we derive the relevant $y$-distribution of events generated through production and decay of a scalar leptoquark under the assumption of diagonal generation coupling. Armed with this distribution, in Sec.~\\ref{3} we present a statistical method for assessing the significance of discovery criteria. Our conclusions are collected in Sec.~\\ref{4}. ",
        "conclusions": "\\label{4}  In this paper we have introduced the measurement of inelasticity as a powerful tool for probing new physics in cosmic neutrino interactions. As an illustrative example, we have discussed the possibility of detecting leptoquark production at the IceCube neutrino telescope~\\cite{Robinett:1987ym}. We estimated the expected event rate at IceCube to be comparable to the one predicted for cosmic ray facilities that make use of the atmosphere as the detector calorimeter~\\cite{EspiritoSanto:2005mx}. However, the ability of IceCube to accurately measure the inelasticity distribution of events provides a unique method for SM background rejection, allowing powerful discrimination of resonant processes: we have shown that production of leptoquarks with masses $\\agt 250$~GeV and diagonal generation couplings of ${\\cal O}(1)$ can be directly tested at the Antarctic ice-cap.   In closing, some comments are in order. First, we have not taken account of any systematic considerations concerning the detector -- these are beyond the scope of the present work. Second, for reasons of simplicity, we have not included upcoming events close to the horizon. Although these Earth-skimming neutrinos have the potential of nearly doubling our signal event background (and thus nearly halving the required observation time scale), their proper consideration will require a full Monte Carlo simulation."
    },
    "0609/nlin0609061_arXiv.txt": {
        "abstract": "\\noindent    Solving the Boltzmann - BGK equation, we investigate a flow generated by an infinite plate oscillating with frequency $\\omega$.  Geometrical simplicity of the  problem allows a solution in the entire range of  dimensionless frequency variation $0\\leq \\omega \\tau\\leq \\infty$, where $\\tau$ is a  properly defined  relaxation time. A  transition from viscoelastic behavior of Newtonian fluid ($\\omega\\tau\\rightarrow 0$)  to purely elastic dynamics in the limit $\\omega\\tau\\rightarrow \\infty$ is discovered. The relation of the derived solutions to microfluidics  (high-frequency micro-resonators) is demonstrated on an example of a  \"plane oscillator\". ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609581.txt": {
        "abstract": "We present detailed time-averaged X-ray spectroscopy in the 0.5--10~keV band of the Seyfert~1.9 galaxy NGC~2992 with the \\suzaku X-ray Imaging Spectrometers (XIS). The source had a factor $\\sim 3$ higher 2--10~keV flux ($\\sim 1.2 \\times 10^{-11} \\rm \\ erg \\ cm^{-2} \\ s^{-1}$) than the historical minimum and a factor $\\sim 7$ less than the historical maximum. The XIS spectrum of NGC~2992 can be described by several components. There is a primary continuum, modeled as a power-law with a photon index of $\\Gamma = 1.57^{+0.06}_{-0.03}$ that is obscured by a Compton-thin absorber with a column density of $8.0^{+0.6}_{-0.5} \\times \\rm 10^{21} \\ cm^{-2}$. There is another, weaker, unabsorbed power-law component (modeled with the same slope as the primary), that is likely to be due to the primary continuum being electron-scattered into our line-of-sight by a region extended on a scale of hundreds of parsecs. We measure the Thomson depth of the scattering zone to be $\\tau_{\\rm es} = (0.073 \\pm 0.021)/[\\Omega/4\\pi]$, where $\\Omega/4\\pi$ is the fraction of the sky covered by the zone (as seen from the X-ray source) that is visible to the observer. An optically-thin thermal emission component, which probably originates in the same extended region, is included in the model and yields a temperature and luminosity of $kT=0.656^{+0.088}_{-0.061}$~keV and $\\sim 1.2 \\pm 0.4 \\times 10^{40} \\rm \\ erg \\ s^{-1}$ respectively. We detect an Fe~K emission complex which we model with broad and narrow lines and we show that the intensities of the two components are decoupled at a confidence level $>3\\sigma$. The broad \\fekalfa line has an equivalent width of $118^{+32}_{-61}$~eV and could originate in an accretion disk (with inclination angle greater than $\\sim 30^{\\circ}$) around the putative central black hole. The narrow \\fekalfa line has an equivalent width of $163^{+47}_{-26}$~eV and is unresolved (FWHM~$<4090 \\rm \\ km \\ s^{-1}$) and likely originates in distant matter. The absolute flux in the narrow line implies that the column density out of the line-of-sight could be much higher than measured in the line-of-sight, and that the mean (historically-averaged) continuum luminosity responsible for forming the line could be a factor of several higher than that measured from the data. We also detect the \\fekbeta line (corresponding to the narrow \\fekalfa line) with a high signal-to-noise ratio and describe a new robust method to constrain the ionization state of Fe responsible for the \\fekalfa and \\fekbeta lines that does not require any knowledge of possible gravitational and Doppler energy shifts affecting the line energies. For the distant line-emitting matter (e.g. the putative obscuring torus) we deduce that the predominant ionization state is lower than Fe~{\\sc viii} (at 99\\% confidence), conservatively taking into account residual calibration uncertainties in the XIS energy scale and theoretical and experimental uncertainties  in the Fe~K fluorescent line energies. From the limits on a possible Compton-reflection continuum it is likely that the narrow \\fekalfa and \\fekbeta lines originate in a Compton-thin structure. ",
        "introduction": "\\label{intro}  The Fe~K emission line profiles in active galactic nuclei (AGN) generally consist of a prominent core centered at $\\sim 6.4$~keV (typically with FWHM~$<10,000  \\rm \\ km \\ s^{-1}$ or so) and in some sources an additional underlying broader component. The relatively narrow core at $\\sim 6.4$~keV is found to be ubiquitous in nearby, low to moderate luminosity AGN (e.g. Nandra \\etal~1997a,b; Sulentic \\etal~1998; Weaver, Gelbord, \\& Yaqoob 2001; Reeves 2003; Page \\etal~2004; Yaqoob \\& Padmanabhan 2004; Jim\\'{e}nez-Bail\\'{o}n \\etal~2005; Zhou \\& Wang 2005; Jiang, Wang, \\& Wang 2006; Levenson \\etal~2006). The broad Fe~K line component potentially carries signatures of the inner accretion disk in the vicinity of the putative black hole by virtue of gravitational and Doppler energy shifts (e.g. Tanaka \\etal~1995; Nandra \\etal~1997a,b; Reynolds 1997; Turner \\etal~1998; Iwasawa \\etal~1999; Yaqoob \\etal~2002; Reynolds \\& Nowak 2003; Fabian \\& Miniutti 2005). The bulk of the Fe~K line core is likely to originate in matter located much further from the central black hole than the outer disk (e.g. the outer parts of the optical broad line region, the putative parsec-scale torus structure, or the optical narrow line region). In addition to the two principal components, there are an increasing number of reports of cases of narrow lines at energies higher than 6.4~keV (from ionized and/or outflowing Fe -- see Bianchi \\etal~2005, and references therein) and lower than 6.4~keV (from apparently redshifted localized emission -- Turner \\etal~2002; Iwasawa, Miniutti, \\& Fabian 2004; Turner \\etal~2006, and references therein).  Deconvolving the disk and distant-matter components of the Fe~K line is extremely challenging because the disk can also contribute to the line core emission (from radii greater than $\\sim 100r_{g} \\equiv 100GM/c^{2}$ or so). In other words, line emission from disks observed at low inclination angles (relative to the disk normal) is partially degenerate with line emission from distant matter. Line emission in the wings of the profile from regions close to the black hole is smeared out over energy and therefore difficult to detect against the continuum. In order to derive the most reliable parameters and constraints for the disk and distant-matter Fe~K line components, {\\it one must simultaneously model both}, otherwise the model parameters will be biased and misleading (e.g. see discussions in Weaver \\& Reynolds 1998, Yaqoob \\& Padmanabhan 2004). So far only some of the brightest AGN have sufficient signal-to-noise ratio to produce useful constraints from dual-line models (e.g. see Yaqoob \\etal~2002, Yaqoob \\& Padmanabhan 2004, and references therein). For the brightest sources the Fe~K line core is usually unresolved even with the \\chandra gratings because high spectral resolution {\\it and} high throughput are required. Therefore, the location of the ``distant'' Fe~K line emitter remains unknown. Nevertheless, decoupling the intensities of the broad and narrow lines (in the sense of showing that the data require both to be non-zero at some confidence level) would be an important step in constraining models of the central engine in AGN. Such decoupling has only been possible for a few of the brightest AGN (e.g. see Fabian \\& Miniutti 2005). We also note that detection of the broad Fe~K line in type~1.5--2 AGN is even more difficult than in type~1 objects because the X-ray spectra of the former class are more complex and are usually affected by significant intrinsic X-ray absorption (e.g. see Turner \\etal~1998; Weaver \\& Reynolds 1998; Matt \\etal~2000; Guainazzi \\etal~2001; Risaliti 2002).  Almost all of our knowledge of the Fe~K line emission in AGN comes from the \\fekalfa line. Although \\fekbeta line emission has been reported in a few type~1.5--2 AGN (Sambruna \\etal~2001; Molendi, Bianchi, \\& Matt~2003; Matt \\etal~2004), constraints on the line energy and equivalent width (EW) have been loose because the branching ratio, \\fekbetap/\\fekalfap, is less than 0.145 (Palmeri \\etal~2003, and references therein). Yet the \\fekbeta line carries {\\it additional} information on the ionization state of line-emitting Fe that would supplement the information carried by the \\fekalfa line.  In this paper we present the results of detailed X-ray CCD spectroscopy of the nearby ($z=0.00771$, Keel 1996) Seyfert~1.9 galaxy NGC~2992 using data from \\suzaku observations. We show that the intensities of the broad and narrow \\fekalfa lines are decoupled for the first time in this source. We also detected the \\fekbeta line corresponding to the \\fekalfa line core and present measurements of its centroid energy, intensity, and EW. We describe a new method that utilizes the \\fekbeta line to deduce tight limits on the ionization state of Fe in the matter responsible for emitting the line core. For the \\fekalfa and \\fekbeta line analysis we modeled the 0.5--10~keV spectrum self-consistently so we also report the broadband results.  NGC~2992 has been observed by every X-ray astronomy mission since the time it was discovered by HEAO-1 to be one of the brightest hard X-ray AGN in the sky, with a 2--10~keV flux of $\\sim 7.2-8.6 \\times 10^{-11} \\rm \\ erg \\ cm^{-2} \\ s^{-1}$ (Piccinotti \\etal~1982). In more than a quarter of a century of X-ray observations, the hard X-ray flux has varied by over a factor of 20, corresponding to a range in the intrinsic 2--10~keV luminosity of $\\sim 0.55-11.8 \\times 10^{42} \\rm \\ erg \\ s^{-1}$ (assuming $H_{0} = 70 \\ \\rm km \\ s^{-1} \\ Mpc^{-1}$, $\\Lambda = 0.73$, $\\Omega = 1$). During an \\asca observation in 1994, NGC~2992 was in its lowest continuum flux state thus far observed, with a 2--10~keV flux of $\\sim 4 \\times 10^{-12} \\rm \\ erg \\ cm^{-2} \\ s^{-1}$ (Weaver \\etal~1996). During that observation the Fe K line equivalent width (EW) was very high ($\\sim 500-700$~eV), indicating that the line intensity had not responded to the declining continuum. In 1997 a \\bsax observation found NGC~2992 to be still in a low continuum state, and the Fe~K line EW was even higher, but then a second \\bsax observation in 1998 revealed that the source continuum had fully recovered to its bright HEAO-1 state (see Gilli \\etal~2000). In this state, the EW of the Fe~K line was very low ($<100$~eV) since the line intensity again did not respond to the change in continuum level. These measurements were mostly sensitive to the Fe~K line core, which from the lack of variability, likely arises in a distant, parsec-scale reprocessor, as suggested by Weaver \\etal~(1996).  The results from an \\xmm observation of NGC~2992 in May 2003 are unpublished. We examined those data and found that they suffered from heavy pile-up. A detailed analysis of the \\xmm data is beyond the scope of this work and will be reported elsewhere.  NGC~2992 was the subject of an \\rxte monitoring campaign consisting of 24 observations in the period 2005 March to 2006 January. The \\suzaku observations that are the subject of the present paper were made during the \\rxte campaign and were quasi-simultaneous with some of the \\rxte observations. The scope of the present paper is restricted to the \\suzaku data only, but a detailed analysis of the data from the \\rxte campaign will be presented elsewhere.  The present paper is organized as follows. In section~\\ref{data} we describe the observations and details of the data reduction. In section~\\ref{spfit} we describe the spectral fitting analysis and in section~\\ref{results} we discuss the detailed results of that analysis, including brief comparisons with historical data and results. In section~\\ref{kakbanal} we describe a new method to tightly constrain the ionization state of Fe using the narrow \\fekalfa and \\fekbeta emissions lines. In section~\\ref{concl} we summarize our results and give our conclusions. ",
        "conclusions": "\\label{concl}  We have presented detailed X-ray spectroscopy in the 0.5--10~keV band of the Seyfert~1.9 galaxy NGC~2992 with the \\suzaku XIS. Here we summarize the results.  \\begin{enumerate}  \\item The mean observed 2--10~keV flux of NGC~2992 during the three \\suzaku observations in 2005 November and December was $\\sim 1.2 \\times 10^{-11} \\rm \\ erg \\ cm^{-2} \\ s^{-1}$, a factor of $\\sim 3$ higher than the historic minimum flux, and a factor of $\\sim 6$ lower than the historic maximum. The absorption-corrected intrinsic 2--10~keV luminosity was $1.6 \\times 10^{42} \\rm \\ erg \\ s^{-1}$.  \\item We detected the extended, optically-thin soft X-ray thermal emission found with previous X-ray astronomy missions, and obtained a temperature of $kT=0.656^{+0.088}_{-0.061}$~keV. This is the most precise measurement of $kT$ for NGC~2992 to date although the errors are purely statistical and do not include any possible systematic effects. We obtained a luminosity for the soft emission component of $\\sim 1.2 \\pm 0.4 \\times 10^{40} \\rm \\ erg \\ s^{-1}$ which is consistent with that deduced from \\chandra data (Colbert \\etal~2005).  \\item We modeled the intrinsic X-ray continuum with a power law and obtained $\\Gamma = 1.57^{+0.06}_{-0.03}$ for the power-law photon index.  \\item We derived a column density of $8.0^{+0.6}_{-0.5} \\times 10^{21} \\rm \\ cm^{-2}$ for the line-of-sight X-ray obscuration, which corresponds to a Thomson depth of $\\sim 0.006^{+0.005}_{-0.004}$ (i.e. the absorber is Compton-thin, at least in the line-of-sight). This column density is consistent with values obtained from previous \\bsax and \\asca observations.  \\item We estimated $\\tau_{\\rm es} = 0.10 \\pm 0.03$ for the optical depth of the extended, warm, electron scattering zone which is responsible for scattering $(7.3 \\pm 2.1)\\%$ of the intrinsic hard X-ray continuum into the line-of-sight.  \\item We detected broad and narrow Fe~K line emission. When we modeled the broad and narrow lines with line emission from a relativistic disk around a Schwarzschild black hole and a Gaussian profile respectively, we found that the intensities of the broad and narrow lines were decoupled at a confidence level of greater than $3\\sigma$. We obtained an EW of $118^{+32}_{-61}$~eV, a radial line-emissivity law of $r^{-1.5 (-0.8,+5.1)}$, and a lower limit of $31^{\\circ}$ on the disk inclination angle from the disk-line model. The actual values are model-dependent, however.  \\item The narrow Fe~K line core EW was $163^{+47}_{-26}$~eV. The line core was unresolved by the XIS and we obtained a 99\\% confidence, two-parameter upper limit on the FWHM of $4090 \\rm \\ km \\ s^{-1}$, after taking into account the time-dependent degradation in the XIS spectral resolution (as discussed in section~\\ref{meanxis} and section~\\ref{deconv}).  \\item The absolute flux in the narrow line is a factor of more than $20$ larger than that expected from a fully-covering spherical shell illuminated by the measured intrinsic 2--10~keV luminosity, having a cosmic abundance of Fe and a uniform column density equal to that measured in the line-of-sight. Therefore, roughly speaking, the product of the solid angle subtended by the line-emitting structure at the source, its angle-averaged column density, its Fe abundance, and the mean historical 2--10~keV luminosity averaged over timescales of the order of the structure light-crossing time, must be $>20$.  \\item We detected the \\fekbeta fluorescent line corresponding to the \\fekalfa line core. Taking into account residual uncertainties in the XIS energy scale calibration and conservative uncertainties on the theoretical energies of the \\fekalfa and \\fekbeta transitions for different ionic species of Fe, we were able to constrain the predominant ionization state of Fe responsible for producing the narrow Fe~K line core to be Fe~{\\sc i}--Fe~{\\sc vii}. Ionization states of Fe~{\\sc viii} or higher are ruled out at 99\\% confidence. This is the most precise determination to date of the ionization state of the matter emitting the Fe~K line core and was achieved using a method that {\\it requires no knowledge of any uncertainties in the XIS gain and any possible gravitational and Doppler shifts}.  \\item We were able to decouple the disk and distant-matter \\fekalfa emission lines because the continuum flux was low, giving a large EW for the line core, with good signal-to-noise ratio. The large core EW also enabled the \\fekbeta line to be measured with a good signal-to-noise ratio, giving redundant information on the width of the distant-matter lines. This demonstrates the importance of measuring the distant-matter Fe~K line emission parameters in order to measure those of the relativistic Fe~K line.  \\item Although the XIS data do not have the bandpass to measure the Compton-reflection continuum associated with the broad Fe~K line, we included it in all of the spectral analysis with a parameterized normalization, $R$, relative to that expected from a steady-state illumination of a disk subtending a solid angle of $2\\pi$ at the X-ray source. We obtained upper limits on $R$ as a function of the disk inclination angle (constrained by the relativistic Fe~K line). We found (at 90\\% confidence for two parameters), $R<0.9$ for an inclination angle of  $80^{\\circ}$, going down to $R<0.6$ when the inclination angle is at its lower limit of $31^{\\circ}$. If all of this reflection is associated with the broad Fe~K line, which presumably originates in a Compton-thick disk, there is no room left for reflection from distant matter. This implies that the narrow Fe~K line core is likely to originate in Compton-thin matter, supported by the lack of a Compton shoulder on the Fe~K line core. Although the EW of the Fe~K line core is too large for steady-state illumination of Compton-thin matter, several possible factors can account for it (see [8] above).  \\end{enumerate}  In the future we will refine our modeling by utilizing the \\suzaku HXD data when the background model has achieved the goal of 1\\% systematic uncertainties. The HXD data will better constrain the Compton-reflection continuum components. We will also improve model constraints by including data from XIS0 and XIS1 when the calibration has advanced further. Future observations with higher spectral resolution, such as gratings or calorimeters, will be able to resolve the narrow core of the \\fekalfa line and thereby better constrain the location of the line-emitter. In addition, improved throughput will enable better measurement of the profile of the broad Fe~K line. Further monitoring of NGC~2992  to measure variability in the Fe~K line components and in the reflection continuum as the continuum undergoes large excursions in amplitude promise to reveal important clues on the accretion process and on the structure of the central engine beyond the accretion disk.  %"
    },
    "0609/astro-ph0609451_arXiv.txt": {
        "abstract": "We describe the angular power spectrum of resolved sources at 3.6 $\\mu$m ($L$-band) in {\\it Spitzer} imaging data of the GOODS HDF-N, the GOODS CDF-S, and the NDWFS Bo\\\"otes field in several source magnitude bins.     We also measure angular power spectra of resolved sources in the Bo\\\"otes field at $K_S$ and $J$-bands using ground-based IR imaging data. In the three bands, $J$, $K_S$, and $L$, we     detect the clustering of galaxies on top of the shot-noise power spectrum at multipoles between $\\ell \\sim 10^2$ and $10^5$. The angular power spectra range from the large, linear scales to small, non-linear scales of galaxy clustering,  and in some magnitude ranges,  show departure from a power-law clustering spectrum. We consider a halo model to describe clustering measurements and  to establish the halo occupation number parameters of IR bright galaxies at redshifts around one. The typical halo mass scale at which two or more IR galaxies with $L$-band Vega magnitude between 17 and 19 are found in the same halo is between $9 \\times 10^{11}$ M$_{\\sun}$ and $7 \\times 10^{12}$ M$_{\\sun}$ at the 1$\\sigma$ confidence level; this is consistent with the previous halo mass estimates  for bright, red galaxies at $z\\sim1$. We also extend our clustering results  and completeness-corrected faint source number counts in GOODS fields to understand the underlying nature of unresolved sources responsible for IR background   (IRB) anisotropies that were detected in deep {\\it Spitzer} images.  While these unresolved  fluctuations were measured at sub-arcminute angular scales, if a high-redshift diffuse component associated with first galaxies exists in the IRB, then it's clustering properties are best studied with shallow, wide-field images that allow a measurement of the clustering spectrum from a few degrees to arcminute angular scales. ",
        "introduction": "The intensity of the cosmic near-infrared background (IRB) is a measure of the total light emitted by stars and galaxies in the Universe. While the absolute background has been estimated by space-based experiments, such as the  Diffuse Infrared Background Experiment (DIRBE; Hauser \\& Dwek 2001) and the Infra-Red Telescope in Space (IRTS; Matsumoto et al. 2005), the total IRB intensity measured still remains fully unaccounted for by sources: only 13.5 $\\pm$ 4.2 nW m$^{-2}$ sr$^{-1}$ is resolved to point sources at 1.25 $\\mu$m (Cambr\\'esy et al. 2001), while current direct measurements range from 25-70 nW m$^{-2}$ sr$^{-1}$. At wavelengths greater than 3 $\\mu$m, with a total IRB intensity of 12.4 $\\pm$ 3.2 nW m$^{-2}$ sr$^{-1}$ (Wright \\& Reese 2000) at 3.6 $\\mu$m, the ``missing source'' problem is less significant with 5.4 nW m$^{-2}$ sr$^{-1}$ resolved to point sources in {\\it Spitzer} data (Fazio et al. 2004a; Franceschini et al. 2006), while fluctuation analyses of {\\it Spitzer} data yield 10.6$^{+0.63}_{-1.95}$ nW m$^{-2}$ sr$^{-1}$ (Savage \\&  Oliver 2005).       We refer the reader to Kashlinsky (2005) for a recent review on the IRB including a summary of past attempts to understand the intensity excess relative to the background predicted by resolved source number counts.  Primordial galaxies at redshifts 8 and higher, especially those involving Population III stars, are generally invoked to explain the missing IR flux between 1 $\\mu$m and 2 $\\mu$m, with most of the intensity associated with redshifted Lyman-$\\alpha$ emission during reionization (e.g., Santos, Bromm \\& Kamionkowski 2002; Salvaterra \\& Ferrara 2003; Cooray \\& Yoshida 2004; Fernandez \\& Komatsu 2006). While models of high-redshift Pop III populations can explain the ``missing'' IRB, these models run into several difficulties if such sources were to account for all of the missing IR intensity. These include the high efficiency required to convert baryons to stars in first galaxies (e.g., Madau \\& Silk 2005) and limits from deep IR imaging data that suggest a lack of a large population of high-redshift dropouts (Salvaterra \\& Ferrara 2006). The recently revised optical depth to reionization (Page et al. 2006), with reionization around a redshift of 10 instead of the previous estimates of 20, decreases the number density of first sources that are required for reionization. In return, the implied fractional contribution from first sources to the total intensity of the IRB  is lowered.  Still, one does expect some contribution to the IRB from sources that reionized the Universe, though the exact intensity of the IRB from such sources is yet unknown both theoretically and observationally. As pointed out in Cooray et al. (2004; also, Kashlinsky et al. 2004), if a high-redshift population contributes significantly to the IRB, then these sources are expected to leave a distinct signal in the anisotropy fluctuations of the near-IR intensity, when compared to the anisotropy spectrum associated with low-redshift sources. Using results from a fluctuation analysis in deep {\\it Spitzer} imaging data with resolved point sources removed from the image down to a deeper magnitude than previous studies on this topic, Kashlinsky et al. (2005) claimed a potential detection of the clustering signature of high-redshift sources at wavelengths of 3.6, 4.5, and 5.8 $\\mu$m. A previous attempt to understand the nature of this excess clustering in {\\it Spitzer} data, when resolved sources are removed down to a magnitude level of 22.5 in the $L$-band,     suggested that it could be the clustering signature of galaxies at redshifts greater than 5, with a total contribution to the IRB at the level of $\\sim$ 1 to 2 nW m$^{-2}$ sr$^{-1}$ in  the $L$-band (Salvaterra et al. 2006); at shorter wavelengths, the intensity of this background remains uncertain as the spectrum of this excess fluctuation component is only established with {\\it Spitzer} for $\\lambda \\geq$ 3.6 $\\mu$m.  While this suggestion was simply based on a model description of the fluctuation spectrum, whether such a  scenario is consistent with faint source counts in deep {\\it Spitzer} images is yet to be considered.  If the 1 to 2 $\\mu$m IRB is not fully resolved by extragalactic sources,  then it could be that most of the missing source problem is due to deficiencies in accounting for the zodiacal light (Dwek, Arendt \\& Krennrich 2005). This possibility may be further consistent with some suggestions that, subject to a priori assumptions on their intrinsic spectral shape, the observed blazar TeV spectra rule out a significant IRB from extragalactic sources (e.g., Aharonian et al. 2006).      While the point source photometry may underestimate the total IR flux in each resolved source by missing the outer regions when constructing source catalogs, it is unlikely that this effect alone can explain the factor of 2 to 3 difference between predicted and measured IRB intensity given that point source photometry is unlikely to miss more than 5\\% to 10\\% of the total flux (see Totani et al. 2001 for a discussion). Such a small residual, however, could produce an excess in fluctuations of the IRB, with the angular spectrum tracing that of the bright resolved sources.  To separate various possibilities related to the nature of intensity as well as anisotropy excesses, further analyses of near-IR background data is required. If the excess background is due to missing flux in low-z galaxies, then one would expect fluctuations in the background to strongly correlate with resolved point sources. The contribution from zodiacal light is known only to within an upper limit, but if residual zodiacal light is the reason, then spatial clustering of the IRB will peak at degree or more angular scales due to its smoothness. Thus, to distinguish between these possibilities, it is necessary to further understand spatial clustering at IRB wavelengths. Here, we concentrate on the fluctuations in IRB due to resolved sources. We study both the shape and the amplitude of the angular power spectrum and use these measurements to make a relative comparison between resolved and unresolved IRB clustering, with the latter from measurements in  Kashlinsky et al. (2005). We also connect our clustering measurements with faint number counts and an underlying halo model of the IR galaxy distribution within the large-scale structure to study if faint galaxies could be responsible for the IRB fluctuations seen with {\\it Spitzer}.  In this paper, we make clustering measurements at 3.6 $\\mu$m in the Great Observatories Origins Deep Survey (GOODS; Dickinson et al. 2003) HDF-N field, the GOODS CDF-S field,  and the NOAO Deep Wide-Field Survey (NDWFS; Jannuzi \\& Dey 1999) Bo\\\"otes field with the source sample divided into several magnitude bins in each of these catalogs. The data come from {\\it Spitzer} IRAC imaging of these fields. We refer the reader to Dickinson et al. (in preparation) for IRAC observations of the GOODS field and Eisenhardt et al. (2004; also see Fazio et al. 2004a) for the Bo\\\"otes-field IRAC shallow survey. In addition to {\\it Spitzer} images, we also measure clustering of resolved sources in the Bo\\\"otes field in ground-based images at $K_S$ and $J$-bands with data taken by the FLAMINGOS Extragalactic Survey (FLAMEX; Elston et al. 2006).  Here we present clustering measurements in terms of an angular power spectrum in multipole space. Due to simplicity in the measurement, galaxy clustering is also measured in real space in terms of the angular correlation function or $w(\\theta)$. As discussed in Tegmark et al. (2002), Fourier or multipole-space measurements should be preferred over the real space due to the simple behavior of the covariance matrix and the fact that multipole-space measurements can be easily related to the underlying three-dimensional power spectrum without a complex window function that maps fluctuations in Fourier space to that of the angular space. In the case of the real space correlation function,  clustering at different angular scales are also correlated (Eisenstein \\& Zaldarriaga 2001) and this also lead to a complex covariance matrix that must be accounted for when model fitting the measurements.  Despite these considerations, real-space clustering measurements are easy to make and have been studied in the literature at 3.6 $\\mu$m and other {\\it Spitzer} IRAC bands (Oliver et al. 2004; Fang et al. 2004).  While model fitting was considered in terms of a power-law, a clear departure from a power-law model fit is present in angular clustering of 3.6 $\\mu$m galaxies in the   {\\it Spitzer} Wide-Area Infrared Extragalactic (SWIRE) as described by Oliver et al. (2004; for example, see their Figure~1 for clustering in Lockman and EN1 fields that indicates a rise above a power-law at small angular scales). On the other hand, clustering measurements in the First Look  Survey (FLS) at 3.6 $\\mu$m by Fang et al. (2004) do not show clear evidences for a departure from the power-law (see, Figure~1 in Fang et al. 2004). This could be due to the smaller number of galaxies used for clustering measurements, when compared to statistics in Oliver et al. (2004). Here, using the Bo\\\"otes-field IRAC shallow survey over 6 deg.$^2$, we find a departure from the power-law description and model that departure in terms of the halo model for galaxy clustering (Cooray \\& Sheth 2002). A departure from a power-law description, however, is not surprising for galaxy clustering at high redshifts due to non-linear clustering at angular scales less than a few arcminutes (Ouchi et al. 2005; Lee et al. 2005). Such departures have been detected for Lyman-break galaxies at $z \\sim 3$ to 5 (Ouchi et al. 2005; Lee et al. 2005; Cooray \\& Ouchi 2006), as well as at low redshifts    with surveys such as SDSS (Zehavi et al. 2005), DEEP2 (Coil et al. 2004), and COMBO-17 (Phleps et al. 2004). These departures were expected, and are described, in terms of analytical approaches such as the halo model (Cooray \\& Sheth 2002) and conditional luminosity functions (Cooray  2006a; van den Bosch et al. 2006).  The {\\it Paper} is organized as following: in the next section, we briefly summarize our   data and in \\S~3 we describe statistics in our catalogs. In \\S~4, we describe the procedure we used to measure clustering of IR sources and in \\S~5 we present  a summary of the analytical model used to describe clustering. Based on model fits, we extract quantities related to how IR galaxies occupy dark matter halos. In \\S~6, we discuss our results in the context of recent IRB anisotropy measurements with {\\it Spitzer} and future wide-field surveys with {\\it Akari} (Matsuhara et al. 2006), Cosmic Infrared Background Explorer (CIBER; Bock et al. 2006), and {\\it Spitzer} to resolve the ``missing source'' problem in the near-IR background intensity. We conclude with a summary in \\S~7.   Throughout this paper, we take a reference cosmology model with parameters $\\Omega_m=0.3$, $\\Omega_\\Lambda =0.7$, $h=0.7$ (unless left as a free parameter, such as in the luminosity function), and a normalization to the matter power spectrum at 8 h$^{-1}$ Mpc scales $\\sigma_8=0.84$. ",
        "conclusions": "\\subsection{Clustering}  In Figure~4 (left panel), we show the clustering measurements and a model description of the data. GOODS data, due to low sky area of 0.092 deg.$^2$ between the two HDF-N and CDF-S fields, measure only non-linear clustering, but between magnitudes 17 and 19 (45 and 7 $\\mu$Jy), Bo\\\"otes data allow the angular power spectrum to be extended to large angular scales related to the linear part of the angular power spectrum. The combination, Bo\\\"otes and GOODS, shows evidence for the transition between linear to  non-linear clustering with a departure from a single power-law clustering spectrum. With counts corrected for completeness, over the angular scales of overlap, we find a good agreement between the clustering spectrum determined in the large-area Bo\\\"otes field and the ones determined with GOODS images. If not completeness corrected, we find a $\\sim 10$\\% offset between Bo\\\"otes clustering spectrum and the GOODS clustering spectrum at multipoles of $10^4$, with Bo\\\"otes field underestimating power relative to the GOODS measurements.  In addition to the combination of GOODS and Bo\\\"otes in the $L$-band, we also detect a departure from a power-law in the ground-based FLAMEX Bo\\\"otes $J$-band data though $K_S$-band clustering measurements could be fitted with a power-law (Figure~4 right panel). In Figure~4, the lines show model descriptions of the angular power spectrum of source clustering using the CLF model for 3.6 $\\mu$m galaxies in the bin of 17 to 19 magnitudes and 21 to 21.5 magnitudes separately. These predictions make use of the statistical redshift distribution for $L$-band galaxies estimated based on photometric redshift estimates shown in Figure~3 (left panel). Our model fits are not strongly sensitive to assumptions about the redshift distribution as long as we do not take all sources to be either at very low or very high redshift. Note the factor of 100 difference in the y-axis scale between the left and the right panels of Figure~4; for sources in the same magnitude bin of 17 to 19 (255 to 40 $\\mu$Jy $J$-band, and 107 to 17 $\\mu$Jy $K_S$-band), $J$-band sources produce a factor of $\\sim$ 7.5 larger rms fluctuations, $\\sqrt{l^2 C_l/2\\pi}$, in the resolved IRB relative to sources in the same magnitude ranges in the $L$-band. This large increase is consistent with the difference in the mean intensity of the     IR background between these two different wavelengths and differences in the average clustering bias of $J$-band and $L$-band galaxies.  The CLF model-based descriptions of IR galaxy clustering spectra, as shown in Figure~4, require knowledge on the relation between both central galaxy luminosity and halo mass and the total luminosity and halo mass. As outlined in Section~5.2, we obtain these relations based on model descriptions to the measured LFs at 3.6 $\\mu$m in the {\\it Spitzer} SWIRE survey (Babbedge et al. 2006). This procedure is similar to the ones used to described galaxy clustering in SDSS at low redshifts and DEEP-2 and Lyman-break galaxies at high redshifts (Cooray 2006a; Cooray \\& Ouchi 2006). A different modeling technique, but with the same underlying approach, is available in van den Bosch et al (2006) where CLFs are  a priori assumed to be Schechter-function shapes. Similarly, we also use the simple halo model (in Section~5.1) to model fit the data as well. Here, the models require three parameters, $M_{\\rm min}$,  the minimum  dark matter halo mass at which IR galaxies begin to appear, $\\beta$, the power-law slope of the satellite occupation with halo mass, and $A$, the normalization of the satellite occupation number relative to the central galaxy occupation of unity in halos with mass above $M_{\\rm min}$. Instead of model fits to the CLFs, which involve a large number of parameters, here we will directly constrain parameters such as  $A$ and $\\beta$. The right panel of Figure~4 shows clustering measurements at $J$- and $K_S$-bands using the source catalogs from the ground-based imaging data of the Bo\\\"otes field. Again, in addition to clustering measurements, we also show model descriptions based on the halo model for galaxies with magnitude brighter than 19 in both $J$- and $K_S$-bands.     We have not attempted to model fit fainter galaxy samples due to issues related to the completeness of the catalogs.   \\begin{figure*}[!t] \\centerline{\\psfig{file=Fig7.eps,width=7in,angle=0}} \\caption{(a) Model fits to the $L$-band Bo\\\"otes and GOODS HDF-N clustering measurements in the magnitude bin of 17 to 19 (45 to 7.0 $\\mu$Jy), and (b) $J$-band Bo\\\"otes field clustering in the magnitude bin of 17 to 19 (255 to 40 $\\mu$Jy). We show the constraints on parameters that describe the satellite occupation number, following Section~5.1 (see text for details). In (a), for the $L$-band galaxies with magnitudes between 17 and 19, model fits suggest $\\beta < 0.98$  at the 1$\\sigma$ level, while $A$ is between 0.4 and 1.5.  In (b), for $J$-band galaxies with magnitudes between 17 and 19, $0.27 <  \\beta < 0.89$ at the 1$\\sigma$ level, while $A$ is between 0.4 and 1.1. In (b), the dashed lines show constant $\\bar{n}_g = \\int dz w(z) \\int dN/dM \\langle N_{\\rm gal} \\rangle/\\int dz w(z)$; the degeneracy direction in $\\beta$-$A$ plane traces the number density of IR sources, normalized to the redshift distribution of the sources, $w(z)$.} \\label{fig:contour} \\end{figure*}  In Figure~7, we consider a likelihood model fit to the angular clustering measurements only by taking two parameters, $\\beta$ involving the slope of the satellite counts as a function of halo mass and $A$ the overall normalization of the central-to-satellite galaxy occupation number to be free  while $M_{\\rm min}$ is taken to be fixed at  $5 \\times 10^{11}$ M$_{\\sun}$ for galaxies with magnitudes in the range 17 to 19. This $M_{\\rm min}$ value is not the minimum mass for galaxies with any luminosity to appear in halos but rather those in the appropriate magnitude ranges.  Since we are considering IR bright galaxies in the magnitude range between 17 and 19, this minimum mass is consistent with the expected value  with the CLF model, though the cut off is not sharp at a fixed mass scale given the dispersion we allow between luminosity and halo mass . With large catalogs, the accuracy of clustering measurements can be improved and further parameters determined, but at this stage, it is unlikely that more than 2 parameters can be extracted from the data. We only model fit the range of 17 to 19 magnitude in both the $L$-band and the $J$-band, since at deeper bins with either the GOODS data or ground-based Bo\\\"otes data alone, the clustering measurements are mostly limited to non-linear scales only leading to large degeneracies between parameters in the occupation number, and the photometric estimates for the source distribution of these galaxies are less certain than bright galaxies in the magnitude bin between 17 and 19.  Based on the constraints on $A$ and $\\beta$ shown in Figure~7, we estimate the halo mass scale at which two IR sources with a magnitude between 17 and 19 in the $L$-band appear in the same halo, or when $\\langle N_{\\rm gal}(m) \\rangle =2$, is $(0.9-7)\\times 10^{12}$ M$_{\\sun}$. This is consistent with the number in the CLF-based model with a value of $\\sim 2\\times 10^{12}$ M$_{\\sun}$. Note that in CLF-based models, due to the large number of parameters involved, we do not vary them to fit them individually, but obtain an overall description consistent with 3.6 $\\mu$m LFs between redshifts of 0 and 1 and also use a single set of values for the input parameters when comparing to measurements. In terms of CLF models, this is the best we can do at this stage. The mass scale we estimate for galaxies at $z$ of 1 can be compared to previous estimates in the literature. For example, using conditional luminosity function based  models of galaxy clustering (Cooray 2006a) compared to DEEP2 (Coil et al. 2004) and COMBO-17 (Phleps et al. 2003)     measurements of galaxy clustering at redshifts between 0.7 and 1.3 suggests that the halo mass scale at which bright galaxies with luminosities above $L_\\star$ at $z \\sim 1$ start to appear as satellites is  $\\sim (1-8)\\times 10^{12} M_{\\sun}$ at the 1$\\sigma$ confidence range, in good agreement with the same scale we have established for $z \\sim 1$ IR galaxies.  The satellite halo occupation number for IR sources in the $L$-band is such that the power-law slope with mass is below 0.98 at the 1$\\sigma$ confidence level.  For $J$-band galaxies with magnitudes between 17 and 19, the power-law slope with mass of the satellite occupation number is $ 0.27 < \\beta < 0.89$ at the 1$\\sigma$ confidence level, which is consistent with the expected value for the slope 0.72 with halo mass from out CLFs. As shown in Figure~7 right panel, we illustrate the degeneracy direction in the $\\beta$-$A$ parameter plane by plotting lines of constant average density of galaxies calculated as $\\bar{n}_g = \\int dz w(z) \\int dM dN/dM \\langle N_{\\rm gal}\\rangle$, where $w(z)$ is the normalized redshift distribution of sources used for clustering measurements. The degeneracy in $\\beta$-$A$ is such that one traces constant $\\bar{n}_g$ values.     This number density,  as a function of redshift, is directly available with LFs and was the basis to normalize our CLF description. In future, if LF measurements can be improved (the measurements have low signal-to-noise at $z > 0.5$ as can be seen in Fig.~5) and also be extended to higher redshifts, then we can hope to improve constraints on the occupation number, or more importantly on the conditional luminosity function. An analysis based on such an approach for $z \\sim 3$ to 4 Lyman-break galaxies are available in Cooray \\& Ouchi (2006) and we hope to return to this topic again in a later paper using additional measurements we have now begun to make from the same set of data as used for the clustering measurements here (Cooray et al., in preparation).   \\subsection{Anisotropies in the IRB}  To understand how these fluctuations of resolved sources compare with that of the diffuse IRB, in Figure~8 left panel we compare our clustering measurements with the anisotropy power spectrum presented in Kashlinsky et al. (2005) for the unresolved component with all sources brighter than 22.5 magnitude (0.28 $\\mu$Jy) removed.  In the case of these unresolved anisotropies, the small scale structure is that of a shot-noise with the power spectrum scaling as $\\ell^2$; In fact, it is this shot-noise that Kashlinsky et al. (2005) used to determine that they had removed resolved sources down to a magnitude limit of 22.5. The flattening of the fluctuations at the two lowest $\\ell$ bins is taken to be indicative of an excess component in the IRB,      with that component ascribed to first galaxies.     The predictions shown in lines come from the CLF model with parameters that reproduce the LF.  The right panel of Figure 8 shows the completeness corrected number counts.  The lines here are expectation for number counts using the CLF by integrating over LFs as a function of volume element. Note that at magnitudes below 18, the models slightly under predict the counts despite a reasonable agreement with LFs. This difference, however, comes from stellar contamination to source counts as our counts are not corrected for stars, though based on the stellar model of Fazio et al. (2004), we do correct for stars in our clustering measurements.  We show three specific CLF models for counts which are all normalized to LFs at $z < 1$, but with differences at the faint end not probed by the LFs. The corresponding clustering spectra associated with those counts are shown in the left panel of Fig.~8. In Fig.~8 right panel, the black lines involve a model where we set the minimum mass for appearance of at least one galaxy regardless of the luminosity in a dark matter halo to be  $3 \\times 10^{10}$ M$_{\\sun}$ (or effectively a luminosity with a corresponding magnitude below 30 in $L$-band since our discussion below is for galaxies brighter than 30th magnitude). That is, at halo masses below this value galaxies are not expected to form due to a variety of reasons from photoionization during reionization, preheating of gas from AGNs, and/or expulsion of gas from first generation of supernovae. This number is consistent with the value generally quoted in the literature based on semi-analytical models of galaxy formation (e.g., Kauffmann et al. 1993; Benson et al. 2001). Furthermore, this value is constrained to be below $3 \\times 10^{10}$ M$_{\\sun}$ using direct likelihood fits to clustering measurements in SDSS (Cooray 2006a), though this is for galaxies in the $r'$-band. With this cut-off in halo mass, number counts peak at $L$-band magnitudes around 24.  To obtain an excess of galaxies at the faint-end, we arbitrarily lower this low mass cut off for galaxy formation and set it at a mass scale of  $10^{8}$ M$_{\\sun}$ and take the central galaxy luminosity-halo mass relation to continue down to such low masses. This introduces an excess of faint galaxies with an increase in the number counts at magnitudes starting from a magnitude of 22. If we use these counts down to a magnitude limit of 26, the sources between the magnitude range of 22.5 and 26 produce a shot-noise of $\\sim 5 \\times 10^{-11}$ nW$^2$ m$^{-4}$ sr$^{-1}$, which is roughly a factor of 8 higher than the shot-noise seen in the fluctuation measurements ($\\sim 7.2 \\times 10^{-12}$ nW$^2$ m$^{-4}$ sr$^{-1}$). Note that such a large slope at the faint-end has been suggested based on a study of the one-point probability distribution function of the unresolved IR intensity (Savage \\& Oliver 2005), the large difference between the expected and measured shot-noise level suggests that $L$-band source counts at the faint-end do not continue to increase down to magnitude limits of 26 with the same slope as the one suggested with completeness-corrected counts between 19 and 22. The order of magnitude lower shot-noise measured by Kashlinsky et al. (2005) suggests that the counts either flatten or turn over  consistent with our basic CLF model with a minimum mass cut off around  $3 \\times 10^{10}$ M$_{\\sun}$ for appearance of IR galaxies in dark matter halos.  The integrated background associated with the default description for number counts (in black lines) is $\\sim$ 0.08 nW  m$^{-2}$ sr$^{-1}$ between the magnitude range of 22.5 and 30, suggesting that if the excess clustering suggested by Kashlinsky et al. (2005) is correct, then these sources are unlikely to be a significant fraction of the total IRB intensity at 3.6 microns. If sources were to continue with the large positive slope to the faint-end, then the background between magnitude limits of 22.5 and 30 would have been close to 1.2 nW  m$^{-2}$ sr$^{-1}$, which is an appreciable fraction given that the total background is about 10 nW  m$^{-2}$ sr$^{-1}$ at 3.6 $\\mu$m. Note that the background intensity of $\\sim$ 0.08 nW m$^{-2}$ sr$^{-1}$ we suggest for faint sources to describe unresolved source clustering shot-noise is also below the level of 1 to 2 nW m$^{-2}$ sr$^{-1}$ background intensity from a population of $z > 5$ galaxies that was considered to reproduce clustering measurements of Kashlinsky et al. (2005) by Salvaterra et al. (2006). Salvaterra et al. (2006) models assume that shot-noise is dominating from $\\ell$ of 10$^5$ and only consider the contribution coming from galaxies at high redshifts.  To emphasize the redshift ranges where the contributions are coming from, in our CLF, we put a low redshift cut at $z \\sim 5$. The expected counts are show in green in the right panel with corresponding clustering spectrum for those galaxies in the left panel of Figure~8. The clustering level is well below the measured level at two  orders of magnitude below the amplitude of Kashlinsky et al. (2005). Since the clustering spectrum is proportional to square of the background, see Section ~5, one can  arbitrarily renormalize the counts, as in the case of models in Salvaterra et al. (2006), to fit Kashlinsky et al. (2005) data, but then the observed excess lead to a total intensity of IRB an order of magnitude higher. Similarly, counts are also increased by an order of magnitude. This renormaization, however, is not correct since it leads to a large surface density of galaxies at redshifts more than 5 at a level more than allowed by the expected density of dark matter halos at these redshifts. This expected high density of galaxies has been already realized in Salvaterra \\& Ferrara (2006) who have argued against a large density of first galaxies that dominate the IR background given constraints on the number counts. This arbitrary normalization is not an issues for our CLF models  since the galaxy density is constrained given the connection between dark matter halos and galaxies  they host. Unlike the simple model in Salvaterra et al. (2006) the amplitude is not a free parameter.  More importantly, Figure~8 left panel also shows that the clustering predicted for galaxies in the magnitude range between 22.5 and 26 is consistent with measurements of Kashlinsky et al. (2005) for unresolved background and for clustering of resolved galaxies. Here, non-linear clustering of galaxies within dark matter halos continue to be important out to multipoles of 10$^6$. For comparison, in Figure~8 left panel, we also show the clustering measurements in  our faintest bin between 21 and 22.5 in GOODS. The excess seen in Kashlinsky et al. (2005) can be described as simply an extension of resolved component that we have measured and extending to a larger angular scale than for measurements in Kashlinsky et al. (2005). In combination, there is a consistent description to the data from our CLFs suggesting that the unresolved component seen by Kashlinsky et al. (2005)  can be described by galaxy counts with magnitudes between 22.5 and 26 at redshifts below 5 with contributions mostly from redshifts 1 to 3, and that there is no need to invoke a large surface density of  galaxies at redshifts above 5.  The limited range of multipoles probed by Kashlinsky et al. (2005) limits us from making an exact statement given that the data can also be modeled from a clustering spectrum that trace the linear density field at redshifts above 8 (dot-dashed lines following Cooray et al. 2004). This diffuse background is taken to be produced by Pop-III stars and their spectra are completely different from known galaxies given the dominant energy in the Lyman-$\\alpha$ line, which is redshifted to near-IR wavelengths today (see, Santos et al. 2002). Thus, our modeling based on CLFs described above will not contain such a separate component especially if the Pop-III formation happens at $z > 8$. The 3.5 $\\mu$m IRB intensity associated with  models in the top and the bottom curve is $\\sim$ 2.5 and 0.3 nW m$^{-2}$ sr$^{-1}$, respectively. The angular power spectrum of clustering  peaks at a multipole of $\\sim 10^3$, corresponding to the peak of the linear clustering when projected at these     high redshifts. If the IRB excess is due to this component,  the the clustering excess should peak at multipoles around $\\ell^2$ $\\sim 10^3$; this is simply a reflection of the projection associated with the peak of the linear mass power spectrum. As discussed in Cooray et al. (2004; Bock et al. 2006), in the H-band, the IRB intensity is roughly a factor of 10 higher given that the Pop-III spectrum is sharply rising towards lower wavelengths as the  intensity there is dominated by the Lyman-$\\alpha$ emission associated with recombinations during the reionization of the Universe. The two models shown in Figure~8 left panel bracket two extreme regions in terms of the reionization history, the density of first galaxies, their clustering     strength or bias, and the Lyman-$\\alpha$ photon production. It is clear that one should focus on degree to arcminute angular scales for anisotropy studies in the IRB instead of sub-arcminute angular scales considered by Kashlinsky et al. (2005) at multipoles above 10$^4$. As is clear from Figure~8 left panel, such imaging need not be very deep since degree-scale shallow images that allow removal of sources down to a magnitude of 21 in the $L$-band can be easily used to probe the range suggested by the two extreme models.  To briefly summarize our discussion so far, the excess clustering power spectrum measured in Kashlinsky et al. (2005) is consistent with faint galaxies with magnitudes between 22.5 and 26 and predominantly at redshifts between 1 and 3 as extrapolated from measured LFs out to $z \\sim 1$. These faint galaxies are unresolved in the deep {\\it Spitzer} images, but not a significant contributor to the total intensity of the IRB. If we were to explain the missing IRB with a new population of sources, such as Pop-III stars, then the clustering spectrum can be explained, but there is an overall increase in fluctuations at multipoles of $\\sim$ 10$^3$ which is not currently probed by unresolved clustering.  \\begin{figure*}[!t] \\centerline{\\psfig{file=Fig8a.eps,width=2.8in,angle=-90} \\psfig{file=Fig8b.eps,width=3.0in,angle=-90} } \\caption{ IRB fluctuations from a $z> 8$ diffuse component that traces the linear density field (dot-dashed lines), galaxy clustering (solid lines), and shot-noise from galaxies (dashed lines that scale as $l^2$ with increasing multipole). The measurements are from Kashlinsky et al. (2005) in the $L$-band with fainter symbols showing the background fluctuations after rebinning the original measurements. We consider three model descriptions for clustering using the CLF models (see text for details). The corresponding counts are shown and compared to completeness corrected measured counts in the right panel. In the right panel, in addition to total counts (solid lines), we also show counts of central (dashed lines) and satellite (dot-dashed lines) galaxies. These predictions for number counts are obtained by simply integrating the model LFs over the volume with  luminosity converted to an observed flux as a function of redshift. The black lines show our default model and the shot-noise component can be roughly described as due to point sources with $L$ band magnitudes between 22.5 and 26. This model also reproduces clustering suggesting that unresolved clustering spectrum measured by Kashlinsky et al. (2005) can be described by galaxies in the magnitude range between 22.5 and 26. If we populate IR galaxies down to a halo mass of 10$^8$ M$_{\\rm sun}$, we find counts shown by the blue line, with a slope consistent with the faint-end slope suggested in Savage \\& Oliver (2005). These counts lead to an order of magnitude higher shot-noise than measured with the clustering spectrum of IR anisotropies and a larger clustering amplitude than measured. Finally, with green lines, we only consider counts at $z > 5$ for the same model, but in this case, we fail to explain the measurements given that these faint galaxies are not strongly clustered even at high redshifts. The two dot-dashed lines show a model expectation for clustering in the diffuse IRB from unresolved sources at $z > 8$ prior to complete reionization (Cooray et al. 2004). The 3.5 $\\mu$m IRB intensity associated with models in the top and the bottom curve is $\\sim$ 2.5 and 0.3 nW m$^{-2}$ sr$^{-1}$, respectively. The angular power spectrum of clustering  peaks at a multipole of $\\sim 10^3$, corresponding to the peak of the linear clustering when projected at redshifts greater than 8. } \\label{fig:irb} \\end{figure*}      \\subsection{Wide-Field Surveys for IRB Anisotropy Measurements}  While degree to arcminute angular scale fluctuation measurements allow first-stars at $z>8$  to be more easily identified, extending fluctuation measurements to low multipoles is challenging with {\\it Spitzer} IRAC images alone due to both the small field-of-view and to problems such as uncertainties in relative calibration associated with stitching independent fields together to make a larger map. The Bo\\\"otes field considered here comes from a large number of independent images, but we are able to make clustering measurements of resolved sources since the catalogs are less affected by issues related to sensitivity variations from field to field. Nevertheless, we are also studying the possibility to use the 6 deg.$^2$ central region of the Bo\\\"otes mosaic to at least place limits on the amplitude of unresolved source clustering at 10 arcminute angular scales to a degree angular scales. These results will be presented in an upcoming paper (Cooray et al. 2006).  As mentioned in above, to separate various possibilities, the need for large-area clustering measurements is clear from the left panel of Figure~8 (see, also Figure~2 of Bock et al. 2006). While at multipoles of $10^4$ to 10$^5$ the model clustering of a $z> 8$ diffuse component overlaps with faint galaxy clustering, at multipoles of $10^3$ the clustering is well separated. Thus, at degree angular scales, instead of deep images one can use shallow images down to a brighter magnitude limit to search for an excess component than the magnitude limit one has to image when probing the same diffuse component with images that allow clustering studies only at arcminute angular scales and below. In consideration of the two issues discussed so far, large-angular scale excess and the large difference between measured IRB intensity and the predicted intensity based on counts at lower wavelengths than studied by {\\it Spitzer}, we suggest observations that attempt to resolve the nature of the IRB excess should concentrate on shallow, wide-field imaging at wavelengths below 3.6 $\\mu$m.  The Cosmic Infrared Background Explorer (CIBER; Bock et al. 2006) is designed directly to address the ``missing source'' problem in IRB that is significant between 1 $\\mu$m and 2 $\\mu$m through shallow, but wide-field images. CIBER uses wide-field imagers with a field-of-view of 4 deg.$^2$ in both I- and H-bands, supported by spectrometers to determine the contribution from zodiacal light and establish the spectrum of IRB intensity from 0.7 to 1.8 $\\mu$m. With the wide-field coverage, CIBER can measure clustering of IRB light from multipoles below $10^2$ to $10^4$ covering the linear to non-linear regime from a single image. Furthermore, CIBER will image the Bo\\\"otes field so that existing source counts from {\\it Spitzer} and other ground-based imaging data can be used to remove point sources. The combination of resolved and unresolved source clustering by extending studies such as the one presented here should allow us to pin down the sources responsible for the missing IRB and directly address the ``missing source'' problem of the IRB.  In addition to CIBER, at wavelengths above 2.2 $\\mu$m, the North Ecliptic Pole wide-field survey with the IR Camera on {\\it Akari} (Matsuhara et al. 2006) will provide an additional dataset to study large-scale fluctuations in the near-IR background. The shallow survey useful for clustering measurements will span over 6.2 deg.$^2$ and will allow  clustering measurements from the linear to non-linear scales similar to the angular power spectrum established here for 3.6 $\\mu$m resolved sources in the Bo\\\"otes field.  Finally, it could be that after the depletion of cryogen necessary for longer wavelength instruments, {\\it Spitzer} will continue to operate only the first two channels of IRAC. Such a scenario could be exploited for a very wide-field survey which in return could be used to further improve the clustering measurements of both resolved and unresolved components. While the deep imaging data with {\\it Spitzer} IRAC of the GOODS HDF-N and CDF-S fields have now allowed clustering measurements down to very faint magnitudes, these measurements are restricted to the non-linear regime. It will be very useful to measure clustering for sources at magnitude levels of 21 and fainter so that background anisotropy measurements at degree angular scales can be properly combined with resolved source clustering at these faint magnitude levels to obtain a complete picture of the IRB anisotropies."
    },
    "0609/astro-ph0609667_arXiv.txt": {
        "abstract": "We have modelled the rotational modulation of X-ray emission from T Tauri stars assuming that they have isothermal, magnetically confined coronae.  By extrapolating surface magnetograms we find that T Tauri coronae are compact and clumpy, such that rotational modulation arises from X-ray emitting regions being eclipsed as the star rotates.  Emitting regions are close to the stellar surface and inhomogeneously distributed about the star.  However some regions of the stellar surface, which contain wind bearing open field lines, are dark in X-rays. From simulated X-ray light curves, obtained using stellar parameters from the Chandra Orion Ultradeep Project, we calculate X-ray periods and make comparisons with optically determined rotation periods.  We find that X-ray periods are typically equal to, or are half of, the optical periods.  Further, we find that X-ray periods are dependent upon the stellar inclination, but that the ratio of X-ray to optical period is independent of stellar mass and radius. ",
        "introduction": "One of the recent results from the Chandra Orion Ultradeep Project (COUP) was the detection of significant rotational modulation of X-ray emission from low mass pre-main sequence stars.  The detection of such modulation suggests that the coronae of T Tauri stars are compact and clumpy, with emitting regions that are inhomogeneously distributed across the stellar surface, and confined within magnetic structures that do not extend out to much beyond a stellar radius \\citep{fla05}. There is also evidence for much larger magnetic loops, possibly due to the interaction with a surrounding circumstellar disc \\citep{fav05,gia06}.  A model already exists for T Tauri coronae, where complex magnetic field structures contain X-ray emitting plasma close to the stellar surface, whilst larger magnetic loops and open field lines are able to carry accretion flows (\\citealt{jar06}; \\citealt{gre06}).  In this paper we use surface magnetograms derived from Zeeman-Doppler imaging to extrapolate the coronae of T Tauri stars using stellar parameters taken from the COUP dataset \\citep{get05}.  By considering isothermal coronae in hydrostatic equilibrium we obtain simulated X-ray light curves, for a range of stellar inclinations, and then X-ray periods using the Lomb Normalised Periodogram (LNP) method. We compare our results with those of \\citet{fla05}, who demonstrate that those COUP stars which show clear evidence for rotationally modulated X-ray emission appear to have X-ray periods which are either equal to the optically determined rotation period ($P_X=P_{opt}$) or are half of it ($P_X=0.5P_{opt}$).   \\begin{figure*} \\def\\subfigtopskip{4pt} \\def\\subfigbottomskip{4pt} \\def\\subfigcapskip{2pt} \\centering \\begin{tabular}{cc} \\subfigure[]{ \\label{lqhya_extrap} \\psfig{figure=gregory_figure1a.eps,width=80mm} } & \\subfigure[]{ \\label{abdor_extrap} \\psfig{figure=gregory_figure1b.eps,width=80mm} }  \\\\ \\end{tabular} \\caption[]{Closed coronal structures showing field lines which contain X-ray emitting plasma, extrapolated from surface magnetograms of (a) LQ Hya and (b) AB Dor, for COUP source number 6 ($M_{\\ast}=0.23M_{\\odot}, R_{\\ast}=1.6R_{\\odot}, P_{rot}=9.81d$). The magnetic structures are compact and inhomogeneously distributed about the stellar surface.} \\label{extrap} \\end{figure*} ",
        "conclusions": "By extrapolating surface magnetograms of young main sequence stars we have shown that magnetic fields with a realistic degree of complexity reproduce the observed rotational modulation of X-ray emission from T Tauri stars \\citep{fla05}.  This model has already been used to correctly predict X-ray EMs, mean coronal densities and the observed increase in X-ray EM with stellar mass \\citep{jar06}.  We find that X-ray emitting regions are inhomogeneously distributed around the stellar surface and are typically compact ($\\la R_{\\ast}$).  This agrees with the findings of \\citet{fla05} who point out that in order to explain the amplitude of the modulation of X-ray emission which they detect, emitting regions must be compact in order to undergo eclipse.  We found that the ratio $P_X/P_{opt}$ is independent of stellar mass and radius, but that X-ray periods are dependent on stellar inclination.  For example, in the case where there are two dominant X-ray emitting regions in opposite hemispheres, we found that the X-ray period could be either $0.5P_{opt}$ or $P_{opt}$ depending on how the star is inclined to the observer.  In order words, the presence of two bright emitting regions in opposite hemispheres does not immediately imply that $P_X=0.5P_{opt}$.  Further, some coronal fields with compact emitting regions produce only small amplitude modulation of X-ray emission when viewed at unfavourable inclinations.  We do not find X-ray periods of $P_X>2P_{opt}$, unlike \\citet{fla05}.  However, they argue that those period detections are likely to spurious, given the arduous task of removing flares from observed COUP X-ray light curves. Also the apparent lower limit of $P_X=0.5P_{opt}$ may not be of true physical origin.  We have demonstrated this by constructing a corona which resulted in $P_X=0.333P_{opt}$; however this sort of field structure is perhaps uncommon. Instead the majority of X-ray periods cluster around $P_X=[0.5,1]P_{opt}$.  We found that $P_X$ (for a fixed inclination) was the same for both the LQ Hya and AB Dor-like coronal fields, despite the distribution of emitting regions being different for the two cases.  This suggests that for complex magnetic fields, where there are many emitting regions distributed across the stellar surface, it becomes difficult to disentangle the contribution to the X-ray emission from any particular emitting region. Consequently, even though different stars are likely to have different coronal structures, the X-ray periods of T Tauri stars appear to cluster around $P_{opt}$ and $0.5P_{opt}$.  In other words, although different coronal structures give rise to different amounts of modulation at a fixed inclination, the X-ray period is likely to be the same.  Throughout we have only considered the lower coronal temperatures derived from the COUP dataset. If instead we had used the higher coronal temperatures our results and conclusions remain the same.  For a higher coronal temperature more field lines are unable to contain the coronal gas and are blown open to form a stellar wind, meaning that coronae are (slightly) more compact.  This leads to small, but not important, differences in the resulting X-ray periods, however the amplitude of modulation is larger (typically 60\\% compared with 50\\% for the lower coronal temperatures).  The amplitudes of modulation which we determine compare well with the observed values of \\citet{fla05} of 20-70\\%.  Although we have not considered it in this paper, the coronal structure of T Tauri stars may evolve on timescales that are comparable with the COUP observing time.  For example a dominant X-ray emitting region may be lost due to a large coronal mass ejection which would open up a previously-closed region of field.  This may then change the resulting X-ray period.  However, it is still interesting that even with two different field structures such as those in Fig. \\ref{extrap}, the X-ray periods are still typically equal to the optical period or are half of the optical period.  We have not considered how the presence of a circumstellar disc and active accretion influences coronal structures, and the possible resulting effects on the X-ray period/amplitude of modulation. A comprehensive and detailed discussion of how accretion processes may influence X-ray emission in T Tauri stars is provided by \\citet{pre05}.  The growing consensus, as was confirmed by COUP, is that X-ray emission from T Tauri stars does not depend upon the presence of a disc, but is influenced by active accretion. For example, \\citet{fei02} found X-ray activity levels were the same for stars with and without discs\\footnote{It is important to remember that the presence of a disc does not immediately imply that accretion is taking place.  Indeed many T Tauri stars show excess infrared emission, indicating the presence of a disc, but lack strong H$\\alpha$ or CaII emission, characteristic of active accretion.}.  However, \\citet{fla03} found a difference in the X-ray activity levels of accreting and non-accreting stars.  The process of accretion therefore influences the amount of X-ray emission detected, with stars that show evidence for accretion having X-ray luminosities which are a factor of 2-3 smaller than those stars without discs and those stars which show evidence for discs, but not for accretion \\citep{pre05}.  Accretion may therefore affect the coronal structure of T Tauri stars, which in turn may have implications for rotational modulation of X-ray emission, and X-ray periods.  It is important to note however that \\citet{fla05} have searched for correlations between observed modulation amplitudes and the $\\Delta$(I$-$K) excess, a disc indicator, and EW(CaII), an accretion indicator, but do not find any.  Furthermore, \\citet{fla05} also search for dependencies on the X-ray period detection fraction with the indicators of discs and accretion, but again find none.  This suggests that there is no difference in the spatial distribution of X-ray emitting regions in both classical and weak line T Tauri stars (with the former having both a disc and active accretion, and the latter neither). However, \\citet{fla05} did find a suggestion (but with a low significance of only about $1\\%$) that stars with $P_X=0.5P_{opt}$ are preferentially active accretors.  \\citet{jar06} have shown that stars (typically of lower mass) which have coronae that would naturally extend to beyond the corotation radius would have their outer corona stripped by the presence of a disc.  In order to check if this would affect the resulting X-ray period we determined the coronal structure of all the stars considered in \\S3 assuming that they were surrounded by a disc.  Any field line which passed through the disc at, or within the corotation radius was assumed to have the ability to carry an accretion flow and was therefore considered to be ``mass-loaded'' and set to be dark in X-rays - a process which is discussed in detail by \\citet{jar06} and \\citet{gre06}.  We then calculated X-ray emission measures with rotational phase and determined the amplitude of modulation of X-ray emission and X-ray periods, using the same methods as discussed in \\S3.  We found no difference in the values of $P_X$ and very little difference in the amplitudes of modulation.  We therefore conclude, exactly as the observations suggest, that the presence of discs does not influence rotational modulation of X-ray emission. However, active accretion might.  We have yet to consider how accretion could influence rotational modulation of X-ray emission.  As already discussed above stars which are actively accreting show lower levels of X-ray activity than those which are non-accretors.  Again, the possible explanations for this are discussed by \\citet{pre05}, with the most likely being that magnetic reconnection events in the magnetosphere cannot heat the dense material within accretion columns to a high enough temperature to emit in X-rays.  Some models suggest that accretion will distort magnetic field lines giving rise to instabilities and reconnection events (e.g. \\citealt{rom04}).  The energy liberated in such reconnection events may not be able to heat the high density material in accretion columns to a sufficiently high temperature in order to emit in X-rays.  Thus accretion columns rotating across the line-of-sight to the star may be responsible for the observed reduction in X-ray emission in accreting T Tauri stars compared with the non-accretors.  The effect of this on X-ray periods and modulation amplitudes will be considered in future work. A detailed accretion flow model will have to be constructed in order to account for the attenuation of coronal X-rays by accretion columns rotating across the line-of-sight.  Also, inner disc warps may have a similar effect at high stellar inclinations. This could change the shape of the X-ray variability curves shown in Figs. \\ref{lightcurve} and \\ref{lnp_curves}, with a reduction in X-ray emission at particular rotational phases, depending on the accretion flow geometry.  Our model may allow for the simultaneous confinement of X-ray emitting plasma in closed loops close to the stellar surface, but also for other field lines to stretch out into the inner disc and carry accretion flows.  \\citet{gre06} find that often open field lines can carry accretion flows, which may allow accretion to influence the magnetic field structure on the large scale, with field lines being ``mass-loaded'' with disc material,  but leaving compact coronae relatively unaffected by the process of accretion.  This model would allow for the existence of large extended magnetic structures, inferred to exist from the detection of large flaring loops by \\citet{fav05}, with the coexistence of compact coronae.  However, much quantitative work remains to be done here.  Another new accretion model may offer insights into the connection between X-ray emission and accretion in T Tauri stars. An MHD model by \\citet{von06} suggests that rather than material being loaded onto field lines at the inner edge of the disc, it flows onto the star only at times when, and into regions where, the accretion flow is able to overcome the stellar wind.  In such a way this model predicts that accretion only occurs into regions of low field strength.  It is also worth noting the observational results of \\citet{fla05} that indicate that rotational modulation of X-ray emission is not preferentially detected in either accreting or non-accreting T Tauri stars and that therefore the spatial structure of the X-ray emitting plasma is the same in both cases.  Hence, although accretion does affect X-ray emission in accreting stars, and may distort the structure of stellar coronae, it may not affect the existence of rotational modulation."
    },
    "0609/astro-ph0609384_arXiv.txt": {
        "abstract": "{We present NIR surface photometry of 11 edge-on galaxies obtained in the course of a long term project aimed at analysing the occurrence and type of the truncation of the outer disks. Observations were carried out at the 1.5 m CST (Carlos S\\'anchez Telescope) in Tenerife (Spain) using the CAIN infrared camera. 7 galaxies exhibit clear truncation on their disk profiles and 4 galaxies were observed to be clearly untruncated within observational limits. We describe the truncations as real, smooth and complete (as suggested by extrapolation and in the sense that the measured truncation curve goes into the noise at a truncation radius $R_{tr}$), following a decline proportional to $(R_{tr} -R)^{-n}$  (where $R$ is the radius). Despite its deep photometric reach, the data presented do not permit a detailed exploration of the region where optical data show a second slope. Special care  was taken concerning the surface brightness deprojection of edge-on galaxies, which was carried out by two methods, one comprising the inversion of Abel's integral equation and the other following a numerical method. These methods gave nearly identical results. NIR observations of truncations could differ from observations in the optical, since the two domains trace different stellar populations. ",
        "introduction": "It is widely accepted that the galactic disk brightness profiles, at least in the inner regions where the contribution of the bulge to the observed emission is still low, can be reproduced with an exponential function. At large radii, it is found that the stellar distribution often departs from this simple description and decreases much faster. One of the goals of this paper is to contribute to determining the functional form of such a rapid decline.  This decline of the stellar density was discovered by van der Kruit (1979), who denoted it as truncation. In this and subsequent papers (van der Kruit \\& Searle, 1981a, 1981b) the basic properties of truncations were described as well as the first theoretical ideas to explain them. Some recent studies have found that there is no complete truncation, but rather a sharp change in the exponential constant, called a ``break'' (Pohlen et al. 2002a, Pohlen et al. 2004 and references therein).  Let us define the truncation curve as \\begin{equation} \\tau(R) = \\mu(R) - \\mu_D(R) \\end{equation} as in Florido et al. (2001) (hereafter F01). Here, $R$ is the radius, $\\mu$ the actual surface brightness (in mag arcsec$^{-2}$), and $\\mu_D$ the same quantity obtained by the extrapolation of $\\mu(R)$ in the inner disk, where the exponential decline can be followed unambiguously.  What is $\\tau(R)$? Does it depend on the observed wavelength? These questions, of fundamental importance to understanding this phenomenon, motivated the present paper and are only partially answered. Nevertheless, the observations reported here will contribute to achieving a final description, which is difficult to assess because truncations (or ``breaks'') take place at a very low brightness in regions with a very low S/N ratio. The different theoretical models currently proposed are still far from obtain the functional form of $\\tau(R)$, but they should be able to do so eventually. Excellent reviews of the question have already been made, such as those by van der Kruit (2001) and Pohlen et al. (2004).  We opted to observe at NIR wavelengths, for the following reasons: firstly, the old stellar population, which makes up the great majority in the galactic disk of spiral galaxies, can be better traced, as it emits the bulk of its flux at NIR. There is also a contribution to the observed flux from the young population and from AGB stars, but this tends to be restricted  spatially to specific areas within the disk (e.g. spiral arms and bulges) and does not change the basic trend in the stellar distribution throughout the disk. Another important argument is that the extinction effects are minimised and so the observed brightness distribution more closely resembles the true stellar distribution, while the stellar flux still largely dominates the measured flux. This is not the case at longer wavelengths, at which the extinction is even lower. Also, the instrumental set-up allows us to reach outer regions on the galactic disks, where truncations are normally found, due to a combination of an excellent dry site, a telescope designed and maintained for NIR observations and cryogenic instrumentation with background contamination controlled via a cold stop.  Most previous studies (for example, Barteldrees \\& Dettmar 1994; de Grijs et al. 2001; Kregel et al. 2002; Pohlen et al. 2002a) have been carried out at optical wavelengths. To be complete, we mention the paper by Schwarzkopf \\& Dettmar (2000) as it contains a sample of 110 highly inclined galaxies, containing also K$_s$ data for 41 of the galaxies, but most of these are probably not deep enough.  As a complement to this large amount of statistical data, detailed studies of the truncation curve are required. The observations reported in the present paper and in F01 are considerably deeper than those of any previous study dealing with NIR photometry on the outskirts of spirals, here reaching values of 20-23 mag arcsec$^{-2}$ in J and 18.5-21 mag arcsec$^{-2}$ in K$_s$, which correspond to a face-on surface brightness in the range of 25-26 mag arcsec$^{-2}$ in J and 24-25 mag arcsec$^{-2}$ in K$_s$. The increase of limiting magnitudes from edge-on to face-on views depends on the galaxy, mainly due to its inclination and thickness. These sensitivity limits enable us to extend the structural analysis to the outer areas of the disks, at galactocentric radii of comparable size to those achieved in optical studies.  F01 presented NIR observation dealing with truncations, and the present paper can be considered as a continuation of the former. The telescope and the infrared camera proved to be very efficient for investigating the low surface brightness at the outskirts of disks.  More NIR observations are necessary, especially taking into account that optical and NIR observations may differ substantially and that this difference can be crucial. Recent star formation beyond the truncation (or break) radius would have a larger influence in the blue but very little influence in the NIR range.  Different studies of truncations have used different techniques and so the results presented are not easily comparable. Some groups observe edge-on galaxies, while others consider face-on galaxies. Comparison is especially difficult when the data used are obtained at different wavelengths. For correct identification of the differences between optical and NIR observations a parallel study  of the same galaxies would be preferable, using the same analysis. This would be the best way to determine, definitively, whether there is a clear difference between the results obtained using the different types of wavelengths. We present here the results of our NIR survey which are noticeably deep and interesting in themselves. Observations at these wavelengths can provide vital clues to our understanding of the truncation phenomenon. ",
        "conclusions": "In order to obtain more information about the truncation of the stellar disk, it is necessary to incorporate observations in the NIR, which is the best tracer of the old population and relatively free of extinction problems, and to use larger telescopes. On the other hand, observations in the optical and in the NIR should be considered complementary, as both wavelength ranges may give information about different aspects of the feature. In fact, they may differ not only in the value of certain fitting parameters, but in the functional form of the truncation curve itself. For example, beyond the truncation (or break) radius, non negligible star formation could affect the young population distribution but not that of the old population. If stars did not move from their birthplace, a difference in the blue and the NIR profiles would not be expected, but stars can move.  If there is a (smooth and complete) truncation in the IR, as shown by F01 and this study, and a sharp break in the optical, as shown by Pohlen et al. (2004), this is not necessarily a contradiction; it could be the sign of  differing behaviour patterns in the optical and in the near infrared, for dynamic reasons. We choose the traditional description introduced by van der Kruit (1987), introducing a taring function to smooth out the sharp cut-off, also called a soft cut-off, and propose a truncation curve proportional to $(R_{tr} - R)^{-n}$, with $n= 0.8 \\pm 0.3$, a value close to 1, i.e. it becomes complete at $R=R_{tr}$. This is the profile of type 2, labelled ``smooth'' in Fig. 1.  The type 3 profile is clearly different from that of type 2, either due to the different information provided by the NIR and the optical, or due to insufficient coverage of the NIR observations. In the light of the former consideration, we have chosen the simplest fit suggested by our data, but they are not deep enough to conclude that an external region with a second slope does not exists.  Unfortunately, there are not many galaxies measured in both the optical and the NIR. One of them is NGC 522, which was measured in the optical by Pohlen et al. (2002b). This galaxy is a very good example of a two-slope with a narrow break region in between, when measured in the optical. However, as can be seen in Fig. 5, the J-profile on the right hand side seems to be smooth. The left side as well as the right Ks side could be described as a two-slope. The data does not allow us to clearly distinguish between these two types. Hence, no clear conclusions can be obtained from this comparison. This important point should be addressed in a future study.  As the radial scale height is a natural unit of length in a spiral galaxy, $R_{tr}/h$ is a parameter characterising truncations. For these galaxies, we found the following values for  $R_{tr}/h$: 3.2 for J and 2.9 for K$_s$, with $\\sigma=0.9$ and 0.8 respectively for each filter. The colour dependences of $R_{tr}$ will be discussed in Paper II. These values are in reasonable agreement with previous estimations. For  comparison, it should be taken into account that the difference between the radius at which the truncation curve leaves the inner exponential and the radius at which the truncation finishes is of the order of 20-30$\\arcsec$, and that in the break description only the first value is meaningful.  It is not easy to establish how often galaxies are truncated, due to the difficulty in determining whether a given galaxy is non truncated or whether $R_{tr}$ is too distant to be observed. If we consider 7 out of 11 galaxies truncated, this is equivalent to approximately 2/3. However, this figure should be considered with some caution as our sample is small. However, we confirm that the frequency of truncated galaxies is high, which makes this topic of crucial importance in our understanding of the evolution of disk galaxies."
    },
    "0609/astro-ph0609721_arXiv.txt": {
        "abstract": "We studied the orbital evolution of Jupiter-family comets (JFCs), Halley-type comets (HTCs), and long-period comets, and probabilities of their collisions with planets. In our runs the probability of a collision of one object with the Earth could be greater than the sum of probabilities for thousands of other objects. Even without a contribution of such a few bodies, the  probability of a collision of a former JFC with the Earth was greater than 4$\\cdot$10$^{-6}$. This probability is enough for delivery of all the water to Earth's oceans during formation of the giant planets. The ratios of probabilities of collisions of JFCs and HTCs with Venus and Mars to the mass of a planet usually were not smaller than that for Earth. Among 30,000 considered  objects with initial orbits close to those of JFCs, a few objects got Earth-crossing orbits with semi-major axes $a$$<$2 AU and aphelion distances $Q$$<$4.2 AU, or even got inner-Earth ($Q$$<$0.983 AU), Aten, or typical asteroidal orbits, and moved in such orbits for more than 1 Myr (up to tens or even hundreds of Myrs). From a dynamical point of view, the fraction of extinct comets among near-Earth objects can exceed several tens of percent, but, probably, many extinct comets disintegrated into mini-comets and dust during a smaller part of their dynamical lifetimes if these lifetimes were large. ",
        "introduction": "Farinella et al. (1993), Bottke et al. (2002), Binzel et al. (2002), and Weissman et al. (2002) believe that asteroids are the main source of near-Earth objects (NEOs). Wetherill (1988) supposed that half of NEOs are former short-period comets. Trans-Neptunian objects (TNOs) can migrate to the near-Earth space. Duncan et al. (1995) and Kuchner et al. (2002) investigated the migration of TNOs to Neptune's orbit, and Levison \\& Duncan (1997) studied their migration from Neptune's orbit to Jupiter's orbit. Levison et al. (2006) studied formation of Encke-type objects. More references on papers devoted to the migration of bodies from different regions of the solar system to the near-Earth space can be found in our previous publications on this problem (Ipatov 1995, 1999, 2000, 2001, 2002; Ipatov \\& Hahn 1999, Ipatov \\& Mather 2003, 2004a-b, 2006a). As migration of TNOs to Jupiter's orbit was considered by several authors, Ipatov (2002), and Ipatov \\& Mather (2003, 2004a-b, 2006a) paid particular attention to the orbital evolution of Jupiter-crossing objects (JCOs), considering a larger number of JCOs than before.  In the present paper, we summarize our studies of migration of cometary objects into NEO orbits, paying particular attention to the probabilities of collisions of cometary objects with the terrestrial planets. These studies are based on our previous runs and on some new runs. Earlier we did not consider the evolution of orbits of Halley-type comets and long-period comets and did not study the probabilities of collisions of different comets with the giant planets. Though some used runs are the same as earlier, the discussion on migration of small bodies based on these runs is different. ",
        "conclusions": "Some Jupiter-family comets can reach typical NEO orbits and remain there for millions of years. From the dynamical point of view (if comets didn't disintegrate) there could be (not 'must be') many (up to tens of percent) extinct comets among the NEOs, but, probably, many extinct comets disintegrated into mini-comets and dust during a smaller part of their dynamical lifetimes if these lifetimes were large. Disintegration of comets can provide a considerable fraction of cometary dust among the zodiacal dust particles. The probability of a collision of one object moving for a long time in Earth-crossing orbits, with the Earth could be greater than the sum of probabilities for thousands of other objects, even having similar initial orbits. Even without a contribution of such a few bodies, the probability of a collision of a former JFC (during its dynamical lifetime) with the Earth was greater than 4$\\cdot$10$^{-6}$. This probability is enough for delivery of all the water to Earth's oceans during formation of the giant planets. The ratios of probabilities of collisions of JFCs and HTCs with Venus and Mars to the mass of a planet usually were not smaller than that for Earth."
    },
    "0609/astro-ph0609735_arXiv.txt": {
        "abstract": "Wide-field ($\\sim$8$'$ $\\times$ 8$'$) and deep near-infrared ($JHKs$ bands) polarization images of the Orion Nebula are presented. These data revealed various circumstellar structures as infrared reflection nebulae (IRN) around young stellar objects (YSOs), both massive and low-mass. We found the IRN around both IRc2 and BN to be very extensive, suggesting that there might be two extended ($>$0.7 pc) bipolar/monopolar IRN in these sources. We discovered at least 13 smaller-scale ($\\sim$0.01-0.1 pc) IRN around less-massive YSOs including the famous source $\\theta${$^{2}$} Ori C. We also suggest the presence of many unresolved ($<$690 AU) systems around low-mass YSOs and young brown dwarfs showing possible intrinsic polarizations. Wide-field infrared polarimetry is thus demonstrated to be a powerful technique in revealing IRN and hence potential disk/outflow systems among high-mass to substellar YSOs. ",
        "introduction": "The Orion Nebula and its associated population of stars are amongst the best studied objects in the sky. As a component of the ridge of molecular cloud (OMC-1) extending north-south, M42 is the HII region excited by the Trapezium, a group of massive (OB-type) young stars, near the surface of the molecular cloud. Of the Trapezium stars, $\\theta${$^1$} Ori C is the most dominant photo-ionizing source. One arcmin to the north-west on the sky are the two massive YSOs, IRc2 and BN, whose masses are suggested to be 25 and {$\\gtrsim$}7 {$M$}$_{\\sun}$, respectively (see Genzel \\& Stutzki 1989; Jiang et al. 2005). In fact IRc2 is a cluster of sources seen at infrared wavelengths, with additional nearby sources detected at different wavelengths (see e.g., Beuther et al. 2004) such as the radio source (source I) and the submillimetre source (SMA1). For convenience we refer to the sources responsible for the major outflows as the ``IRc2 prtostars''. At least two outflows are known to originate from the region near the IRc2 protostars; one high-velocity outflow in the SE-NW direction, observed in radio molecular lines, and in optical and near-infrared shocked lines (Allen \\& Burton 1993; Chernin \\& Wright 1996), and one lower-velocity outflow in the NE-SW direction best seen in thermal SiO and H$_{2}$O maser emission, as well as some H$_{2}$ bow shocks (Genzel \\& Stutzki 1989; Chrysostomou et al. 1997). The driving sources of these outflows are uncertain. Situated to the south is the Bright Bar, which corresponds to the tangential region of the wall of the ionization front sculpting OMC-1. Distributed over the entire nebula is the Orion Nebular Cluster (ONC; aka. Trapezium cluster), which is composed of some 3500 low-mass YSOs and about 150 stars with protoplanetary disks revealed by $HST$ imaging (O'Dell 2001).  Many of these features are clearly seen with near-infrared intensity imaging (see e.g., Hillenbrand \\& Carpenter 2000; see also Figure~1). Near-infrared wavelengths are useful because of a combination of the low temperature of the YSOs and the extensive dust-extinction over the star forming region. However, near-infrared radiation from massive star forming regions is varied in its origin: not only direct radiation from YSOs but also gas free-free continuum emission, dust thermal emission, and dust scattering from various illuminating sources. The last scattered component from both diffuse nebulae over the whole region and the local environs around YSOs are best studied with polarized radiation.  However, polarimetric studies covering the central few arcmin of M42, where a variety of activities are seen have been limited so far. Imaging polarimetry has been conducted either only at optical wavelengths (Pallister et al. 1977) or toward a small region near IRc2 or BN at near-infrared wavelengths (Minchin et al. 1991; Jiang et al. 2005; Simpson et al. 2006). In this {Letter}, we present for the first time the near-infrared polarization images of a $\\sim$1 pc region of M42 and describe newly detected features seen as IR nebulae on various scales. ",
        "conclusions": "\\subsection{Polarized Intensity Images of M42 and Large Scale Infrared Nebulae} The outflows produced by both massive and low-mass YSOs are or were powerful enough to open a cavity within their parent cloud or core, in a direction that tends to be perpendicular to dense disk-like structures around the star. The radiation from the star preferentially escapes in the polar direction of the cavity, and is then scattered by the dust at the wall of the cavity, forming so-called polarized infrared reflection nebulae (IRN; see e.g., Hodapp 1984; Sato et al. 1985; Weintraub, Goodman, \\& Akeson 2000)  Figure~1 shows the three-color composite, intensity and polarized intensity images in the $J$, $H$, and $Ks$ bands. The IRN revealed by our polarization images is a signature of circumstellar structures of various scales (from $\\sim$1 pc down to less than the seeing size, $<$690 AU). The infrared polarization has been used to derive the morphology of the IRN associated with past or present outflow phenomena (Tamura et al. 1991; Sato et al. 1985; Hodapp 1984), either via spatially resolving the IRN or via detecting significant levels of integrated (i.e., unresolved) polarization.   Our polarization images reveal a new picture of the geometry in this region. The most dominant feature is the bipolar IRN centered on the IRc2 protostars, seen as red nebulae extending predominantly to the east and west in the color composite image (note that an arc to the north belongs to another IRN illuminated by the BN object as described below) and the $Ks$ band polarization vector map in Figure~2 (see also a sketch in Figure~3). This most likely traces the cavity carved out by the molecular outflow from the IRc2 protostars. First, the highly polarized (up to $p$ $\\sim$ 40 \\% at $Ks$) IRN is more extended than previously thought; the extent of the outflow traced by our polarization images is $\\sim$0.7 pc east-west, while the extent of the H$_{2}$ outflow has been suggested to be $\\sim$0.3 pc to the NNW and only $\\sim$0.1 pc to the west (Allen \\& Burton 1993; Schultz et al. 1999). Therefore, the molecular cloud near the IRc2 protostars might already have had a cavity formed with a size-comparable to the diameter of the ionized region around the Trapezium. Second, the morphology is different between the H$_{2}$ emission and the IRN; in H$_{2}$ emission the well known ``fingers'' (Allen \\& Burton 1993) most prominently extend to the NW, while in polarized emission the extension is to the west. The western nebula clearly shows a V-shaped cavity structure with a wide opening angle ($\\sim$50$\\arcdeg$), which is typical for IRN associated with molecular outflows (Tamura et al. 1991; Staude \\& Els{\\\"a}sser 1993). This, together with the extension to the east, makes the total outflow quite symmetric around the IRc2 protostars in both direction and size. Note that the cavity direction is perpendicular to the dense ridge of OMC-1. Once UV radiation from the IRc2 protostars becomes powerful enough, the cavity evacuated by the outflow might become a circular-symmetric HII region, as is seen around the Trapezium.  In addition to the IRN illuminated by the IRc2 protostars, there is another large IRN illuminated by BN (Simpson et al. 2006) as also seen in the polarization vector pattern near BN (Figure~2). This IRN extends to the north (seen as a reddish arc with its apex coincident with BN), but no clear counterpart lobe is seen in our images. Although the IRN could simply be due to dust in the cloud that is illuminated by BN, we suggest that the IRN associated with BN might represent an independent outflow, whose direction is distinct from that of the IRc2 protostars. Such a large (0.3 pc) and curved morphology starting from BN whose direction is consistent with the polar region of the compact polarization disk (Jiang et al. 2005) could be carved out by a massive outflow associated with BN. The misalignment of the BN and IRc2 ``outflows'' is not surprising if BN is a ``run-away'' star from the Trapezium region (Tan 2004).  \\subsection{Medium Scale Infrared Nebulae around YSOs} A number of smaller-scale ($\\sim$0.01-0.1 pc) IRN are also seen in our polarization images; there are at least 13 ``resolved'' IRN (see Figure~3). They are around low- to intermediate-mass YSOs. Most prominent are those near OMC-1 S and $\\theta${$^2$} Ori C. The former is seen as a reddish and irregular pattern in the polarized images (see Hashimoto et al. 2006 for $JKs$ data and more discussion). The latter shows a bipolar-pattern in the polarized intensity and a centro-symmetric polarization vector pattern. Although $\\theta${$^2$} Ori C is a well-known object and has been observed at many times, no circumstellar structure has been suggested previously. Note, however, that care must be taken in determining the exact geometry of IRN from the polarized intensity images in Figure~1 alone, because a combination of circumstellar scattering of light from the central star and illumination by the Trapezium stars, could systematically change the apparent polarization pattern. These (as well as unresolved, highly polarized sources described below) might serve as the best targets for future high-resolution imaging with adaptive optics (e.g., Perrin et al. 2004).  \\subsection{Unresolved Infrared Nebulae around Low-Mass YSOs and Young Brown Dwarfs} There are numerous ``unresolved'' (point-like, $<$690 AU) sources detected in our near-infrared images, most of which are low-mass YSOs and comprise part of the ONC cluster. The circumstellar structure is not resolved in our 1$\\farcs$5 resolution images. However, software aperture polarimetry of these sources can provide geometrical evidence for circumstellar structures (Sato et al. 1985; Tamura \\& Sato 1989). Interestingly, 86 (17\\%) sources show an aperture polarization larger than 6\\%, while 58 have a larger $p(H)/(H-K)$ value than that for BN. Such a large near-infrared polarization is likely to arise from scattering in the small-scale IRN system around each source, rather than originating from interstellar or intra-cloud extinction, except for extremely red sources. The details of these polarized sources are described elsewhere (N. Kusakabe et al. in preparation), but in this {Letter}, we just point out that some young brown dwarf candidates, whose masses have been estimated to be less than 0.08 {$M$}${_{\\odot}}$ from {\\it both} photometry and spectroscopy (Slesnick, Hillenbrand, \\& Carpenter 2004), appear to have intrinsic polarizations because their $p(H)/${$\\tau$}$(H)$ values are greater than that for the dichroic polarization due to extinction in the molecular clouds (Jones, Klebe, \\& Dickey 1992). In particular, the $H$ band polarizations of HC022 (0.05 $M$$_{\\sun}$, $A_{V}$ = 2 mag) and HC064 (0.03 $M$$_{\\sun}$, $A_{V}$ = 4 mag) are 1.49 $\\pm$ 0.06 \\% and 5.82 $\\pm$ 0.09 \\%, respectively, and are most likely due to scattering in unresolved IRN. It has been suggested that young brown dwarfs are associated with disks based on indirect evidence such as excess IR emission (e.g., Muench et al. 2001). Besides a few examples of disk imaging by $HST$ (Bally, O'Dell, \\& McCaughrean 2000) and a detection of an outflow by spectro-astrometry (Whelan et al. 2005), our results serve as direct ``geometrical'' evidence for IRN associated with disk-outflow systems around young brown dwarfs.  \\subsection{Polarized Orion Bar} The polarized radiation at shorter wavelengths (our $J$ band data and previous optical data by Pallister et al. 1977) is dominated by scattered light from the Trapezium stars rather than that from the IRc2 protostars, as evidenced from the polarization vector patterns (Figure 2). Our polarization images show exactly where the scattering of this diffuse nebula (seen as white) occurs; it originates from the boundary between the ionization front and the background molecular cloud, OMC-1. This is evidenced by the systematic displacement of the position of the Orion Bar seen in the intensity image (corresponding to the ionization front) and that seen in the polarized intensity image (corresponding to the scattering region), and the very low polarization near the Trapezium. The distance between the ionization front and the dust scattering region is $\\sim$0.02 pc. The typical polarization level in the Bar region is 5-10\\% in the $J$ band. The scattering dust is mostly concentrated at the boundary because the dust is swept-up by stellar radiation pressure from the Trapezium stars and slowed down quickly by gas drag in the molecular cloud (Ferland 2001)."
    },
    "0609/astro-ph0609029_arXiv.txt": {
        "abstract": "{ We present new grids of evolutionary models for the so-colled ``Anomalous'' Cepheids (ACs), adopting $Z$=0.0001 and various assumptions on the progenitor mass and mass-loss efficiency. These computations are combined with the results of our previous set of pulsation models and used to build synthetic populations of the predicted pulsators as well as to provide a Mass-Luminosity relation in the absence of mass-loss. We investigate the effect of mass-loss on the predicted boundaries of the instability strip and we find that the only significant dependence occurs in the Period-Magnitude plane, where the synthetic distribution of the pulsators is, on average, brighter by about 0.1 mag than the one in absence of mass-loss. Tight Period-Magnitude relations are derived in the $K$ band for both fundamental and first overtone pulsators, providing a useful tool for distance evaluations with an intrinsic uncertainty of about 0.15 mag, which decreases to $\\sim$ 0.04 mag if the mass term is taken into account. The constraints provided by the evolutionary models are used to derive evolutionary (i.e, mass-independent) Period-Magnitude-Color relations which provide distance determinations with a formal uncertainty of the order of $\\sim$ 0.1 mag, once the intrinsic colors are well known.  We also use model computations from the literature to investigate the effect of metal content both on the instability strip and on the evolutionary Period-Magnitude-Color relations. Finally, we compare our theoretical predictions with observed variables and we confirm that a secure identification of actual ACs requires the simultaneous information on period, magnitude and color, that also provide constraints on the pulsation mode. }  \\authorrunning{Fiorentino \\it{et al.}} \\titlerunning{Synthetic properties of bright metal-poor variables. I. } ",
        "introduction": "While the majority of metal-poor radial pulsating variables is represented by the RR Lyrae stars, other classes of variables are observed in metal-poor stellar systems.  According to the current literature, one finds the ``Anomalous'' Cepheids (AC) and the ``Population II'' Cepheids (P2C), the former with periods ($P$) from $\\sim$ 0.5 to $\\sim$ 2 days, the latter from $\\sim$ 1 to 25 days. Both types of variables are brighter than RR Lyrae stars and belong to the central He-burning evolutionary phase, as are RR Lyrae stars. However. ACs are more massive whereas P2Cs are less massive than RR Lyrae stars with similar metal content. The purpose of the present paper is to perform a theoretical analysis of the ACs, which are observed in the majority of the Local Group dwarf galaxies which have been surveyed for variable stars, but that are almost absent in other metal-poor stellar systems such as Galactic Globular Clusters. The origin of this ``anomaly'' is the evidence that they do not follow the Period-Luminosity relation of the P2Cs observed in these old metal deficient clusters, being significantly brighter at fixed period (see Norris \\& Zinn 1975; Zinn \\& Searle 1976; Smith \\& Stryker 1986; Nemec et al.  1994).   Several authors (Dolphin et al. 2002, 2003; Clementini et al. 2003; Cordier, Goupil \\& Lebreton 2003) have suggested that ACs are the natural extension of the Population I Classical Cepheids to lower metal contents and smaller masses. This suggestion is well supported by theoretical investigations (Marconi, Fiorentino \\& Caputo 2004 [MFC]; Caputo et al. 2004 [C04]) where, based on the constraints provided by pulsation and evolutionary models, it is shown that Anomalous and Classical Cepheids define a common region in the $M_V$-log$P$ plane, with the former ones at lower luminosities and shorter periods, as actually observed in the dwarf irregular galaxy Leo A (Dolphin et al. 2002). Moreover, it is also shown that this region is well separated from that populated by RR Lyrae stars and P2Cs, in full agreement with observations.  Concerning the evolutionary phase, there is a general consensus that ACs are central He-burning stars with a mass around 1.5\\msun. Previous studies (see Castellani \\& Degl'Innocenti 1995 [CD95]; Bono et al. 1997 [B97]; Caputo 1998; C04 and references therein) have shown that for metal abundances $Z \\le$ 0.0004 the central He-burning models more massive than $\\sim$ 1.2\\msun evolve into the pulsation region at a luminosity not dramatically higher than the RRL level, predicting pulsators with AC-like periods and luminosities.  In this study, we will extend those investigations in order to define a sound theoretical scenario for the analysis of the AC properties. In particular, our intention is to verify their use as distance indicators and to study the effects of mass-loss on the various relations connecting evolutionary and pulsational quantities.  The paper is organized as follows: in Section 2, we present new evolutionary tracks computed without mass-loss, as well as the central He-burning models obtained from a given progenitor star assuming different amounts of mass-loss. Using also the constraints provided by pulsation models, synthetic populations of the predicted pulsators are obtained. Section 3 presents the results for the Color-Magnitude diagram as well as the Period-Magnitude, Period-Color, and Period-Magnitude-Color relations. Section 4 deals with the comparison between theoretical predictions and observations and the Conclusions close the paper.  \\begin{figure} \\includegraphics[width=8cm]{ACf1.ps} \\caption{\\small{Canonical (no mass-loss) evolutionary tracks of central He-burning models with $Z$=0.0001, $Y$=0.24 and the labeled masses. In order to point out the connection of evolution with pulsation, we subtract the value at the red edge (FRE) of the pulsation region from the effective temperature of each model. Dots indicate the initial central He-burning phase, crosses indicate the central He-exhaustion phase corresponding to t$_{He}$ (see Table 1), as the vertical lines represent the predicted boundaries of the instability strip (see text for details).}} \\end{figure} ",
        "conclusions": "We have presented an homogeneous and updated theoretical scenario for the study of AC properties based on a new set of evolutionary tracks with $Z$=0.0001, progenitor mass in the range 1.3-2 \\msun and  several assumptions on the efficiency of mass loss. This evolutionary framework is related to the pulsation model predictions by MFC. We find that predicted pulsators with mass 1.3-2.0 $M_{\\odot}$ follow a ML relation as given by $\\langle$log$L\\rangle \\sim$ 1.77($\\pm$0.05)+2.07log$M$, in absence of mass loss. Moreover, the evolutionary and pulsational models have been used to build synthetic populations of AC pulsators, each characterized by a progenitor mass, a mean mass and a percentage of central He-burning stars populating the IS. We find that  pulsators with a given mass originating in the presence of mass loss are systematically fainter than the ones in the absence of mass loss, with the decrease in luminosity following the variation in the He-core mass. We have also investigated the effect of mass loss on the predicted IS boundaries in the $CM$, $PC$ and $PM$ planes and we find that the only significant dependence occurs  in the $PM$ plane where the synthetic distribution in the presence of mass loss is, on average, brighter by about 0.1 mag than the one in the absence of mass loss.  Concerning the existence of a $PM$ relation for these stars, we confirm that in the case of optical magnitudes it depends on the pulsator distribution within the IS, whereas tight near-infrared $PM_K$ relations can be derived for both fundamental and first overtone pulsators, providing a reliable tool for distance evaluations with an intrinsic uncertainty of about 0.15 mag. If the mass term is taken into account a tighter (mass-dependent) PM$_K$ relation (r.m.s=0.04 mag.) is obtained. Conversely, the predicted $PC$ relations dealing with $B-V$ and $V-I$ colors appear useful for reddening determinations with an intrinsic uncertainty of less than 0.1 mag.  Mass-dependent $PMC$ relations have been derived by using all the models, as well as evolutionary $PMC_e$ relations where the mass-term is removed. The latter relations allow distance determinations with a formal uncertainty of the order of 0.1 mag, provided that the intrinsic colors are well known.  To take into account the metallicity effect, we have used other evolutionary tracks from the literature with $Z$ in the range 0.0001-0.001. In particular, we find that an increase of the metal content from 0.0001 to 0.0004 yields that, at fixed mass, the average luminosity of predicted AC pulsators decreases by $\\sim$ 0.1 dex, producing an effect on the IS boundaries similar to the one due to mass loss, whereas the zero-points of the predicted $PMC_e$ relations becomes brighter  by $\\sim$ 0.1 mag.  By comparing the predicted edges of the IS with observed pulsators in the $CM$, $PC$ and $PM$ planes, we find discordant results in the properties of individual stars, thus confirming that a sure identification of actual ACs requires simultaneous information on period, magnitude and color through the application of the predicted $PMC_e$ relations. In this case, we are also able to provide constraints on the pulsation mode.  We show that AC candidates that satisfy the predicted edges of the IS in the $CM$, $PC$ and $PM_V$ diagrams can be used to obtain empirical $PM_V$ relations for F and FO pulsators."
    },
    "0609/astro-ph0609359_arXiv.txt": {
        "abstract": "Population synthesis is used to study the contribution from unresolved sources to the Galactic ridge emission measured by EGRET. Synthesized source counts are compared with the 3rd EGRET catalogue at low and high latitudes. For pulsar-like populations, 5-10\\% of the emission $>100$ MeV comes from sources below the EGRET threshold. A steeper luminosity function can increase this to  20\\%  without violating EGRET source statistics. Less luminous populations can produce much higher values without being detected. Since the unresolved source spectrum is different from the interstellar spectrum, it could  provide an explanation of the observed MeV and GeV excesses above the predictions, and we give an explicit example of how this could work. ",
        "introduction": "The Galactic plane is known to be an intense emitter of X-rays and \\grays from keV to at least 100 GeV. At X-ray energies, it has recently been claimed \\citep{2006A&A...452..169R} % that the 2-10 keV emission can be explained entirely by a population of weak  sources, mainly CVs. Above 50 keV, unresolved sources also appear to be required to explain the hard-power law emission observed by INTEGRAL \\citep{2005A&A...444..495S,  % 2005ApJ...635.1103B}; % AXPs and/or pulsars are potential candidates. This   raises the question of the situation for \\grays, in particular the range observed by EGRET, 30 MeV - 100 GeV. \\gray telescopes  are relatively  insensitive and reveal only the `tip of the iceberg' of the sources, so most of them will  go undetected unless it happens that only strong sources exist in the Galaxy. Conventional wisdom is that the source contribution to the unresolved ridge emission is at the few percent level, The problem of the GeV excess in the diffuse emission compared to the expected spectrum from interstellar processes \\citep{2004ApJ...613..962S},  % is however a hint that the emission may have other  components whose contribution is also energy-dependent. The failure to explain the 1-30 MeV emission measured by COMPTEL from interstellar components is another pointer to a source  population.      This topic has wide ranging implications: for example it has been claimed that the GeV excess is a signature of dark matter decay \\citep{2005A&A...444...51D}% \\footnote{but see  critique by \\citet{2006JCAP...05..006B}.} % , but this is only plausible if  all  viable alternatives are excluded, and source populations provide one important such candidate. ",
        "conclusions": "1. Modelling the contribution from unresolved sources  is  essential to understanding  the diffuse Galactic emission .  \\noindent 2. The contribution from unresolved sources to the EGRET low-latitude emission is at least 5-10\\%, and can be 20\\% for steep luminosity functions.  \\noindent 3. An arbitrarily large fraction of the diffuse emission could come from sources  $\\Lg$($>$100 MeV) $<10^{35}$ \\Lgunits from sources without violating EGRET source counts anywhere on the sky.  \\noindent 4. The GeV excess can be produced naturally by a sufficient population of  sources like Geminga and Vela, but this has to be studied with more detailed models.  \\noindent 5. Crab-like sources can produce the COMPTEL excess but not the GeV excess.  \\noindent 6. A combination of source populations combined with the conventional model of interstellar emission could explain the full  COMPTEL/EGRET Galactic ridge spectrum.  \\noindent 7. Whether  GLAST can settle these issues depends critically on the source luminosity function."
    },
    "0609/astro-ph0609390_arXiv.txt": {
        "abstract": "{ The results obtained in paper I are used to study possible relationships between the truncation radius of stellar discs in the NIR and structural parameters of the galaxies. The NIR truncation radius is larger for brighter galaxies, being proportional to  $V_m^c$ with $c \\approx 3/2$, and with $V_m$ being the asymptotic rotation velocity at large radii (when the rotation curve becomes flat), and is lower for higher wavelengths. When it is normalized to the scalelength, the truncation is an increasing function of the central surface brightness and is lower for late type galaxies, although these correlations are weaker. These relations are in agreement with the scenario of magnetically driven truncations. ",
        "introduction": "Paper I (Florido et al. 2005) presented the result of the observation of 18 edge-on spiral galaxies in the NIR in order to study the truncation of the stellar disc, the main objective being to describe the truncation curve. This curve is defined as the progressive decrease of the photometric profile taking place at large radii with respect to the mean exponential profile obtained at moderate radii. In this paper we study the relationship of truncation to other parameters of the galaxy. We investigate some key properties that could restrict some of the theoretical scenarios that have been proposed: what types of galaxy show truncation and what relation exists between the truncation, radius, mass, luminosity, rotation velocity at large radii, type, central surface brightness and  age of stellar population. The results obtained from this study will be merged with those from a previous one (Florido et al. 2001) in order to obtain a larger sample. No merger with other observations in the optical has been attempted, as the assumption of the same behaviour for optical and NIR truncations is not theoretically justified.  In early works by van der Kruit (1979) truncations were found to have sharp cut-offs. In paper I and in Florido et al. (2001) truncations were found to be smooth. In these latter two papers it was shown that the smooth behaviour of truncations was not due to projection effects, but was real, existing even in deprojected profiles. In some recent papers (e.g. Pohlen et al. 2002, P\\'erez 2004) this description is replaced by a two-exponential profile with a sharp elbow. There is increasing evidence favouring the two-slope description, at least in the optical. We are not interested here in the shape of the profile (see paper I); it may even be that the differences between authors are due to a true difference between the NIR and the optical. For the present, it is not our goal to determine whether this discrepancy is real or is a result of the low signal/noise at these very large radii, especially important in the NIR, and we refer to a ``truncation (break)'' or simply ``truncation''.  In paper I it was proposed that most galaxies in the NIR seem to have a truncation curve that is proportional to $(R-R_{tr})^{-n}$ where $n \\approx 1$, $R$ is the radius and $R_{tr}$ is the truncation radius at which the truncation seems to be complete, as obtained by extrapolation of the data. ",
        "conclusions": "The magnetic model satisfactorily explains all the statistical properties reported in this paper. We thus propose this scenario as responsible for producing truncations. The agreement is particularly noticeable concerning the relation between truncation radius and the power 3/2 of the asymptotic rotation velocity.  Magnetic forces are important in the rotation dynamics of the gas, as deduced from observed magnetic field strengths. These forces no longer act when gas is converted into stars, which must have a large influence on the stability of the orbits of young stars, even leading to massive escape.  Large galaxies of type Sa and Sb have less need for dark matter, hence, in the magnetic interpretation, they require lower magnetic fields, i.e. the truncation would be a less important effect in large (massive and luminous) galaxies, especially those with a prominent bulge. The fact that late-type galaxies, which are  bulge-poor, could be more affected by truncations is in agreement with the magnetic scenario.  Pohlen et al. (2004) found several examples of truncations in edge-on SO galaxies. Even if these galaxies are at present gas-poor, they had both gas and star formation in the past. Thus, the mechanisms suggested by the magnetic hypothesis would have been operating in the past, and they would have inherited truncations.  The fact that the truncation in the optical could be  better described by a two-exponential profile could be compatible with a smooth, complete truncation in the near infrared. The newly formed stars would be, under the magnetic interpretation, in the process of escaping. If the escape speed is roughly equal to $V_m$, i.e. of the order of 100 km s$^{-1}$, and if the length of the escape path is of the order of $\\sim$ 10 kpc, the escape time would be of the order of $10^8$ years, a time that differentiates clearly between new-born and old stars. The fact that $R_{tr}(K_s) <R_{tr}(J)$ also fits this interpretation.  In the magnetic interpretation, the truncation corresponds to a true decrease in the total mass (stars + gas) and a connection with a decrease in the rotation velocity should be expected. This decrease has been observed (Casertano 1983, van der Kruit 2001, Bottema 1995) although this fact requires more observational support and may be explained by other means."
    },
    "0609/astro-ph0609673_arXiv.txt": {
        "abstract": "We present recent results for galactic WNL stars, obtained with the new Potsdam Wolf-Rayet (PoWR) hydrodynamic model atmospheres. Based on a combination of stellar wind modeling and spectral analysis we identify the galactic WNL subtypes as a group of extremely luminous stars close to the Eddington limit. Their luminosities imply progenitor masses around 120\\,$M_\\odot$ or even above, making them the direct descendants of the most massive stars in the galaxy. Because of the proximity to the Eddington limit our models are very sensitive to the L/M ratio, thus allowing for a direct estimate of the present masses of these objects. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609445_arXiv.txt": {
        "abstract": "We discuss the importance for the long-term cluster evolution of the mass loss from intermediate-mass stars (0.8--8 M$_\\odot$). We present constraints on the mass loss from red giants in clusters in the Magellanic Clouds, a search for the intra-cluster medium in galactic globular clusters, and a simple estimate for the cluster evolution due to red giant mass loss compared to stellar escape. ",
        "introduction": "The evolution of a star cluster depends critically on its ability or failure to retain its mass. Dynamical processes cause mass segregation and preferential escape of low-mass stars from clusters. Besides losing entire stars, clusters also lose mass due to envelope mass loss from its member stars, exacerbating cluster dispersal. We here briefly discuss the importance for the long-term cluster evolution of the mass loss from intermediate-mass stars in the mass range 0.8--8 M$_\\odot$. ",
        "conclusions": ""
    },
    "0609/astro-ph0609323_arXiv.txt": {
        "abstract": "The nature of the Local Interstellar Cloud (LIC) is highly constrained by the combination of \\emph{in situ} heliospheric and line-of-sight data towards nearby stars. We present a new interpretation of the LIC components of the absorption line data towards \\epsC, based on recent atomic data that include new rates for the Mg$^+$ to Mg$^0$ dielectronic recombination rate, and using \\emph{in situ} measurements of the temperature and density of neutral helium inside of the heliosphere.  With these data we are able to place interesting limits on the gas phase abundance of carbon in the LIC.  If the C/S abundance ratio is solar, $\\sim 20$, then no simultaneous solution exists for the $N($\\ion{Mg}{1}$)$, $N($\\ion{Mg}{2}$)$, $N($\\ion{C}{2}$)$ and $N($\\ion{C}{2}$^*)$ data. The combined column density and \\emph{in situ} data favor an abundance ratio \\ACAS\\ $= 47^{+22}_{-26}$.  We find that the most probable gas phase C abundance is in the range 400 to 800 ppm with a lower limit of $\\sim 330$.  We speculate that such a supersolar abundance could have come to be present in the LIC via destruction of decoupled dust grains.  Similar enhanced C/H ratios are seen in very low column density material, \\NH$ < 10^{19}$ \\scm, towards several nearby stars. ",
        "introduction": "\\label{sec:intro}  The interstellar cloud that surrounds our solar system, the Local Interstellar Cloud (LIC), is a tiny, radius $\\lesssim 1$ pc, low density interstellar cloud that has been detected in absorption lines towards many towards nearby stars \\citep[e.g.,][]{McClintocketal:1978,Bruhweiler:1984,Genovaetal:1990, Lallement+Ferlet_1997,Hebrard_etal_1999,Gry+Jenkins_2001, Wood+Linsky_1997}.  The cloud is partially ionized, $\\sim 20-40$\\%, apparently due to the low intensity, relatively hard ionizing radiation field \\citep[][hereafter SF02]{Slavin+Frisch_2002,FrischSlavin:2003}. Direct observations of interstellar He$^0$ which penetrates the heliosphere gives us tight constraints on its temperature and density: $T = 6300\\pm340$ K, $n($He$^0) = 0.015\\pm0.0015$ cm$^{-3}$ \\citep{Witte_2004,Moebius_etal_2004}.  The star $\\epsilon$ CMa (HD 52089) is located in the downwind direction of the flow of ISM past the Sun ($\\ell = 239.8$\\degr, $b = -11.3$\\degr, $d = 132$ pc).  The absorption line data set for the line of sight towards \\epsC\\ is particularly complete \\citep[][hereafter GJ]{Gry+Jenkins_2001}.  Using this dataset, in particular the ratios $N($\\ion{Mg}{2}$)/N($\\ion{Mg}{1}$)$ and $N($\\ion{C}{2}$^*)/N($\\ion{C}{2}$)$, GJ put constraints on the electron density, $n(e)$, and temperature of the LIC and of the other clouds seen in this line of sight.  As we discuss, their method involved assuming a gas phase abundance ratio for C/S of 20, appropriate to Solar abundances.  However, we show that when updated atomic rates are used for the Mg$^+$ recombination coefficient, in particular for dielectronic recombination, this solution becomes no longer viable.  This leads us to the conclusion that C has a supersolar abundance in the LIC, $A_\\mathrm{C} =  400 - 800$ ppm (parts per million).  This result, which is consistent with detailed models of the photoionization of the LIC \\citep[][hereafter SF06]{Slavin+Frisch_2006}, then allows for a solution for $n(e)$ which is consistent with all other data. High carbon abundances are also seen towards several other stars that sample the LIC. ",
        "conclusions": "We have shown that using updated atomic data and the \\emph{in situ} data on the temperature of the gas in the LIC, we can put interesting constraints on the C abundance in the LIC.  Unless a number of observed or inferred data values are at the far ends of their uncertainty ranges, it appears that C is not only undepleted in the LIC, but has an abundance of $A_\\mathrm{C} \\approx 400-800$ ppm, considerably larger than the solar value.  In addition, the ratio of C to S abundance appears to be necessarily larger than the solar ratio of $\\sim 20$ with values in the range of $40 - 50$ favored by the constraints on $n(e)$ in the LIC.  We speculate that this supersolar abundance of C, if it truly exists, may be indicative of a local enhancement in dust in the past.  In this picture the gas phase C abundance was enhanced by the passage of an interstellar shock that destroyed the carbon bearing grains more effectively than other, e.g.\\ silicate, grains. This could be related to the unusual dust size distribution in the LIC as indicated by Ulysses observations of large grains penetrating into the heliosphere \\citep{Grun_etal_1994,Frisch_etal_1999}.  Further investigations into the C abundance in the LIC should be undertaken to confirm our findings, especially studies of other lines of sight that show the LIC velocity component in the absorption lines of C, S, Mg and other elements. The LIC is the best observed low density interstellar cloud and thus can provide us with valuable information on this phase of the interstellar medium. The C abundance and its evolution in the warm diffuse ISM has many implications for the nature of interstellar dust and metallicity in the Galaxy."
    },
    "0609/astro-ph0609115_arXiv.txt": {
        "abstract": "We study the dependence of the exposure required to directly detect a WIMP directional recoil signal on the capabilities of a directional detector. Specifically we consider variations in the nuclear recoil energy threshold, the background rate, whether the detector measures the recoil momentum vector in 2 or 3 dimensions and whether or not the sense of the momentum vector can be determined. We find that the property with the biggest effect on the required exposure is the measurement of the momentum vector sense. If the detector cannot determine the recoil sense, the exposure required is increased by an order of magnitude for 3-d read-out and two orders of magnitude for 2-d read-out. For 2-d read-out the required exposure, in particular if the senses can not be measured, can be significantly reduced by analyzing the reduced angles with the, time dependent, projected direction of solar motion subtracted. The background rate effectively places a lower limit on the WIMP cross-section to which the detector is sensitive; it will be very difficult to detect WIMPs with a signal rate more than an order of magnitude below the background rate. Lowering the energy threshold also reduces the required exposure, but only for thresholds above 20 keV. ",
        "introduction": "Weakly Interacting Massive Particle (WIMP) direct detection experiments aim to directly detect non-baryonic dark matter via the elastic scattering of WIMPs on detector nuclei~\\cite{DD}, and are presently reaching the sensitivity required to detect the lightest neutralino (which in most supersymmetry models is the lightest supersymmetric particle and an excellent WIMP candidate). Since the expected event rates are very small ( ${\\cal O} (10^{-5} - 1)$ counts ${\\rm kg^{-1} day^{-1}}$) distinguishing a putative WIMP signal from backgrounds due to, for instance, neutrons from cosmic-ray induced muons or natural radioactivity, is crucial. The Earth's motion about the Sun provides two potential WIMP smoking guns: i) an annual modulation~\\cite{amtheory} and ii) a strong direction dependence~\\cite{dirndep} of the event rate. The dependence of the differential event rate on the atomic mass of the detector (see e.g. Refs.~\\cite{jkg,ls,kraus}) is a third possibility, however this would require good control of systematics for detectors composed of different materials.  The direction dependence of the WIMP scattering rate has several advantages over the annual modulation. Firstly, the amplitude of the annual modulation is typically of order a few per-cent~\\cite{amtheory,ls} while the event rate in the direction of solar motion is roughly an order of magnitude larger than that in the opposite direction~\\cite{dirndep,ls}. Secondly, it is difficult for the directional signal to be mimicked by backgrounds. In most cases (a point source in the lab is a possible exception) a background which is anisotropic in the laboratory frame will be isotropic in the Galactic rest frame as the time dependent conversion between the lab and Galactic co-ordinate frames will wash out any lab specific features. Low pressure gas time projection chambers (TPCs), such as DRIFT ({\\bf D}irectional {\\bf R}ecoil {\\bf I}dentification {\\bf F}rom {\\bf T}racks)~\\cite{drift,sean:drift} and NEWAGE~\\cite{newage}, see also Ref.~\\cite{martoff}, seem to offer the best prospects for a detector capable of measuring the directions of sub-100 keV nuclear recoils.   Early studies found that in principle as few as 30 events would be required to distinguish a WIMP induced signal from isotropic backgrounds~\\cite{copi:krauss,lehner:dir}. In reality the number of events, and more importantly the required exposure, will depend on the detector properties including the energy threshold, background event rate, whether the read-out measures the recoil momentum in 2 or 3 dimensions (and if 2-d in which plane)~\\cite{pap1,pap2,copi2d} and whether the sense (i.e. the absolute sign $+ {\\vec{p}}$ or $- { \\vec{p}}$) of the recoil momentum vector can be measured~\\cite{pap1,pap2,copi2d}. In this paper we extend our previous work~\\cite{pap1,pap2} by considering the effects of non-zero background and varying energy threshold and also making a systematic comparison of the number of events, and hence the exposure, required to reject isotropy for benchmark detector configurations covering the range of possibilities. Our aim is to provide guidance for experimentalists wishing to optimize the detection potential of their experiment. The performance of a real detector will of course be more complex than our simulated detector and the cost (both financial and man-hours) of improvements to the detector performance will in reality have to be taken into account when formulating an optimization strategy. None the less we believe that our results will provide useful indications of which aspects of the detector performance have most effect on the ability to detect a WIMP directional signal. ",
        "conclusions": "We have examined the effects of the detector properties on the number of events, and hence exposure, required to reject isotropy of recoil directions, and hence identify a putative WIMP signal as Galactic in origin. Specifically we have considered the nuclear recoil energy threshold, the background rate, 2-d and 3-d read-out and whether or not the sense of the recoils can be measured.  Our simulations are for a TPC detector filled with $CS_2$. We expect that results for other directional detectors would be qualitatively similar.   We took a detector with 3-d vector read-out, energy threshold $E_{\\rm TH}= 20$ keV, no background ($S/N=\\infty$), with the uncertainty in reconstructing the recoil directions taken into account, as our baseline realistic detector configuration and then varied one or two of these properties at a time. The property which has the largest effect on the exposure required is whether or not the sense of the recoils can be measured. If the data is axial (i.e. the sense can not be measured) the exposure is increased by one order of magnitude for 3-d read-out and two orders of magnitude for 2-d read-out compared to the baseline detector. With 2-d read-out this can be improved to a factor $\\sim 30$ if the reduced angles (with the direction of solar motion subtracted) are calculated and analyzed. If the read-out is 2-d, and the senses can be measured, the exposure is increased by a factor of $\\sim 3$ over the baseline. Once more this can be reduced, to a factor of $\\sim 2$, by using the reduced angles. These figures are for 2-d read-out in the optimal plane, which has normal perpendicular to the Earth's spin axis. For other read-out planes the required exposures are larger, and location dependent (see section 3 and Ref.~\\cite{pap2} for further information). We also caution that the results for 2-d read-out assume perfect angular resolution, and projection effects will make this a more significant factor than for 3-d read-out. Estimates based on the projected length of recoil tracks in the planes indicate that the required exposure would increase by at least a factor 2~\\cite{pap2}.    The effect of non-zero background depends on the underlying signal rate.  If the signal and background rates are the same the exposure, relative to that for zero background, is increased by a factor of $\\sim 2.5$. If the background rate is an order of magnitude larger than the signal rate the increase is a factor of $\\sim 15$. As the signal to noise ratio is reduced below 0.1 the detection of the signal anisotropy becomes extremely challenging.  As the energy threshold is increased the WIMP recoil signal becomes increasingly anisotropic and the number of events required decreases. However the event rate decreases and the net effect is that the exposure required increases with increasing energy threshold above $E_{\\rm TH} \\approx 20 \\, {\\rm keV}$. The exposure increases by a factor $\\sim 1.3$ for every $10 \\, {\\rm keV}$ increase in the energy threshold. Below $E_{\\rm TH} \\approx 20 \\, {\\rm keV}$ the exposure is roughly independent of the threshold energy. In practice achieving such low energy thresholds is not possible; the path lengths of low energy recoils are too short for the directions to be measured. We also caution that the actual `energy threshold' of a directional detector is the maximum of the detection, discrimination and direction reconstruction thresholds.   In summary the property with the biggest effect on the detector performance is whether or not the sense of the recoils can be measured. For 2-d read-out the performance, in particular if the senses can not be measured, can be significantly improved by recording and analyzing the reduced angles (with the projected direction of solar motion subtracted). The background rate effectively places a lower limit on the WIMP cross-section to which the detector is sensitive; it will be very difficult to detect a WIMP directional signal with a recoil rate more than an order of magnitude below the background rate. If the energy threshold is larger than about 20 {\\rm keV} significant improvements in the exposure required can be made by reducing the energy threshold (i.e. a factor of 2 if $E_{\\rm th}$ is reduced from $50 \\, {\\rm keV}$ to $20 \\, {\\rm keV}$ and a factor of 5 if it is reduced from $100 \\, {\\rm keV}$ to $50 \\, {\\rm keV}$)."
    },
    "0609/astro-ph0609786_arXiv.txt": {
        "abstract": "{Using high spatial resolution {\\sl Hubble Space Telescope}/Advanced Camera for Surveys archival imaging observations of \\object{Arp 116}, centred on the elliptical galaxy \\object{NGC 4649}, we explore the novel technique of pixel-by-pixel analysis of the galaxy's colour-magnitude diagramme to search for any evidence of recent enhanced star formation due to the apparent tidal interaction with its spiral companion, \\object{NGC 4647}. From a detailed analysis of the system's geometry, and based on additional circumstantial evidence from extant multi-wavelength observations, we conclude that, while there may be grounds for the tidal-interaction assumption for this system, any interaction has thus far been of insufficient strength to trigger an enhanced level of recent star formation in the elliptical component, although close inspection of our colour images shows a faint excess of bluer pixels (a $\\sim 0.20$ mag bluer ``loop'') in the elliptical galaxy on the side of the spiral companion. Given that there appears to be a moderate reservoir of available gas for ongoing star formation (although at low column density), this suggests that we are currently witnessing the {\\it onset} of the tidal interaction between \\object{NGC 4647} and \\object{NGC 4649}. In addition, the triggering of new star formation in \\object{NGC 4649} may be significantly impeded due to the much lower mass of the spiral component.} ",
        "introduction": "One of the major achievements of the {\\sl Hubble Space Telescope (HST)} to date has been the discovery that the progenitors of globular clusters (GCs), once thought to be the oldest building blocks of galaxies, continue to form until today (e.g., de Grijs et al. 2003b, 2005). {\\sl HST} observations allow us to resolve and study GC systems far beyond the well-studied GC systems in the \\object{Milky Way} and the \\object{Magellanic Clouds}.  Young massive star clusters (YMCs; $M_{\\rm cl} \\ga 10^5$ M$_\\odot$) have been detected in a great variety of actively star-forming galaxies. Systems containing a large number of YMCs are most often interacting spiral-spiral pairs (e.g., the \\object{``Antennae''} galaxies, Whitmore \\& Schweizer 1995; \\object{NGC 7252}, Whitmore et al. 1993; the \\object{``Mice''} and \\object{``Tadpole''} systems, de Grijs et al. 2003d), since these naturally provide large reservoirs of gas for star and cluster formation. In fact, massive star {\\it cluster} formation is likely the major mode of star formation in such extreme starburst environments (cf. de Grijs et al. 2003d). Therefore, we can use young and intermediate-age massive star clusters as efficient tracers of the recent violent star formation and interaction history of galaxies by determining accurate ages for the individual clusters, even at those ages when morphological interaction features in their host galaxies have already disappeared.  The formation of YMCs and field stars alike is limited by the supply of gas. The colliding gas masses associated with interacting galaxies determine the violence of the star and cluster formation, as well as the overall gas pressure. Hence the number, and perhaps also the nature (such as their masses and compactness), of the YMCs, and the ratio of YMC to field-star formation in particular, is expected to depend on Hubble type. While observations of gas-rich spiral-spiral mergers are numerous, star-formation tracers in mixed-pair mergers have thus far largely been overlooked, although they are undoubtedly of great interest in the context of the parameter space covered by the star cluster and field star formation processes.  One of the most important open questions in this field, and one that we address in this {\\it Research Note}, is whether the smaller amount of gas available for star formation in early-type galaxies might act as a threshold for star formation in general, and for massive and compact cluster formation in particular, or possibly result in enhanced field star formation (including star formation in small clusters, $M_{\\rm cl} \\la 10^3$ M$_\\odot$) with respect to the mode of star formation in massive clusters.  The present analysis is of particular current relevance as ever more sensitive observations increasingly reveal neutral and/or molecular hydrogen in elliptical and early-type spiral galaxies (e.g., Wiklind \\& Henkel 1989; Cullen et al. 2003, 2006) and in late-stage merger remnants, which might be the progenitors of present-day elliptical galaxies (cf. Georgakakis et al. 2001). Star (cluster) formation in early-type galaxies, long thought to be essentially dust- and gas-free environments, has therefore returned to the forefront of interest. ",
        "conclusions": ""
    },
    "0609/astro-ph0609265_arXiv.txt": {
        "abstract": "Recent observations of quasi-periodic oscillations in the aftermath of giant flares in soft gamma-ray repeaters suggest a close coupling between the seismic motion of the crust after a major quake and the modes of oscillations in a magnetar. In this paper we consider the purely elastic modes of oscillation in the crust of a neutron star in the relativistic Cowling approximation (disregarding any magnetic field). We determine the axial crust modes for a large set of stellar models, using a state-of-the-art crust equation of state and a wide range of core masses and radii. We also devise useful approximate formulae for the mode-frequencies. We show that the relative crust thickness is well described by a function of the compactness of the star and a parameter describing the compressibility of the crust only. Considering the observational data for SGR 1900+14 and SGR 1806$-$20, we demonstrate how our results can be used to constrain the mass and radius of an oscillating neutron star. ",
        "introduction": "\\label{sec:intro} Asteroseismology aims to probe stellar physics via various observed modes of vibration. To be successful in this endeavour one needs both accurate observations and detailed theoretical models to test the observations against. In the case of neutron stars we have until very recently had neither reliable observations nor detailed theoretical models. The situation appears to have changed with the observations of quasiperiodic oscillations (QPOs) following giant flares in three soft gamma-ray repeaters (SGRs) \\citep{israel:qpo,sw:qpo,sw:flare3,sw:flare2}. Analysis of the relevant X-ray data has unveiled a number of periodicities, with frequencies that agree reasonably well with the expected torsional modes of the neutron star crust [\\eg\\ \\citep{duncan:1998A}]. These observations are tremendously exciting because they will allow us, for the very first time, to test our neutron star oscillation models. They also provide strong motivation to improve these models. After all, in order to be accurate enough to make a comparison to the observations meaningful, a model must be fully relativistic. It should also allow for the presence of a strong magnetic field and possible superfluid components.  In this paper we take some modest steps towards more realistic models of pulsating neutron stars by including the crust elasticity in a relativistic calculation of axial oscillations. To simplify the problem we work within the Cowling approximation, \\ie\\ we neglect the dynamical nature of spacetime. This is expected to be an accurate approximation for axial modes in the crust since they originate in a region of relatively low density and should not lead to significant variations in the gravitational field. We are not considering the class of axial gravitational-wave $w$-modes at this point \\citep{kokkotas:axialmodes}, apart from formally noting how they would appear in our formulation (see Appendix~A). We calculate axial modes for an up-to-date crust equation of state (not very controversial) for a large set of core masses and radii, corresponding to different (unspecified) supranuclear equations of state in the core and (similarly unspecified) central pressures. Our approach allows us to largely ignore the detailed structure of the core which is only poorly constrained by observations and theoretical considerations. We show how, when used in conjunction with observed mode frequencies, our numerical results can be used to put constraints on the global equation of state. We also determine an approximate solution to the mode-problem. This leads to useful formulae that show how the axial modes depend on the key parameters, the star's mass and radius, the crust compressibility and the shear modulus. When compared to our numerical data these approximations are surprisingly good. We show how they can be used to deduce the main stellar parameters from a set of observed frequencies, thus outlining a complete asteroseismology analysis.  Our study does not at this point account for the presence of a magnetic field. This obviously means that one should be careful before assuming that our results can be used to interpret the magnetar observations. However, we feel that the magnetic problem is still somewhat beyond reach. The main reason for this is the strong coupling between any motion in the crust and Alfv\\'en waves in the core. As we have recently argued \\citep{gsa:mhd} [see also \\citet{levin:magnetars}], this coupling is likely to lead to global oscillations which involve significant core motion. This problem is computationally orders of magnitude more difficult because the magnetic field plays an active role. Having said that, it may well turn out to be the case that the modes that dominate the observed signal remain quite close to the pure crust modes, despite the presence of the magnetic field. One can argue that this should be the case if the modes are excited by crust ``cracking'' \\citep{duncan:1998A}, where the bulk of the energy is deposited into elastic motion \\citep{gsa:mhd}. This is, of course, a qualitative argument and better calculations are needed to establish whether or not it is quantitatively useful.  Before we proceed, it is worth noting that the approach to neutron star asteroseismology described in this paper is directly applicable to weaker magnetic field pulsars provided that they are not too rapidly rotating. Glitches in these objects may excite torsional oscillations in the crust, even though such oscillations are yet to be detected. ",
        "conclusions": "In this paper, we have considered torsional oscillation modes in neutron star crusts. Using a state-of-art EOS in the crust but otherwise simplifying the problem as far as possible we have shown how global properties of the star may be deduced from observations of such modes. We have done this using both approximate formulae and a large sample of fully numerical mode-results.  As an example we applied our methods to SGRs assuming that the post-flare QPOs are connected with pure crustal oscillations. The results show that the devised strategy for asteroseismology is, in principle, quite powerful. Of course, the basic model needs to be improved before the inferred stellar parameters should be taken too seriously. In particular, we expect systematic effects due to our neglect of the magnetic field, which obviously should be important for magnetars. One can immediately think of several ways to improve upon our calculations. Below we list the most important.  For a magnetic star, one would not expect the QPOs to be \\emph {exactly} at the pure crustal frequencies. The toy model of \\cite{gsa:mhd} suggests that the QPOs are really global oscillations of the entire star, but that the modes that are predominately excited lie close to the pure crust modes.  In addition one would expect the geometry and strength of the magnetic field to have a key role in determining which modes are excited and how strongly. Clearly, this question can only be settled in a realistic spheroidal model. Work on that problem should be strongly encouraged.  The magnetic field will not simply lead to an offset of the mode frequency with respect to that of a pure crust mode. The nature of the modes will also be affected by the presence of the field. The considerations of \\cite{piro:flares} show that (in his plane-parallel geometry) the fundamental modes are only weakly affected by the magnetic field, whereas the overtones are strongly dependent on the field. This is to be expected from our expressions \\refeq{f0est} and \\refeq{f1est} where we see that the fundamental modes are determined by the tangential shear velocity $v_t$ while the overtones instead depend on the radial velocity $v_r$. A magnetic field will boost the elasto-MHD speeds along the field direction and it therefore seems possible that the toroidal part of the field will affect the fundamental modes whereas the poloidal part will influence the overtones. Should this prove to be true for the spherical case it would be very interesting since it would offer a possibility of extracting information concerning the field geometry. That the geometry of the magnetic field configuration is essential to the behaviour of the torsional modes was emphasised by \\cite{mps:qpo} and recently by \\cite{sks:torsional}.  In our model we have assumed zero temperature. A non-zero temperature will lead to a systematic effect since the shear modulus will decrease \\citep{oi:shearmod}. It will also mean that the crust will be slightly thinner as the ocean is deeper. In our calculations we have integrated the equations all the way to the density of iron at zero pressure.  In the deep crust matter is not expected to be microscopically isotropic, but instead form so-called pasta phases. If these anisotropies extend to macroscopic (fluid-dynamic) scales, \\ie\\ if they do not average out over large distances compared to the lattice spacings, the radial and tangential shear speeds will be different, perhaps by a factor of two \\citep{pp:liquid}. Additionally, the crust may not be in a relaxed state, causing the background strain to influence the torsional oscillations. We have not taken this into account in our model.   In the inner crust we expect that at least a portion of the dripped neutrons are in a superfluid state.  It is, in fact, likely that a significant portion of the mass does not directly participate in the oscillations. Of course, the superfluid will nevertheless be set in motion via the effect of entrainment [for discussion and relevant references see \\cite{ac:sfreview}]. It would be very interesting the examine the effects of this mechanism on the mode spectrum using the recently developed formalism of \\cite{cs:superelastic} (which incidentally also allows for a MHD-type magnetic field).  Our models are static, \\ie\\ non-rotating.  Although this is a very good approximation for the slowly rotating magnetars it would be useful to extend our treatment if toroidal oscillations were to be observed in faster spinning pulsars.  Our model is non-dissipative. In reality a number of processes, including the emission of gravitational and electromagnetic radiation, will damp the oscillations. Mutual friction between superfluid and normal components in the inner crust and viscosity will also be important. Observations of the QPOs suggest a damping time of a few tens of seconds in the SGRs. It would be interesting to estimate the characteristic damping times for each damping mechanism to see if they are consistent with the observations. In addition, if the emission mechanism of the observed X-ray signal was known, a lower limit of the amplitude of the motion of the crust could be derived. This might lead to an estimate of the maximum strain in the crust and hence an assessment of the viability of elastic motion as the origin of the QPOs. Clearly, if the maximum strain exceeds the breaking strain, the oscillations cannot be of elastic origin."
    },
    "0609/astro-ph0609053_arXiv.txt": {
        "abstract": "We present the results of a radial velocity (RV) survey of 14 brown dwarfs (BDs) and very low-mass (VLM) stars in the Upper Scorpius OB association (UScoOB) and 3 BD candidates in the $\\rho$~Ophiuchi dark cloud core. We obtained high-resolution echelle spectra at the Very Large Telescope using Ultraviolet and Visual Echelle Spectrograph (UVES) at two different epochs for each object, and measured the shifts in their RVs to identify candidates for binary/multiple systems in the sample. The average time separation of the RV measurements is 21.6~d, and our survey is sensitive to the binaries with separation $<0.1$~au.  We found that  4 out of 17 objects (or $24^{+16}_{-13}$~per~cent by fraction) show a significant RV change in 4--33~d time scale, and are considered as binary/multiple `candidates.'  We found no double-lined spectroscopic binaries in our sample, based on the shape of cross-correlation curves. The RV dispersion of the objects in UScoOB is found to be very similar to that of the BD and VLM stars in Chamaeleon~I (Cha~I). We also found the distribution of the mean rotational velocities ($v\\,\\sin i$) of the UScoOB objects is similar to that of the Cha~I, but the dispersion of $v\\,\\sin i$ is much larger than that of the Cha~I objects. ",
        "introduction": "\\label{sec:Introduction}  Most stars are member of binary systems and it is therefore important that a complete star formation theory be able to predict the binary fraction, period distribution, and mass-ratio distribution of  stellar objects across a wide range of masses. Furthermore, the study of individual binary systems is the only direct means to determine fundamental stellar properties such as stellar masses and radii.  Recent high-resolution imaging studies of brown dwarfs (BDs) and very low-mass (VLM) stars have placed strong constraints on binaries with separations of $\\sim1-100\\,\\mathrm{au}$. For example, Hubble Space Telescope (\\emph{HST}) observations of $\\alpha$~Per and the Pleiades indicates a binary fraction ($f$) of $>10$~per~cent with a bias towards separations ($a$) of less than $15$~au, and a mass-ratio ($q$) of $>0.7$ \\citep{martin:2003} for objects around and below the hydrogen burning limit (see also \\citealt{bouy:2006}). A similar lack of wide binaries was found in the field T-dwarf study \\citep{burgasser:2003}, while $f\\approx15$~per~cent (for the field objects with 0.03--0.1~$\\mathrm{M_{\\sun}}$ and M8.0--L0.5) was  determined by \\citet{close:2003} using the adaptive optics at Gemini North. They also found the vast majority of systems have a semimajor axis $<20\\,\\mathrm{au}$ (see also \\citealt{siegler:2005} but note too the results of \\citealt{luhman:2004}). An \\emph{HST} study of more than 80 field late M and L dwarfs \\citep{gizis:2003} indicated $f\\approx15$~per~cent with separations in the range of 1.6--16~au. For a small (12) sample of BDs and VLM stars (0.04--0.1~$\\mathrm{M_{\\odot}}$) in Upper Scorpius OB association (UScoOB), \\citet*{kraus:2005} found $f=25_{-8}^{+16}$~per~cent for $5\\,\\mathrm{au}<a<18\\,\\mathrm{au}$ by using a similar imaging technique. More recently, \\citet{basri:2006}, combined with the results of earlier works, found the upper limit of the overall binary fraction for VLM stars of $26\\pm10$~per~cent.  Using a Monte Carlo simulation, the data from radial velocity surveys available in the literature, and by carefully considering the sensitivity and sampling biases, \\citet{maxted:2005}  found an overall BD/VLM binary frequency of 32--45 per~cent assuming $f=15$~per~cent for $a>2.6$~au. A recent photometric study \\citep{pinfield:2003} of low-mass objects in Pleiades and Praesepe suggested, albeit indirectly, $f$ as large as 50~per~cent. which would only be compatible with direct imaging studies if 70--80~per~cent of those binaries have $a<1\\,\\mathrm{au}$. For a more comprehensive review of the current status of BD/VLM binary fraction and the separation distribution, readers are refer to a recent review of multiplicity studies by \\cite{burgasser:2006}.  The extensive imaging surveys provide excellent observational constraints on wider BD+BD binaries, but it is now necessary to search for shorter period BD+BD binaries systematically. Binaries with the separation of less than 1~au are not resolved by current imaging techniques, but will be detectable as spectroscopic binaries, providing the mass-ratio is not too extreme, and velocity separation is large enough. The first BD+BD spectroscopic binary, PPl~15  \\citep{basri:1999}, showed a double-peaked cross-correlation function with a maximum velocity separation of $>70\\,\\kmps$. The binary was found to have an eccentric orbit ($e=0.4$) with a period of $\\sim5.8\\,\\mathrm{d}$. \\citet{basri:1999} suggested that the formation process of substellar objects is biased towards smaller separation binaries based on the short period of PPl~15 and the lack of Pleiades BD binaries with separations $>40\\,\\mathrm{au}$. Note that the median separation of binaries with solar-type primaries is 30~au \\citep{duquennoy:1991}. Pioneering work on the RVs of BDs and VLMs are presented by \\citet{guenther:2003}, \\cite{kenyon:2005} and \\citet{joergens:2006} who found a several binary candidates; however, the orbital parameters and masses of binaries remains unknown because the follow-up spectroscopic monitoring is lacking or is still being undertaken. In addition to the follow-up observations, the number of BDs and VLM star binary candidates needs to be increased in order to have better statistics on short-period binary parameters.  The first BD+BD eclipsing binary (2MASS J0532184--J0546085) was discovered by \\citet*{stassun:2006} from the $I$-band photometric monitoring of the system. Combining their light curves and the results from the follow-up radial velocity measurements, they were able to determined the precise orbital and physical parameters of the system.  The projected semimajor axis and the period of the binary are found as $0.0398\\pm0.0010$~au and $9.779621\\pm0.000042$~d respectively.  The separation distribution of BD/VLM binaries is critical to understanding their origin. There are two main models for the formation of BDs and VLM stars: first, they have low masses because they form in low-mass, dense molecular cloud cores (e.g.~\\citealt{padoan:2002}); second, BD/VLM objects have low masses because they are ejected from the dense core in which they form via dynamical interactions in multiple system, cutting off their accretion before they have reached stellar masses (\\citealt{reipurth:2001}; \\citealt*{bate:2002}). Alternatively, there is a third model in which a free-floating BD or planetary-mass object can be formed in the process of the photo-evaporation (e.g.~\\citealt{mccaughrean:2002};~\\citealt{whitworth:2004}) with the outer layers of a pre-stellar core ($\\sim0.2\\,\\mathrm{M_{\\odot}}$) removed by the strong radiation pressure from the nearby massive OB stars before the accretion onto the protostar at core centre occurs.  Due to the dynamical interaction involved in the second model, BD/VLM binaries that survive are generally expected to have small separations. In the first model, wider binaries may be expected to be more common. \\citet*{bate:2002b} suggested that close binaries ($a<10\\,\\mathrm{au}$) do not form directly, but result from hardening of wider systems though a combination of dynamical interactions, accretion and interactions with circum-binary discs. If BD/VLM binaries have formed through such mechanisms, one would not expect to find binaries with 1--10~au separations without also finding many with separation $<1$~au. If an absence/rarity of binaries with 1~au were found, it may support the idea that they are ejected quickly from multiple systems before they have undergone the interactions that shorten their periods.  Our immediate aim is to identify spectroscopic and close BD/VLM binaries using the high-resolution echelle spectroscopy at two epochs. This experiment is sensitive to VLM binaries with separations of $<0.1$~au which corresponds to a period of $\\sim10\\,\\mathrm{d}$. A larger sample of candidates will enable us to measure the binary fraction of these short-period/close binaries (once confirmed), and address whether there is a significant population of `hidden' VLM companions. The long term goal of this project to follow up the binary candidates found in this paper by spectroscopically monitoring them over different time scales, enabling us to obtain the radial velocity curves and their minimum masses.  In Section~\\ref{sec:observation-reduction}, we describe the observations and the data reduction. The results of radial velocity and rotational velocity ($v\\,\\sin i$) measurements are presented in Section~\\ref{sec:Results}. We discuss the binary/multiplicity fraction indicated by our RV survey in Section~\\ref{sec:Binary-fraction}, and give our conclusions in Section~\\ref{sec:Conclusions}.     \\begin{table*}  \\caption{Summary of known properties of the targets from literature: a.~\\citet{luhman:1999} (original list for $\\rho$~Oph), b.~\\citet*{ardila:2000} (original list for UScoOB), c.~\\citet*{wilking:1999}, d.~\\citet{muzerolle:2003}, e.~\\citet{kraus:2005}, and f.~\\citet*{mohanty:2005}. }  \\label{tab:literatures}  \\begin{center}\\begin{tabular}{llcccc} \\hline Object& Sp.& mass& RV& $v\\sin i$& Known multiple?\\tabularnewline & & $\\left[\\mathrm{M_{\\odot}}\\right]$& $\\left[\\mathrm{km\\, s^{-1}}\\right]$& $\\left[\\mathrm{km\\, s^{-1}}\\right]$& \\tabularnewline \\hline GY 5& $\\mathrm{M7^{c}}$& $0.07^{\\mathrm{d}}$& $-6.3\\pm1.9^{\\mathrm{d}}$& $16.8\\pm2.7^{\\mathrm{d}}$& no\\tabularnewline GY 141& $\\mathrm{M8.5^{a}}$& $0.02^{\\mathrm{d}}$& \\ldots{}& $6.0^{f}$& no\\tabularnewline GY 310& $\\mathrm{M8.5^{c}}$& $0.08^{\\mathrm{a,d}}$& \\ldots{}& $10.0^{f}$& no\\tabularnewline USco 40& $\\mathrm{M5^{b}}$& $0.1^{\\mathrm{b}}$& \\ldots{}& $37.5^{f}$& no\\tabularnewline USco 53& $\\mathrm{M5^{b}}$& $0.1^{\\mathrm{b}}$& \\ldots{}& $45.0^{f}$& no\\tabularnewline USco 55& $\\mathrm{M5.5^{b}}$& $0.10+0.07^{\\mathrm{e}}$& \\ldots{}& $12.0^{f}$& $\\mathrm{visual^{e}}$\\tabularnewline USco 66& $\\mathrm{M6^{b}}$& $0.07+0.07^{\\mathrm{e}}$& $-4.4\\pm0.6^{\\mathrm{d}}$& $27.5^{f}$& $\\mathrm{visual^{e}}$\\tabularnewline USco 67& $\\mathrm{M5.5^{b}}$& $0.10^{\\mathrm{e}}$& \\ldots{}& $18.0^{f}$& no\\tabularnewline USco 75& $\\mathrm{M6^{b}}$& $0.07^{\\mathrm{e}}$& $-5.6\\pm1.1^{\\mathrm{d}}$& $63.0^{f}$& no\\tabularnewline USco 100& $\\mathrm{M7^{b}}$& $0.05^{\\mathrm{e}}$& $-8.9\\pm0.6^{\\mathrm{d}}$& $50.0^{f}$& no\\tabularnewline USco 101& $\\mathrm{M5^{b}}$& $0.05^{\\mathrm{b}}$& \\ldots{}& \\ldots{}& no\\tabularnewline USco 104& $\\mathrm{M5^{b}}$& $0.05^{\\mathrm{b}}$& \\ldots{}& $16.0^{f}$& no\\tabularnewline USco 109& $\\mathrm{M6^{b}}$& $0.07+0.04^{\\mathrm{e}}$& $-3.8\\pm0.7^{\\mathrm{d}}$& $6.0^{f}$& $\\mathrm{visual^{e}}$\\tabularnewline USco 112& $\\mathrm{M5.5^{b}}$& $0.1^{\\mathrm{e}}$& \\ldots{}& $8.0^{f}$& no\\tabularnewline USco 121& $\\mathrm{M6^{b}}$& $0.02^{\\mathrm{b}}$& $-38.9\\pm1.0^{\\mathrm{d}}$& \\ldots{}& no\\tabularnewline USco 128& $\\mathrm{M7^{b}}$& $0.05^{\\mathrm{e}}$& $-3.0\\pm1.6^{\\mathrm{d}}$& $0.0^{f}$& no\\tabularnewline USco 130& $\\mathrm{M7.5^{e}}$& $0.04^{\\mathrm{e}}$& \\ldots{}& $14.0^{f}$& no\\tabularnewline USco 132& $\\mathrm{M7^{b}}$& $0.05^{\\mathrm{e}}$& $-8.2\\pm1.1^{\\mathrm{d}}$& \\ldots{}& no\\tabularnewline \\hline \\end{tabular}\\par\\end{center}  \\end{table*} ",
        "conclusions": "\\label{sec:Conclusions}  We have presented two-epoch RV survey of 18 young BDs and VLM stars ($0.02~M_{\\sun} < M_{*} < 0.1~M_{\\sun}$) in UScoOB and $\\rho$~Oph dark cloud core using the high resolution UVES echelle spectroscopy at VLT.  The average time separation of RV measurements are 21.6~d, and our RV survey is sensitive to binaries with separation smaller than 0.1~au.  One of our targets, USco~121, is most likely a non-member of the UScoOB association based on the deviation of the RV from the result of the population in the association. A similar conclusion was found by \\citet{muzerolle:2003} from their RV study and the low lithium abundance.  We found 4 (USco~112, USco~128, USco~40 and USco~55) out of 17 objects as our binary/multiple candidates. This corresponds to the binary fraction of $24^{+16}_{-13}$~per~cent for the binary separation $a<0.1$~au. The recent high-resolution imaging survey of brown dwarfs and very low-mass objects (M5.5--M7.5 ) in the UScoOB by \\citet{kraus:2005} (which was not known to authors at the time of our observation: April--May, 2004) confirms that USco~55 and USco~66 are multiple systems, and USco~109 is most likely a multiple system.  Their projected separations are $a>4.0$~au which is well above our sensitivity limit of 0.1~au; therefore, it is not surprising that we did not identify USco~66 and USco~109 as multiples. Interestingly, we also found USco~55 as a candidate for a multiple system. This may indicate that it is a triple system with one of the objects located within 0.1~au of the primary.  In addition, we identify USco~112 and 128 as binary candidates, but they did not find them as multiples. Further they found USco~67, 75, 130 and 132 as non-multiple, as did we.  We found the RV dispersion ($1.0\\,\\kmps$) of the objects in UScoOB is very similar to that of the BDs and VLM stars in Chamaeleon~I (Cha~I) previous study by \\citet{joergens:2006b}. The rotational velocities ($v\\,\\sin i$) of the samples were also measured. The distribution of $v\\,\\sin i$ for the UScoOB objects peaks around $16.9\\,\\kmps$ which is also similar to that of the Cha~I population found by \\citet{joergens:2006b}; however, the dispersion of $v\\,\\sin i$ for the UScoOB objects ($27.8\\,\\kmps$) is found to be much larger than that of the Cha~I objects ($8.0\\,\\kmps$)  Follow-up spectroscopic observations of the binary candidates presented here are planned in near future. There are only a few RV variable binary candidates identified in earlier surveys (\\citealt{guenther:2003}; \\citealt{kenyon:2005}; \\citealt{joergens:2006}). As \\citet{burgasser:2006} points out most of the current RV and imaging surveys use samples from magnitude-limited survey, but one should attempt to use the samples from volume-limited survey in order to find a correct statistics on binary parameters more straightforwardly, i.e. without correcting for bias. \\textcolor{black}{ It is possible that our binary fraction is an overestimate since we are biased to brighter objects. However the lack of a double-line system suggests that a primary-to-secondary luminosity ratio of $\\gg 1$ for the RV variables found in this paper. }"
    },
    "0609/astro-ph0609579_arXiv.txt": {
        "abstract": "We have measured the contribution of submillimeter and mid-infrared sources to the extragalactic background radiation at 70 and 160$\\mu$m. Specifically, we have stacked flux in 70 and 160$\\mu$m Spitzer Space Telescope ({\\sl Spitzer}) observations of the Canada-UK Deep Sub-millimeter Survey 14h field at the positions of 850$\\mu$m sources detected by SCUBA and also 8 and 24$\\mu$m sources detected by {\\sl Spitzer}.  We find that per source, the SCUBA galaxies are the strongest and the 8$\\mu$m sources the weakest contributors to the background flux at both 70 and 160$\\mu$m. Our estimate of the contribution of the SCUBA sources is higher than previous estimates. However, expressed as a total contribution, the full 8$\\mu$m source catalogue accounts for twice the total 24$\\mu$m source contribution and $\\sim 10$ times the total SCUBA source contribution.  The 8$\\mu$m sources account for the majority of the background radiation at 160$\\mu$m with a flux of 0.87$\\pm$0.16 MJy/sr and at least a third at 70$\\mu$m with a flux of 0.103$\\pm$0.019 MJy/sr. These measurements are consistent with current lower limits on the background at 70 and 160$\\mu$m.  Finally, we have investigated the 70 and 160$\\mu$m emission from the 8 and 24$\\mu$m sources as a function of redshift. We find that the average 70$\\mu$m flux per 24$\\mu$m source and the average 160$\\mu$m flux per 8 and 24$\\mu$m source is constant over all redshifts, up to $z\\sim 4$. In contrast, the low-redshift half $(z<1)$ of the of 8$\\mu$m sample contributes approximately four times the total 70$\\mu$m flux of the high-redshift half. These trends can be explained by a single non-evolving SED. ",
        "introduction": "Excluding the microwave background, approximately half of the entire extragalactic background radiation is emitted by dust at far infra-red (IR) and sub-millimeter (submm) wavelengths \\citep[e.g.,][]{fixsen98,hauser01,dole06}.  This cosmic IR background (CIB) radiation peaks at a wavelength of $\\sim 200\\mu$m, yet compared to the optical, relatively little is known about the sources responsible.  Surveys conducted by the Sub-millimeter Common User Bolometer Array (SCUBA) and the Max-Planck Millimeter Bolometer (MAMBO) over the last decade have directly resolved up to two-thirds of the CIB at $850\\mu$m and 1.1mm into discrete, high redshift sources (although this fraction is uncertain due to the uncertainty in measurements of the CIB at these wavelengths). Discovery of this population has been extremely important since SCUBA galaxies represent the most energetic star-forming systems at an epoch when the Universe was at its most active. However, the impact this has had on understanding the nature of the CIB is relatively minor since at these wavelengths, the CIB has 30 to 40 times less power than at the peak.  A recent study using a large sample of 73 bright ($\\simgreat 5$mJy) SCUBA sources by \\citet{chapman05} indicated that the population contributes a mere $\\sim 2\\%$ of the CIB at the peak, with an extrapolation of up to $\\sim 6\\%$ for sources down to the fainter limit of 1mJy.  However, this work relied on an assumed spectral energy distribution (SED) for the SCUBA sources constrained only by a redshift, the 850$\\mu$m SCUBA flux and a radio flux at 1.4GHz. Furthermore, redshifts were obtained from optical spectra having identified the optical sources with radio counterparts to the SCUBA sources.  This introduces two selection effects. The first causes SCUBA sources with $z \\simgreat 3$ to be missed by requiring a radio detection, the selection function for radio sources falling off rapidly at $z \\sim 3$ due to the K-correction.  The second causes a paucity of sources around $z \\sim 1.5$ where no emission lines fall within the observable wavelength range of their spectra.  This motivates the first of two main goals of this paper. By stacking the flux in 70 and 160$\\mu$m MIPS images at the positions of SCUBA sources detected in the Canada-United Kingdom Deep Sub-millimeter Survey (CUDSS) 14-hour field \\citep{eales00,webb03}, we directly measure the contribution of the SCUBA sources to the CIB at wavelengths in the vicinity of the peak.  In addition to the submm surveys, space-borne mid-IR surveys conducted by the Infra-red Astronomical Satellite (IRAS) and the Infra-red Space Observatory (ISO) have resolved significant contributions to the CIB from the shorter wavelength side of the peak \\citep[see, for example, the review by][]{lagache05}. The introduction of the Spitzer Space Telescope \\citep[{\\sl Spitzer};][]{werner04} means that such surveys can be carried out over much wider areas and to much greater depths.  In particular, the Multi-band Photometer for {\\sl Spitzer} \\citep[MIPS;][]{rieke04} has been used for a variety of mid- and far-IR surveys to resolve sources contributing to the CIB. \\citet{papovich04} showed that approximately 70\\% of the CIB at 24$\\mu$m can be resolved into IR galaxies with flux $\\ge 60\\mu$Jy. In contrast, \\citet{dole04} found that at the longest two MIPS wavelengths, 70 and 160$\\mu$m, only 20\\% and 10\\% of the CIB can be directly resolved into distinguishable sources brighter than 3.2 and 40mJy respectively. However, the error on these fractional quantities is very large since the absolute flux of the background at 160$\\mu$m is currently unknown to a factor of $\\sim 2$ and at 70$\\mu$m, the uncertainty is even larger.  A major problem with attempting to directly resolve sources in deep 70 and 160$\\mu$m MIPS surveys is source confusion due to the large instrument point spread function (PSF). This problem can be circumvented by measuring the 70 and 160$\\mu$m MIPS flux at the position of objects selected in other wavebands for which there are already accurate positions. This stacking technique has been successfully used by several authors with SCUBA data that also suffer from confusion \\citep[e.g.,][]{peacock00,serjeant04,knudsen05,dye06,wang06}. \\citet{dole06} stack MIPS flux at the positions of 24$\\mu$m sources with fluxes $> 60\\mu$Jy to find that they represent the bulk of the CIB at 70 and 160$\\mu$m respectively (see Section \\ref{sec_full_stack_results}). These contributions are investigated as a function of 24$\\mu$m source flux and, based on external studies of the redshift distribution of MIPS 24$\\mu$m sources, the authors conclude that the majority of the radiation must be emitted at $z \\sim 1$.  This provides the second main motivation for the present paper.  We extend the analysis of \\citet{dole06} in two ways. Firstly, we additionally measure the contribution to the CIB at 70 and 160$\\mu$m from 8$\\mu$m sources observed with {\\sl Spitzer}'s Infra-red Array Camera \\citep[IRAC;][]{fazio04}. Secondly, we investigate how the contribution from the 8 and 24$\\mu$m populations varies with redshift, using photometric redshifts established for these sources in our earlier work \\citep{dye06}.  This paper is set out as follows. In the following section we describe the data. Section \\ref{sec_analysis} outlines our stacking procedure. Our results are presented in Section \\ref{sec_results}, followed by a summary and brief discussion in Section \\ref{sec_summary}. ",
        "conclusions": "\\label{sec_summary}  In this paper, we have quantified the contribution of flux to the CIB at 70 and 160$\\mu$m from SCUBA sources and 24 and 8$\\mu$m {\\sl Spitzer} sources in the CUDSS 14hour field. By stacking flux at the position of the different sources, we have found that the SCUBA sources make the highest contribution per source and that the 8$\\mu$m sources make the lowest. Conversely, the opposite is true of the total contribution from all sources, leading to the conclusion that the bulk of the CIB is most efficiently resolved by sources detected at wavelengths shorter than the peak of the CIB emission at $\\sim 200\\mu$m.  Our stacking suggests a somewhat larger contribution from the CUDSS SCUBA sources to the 200$\\mu$m CIB peak than previously thought. \\citet{chapman05} estimated an upper limit on the contribution of $> 1$mJy SCUBA sources detected at 850$\\mu$m to the CIB emission at 160$\\mu$m of $\\simless 0.04$ MJy/sr.  Despite our SCUBA sources having a brighter sensitivity level of 3.5mJy, we measure a contribution at 160$\\mu$m of 0.125$\\pm$0.040 MJy/sr, without removing any bright MIPS sources, or 0.060$\\pm$0.042 MJy/sr having removed all $\\geq 3 \\sigma$ sources.  Since measurements of the CIB at 70 and 160$\\mu$m are presently very uncertain, the fractional contribution to the CIB made by our sources can only be expressed within a range set by present upper and lower limits. Using the DIRBE estimates as upper limits and the lower limits set by \\citet{dole06}, our strongest contributors, the 8$\\mu$m sources, resolve somewhere between $\\sim 35-75$\\% of the background at 70$\\mu$m and $\\sim 90-100$\\% at 160$\\mu$m.  By combining our results in the present work with our previous stacking of 450 and 850$\\mu$m SCUBA flux \\citep{dye06}, we have established that the 8 and 24$\\mu$m sources detected by {\\sl Spitzer} are on average borderline ULIRGs. The average source we detect is $\\sim 10$ times fainter than the average SCUBA source detected by \\citet{pope06} and \\citet{chapman05} integrating flux over the wavelength range 8 - 1000$\\mu$m. Furthermore, the temperature of $\\sim 40$K of our average source is consistent with temperatures of submm galaxies found in the local Universe \\citep[e.g.,][]{dunne01} and with the median temperature of $(36\\pm7)$K of SCUBA sources in the high redshift Universe found by \\citet{chapman05}. However, our average source is warmer than the median temperature of T$_{\\rm med} \\simeq 30$K of 10 SCUBA sources measured by \\citet{pope06}.  Using photometric redshifts assigned to the 8 and 24$\\mu$m sources, we have investigated how the contribution of 70 and 160$\\mu$m flux to the CIB varies with redshift. We have found that the 8$\\mu$m sources at low redshifts, $z < 1$ (accounting for half of them), are the strongest contributors to the CIB at 70$\\mu$m, their flux amounting to $\\sim 4$ times that of the $z>1$ sources.  The 70$\\mu$m emission per 24$\\mu$m source as well as the 160$\\mu$m emission per 8 and 24$\\mu$m source is consistent with an even distribution over redshift. This verifies the result of \\citet{dole06} that the majority of the emission at 70 and 160$\\mu$m from 24$\\mu$m sources must come from a redshift of $z \\sim 1$ where the redshift distribution of these sources peaks. We have shown how this distribution can be reproduced from our observed 8 and 24$\\mu$m catalogue of fluxes and redshifts using a single non-evolving source SED with a dust temperature of 40K.  As a concluding remark, this study and similar recent studies \\citep[e.g.,][]{serjeant04,knudsen05,dye06,wang06} indicate that the CIB is not predominantly due to a rare population of exceptionally luminous submillimeter sources as hinted at by early SCUBA observations, but that a much more numerous galaxy population of modest average luminosity is responsible instead. However, we have shown here that the SCUBA galaxies probably do make a larger contribution than previously thought, although much larger numbers of SCUBA sources such as those of the SCUBA half degree extragalactic survey \\citep[SHADES; ][]{mortier05} or those resulting from future SCUBA2 surveys\\footnote{see http://www.roe.ac.uk/ukatc/projects/scubatwo} will be required to improve the precision of these measurements.  \\appendix"
    },
    "0609/astro-ph0609635.txt": {
        "abstract": "Medium resolution spectra (100~\\kms) are used to deduce the size and kinematics of the coronal lines in a sample of nearby Seyfert galaxies (Circinus, NGC 1386, NGC 1068, NGC 3227, NGC 3783 and MCG-6-30-15) by means of simultaneous observations of [\\ion{Fe}{11}]~7889 \\AA, [\\ion{Fe}{10}]~6374 \\AA, [\\ion{Fe}{7}]~6087 \\AA, [\\ion{Si}{6}]~2.48~$\\mu$m and [\\ion{Si}{7}]~1.96~$\\mu$m.  The coronal lines extend from the unresolved nucleus up to distances between a few tens to a few hundreds of parsecs.  The region of the highest ionized ions studied, [\\ion{Fe}{11}] and [\\ion{Fe}{10}], is the least spatially extended, and concentrates at the center; intermediate ionization lines extend from the nucleus up to a few tens to a few hundred parsecs; lower [\\ion{O}{3}]-like ions are known to extend in the kpc range. All together indicates a stratification in the ionized gas, usually interpreted in terms of nuclear photoionization as the driving ionization mechanism.  However, coronal line profiles show various peculiarities: they are broader by a factor of two than lower ionization lines, the broadening being in terms of asymmetric blue wings, and their centroid position at the nucleus is blueshifted by a few hundreds of \\kms. Moreover, in two objects --NGC~1386 and NGC~1068, a double peak [\\ion{Fe}{7}]~6087~\\AA\\ line is detected in the nuclear and extended coronal region, this being the first report in the literature of such type of profile in coronal lines in active galactic nuclei. If interpreted as outflow signatures, the total broadening of the lines at zero intensity levels implies gas velocities up to 2000~\\kms. Although the stratification of ions across the coronal region means that photoionization is the main power mechanism, the high velocities deduced from the line profiles, the relatively large spatial extension of the emission, and the results from single cloud photoionization models indicate  that an additional mechanism is at work. We suggest that shocks generated by the outflow could provide the additional required power for line formation. ",
        "introduction": "Coronal lines (CL) are collisionally excited forbidden transitions within low-lying levels of highly ionized species (ionization potential $>$ 100 eV). They can be produced by either a hard UV continuum \\citep{fer97}, a hot collisionally ionized plasma \\citep{vac89} or most plausible, a mix of both processes \\citep{con98}. Because of their high ionization potential (IP), these lines are particularly suitable for getting information on the otherwise difficult to access observationally, extreme UV - soft X-ray region of the ionizing spectrum \\citep{pr00}.  Due to the high energies involved in their production, coronal lines are genuine tracers of the AGN power mechanism.  Pure starbursts when boosted by massive O-B stars in their Wolf-Rayet phase, may  show \\ion{He}{2} lines but not lines of higher ionization potential ions ($\\chi >54$~eV).  Observationally, coronal lines are present with comparable strength in AGN spectra regardless of their class \\citep{pe84,er97,pr00,rod02,reu03}. The peak position is usually blueshifted relative to the systemic velocity of the host galaxy ($\\Delta$V $\\sim$ 500-800~\\kms) and broader than low ionization lines \\citep{pe84,dro86,er97,rod02}. That has led to the idea that CL are associated with outflows, formed in a intermediate region between the classical narrow-line- and broad-line-region (NLR and BLR, respectively).  The broad line region has sizes of less than a few hundred light-days \\citep{pet04} and thus remains unresolved; the coronal region because of its lower velocity dispersion should be placed further out but still relatively close to the BLR, considering the high ionization level of the gas. Accordingly, the coronal region has also remained spatially unresolved.  The physical nature of the CLR has called the attention of many authors starting by the pioneer work by \\citet{pe84}, who discover the systematic blueshifting and broadening of this high ionization gas.  \\citet{er97} found that the CL occurs predominantly in objects with a soft X-ray excess and suggest a possible relationship between the CLR and the absorption edges associated with warm absorbers present in nearly 50\\% of AGN.  \\citet{por99} showed that the CL may, in fact, be formed by the warm absorber.  If that is the case, the CLR should be dense ($n\\sim$10$^{6}$ cm$^{-3}$) and located at comparable distance as the broad line region.  On the other hand, \\citet{fer97} photoionization calculations show that the CLR can form in gas with densities of $ 10^2 - 10^{8.5} $ cm$^{-3}$ and extends up to several hundred parsecs, thus coinciding with the size of the extended NLR. However, they claim that except for the lower ionization coronal lines ([Ne\\,{\\sc v}], [Si\\,{\\sc vii}], [Ca\\,{\\sc viii}] and [Fe\\,{\\sc vii}]) that can form efficiently in gas that is roughly $\\sim$10~pc and beyond, the rest of coronal lines form optimally in gas that is less than 10~pc from the ionization source.  Coronal lines are present all across the electromagnetic spectrum. The highest ionization species are located in the X-rays region. In the optical, the most prominent ones are from iron. Dust extinction and the fact that reliable detection of these lines requires medium to high spectral resolution and large S/N ratios has made measurements of these lines difficult.  The near- and mid-infrared regions, less affected by extinction, are very rich in coronal lines from different species and ionization levels.  Indeed, they are often the dominant lines in the near-IR spectra of Seyfert galaxies \\citep{reu03}.  The current availability of adaptive optics in 8 m telescopes enabled \\citet{pri05} to determine the size and morphology of the coronal line region for the first time. That is found to extend from the nucleus up to 30 to 200~pc radii at most, and be aligned preferentially with the direction of the lower-ionization cones seen in Seyfert galaxies. As expected, the coronal line region (CLR) is primarily produced very close to the active nucleus but is nevertheless resolved over several tens of parsecs.  This, together with its preferential alignment with the ionization cone prompts to the possibility that the nuclear radiation is intrinsically collimated or that some additional in-situ excitation mechanism is needed to explain CL gas at those observed distances from the ionizing source. This paper attempts to address those issues by studying the physical conditions of the coronal line gas and its kinematics.  To that aim, we present  high-sensitive observations with spectral resolution R$\\sim$100 km~s$^{-1}$ of Fe and Si coronal lines located in the visible and near-infrared spectra respectively. As stated above, the goal is determining the kinematics of the emitting region and constraining the physical properties of the coronal gas. To that aim, a sample of six nearby AGN were selected, and spectra centered at the position of a selected group of bright coronal lines are studied. In the optical, we selected [Fe\\,{\\sc vii}]~6087~\\AA\\ (IP=100~eV), [Fe\\,{\\sc x}]~6374~\\AA\\ (IP=240~eV) and [Fe\\,{\\sc xi}]~7889~\\AA\\ (IP=260e~V). In the near-IR, [Si\\,{\\sc vi}]~1.963~$\\mu$m (IP=170~eV) and [Si\\,{\\sc vii}]~2.48~$\\mu$m (IP=205~eV) were chosen. These lines are all located near in wavelength to low-ionization lines, easing the comparison regarding the size of the emitting region and the shape of  the line profiles between different ionization levels. In \\S~\\ref{obs}, we describe the sample, the observations and data reduction process.  The determination of the coronal region size based on the available data is carried out in \\S~\\ref{size}. In \\S~\\ref{comp} we discuss and compare the size of the CLR with that determined for the NLR using our data and the literature.  A comparison of the observed emission line profiles to unravel the kinematics of the coronal gas is in \\S~\\ref{kin}. Photoionization and shock models to explain the physical conditions of the coronal gas are described in \\S~\\ref{mod}. An overview of our main results and conclusions are in \\S~\\ref{con}. ",
        "conclusions": "\\label{con}   \\begin{itemize}  \\item Coronal lines are expected to be presented in all type of AGN. This work surveyed at unprecedented spectral and spatial resolution, five coronal lines covering a large range both in ionization potential, between 100 and 260 eV, and wavelength (optical and NIR regions) in six Seyfert galaxies: three Type~2, two Type~1 and one Type~1.5. We found that they are all present in the galaxies with the exception of NGC~3727, where only ions from the lowest ionization potentials are detected.  Thus, the shape, or hardness, of the ionizing spectrum in Seyfert galaxies is very similar  regardless of the type. This is in the same line as found by \\citep{pr00} in a different sample of Seyfert galaxies, after the analysis of IR-ISO coronal lines spanning a range in ionization potential between 54 and 300~eV. NGC~3727 is the only case showing very soft ionizing spectrum as only lines from the lowest IP are present.   \\item The coronal region is spatially resolved over scales that range from a few tens of parsecs up to a few hundreds of parsecs from the center. The size of the emitting region varies with  the ionization potential: the highest the ionization potential the most compact the  region becomes. This stratification indicates that nuclear photoionization is the principal excitation mechanism.  \\item Coronal line profiles are characterized by two components: a narrow one whose centroid and FWHM values are within the range found for lower ionization lines, and a broader one, whose centroid is systematically shifted to the blue by a few hundreds of \\kms\\ and has a FWHM  a factor 2 at least larger than that  of the lower ionization lines.  The gas velocities implied by this component vary from object to object, and are in the range  from 500 up to 2000~\\kms. This blueshifting is interpreted as an outflowing wind. However, a nuclear wind should have an approaching and a receding component. As the blueshifted component is the one often seen, the redshifted one has to be obscured by dust at the base of the wind.  \\item We are however able to see the receding component of the wind in two objects: NGC~1068 and NGC~1386. This interpretation comes from the detection of a spatially resolved double peak in the [\\ion{Fe}{7} line, each peak tracing respectively blueshifted and redshifted gas at each spatial location. A spatially resolved double peak could be produced by gas moving in circular orbits in a thick toroid with the rotation axis close to the north-south direction, or by radial motions along a collimated cone, with the axis closely oriented to the plane of the sky. The observational evidence supports this latter scenario.  \\item For a luminosity of the ionizing radiation source in the range 10$^{43.5}$ - 10$^{44.5}$ erg s$^{-1}$, single-cloud photoionization models indicate that  to produce detectable coronal emission over a wide range of ionization potentials,   the ionization parameter must be larger than 0.1 and the density should be larger than 10$^3$ cm$^{-3}$. Under these conditions, such simple single-cloud models can account for coronal emission of up to 100~pc from the nucleus if the ionizing radiation luminosity is higher than 10$^{44}$ erg~s$^{-1}$. For lower luminosities the size of the coronal regions can at most reach 50~pc, which is  smaller than the results obtained for some of the galaxies of our sample. In addition,  these models fail to explain the observed Fe line ratios. Furthermore, even  multi-cloud photoionization models, mimicking the stratification of the NLR, can hardly reproduce the  observations, leading us to conclude that another energy source must  be present.   \\item In support to the above result, diagnostic diagrams using the Fe lines show that single-cloud photoionization models largely depart from the trend shown by the data. However, if the contribution of additional shock excitation is included, with shock velocities between 300 and 900~\\kms\\ in agreement with observations, the {\\it combined} photoionization plus shock effect is able to account for the observed ratios.   \\item It is interesting to remark that the systematic presence of H$_2$ lines in the nuclear region of AGN implies very large variations in the physical conditions of the gas, capable of sustaining highly ionized species along with molecules which easily dissociate with the radiation of the central engine.  The molecular H$_{2}$ emission clearly relates to a different type of clouds from that traced by the ionized gas, on the basis of both its kinematic - low rotation, narrow line widths - and spatial distribution which usually extends over several hundreds parsecs to kpc from the nucleus, and preferentially in directions perpendicular to the ionization cone \\citep{reu02,reu03,rod04,rod05}.  \\end{itemize}"
    },
    "0609/astro-ph0609522_arXiv.txt": {
        "abstract": "The Galactic secondary cosmic ray anti-proton (\\pbar) flux calculated with different propagation models is fairly consistent with data, and the associated propagation uncertainty is small. This is not the case for any \\pbar\\ exotic component of the dark matter halo (see also the companion paper;~\\citealt{Maurin:2006ps}). Detailed propagation models are mandatory if the ultimate goal is to explain an excess. However, simpler and faster approximate formulae for \\pbar\\ are an attractive alternative to quickly check that a given dark matter model is not inconsistent with the \\pbar\\ observed flux. This paper provides such formulae. In addition, they could be used to put constraints on new physics in this channel, where an extensive scan of a large parameter space could otherwise be quite expensive in computer ressources. ",
        "introduction": "The first positive detection of anti-protons at the end of the seventies \\citep{1979PhRvL..43.1196G, 1981ApJ...248.1179B,1982ICRC....9..146B} boosted the interest for indirect searches of new physics in the low energy \\pbar\\ spectrum. After 25 years of efforts and improvements both in measurements and theoretical calculations, the data at low energy (up to a few GeV) are now well accounted for \\citep{2001ApJ...563..172D}. However, at higher energy, calculations tend to predict less \\pbar\\ than observed \\citep{2001ApJ...561..787B} but more data are desirable to confirm this possible trend.  The uncertainty of the \\pbar\\ standard secondary flux \\citep{2001ApJ...563..172D} mainly comes from nuclear physics, as that of astrophysical origin is small ---less than a few percents---and is expected to decrease even more as new B/C measurements are available. In contrast, the \\pbar\\ exotic flux---from exotic sources following the dark matter distribution---is plagued by the degeneracy in the propagation parameters, the corresponding uncertainty being as large as two orders of magnitude at low energy \\citep{2004PhRvD..69f3501D}. Until propagation is better understood and constrained, it may not be worth using refined---and sometimes time consuming---models whose result will anyway be crippled by these large uncertainties. This is especially true when one has to scan as efficiently as possible the huge parameter space existing, for example, in {\\sc susy} theories.  The goal of this paper is to provide approximate formulae for the propagated \\pbar\\ exotic fluxes. This is done in the framework of the diffusion/convection model with constant wind already discussed in \\cite{2001ApJ...555..585M,2001ApJ...563..172D,2002A&A...388..676B,2004PhRvD..69f3501D}. These formulae are reasonably accurate, fast to compute and easy to implement. They have to be thought as an easy-to-use tool, for phenomenologists interested in beyond-the-standard-model theories and wishing to quickly check that any new physics model on the market does not violate the \\pbar\\ constraint. Once the interesting regions of the parameter space (for the dark matter candidate) are identified, more elaborate treatments should be implemented (e.g. {\\sc Galprop}---\\citealt{2005AdSpR..35..156M}; {\\sc DarkSusy}---\\citealt{2004JCAP...07..008G}; or \\citealt{2004PhRvD..69f3501D}).  The paper is organized as follows: \\begin{itemize} \\item In Sec.~\\ref{sec:2D}, we remind the salient ingredients of the constant wind/diffusion model for secondary and primary exotic anti-protons, using a Bessel expansion formalism. \\item In Sec.~\\ref{sec:alternative}, two alternative formulations of the \\pbar\\ primary exotic flux are presented. \\item In Sec.~\\ref{sec:results}, all the formulations are compared and their relative merits discussed. \\end{itemize} The formulae given in Sec.~\\ref{sec:alternative} are implemented in numerical routines publicly available\\footnote{http://wwwlapp.in2p3.fr/$\\sim$taillet/mtc/mtc\\_code.tar}. ",
        "conclusions": "\\label{sec:theEnd} Given the present accuracy of propagation parameters, exotic anti-proton fluxes suffer large uncertainties (see also the companion paper). This limits the benefit of a detailed calculation, which can be expensive in term of computational power, when repeated thousands of times.  We showed that discarding energy redistributions (coulomb, ionization, adiabatic losses, reacceleration and tertiary contribution) provides exotic \\pbar\\ fluxes accurate at the level of $\\lesssim 50\\%$ for IS energies greater than $\\sim 1$~GeV (leading to $\\gtrsim 100$ MeV energies once modulated). This opens the possibility to use one of the following simpler formulae: \\begin{itemize} \\item the approximate 1D--model, which gives the correct flux as long as $L\\lesssim 7$~kpc for PBH-like sources, or $L\\lesssim 1$~kpc for SUSY-like sources; \\item the 2D-propagator. \\end{itemize} The advantage of the 1D--model is its compactness and simplicity (no dependence on the dark matter profile), while the second approach is more accurate. The former could be preferred in situations where speed is a decisive factor, e.g. when scanning a large parameter space as in {\\sc susy} studies. Moreover, in order to be conservative in exclusion studies, small halo size should be taken. In particular conservative estimates are made when using the 3$^{\\rm rd}$ line of Tab.~\\ref{tab:param}: this is exactly the regime where 1D--formulae hold. We also supply numerical routines that could be easily implemented in indirect detection codes.  So far, these formulae are valid only in the framework of the constant wind model. Although all the discussions were developed with respect to exotic \\pbar\\ fluxes, all the results can be transposed almost without any modification to the \\dbar\\ case (the total cross section has to be taken accordingly to the species considered)."
    },
    "0609/astro-ph0609714_arXiv.txt": {
        "abstract": "The technique to realize 3D position sensitivity in a two-phase xenon time projection chamber (XeTPC) for dark matter search is described. Results from a prototype detector (XENON3) are presented. \\vspace{1pc} ",
        "introduction": "\\label{} The XENON dark matter search experiment~\\cite{xenon01} uses time projection chamber (TPC) to search for the hypothetical WIMP dark matter particles. The detector (see the most recent status at \\cite{yamashita:ucla06}) consists a bulk of liquid xenon (LXe) as the target for WIMP interactions. The target is also self-shielded from the external background events, mostly gamma rays. The external gamma rays have more chance to interact near the edge and surface of the LXe target. A fiducial volume cut of the edge and surface events will dramatically reduce the background rate and improve the detector's sensitivity for dark matter search. A fiducial volume cut will also help to remove events from the regions (edge or surface) with irregular or non-uniform electric fields. The signals of events from these regions can be similar to the WIMP signals and make them difficult to be rejected without 3D position sensitivity. A position sensitive detector will also have the capability to distinguish neutron interactions by their multiple scatters in the target, while a WIMP is very unlikely to make more than one scatter due to its tiny interaction cross-section. ",
        "conclusions": ""
    },
    "0609/astro-ph0609187_arXiv.txt": {
        "abstract": "During an imaging survey of the Chamaeleon~I star-forming region with the Advanced Camera for Surveys aboard the {\\it Hubble Space Telescope}, we have discovered a candidate substellar companion to the young low-mass star CHXR~73 ($\\tau\\sim2$~Myr, $M\\sim0.35$~$M_\\odot$).  We measure a projected separation of $1\\farcs30\\pm0\\farcs03$ for the companion, CHXR~73~B, which corresponds to 210~AU at the distance of the cluster. A near-infrared spectrum of this source obtained with the Cornell Massachusetts Slit Spectrograph at the Magellan~II telescope exhibits strong steam absorption that confirms its late-type nature ($\\gtrsim$M9.5). In addition, the gravity-sensitive shapes of the $H$- and $K$-band continua demonstrate that CHXR~73~B is a young, pre-main-sequence object rather than a field star. The probability that CHXR~73~A and B are unrelated members of Chamaeleon~I is $\\sim0.001$. We estimate the masses of CHXR~73~B and other known substellar companions in young clusters with a method that is consistent with the dynamical measurements of the eclipsing binary brown dwarf 2M~0535-0546, which consists of a comparison of the bolometric luminosities of the companions to the values predicted by the evolutionary models of Chabrier \\& coworkers and Burrows \\& coworkers. We arrive at mass estimates of 0.003-0.004, $0.024\\pm0.012$, $0.011^{+0.01}_{-0.003}$, and $0.012^{+0.008}_{-0.005}$~$M_{\\odot}$ for 2M~1207-3932~B, GQ~Lup~B, DH~Tau~B, and CHXR~73~B, respectively.  Thus, DH~Tau~B and CHXR~73~B appear to be the least massive companions to stars outside the solar system that have been detected in direct images, and may have masses that are within the range observed for extrasolar planetary companions ($M\\lesssim0.015$~$M_{\\odot}$). However, because these two objects (as well as 2M~1207-3932~B) probably did not form within circumstellar disks around their primaries, we suggest that they should be viewed as brown dwarf companions rather than planets. ",
        "introduction": "\\label{sec:intro}       Since the first discovery of substellar companions outside of the solar system little more than a decade ago \\citep{wol92,mq95,mb96,opp95}, a great deal of work has focused on understanding the origin of these bodies. The three most widely explored theories for the formation of substellar companions involve the accretion of gas by rocky cores in circumstellar disks \\citep{pol96}, the collapse of gravitationally unstable areas of disks \\citep{bos98,bos06}, and the fragmentation of molecular cloud cores \\citep{lod05}. Testing the validity of these models requires observations of companions across a large range of companion mass, primary mass, orbital separation, and age. Searches for substellar companions at young ages are particularly valuable because they can directly constrain the formation timescale and early dynamical evolution of these objects. In addition, substellar objects are brightest when they are young, and thus can be detected at very low masses with direct imaging. For these reasons, young stars and brown dwarfs in the solar neighborhood \\citep[$\\tau\\sim30$-400~Myr,][]{reb98,mh04,mz04,low05}, open clusters and associations \\citep[$\\tau\\sim10$-100~Myr,][]{mar98,mar00,mar03,low00,neu00a,gue01,cha03,cha05b}, and star-forming regions and OB associations \\citep[$\\tau\\lesssim5$~Myr,][]{duc99,bra00,neu02,luh05wfpc,kra05,kra06} have been popular targets for companion searches with high-resolution imaging. The most reliable substellar companions discovered to date at ages of $\\tau<10$~Myr consist of GG~Tau~Bb \\citep{whi99}, TWA~5~B \\citep{low99,neu00b}, 2M~1101-7732~B \\citep{luh04bin}, 2M~1207-3932~B \\citep{cha04,cha05a}, DH~Tau~B \\citep{ito05}, and GQ~Lup~B \\citep{neu05}.  Adding to the small but quickly growing list of known young substellar companions, we report the discovery of a companion to the low-mass star CHXR~73 \\citep[M3.5, $M\\sim0.35$~$M_\\odot$,][]{luh04cha} in the Chamaeleon~I star-forming region ($\\tau\\sim2$~Myr), which was found serendipitously during a survey for free-floating brown dwarfs with the {\\it Hubble Space Telescope} ($HST$).  In this paper, we present our optical images from $HST$ and near-infrared (IR) imaging and spectroscopy from the Magellan telescopes for CHXR~73~B (\\S~\\ref{sec:obs}) and we use these data to assess its youth and membership in Chamaeleon~I (\\S~\\ref{sec:mem}). We then examine the evidence that CHXR~73~B is a companion (\\S~\\ref{sec:bin}), estimate the mass of CHXR~73~B and other young substellar companions in a uniform manner (\\S~\\ref{sec:mass}), and discuss the origin of these companions and whether they should be referred to as planets or brown dwarfs (\\S~\\ref{sec:disc}). ",
        "conclusions": "\\label{sec:disc}        We have estimated masses for CHXR~73~B and other young low-mass companions in a uniform manner, and with a method that produces masses that are consistent with dynamical measurements for the one known young eclipsing binary brown dwarf, which allows us to examine both the relative and absolute masses of these companions. \\citet{neu05} reported that GQ~Lup~B could have a mass as low as 1~$M_{\\rm Jup}$, possibly making it the first planetary-mass companion imaged around a star. However, GQ~Lup~B is only slightly less luminous than 2M~1101-7732~B \\citep{luh04bin} and is significantly brighter than DH~Tau~B \\citep{ito05}. If 2M~1101-7732~B and 2M~1207-3932~A were placed on the diagram of luminosity versus age in Figure~\\ref{fig:lbol}, they would have masses comparable to that of GQ~Lup~B. For added perspective, we note that GQ~Lup~B has a higher luminosity than many of the known late-type members of nearby star-forming regions \\citep[$>$M8,][]{bri02,luh03}. Thus, it appears that GQ~Lup~B was not the least massive companion imaged at the time of its discovery and our mass estimate of $0.024\\pm0.012$~$M_{\\odot}$ suggests that it is unlikely to be a planet by any definition of the term.  DH~Tau~B and CHXR~73~B are the faintest known companions to young stars ($\\tau<10$~Myr). If our mass estimates of $0.011^{+0.01}_{-0.003}$ and $0.012^{+0.008}_{-0.005}$~$M_{\\odot}$ are accurate, then they (and possibly AB~Pic~B) are the least massive companions directly imaged near stars outside the solar system. In fact, these estimates are just below both the maximum mass observed for extrasolar planetary companions \\citep[$M\\sim0.015$~$M_{\\odot}$,][]{mar05} and the deuterium burning limit \\citep[$M\\sim0.014$~$M_{\\odot}$,][]{sau96,cha00}. If DH~Tau~B and CHXR~73~B have planetary masses, then are they planets? We can pose the same question for 2M~1207-3932~B, which has an even lower mass according to the models. In other words, is a planetary-mass companion always a planet?  The definition of a planet can be based on a companion's mass or on its formation mechanism. Following \\citet{cha06}, we advocate the latter definition in which a planet is an object that forms in the disk around a star or a brown dwarf, while a substellar companion that forms in another manner (e.g., core fragmentation) is a brown dwarf. We apply this definition of a planet to DH~Tau~B, CHXR~73~B, and 2M~1207-3932~B by drawing on theoretical work on planet formation in disks. The maximum separations at which giant planet formation can occur via both core accretion \\citep[$\\lesssim10$~AU,][]{pol96,ina03} and disk instability \\cite[$\\lesssim20$-100~AU,][]{bos03,bos06} are much lower than the projected separations of 200 and 300~AU for CHXR~73~B and DH~Tau~B, respectively. In addition, these systems are too young ($\\tau\\sim1-2$~Myr) for the production of planets through core accretion. Thus, CHXR~73~B and DH~Tau~B probably did not form in disks around their primaries. Formation via disk instability cannot be ruled out for 2M~1207-3932~B \\citep{lod05}, but it would require an unusually high ratio of disk mass to brown dwarf mass for 2M~1207-3932~A when it was younger \\citep{kle03,sch06}. Meanwhile, it has been shown that cloud fragmentation can produce isolated objects that reach into the planetary mass regime. For instance, the least massive free-floating brown dwarf in Chamaeleon~I is one magnitude fainter than CHXR~73~B \\citep[$M\\sim0.08$~$M_{\\odot}$,][]{luh05cha}, so it seems plausible that the same process produced CHXR~73~B, but in a binary system instead of in isolation. Based on these various theoretical and observational considerations, we conclude that CHXR~73~B and other known young substellar companions probably did not form in a planetary fashion in disks and instead should be referred to as brown dwarfs. However, because the products of planet formation and core fragmentation may overlap in separation as well as in mass (e.g., 5~$M_{\\rm Jup}$ at 10~AU), future surveys may uncover companions that cannot be classified unambiguously as either planets or brown dwarf companions."
    },
    "0609/astro-ph0609464_arXiv.txt": {
        "abstract": "We report the discovery of T dwarf companions to the nearby stars HN~Peg (G0V, 18.4~pc, $\\tau\\sim0.3$~Gyr) and HD~3651 (K0V, 11.1~pc, $\\tau\\sim7$~Gyr). During an ongoing survey of $5\\arcmin\\times5\\arcmin$ fields surrounding stars in the solar neighborhood with the Infrared Array Camera aboard the {\\it Spitzer Space Telescope}, we identified these companions as candidate T dwarfs based on their mid-infrared colors. Using near-infrared spectra obtained with SpeX at the NASA Infrared Telescope Facility, we confirm the presence of methane absorption that characterizes T dwarfs and measure spectral types of T2.5$\\pm$0.5 and T7.5$\\pm$0.5 for HN~Peg~B and HD~3651~B, respectively. By comparing our {\\it Spitzer} data to images from the Two-Micron All-Sky Survey obtained several years earlier, we find that the proper motions of HN~Peg~B and HD~3651~B are consistent with those of the primaries, confirming their companionship. HN~Peg~B and HD~3651~B have angular separations of $43\\farcs2$ and $42\\farcs9$ from their primaries, which correspond to projected physical separations of 795 and 476~AU, respectively. A comparison of their luminosities to the values predicted by theoretical evolutionary models implies masses of $0.021\\pm0.009$ and $0.051\\pm0.014$~$M_\\odot$ for HN~Peg~B and HD~3651~B. In addition, the models imply an effective temperature for HN~Peg~B that is significantly lower than the values derived for other T dwarfs at similar spectral types, which is the same behavior reported by Metchev \\& Hillenbrand for the young late-L dwarf HD~203030~B. Thus, the temperature of the L/T transition appears to depend on surface gravity. Meanwhile, HD~3651~B is the first substellar companion directly imaged around a star that is known to harbor a close-in planet from radial velocity surveys. The discovery of this companion supports the notion that the high eccentricities of close-in planets like the one near HD~3651 may be the result of perturbations by low-mass companions at wide separations. ",
        "introduction": "\\label{sec:intro}     Direct imaging surveys for resolved substellar companions are vital for studies of brown dwarfs and planets. These surveys provide measurements of the frequency of companions as a function of various physical properties, such as primary mass, companion mass, separation, and age, which in turn can be used to constrain the mechanisms by which substellar companions form. In addition, compared to their isolated counterparts, each newly uncovered companion represents a superior laboratory for studying the atmospheres and evolution of substellar objects because the age, distance, and metallicity of the primary, and hence the companion, are relatively easy to measure.  A variety of criteria can be adopted for selecting primaries to search for companions. Stars for which planets have been discovered through radial velocity monitoring and young nearby stars are particularly attractive categories to consider. Several direct imaging surveys have focused on known extrasolar planetary systems \\citep{luh02,pat02,mug05b,chau06}, producing one companion that was initially described as a brown dwarf \\citep{els01} but later shown to be a white dwarf \\citep{mug05a,lag06}. Similar observations have been performed toward young stars and brown dwarfs within 100~pc of the Sun \\citep[$\\tau\\sim10$-400~Myr,][]{cha03,neu00a,mz04,low05}, resulting in the discovery of several substellar companions that include G196-3~B \\citep{reb98}, TWA~5~B \\citep{low99,neu00b}, HR~7329~B \\citep{low00,gue01}, AB~Pic~B \\citep{cha05b}, HD~49197~B \\citep{mh04}, 2M~1207-3932~B \\citep{cha04,cha05a}, and HD~203030~B \\citep{mh06}. The spectral types of these companions range from late M through L but do not extend into the cooler T types. The only known T dwarf companions to stars consist of Gl~229~B \\citep{nak95,opp95}, Gl~570~D \\citep{bur00}, Gl~337~D \\citep{wil01,bur05}, $\\epsilon$~Ind~B and C \\citep{sch03,mcc04}, and SCR~1845~B \\citep{bil06}, all of which orbit nearby stars that are older than $\\sim1$~Gyr and have no known planets.  Most resolved substellar companions have been identified through deep, high-resolution imaging at small separations ($\\rho\\lesssim10\\arcsec$) with adaptive optics or the {\\it Hubble Space Telescope} or through shallow imaging at arbitrarily large separations with the Two-Micron All-Sky Survey (2MASS) or the Sloan Digital Sky Survey (SDSS). Because of its wide field of view ($\\rho=2\\farcm5$) and unprecedented sensitivity at mid-infrared (IR) wavelengths, the Infrared Array Camera \\citep[IRAC;][]{faz04} aboard the {\\it Spitzer Space Telescope} \\citep{wer04} complements these previous search methods by providing the best available sensitivity to low-mass companions at separations of $\\rho\\sim10$-$100\\arcsec$ ($r\\sim100$-1000~AU) from nearby stars. To take advantage of these capabilities, we have obtained IRAC images of nearby young stars and planet-bearing stars identified in radial velocity surveys. In this paper, we report the discovery of two T dwarf companions to stars in these samples, HN~Peg ($d=18.4$~pc, G0V) and HD~3651 ($d=11.1$~pc, K0V) \\footnote{In a paper submitted after this one, \\citet{mug06} reported an independent discovery of the companion to HD~3651.}. The former is a young solar analog \\citep{gai98} while the latter has a sub-Saturn planetary companion at 0.3~AU \\citep{fis03,fis05}. We begin by describing the IRAC images that were used to identify the T dwarf companions to these stars (\\S~\\ref{sec:phot}) and followup near-IR photometry and spectroscopy (\\S~\\ref{sec:spec}). We then assess the evidence of companionship (\\S~\\ref{sec:bin}), measure spectral types (\\S~\\ref{sec:class}), and estimate the physical properties of these companions (\\S~\\ref{sec:prop}). Finally, we discuss these newly discovered T dwarfs in the context of previously known brown dwarfs (\\S~\\ref{sec:disc}). ",
        "conclusions": "\\label{sec:disc}  We have presented the discovery of two T dwarf companions in the solar neighborhood, HN~Peg~B (T2.5) and HD~3651~B (T7.5). These T dwarfs are the sixth and seventh known T dwarf companions to stars and are the first T dwarfs discovered with the {\\it Spitzer Space Telescope}. In this section, we describe additional notable aspects of these new brown dwarfs.  HN~Peg~A has been the target of high-precision radial velocity monitoring \\citep{cum99,fis05,kon06} and near-IR coronographic imaging \\citep{mz04}, but no companions have been found through those observations. HN~Peg~B was bright enough to be detected by the imaging of \\citet{mz04}, but it was outside of their field of view ($r<15\\arcsec$). This companion was also beyond the field of view of unpublished images of HN~Peg obtained with the Near-Infrared Camera and Multi-Object Spectrometer aboard the {\\it Hubble Space Telescope} during General Observer program 10176. Various data from the literature indicate that HN~Peg~A has an age of 0.1-0.5~Gyr (\\S~\\ref{sec:age}). If the T dwarf S~Ori~70 \\citep{zap02} is a field dwarf rather than a member of the young cluster $\\sigma$~Ori \\citep{bur04}, then HN~Peg~B may be the youngest known T dwarf. HN~Peg~B closely resembles the recently discovered late-L companion to HD~203030 in terms of age and mass \\citep{mh06}.  Several properties of HN~Peg~B do not behave in the manner predicted for young, low-gravity T dwarfs. Although the colors within the $JHK$ bands and within the IRAC bands are normal, all colors containing one of the former and one of the latter are redder than those of typical T dwarfs by $\\sim0.3$~mag (\\S~\\ref{sec:lbol}). In contrast, low gravity should cause these colors to become bluer, not redder, according to the theoretical spectra of \\citet{bur06}. Based on its position in color-magnitude diagrams, HN~Peg~B is probably not an equal-component binary, but it could be a binary with two unequal components\\footnote{One of the T dwarfs that has been unresolved in previous imaging, 2M 0559-1404 \\citep{bur03}, is overluminous in Figure~\\ref{fig:cmd}, which suggests that it could be a binary. This brown dwarf was identified as a possible binary because of its bright $J$ magnitude relative to late L dwarfs \\citep{dah02}, but more recent 1-5~\\micron\\ measurements indicated that it might not be significantly overluminous compared to other T dwarfs \\citep{gol04,vrba04}.}. Because later T dwarfs emit proportionately more flux in the mid-IR bands than at shorter wavelengths, an unresolved system of this kind (e.g., T1+T5) might be responsible for these anomalous colors. Another possible explanation is photometric variability between the dates of the IRAC and $JHK$ observations \\citep{bur02,eno03}.  Our estimates of the stellar parameters of HN~Peg~B also exhibit peculiarities. When we compare the luminosity and age of HN~Peg~B to the predictions of theoretical evolutionary models, we derive an effective temperature of $1130\\pm70$~K, which is lower than the values previously measured for T dwarfs at T2-T3 \\citep[$T_{\\rm eff}\\sim1400$~K,][]{gol04}. In contrast, the spectral modeling of late T dwarfs by \\citet{bur06b} found that temperature should increase rather than decrease with lower gravity at a given spectral type. If this prediction is combined with the models shown in Figs.~\\ref{fig:bur97} and \\ref{fig:bar03}, then a young T dwarf should be much brighter than an older one at the same spectral type. Instead, we find that HN~Peg~B is at the faint end of T0-T4 dwarfs \\citep{gol04} and is at least one magnitude fainter in bolometric luminosity than expected for its assumed age. One possible explanation for these discrepancies is that the HN~Peg system is much older than various data for the primary seem to indicate (\\S~\\ref{sec:age}). However, \\citet{mh06} have recently reported the same anomalous behavior for a late-L companion near another star that appears to be young, HD~203030, and it seems unlikely that the age estimates for both stars have such large errors. Instead, the presence of this phenomenon in both HN~Peg~B and HD~203030~B strongly indicates that the temperature of the L/T transition is dependent on surface gravity, and in a manner that has not been previously predicted.  When combined with the evolutionary models, the luminosity and age of HN~Peg~B translate to a mass $0.021\\pm0.009$~$M_\\odot$, making it one of the least massive known T dwarfs. In addition, like HD~203030~B \\citep[$M=0.023$~$M_\\odot$][]{mh06}, HN~Peg~B is among the least massive known brown dwarf companions and is a more evolved counterpart to the widely-separated companions DH~Tau~B \\citep[$\\tau\\sim1$~Myr,][]{ito05}, CHXR~73~B \\citep[$\\tau\\sim2$~Myr,][]{luh06cha}, and AB~Pic~B \\citep[$\\tau\\sim30$~Myr,][]{cha05b}, which have masses of 0.01-0.02~$M_\\odot$ \\citep{luh06cha}.    For HD~3651, radial velocity monitoring has uncovered a sub-Saturn-mass planet at 0.3~AU \\citep{fis03,fis05} while no companions have been identified in adaptive optics images \\citep{car05}. The latter observations did not detect HD~3651~B because their field of view was too small ($r<15\\arcsec$). \\citet{tak06} suggested that the high eccentricities of planets in old, single-planet systems like HD~3651 might be caused by secular perturbations from distant low-mass companions, and indeed we have found such a companion in the form of HD~3651~B. This T dwarf is the first resolved substellar companion that has been discovered in one of the extrasolar planetary systems identified in radial velocity surveys. Like HN~Peg~B, HD~3651~B is probably not an equal-magnitude binary brown dwarf according to its position in the color-magnitude diagram in Figure~\\ref{fig:cmd}. HD~3651~B is a virtual twin to another T dwarf companion, Gl~570~D, in terms of spectral type, luminosity, mass, temperature, and radius \\citep{bur00,geb01}. As with Gl~570~D, these measurements have relatively high accuracy because of the companionship of HD~3651~B to a nearby, well-studied star, making it a valuable brown dwarf for calibrating methods of characterizing the physical properties of late T dwarfs \\citep{bur06a}."
    },
    "0609/astro-ph0609378_arXiv.txt": {
        "abstract": "Based on simple physical and geometric assumptions, we have calculated the mean surface molecular density of spiral galaxies at the threshold between star formation induced by cloud-cloud collision and spontaneous gravitational collapse.  The calculated threshold is approximately  $\\log \\Sigma_\\mathrm{crit} \\sim 2.5$, where $ \\Sigma \\quad \\mathrm{M_{\\solar}}\\cdot \\mathrm{pc}^{-2}$ is the observed surface mass density of an assumed flat gas disk.  Above this limit, the rate of molecular cloud collisions dominates over spontaneous molecular cloud collapse.  This model may explain the apparent discontinuity in the Schmidt law found recently at $2 \\lesssim \\log \\Sigma \\lesssim 3$. ",
        "introduction": "The onset of star formation within an ensemble of molecular clouds is thought to have two modes, spontaneous (stochastic) gravitational collapse \\citep{bonnell04, krumholtz05}, and collisions between molecular clouds (e.g., \\cite{tan}).  Most studies to date have focused on one of these two possibilities. In actual physical situations which cover a range of physical parameters, these two processes should both have effect on the overall star formation in galaxy disks.  In regions of higher molecular density, collisions should become dominant over spontaneous formation, because collisions will occur on timescales shorter than spontaneous star formation.  In regions with scarse cloud number density, collisions should also be rare and therefore star formation from other stochastic means should dominate.  The transition from spontaneous to collisional star formation may have observed effects on the relation between molecular gas content and star formation rate (SFR).  This expected transition has not been clearly observed yet, however.  There is growing observational evidence that cloud collision can trigger star formation \\citep{loren76, scoville86, hasegawa94, koda06}.  These studies show that cloud collisions are seen mainly at high density regions, and therefore merging and interacting galaxies are candidates for galaxies forming stars mainly by collision. A comparison of these galaxies with normal galaxies have been conducted by \\citet{young86}.  They have found that merging/interacting galaxies exhibit SFRs that are an order of magnitude higher than normal galaxies of the same molecular gas density.  Similarly, central regions and spiral arms of normal galaxies, which generally have high molecular densities, may include star formation induced by collision.  However, quantitative calculations which give direct (observable) parameters on exactly $\\textit{where}$ in parameter space star formation transits from spontaneous to collisional mechanisms have not been proposed. ",
        "conclusions": "We have constructed a simple model to consider the threshold between star formation dominated by spontaneous gravitational collapse and cloud-cloud collision.  Although both cloud-cloud collision and gravitational collapse processes have for decades been thought to cause star formation, the threshold of these two processes has not attracted much attention.  Most studies of collisional star formation have focused on qualitative aspects.  Numerical studies such as that by \\citet{mazzei} have pointed out a threshold, but with the contradictory result that collisions between clouds result in lower star formation rates.  Our model, by assuming that star formation is a bimodal process dominated either by spontaneous gravitational collapse or cloud-cloud collision, does not take into account the possiblity of star formation induced by expanding SNR shells (self-propagating star formation), or complicated structures in temperatures or chemistry.  Simple as our model is, it succeeds in explaining the apparent discontinuity in the Schmidt law. Moreoever, our formulation coupled with the empirical implication that most clouds are virialized and thus GMC size and mean density are anticorrelated, leads to the conclusion that the threshold surface mass density $\\log \\Sigma_\\mathrm{crit}$ is practically independent of cloud parameters (size, density, mass), but that it depends on the mean velocity dispersion $\\sigma$ of clouds.  Further tests of our formulation can be expected from observations of mergers.  In merging spiral pairs whose rotation are oriented in the reverse direction, the relative velocity of the clouds are increased so that according to equation (7), collisional star formation starts at a very low surface mass density.  However, if the rotation is oriented in the same direction, $\\sigma$ will not increase as much as in the reverse orientation.  Thus, the threshold $\\log \\Sigma_\\mathrm{crit}$ starts at a higher density.  Assuming that collisional star formation results in a discontinuity in the SFR relation with surface gas density, we can compare mergers with these two orientations with average normal galaxies which we can assume to be dominated by a spontaneous process.  Then, we should see that mergers with reverse rotation orientation will start its discontinuity in the Schmidt law at a lower surface mass density than mergers with ordered rotation.\\\\[0.5cm]  The authors thank H. Nakanishi, and S. Onodera, and K. Kohno for fruitful discussions.  S.K. and F.E. were financially supported by a Research Fellowship from the Japan Society for the Promotion of Science for Young Scientists."
    },
    "0609/astro-ph0609134_arXiv.txt": {
        "abstract": " ",
        "introduction": "On 4 July 2005 UT, the Deep Impact spacecraft impacted comet 9P/Tempel 1. The goal of the impact experiment was to expose deep layers of the comet nucleus to study the chemical composition of those layers and to ascertain the structural properties of the nucleus.  Comets are the remnant building blocks from the outer solar nebula. The majority of their lifetimes are spent at large heliocentric distances so they undergo little chemical evolution.  However, even far from the Sun, the outer layers are altered via radiation darkening (Strazzula and Johnson 1991).  Once comets enter the inner Solar System, sublimation removes outer layers, while mantling covers the exterior with dust. \\nocite {stjo91}  Comet 9P/Tempel 1 is a Jupiter family comet which has made many passages into the inner Solar System and has not been extremely far from the Sun in many orbits.  Thus, the expectation was that there was an altered outer layer to the comet, but that the interior was pristine.  The outer layer is thought to be too deep for a small drill carried by a spacecraft to penetrate to pristine material. In order to reach fresh, unaltered material, the Deep Impact mission proposed to create a large crater by placing a small (370kg) spacecraft in the path of the comet, setting up an impact. The impact occurred on 4 July 2005 at 05:44:36 UT (Earth-received time of 05:52:02 UT).  The impact speed was 10.3 km/sec resulting in an impact kinetic energy of 19 GJ (A'Hearn {\\it et al.} 2005). \\nocite{ahdeepimpact}  The impactor was carried by a larger flyby spacecraft which was able to monitor the impact with high spatial resolution.  However, while the flyby spacecraft carried an IR spectrometer, it was not able to perform optical spectroscopic observations of the impact, nor was it able to monitor the comet for many days after impact. Thus, a ground-based campaign was established in order to monitor changes in the comet (Meech {\\it et al.} 2005). \\nocite{medeepimpact}  We observed the impact and its aftermath from the Keck 1 telescope using the HIRES spectrograph as part of the ground-based campaign. High resolution spectra could be used to monitor aspects of the impact which could not be observed at lower spectral resolution.  These include monitoring isotope ratios to determine whether the ices in the interior of the comet are isotopically similar to the ices we normally see subliming from the surface.  Jehin {\\it et al.} (2006) reported the isotopic ratios for $^{12}$C/$^{13}$C and $^{14}$N/$^{15}$N, both during the first few hours after impact and at times other than the impact. They showed that the excavated material is isotopically similar to the normal outgassing. High resolution spectra allow for detailed examination of the rotational states of the molecules. This allowed us to determine whether any prompt emissions were seen or whether the impact produced any rotationally hot lines. High resolution also allowed for detection of any new, unexpected species.  High resolution data are more sensitive to very weak lines than low resolution spectra because higher spectral resolving power disperses the underlying continuum to a greater extent, resulting in higher line-to-continuum contrast (our spectral resolution of $\\sim$6 km\\,sec$^{-1}$ was not sufficient to resolve typical cometary lines with widths of $\\sim$1 km\\,sec$^{-1}$). \\nocite{jedeepimpact} ",
        "conclusions": "1) The crater was not deep enough to penetrate the mantle to primitive material, i.e. the mantle is thicker than we had supposed; or 2) The cometary material on the outside of the nucleus is not altered significantly from the interior materials.   Groussin {\\it et al.} (2006) showed that the nucleus has very low thermal inertia. Thus, neither the diurnal heat wave or the heat wave from the extended passage into the inner solar system would penetrate deeply into the nucleus.  This would leave pristine material near the nucleus.  Thus, as we saw with our Keck data, the interior of the comet did not look substantially different than the exterior layers and the outer layers must be very thin. \\nocite{grlpsc06} \\vspace*{0.5in}   \\begin{center}Acknowledgements\\end{center} We thank Dr. Fred Chaffee for making Director's Discretionary time available pre-impact.  Our thanks to Dr. Hien Tran for obtaining the observations of 30 May and 4 July and for help on 5 and 6 July. This work was funded by NASA Grant NNG04G162G (ALC), NASA Grant NNG06A67G (WMJ), NSF Grant CHE-0503765 (WMJ) and by support from the University of Maryland through a subcontract 2667702 (KJM), which was awarded under prime contract NASW-00004 from NASA .  \\newpage"
    },
    "0609/astro-ph0609628_arXiv.txt": {
        "abstract": "In the last few years our knowledge of the physics of massive stars has improved tremendously. However, further investigations are still needed, especially regarding accurate calibrations of their fundamental parameters. To this end, we have constructed a comprehensive grid of NLTE model atmospheres and corresponding synthetic spectra in the massive star domain. The grid covers the complete B type spectral range, extended to late O on the hot side and early A on the cool side, from supergiants to dwarfs and from weak stellar winds to very strong ones. It has been calculated with the latest version of the FASTWIND code. The analysis of an extensive sample of OB stars in the framework of the COROT space mission will lead to accurate calibrations of effective temperatures, gravities, mass loss rates etc. This paper contains a detailed description of the grid, which has been baptised as BSTAR06 and which will be available for further research in the near future. ",
        "introduction": "In the last decade, with the advent of high-resolution, high signal-to-noise spectroscopy, massive stars, and the many astrophysical processes influenced by them, have regained considerable interest. Significant improvements have been achieved in the development and accuracy of models predicting the stellar atmospheres and their winds. Still, it is surprising to note how many things we do not understand, just to name, e.g., the formation of structure (loosely called ``clumping'') in the stellar wind and its consequences. Though a lot of progress has been made in the O and A type regime, the B type regime in between still suffers from many shortcomings.  Since the detailed spectroscopic analysis of individual objects is a rather time-consuming (and boring!) job, scientists refrained from analysing large samples and never tackled more than a few tens of stars at once. Therefore and unfortunately, knowledge in the B type regime still relies on small number statistics and could thus only gain from large sample studies. As an alternative to (automatic) high precision analyses (e.g., by means of genetic algorithms, \\citealt{Mokiem2005}, or neuronal networks) one could profit from a grid-method, which offers a good compromise between effort, time and precision when an {\\it appropriate grid} has been set up.  Since one of our goals is to derive effective temperature calibrations for the complete B type spectral range, for stars in different evolutionary phases, in a way as homogeneous as possible, we need both a comprehensive sample of stars {\\it and} a huge grid of reliable model atmospheres. The first has been supplied to us by the COROT team, and consists of a database of ground-based observations, obtained in preparation of the space mission and for which they need accurate fundamental parameters. The database comprises FEROS, ELODIE and SARG spectra of some 350 massive OBA stars brighter than 9.5 mag. To satisfy our second need, we developed an appropriate and extensive grid of model atmospheres in the way discussed in Section 3. ",
        "conclusions": "We have computed a dense grid of NLTE atmosphere models with and without winds covering all spectral types B. The computations were performed continuously during a period of seven months on a dedicated Linux cluster with 20 CPUs (3800 MHz processors with 4 Gb RAM memory and 8 Gb Swap memory).  Whenever available these 20 CPUs were extended with 40 extra CPUs (among which 8 more of 3800MHz and 32 of 3400MHz).  In order to avoid unnecessary repetition of the computation of such a huge grid, we will offer BSTAR06 to the community for further research in the near future.  By translating the different effects of the individual parameters on the selected line profiles into code, one can finally perform an automated line profile fitting procedure. This process is currently in its phase of development. Once the automated tool does exist, we will be able to analyse B type stars in a fast and effective way, which hopefully leads to a better understanding of this hot star regime. In particular, the interpretation of the COROT grid should lead to a much better calibration of hot stars across the whole B type spectral range."
    },
    "0609/astro-ph0609558_arXiv.txt": {
        "abstract": "We report on observations of several optical emission lines toward a variety of newly revealed faint, large-scale \\ha-emitting regions in the Galaxy. The lines include \\nii$~\\lambda6583$, \\nii$~\\lambda5755$, \\sii$~\\lambda6716$, \\oiii$~\\lambda5007$, and \\hei$~\\lambda5876$\\ obtained with the Wisconsin H-Alpha Mapper (WHAM) toward sightlines that probe superbubbles, high latitude filamentary features, and the more diffuse warm ionized medium (WIM).  Our observations include maps covering thousands of square degrees toward the well-known Orion-Eridanus bubble, a recently discovered $60\\dg \\times 20\\dg$ bipolar superbubble centered in Perseus, plus several classical \\hii\\ regions surrounding OB stars and hot evolved stellar cores. We use the emission line data to explore the temperature and ionization conditions within the emitting gas and their variations between the different emission regions. We find that in the diffuse WIM and in the faint high latitude filamentary structures the line ratios of \\nii/\\ha\\ and \\sii/\\ha\\ are generally high, while \\oiii/\\ha\\ and \\hei/\\ha\\ are generally low compared to the bright classical \\hii\\ regions. This suggests that the gas producing this faint wide-spread emission is warmer, in a lower ionization state, and ionized by a softer spectrum than gas in classical \\hii\\ regions surrounding O stars, the presumed ionization source for the WIM. In addition, we find differences in physical conditions between the large bubble structures and the more diffuse WIM, suggesting that the ionization of superbubble walls by radiation from interior O associations does not account entirely for the range of conditions found within the WIM, particularly the highest values of \\nii/\\ha\\ and \\sii/\\ha. ",
        "introduction": "\\label{sec:intro}  The interstellar medium (ISM) plays a vital role in the ongoing cycle of stellar birth and death and galactic evolution.  However, the role of interstellar matter, from how its properties are influenced by stars to how in turn its properties influence star formation, is poorly understood and is arguably the least understood portion of the cycle. Warm diffuse ionized hydrogen has become recognized as a major phase of the ISM of our Galaxy; see, for example, reviews by \\citet{KH87, Cox89, Reynolds91b, Mathis00}. This phase consists of regions of warm (10$^{4}$ K), low-density (10$^{-1}$ cm$^{-3}$), nearly fully ionized hydrogen that occupy approximately 20\\% of the volume within a 2 kpc thick layer about the Galactic midplane \\citep[e.g.,][]{Reynolds91a, NCT92, TC93, HRT99}. Near the midplane, the rms density of \\hii\\ is less than 5\\% that of \\hi. However, because of its greater scale height, the total column density of interstellar \\hii\\ along high Galactic latitude sight lines is relatively large, with $N_{\\rm{H\\sc{II}}} \\sim 1/3~N_{\\rm{H\\sc{I}}}$. One kiloparsec above the midplane, warm \\hii\\ may be the dominant state of the interstellar medium in the Milky Way.  Widespread, diffuse ionized gas is now firmly established as an important constituent of the ISM in external galaxies as well \\citep[e.g.,][]{RKH90, HG90, Dettmar92,  WB94, Ferguson+96, RD00, CR01, MV03}.  Despite its significance, the origin and physical conditions within the warm ionized medium (WIM) remain poorly understood. In particular, the ubiquitous nature of this gas is difficult to explain. Of the known sources of ionization within the Galaxy, only O stars generate enough power to sustain the WIM \\citep{Reynolds92}. Therefore, it is generally believed that the O stars, confined primarily to widely separated stellar associations near the Galactic midplane, are somehow able to photoionize a significant fraction of the ISM not only in the disk but also within the halo, 1-2 kpc above the midplane. However, the need to have a large fraction of the Lyman continuum photons from O stars travel hundreds of parsecs through the disk seems to conflict with the traditional picture of \\hi\\ permeating much of the interstellar volume near the Galactic plane. It has been suggested that extensive cavities in the neutral gas, created either by ``superbubbles'' of hot gas from supernovae \\citep{Norman91}, or  carved out by O star photons in low-density regions \\citep{MC93}, may extend far above the midplane \\citep{DS94, DSF00}. Although the existence of \\hi\\ superbubbles has long been established \\citep{Heiles84}, direct observational evidence that cavities are actually responsible for the transport of hot gas and ionizing radiation up into the Galactic halo is very limited.  One piece of evidence for such transport has been provided recently by the WHAM \\ha\\ sky survey \\citep[][also see \\S6 below]{RSH01}.  The WHAM \\ha\\ sky survey is a velocity-resolved map of diffuse interstellar \\ha\\ emission at $1\\dg$ angular resolution over the entire northern sky ($\\delta > -30\\dg$) within approximately $\\pm100$ \\kms\\ of the local standard of rest (LSR) \\citep{Haffner+03}. The survey maps show \\ha\\ emission covering the sky, with ionized gas associated with large scale loops, filaments, and bubbles superposed on a fainter, diffuse background, as well as the bright classical \\hii\\ regions near the Galactic plane.  Several of the high latitude structures appear to be associated with hot stars and OB associations; however, the diffuse background and many features superposed upon it have no clear association with known ionizing sources. This survey provides the basis for studies of the physical conditions within these newly revealed emission regions and the source of their ionization.  Even though the primary source of ionization is believed to be O stars, the temperature and ionization conditions within the diffuse ionized gas differ significantly from conditions within classical O star \\hii\\ regions.  These conditions have been inferred by using optical line ratios diagnostic techniques. For example, anomalously strong \\sii$~\\lambda6716$/\\ha\\ and \\nii$~\\lambda6583$/\\ha, and weak \\oiii$~\\lambda5007$/\\ha\\ emission line ratios (compared to the bright, classical \\hii\\ regions) indicate a low state of excitation, with few ions present that require ionization energies greater than 23 eV \\citep{Reynolds85, HRT99, Rand97}.  This is consistent with the small value of \\hei~$\\lambda5876$/\\ha\\ near the midplane, indicating that the ionization fraction of helium is low and suggesting that the spectrum of the diffuse interstellar radiation field that ionizes the hydrogen is significantly softer than that from the average Galactic O star population \\citep{RT95, TuftePhD}. In addition, the elevated \\nii/\\ha\\ and \\sii/\\ha\\ ratios in the WIM suggest that this low density diffuse gas is significantly warmer than traditional \\hii\\ regions and may require spectral processing of the stellar radiation \\citep[e.g.][]{WM04} and/or an additional heating source beyond photoionization \\citep{RHT99}. Recent observations of \\niiblue/\\nii$~\\lambda6583$ have indeed confirmed that the WIM is about 2000 K warmer than \\hii\\ regions \\citep{Reynolds+01}.  Below we present new WHAM observations of \\nii$~\\lambda6583$, \\sii$~\\lambda6716$, \\oiii$~\\lambda5007$, \\hei$~\\lambda5876$ and \\niiblue\\ toward large-scale emission structures as well as individual lines of sight, representing a substantial increase in the number of observations of these diagnostic emission lines in the Galaxy. This is primarily an empirical study. The emission regions examined span a wide range in location, environment, morphology, and scale, and we have compared the line intensity ratios in these different environments in order to explore the variations in physical conditions between them.  A more detailed analysis and deeper understanding of all these different regions, the relationships between them, and the reasons for the observed differences will require combining these observations with photoionization models, and is beyond the scope of this work.  We begin with an overview  of the relationship between the emission line ratios and the temperature  and ionization state of the gas in \\S\\ref{sec:physconds}. Our observational techniques and data reduction procedure are discussed in \\S\\ref{sec:obs}. In \\S\\ref{sec:hii}, we present our results for several classical O-star \\hii\\ regions that form the basis for comparison to the fainter \\ha\\ emission structures. To illustrate the general spectral difference between the \\hii\\ regions and the WIM, we present in \\S\\ref{sec:wim} spectra toward one of the \\hii\\ regions, where the diffuse gas and the classical \\hii\\ region are along the same line of sight, but at separate radial velocities. Two large bubble-shaped features that each span more than 40\\dg, the Orion-Eridanus bubble and the Perseus superbubble, are discussed in \\S\\ref{sec:bubbles} along with comparisons to the \\hii\\ regions and the WIM. Observations of high latitude filamentary structures are presented in \\S\\ref{sec:hlfil}.  A direct measure of the temperature of ionized gas, through observations of \\niiblue\\ and \\nii$~\\lambda6583$, is discussed in \\S\\ref{sec:niiblue}, followed by a summary and conclusions in \\S\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  We have presented a large number of new observations of several optical emission lines toward \\ha-emitting features in the Galaxy that span a wide range in surface brightness, angular scale, environment, and morphology. We have explored the relative intensities of these emission lines to infer the physical conditions of the emitting gas, and we have compared these conditions with those of traditional, O star \\hii\\ regions in an attempt to gain insight into the nature of the WIM and its relationship to hot stars and large-scale bubbles and filaments within the interstellar medium. We found significant variations in the temperature and ionization state among these emission features, revealing that warm ionized gas is heterogeneous in nature. We have strengthened the general assertion that the WIM is warmer and less ionized compared to classical \\hii\\ regions and found significant variations in temperature and ionization state \\emph{within} the diffuse WIM.  An overview of some of these observations is presented in Figure~\\ref{fig:alldiag}, which shows the diagnostic plots of \\nii/\\ha\\ vs. \\sii/\\ha\\ for various emission regions observed in pointed and survey mode.  The dashed vertical and solid sloped lines represent lines of constant temperature and constant S$^+$/S, respectively, and have the same values as in Figures~\\ref{fig:orirationvs}, \\ref{fig:bowtieratiolocalnvs}, and \\ref{fig:bowtieratiopernvs}. The range of physical conditions is indicated by the different distributions and by the scatter in the data points. Different types of emission regions occupy different areas of the diagram, and even within a given catagory of emission region, variations can be significant. The classical O star \\hii\\ regions (except for the two regions ionized by hot stellar cores) show the least variation (Figure~\\ref{fig:alldiag}a), with temperatures between 6000~K and 7000~K and S$^+$/S $\\approx 0.25$.  In contrast, the faint diffuse WIM (Fig.~\\ref{fig:alldiag}b) has the most variation; that is, it occupies the largest area, with temperatures ranging between 7000~K and 10,000~K and S$^+$/S ranging from about 0.1 to 1. Figure~\\ref{fig:alldiag}b also shows that the mean properties of the WIM can be slightly different in different regions of the Galaxy.  For example, in the Perseus arm S$^+$/S in the WIM has an average near 0.25, while in the more nearby gas toward that direction, its average is near 0.6.  From this study, we make the following closing statements:  1. The temperature of diffuse ionized gas is higher in regions of lower emission measure. This is implied by the elevated \\nii/\\ha\\ and \\sii/\\ha\\ ratios toward the relatively faint Perseus superbubble, high latitude filaments, and diffuse background WIM, compared to the relatively bright Orion-Eridanus bubble and the even brighter O star \\hii\\ regions. The relation between temperature (traced by \\nii/\\ha) and \\ha\\ intensity also holds within the diffuse WIM itself (e.g., Figures~\\ref{fig:oriratioplot}, \\ref{fig:bowtieratiolocal}, and \\ref{fig:bowtieratioper}). The elevated temperatures in the WIM are confirmed by the \\niiblue/\\nii$~\\lambda6583$ intensity ratios (Figure~\\ref{fig:bluenii}) and recent observations of [O~II] \\citep{Mierkiewicz+04}. \\\\  2. The ionization state in diffuse ionized gas is generally lower than that in classical \\hii\\ regions. Several new observations of \\oiii\\ and \\hei\\ indicate that, in general, the fraction of O$^{++}$/O and He$^+$/He in the WIM and in the large bubble structures is low compared to \\hii\\ regions, implying a lower ionization state due to a softer ionizing radiation field and probably a lower ionization parameter (e.g., Figures~\\ref{fig:bowtiediag} and \\ref{fig:nfildiag}). The high ionization of sulfur (i.e., the low values of S$^{+}$/S) in the WIM of the Perseus arm (Figs. 15, 16, 21b) appears to be an exception to this trend. The data also suggest that the diffuse low density gas close to the Galactic plane may be more highly ionized than gas at larger distances from the plane. In the inner Galaxy, this trend appears to reverse \\citep{MR05}. \\\\  3. Conditions within the WIM vary significantly. The mean temperature and ionization state can change considerably from one sight line to the next and even along a single sight line.  Moreover, the mean properties of the WIM change from one region of the Galaxy to another (e.g., see Fig.~\\ref{fig:alldiag}b). Within the WIM, values of T and S$^+$/S extend to significantly higher values than are found within classical \\hii\\ regions or even the extended bubbles. \\\\  4. High latitude filaments superposed on the faint WIM have spectral characteristics similar to the WIM. This suggest a close relationship between these high latitude structures and the more diffuse background. \\\\   5. The Perseus superbubble provides strong evidence that a luminous O star cluster near the midplane can produce wide-spread, nearly WIM-like ionization conditions to distances of 1000 pc or more from the ionizing stars. This superbubble has spectral characteristics similar to those observed in portions of the diffuse WIM (Figs.~\\ref{fig:alldiag}b and d) and quite different from the spectral characteristics of the bright \\hii\\ region that immediately surrounds its source of ionization, the Cas OB6 star cluster and other classical \\hii\\ regions (Fig.~\\ref{fig:alldiag}a and d; Table~\\ref{tab:hiimulti}). The fact that the smaller, brighter, denser Orion-Eridanus bubble has low \\nii/\\ha\\ ratios, similar to those in classical \\hii\\ regions, implies that bubble size, gas density within the ionized shell, and/or the flux and spectrum of the radiation escaping O star clusters may be important in setting the conditions within the ionized gas. For example, the presence of an extra source of non-ionizing heat can raise the gas temperature in low density ($\\lesssim 0.1$ cm$^{-3}$) ionized gas \\citep{RHT99}, which could account for the elevated, more WIM-like temperatures in the low density Perseus superbubble, as well as the more general relationship between temperature and emission measure (point 1 above). \\\\   6. Filamentary structures within the Perseus superbubble have physical conditions that differ from conditions in the fainter, more diffuse parts of the bubble. This implies that these filaments are discrete entities, likely regions of higher density, and not just directions of increased pathlength through folds or edge projections of a shell or sheet. The slightly depressed \\nii/\\ha\\ and \\sii/\\ha\\ ratios in these filaments, compared to the fainter superbubble emission adjacent to them, suggest that they are in fact cooler than the gas along adjacent sightlines through the superbubble. \\\\   In conclusion, we believe that high spectral resolution emission line observations at visible wavelengths open a new window on the study of interstellar matter and processes not available through other techniques at other wavelengths. In particular, the application of nebular line diagnostics to the study of the warm ionized component of the interstellar medium provides an opportunity to understand better, through both observations and modelling, the large-scale effects of hot stars on the ionization and morphology of the interstellar medium within the Galaxy's disk and halo."
    },
    "0609/astro-ph0609072_arXiv.txt": {
        "abstract": "{% The Alpha Magnetic Spectrometer (AMS), to be installed on the International Space Station, will provide data on cosmic radiations in a large energy range. The main physics goals in the astroparticle domain are the antimatter and the dark matter searches.  Dark matter should be composed of non baryonic weakly interacting massive particles, a good candidate being the lightest SUSY particle in R-parity conserving models.  As a prototype for the AMS-02 experiment, the AMS-01 particle spectrometer was flown on the Space Shuttle Discovery in near earth orbit for a ten day mission in June 1998. The direct identification of positrons in AMS-01 was limited to energies below 3 GeV due to the vast proton background and the characteristics of the subdetectors, but the sensitivity towards higher energies (up to 40 GeV) was extended by identifying positrons through the conversion of bremsstrahlung photons.  AMS-02 will greatly improve the accuracy on the positron spectrum, which will be measured up to 300 GeV, together with the antiproton and $\\gamma$-ray flux, thus providing a unique chance to measure all relevant neutralino decay channels with the same experiment. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609244_arXiv.txt": {
        "abstract": "We present narrow-band images of the Red Rectangle (RR) nebula which reveal the distinct morphologies of this intriguing nebula in different optical emission bands. The morphology of the RR nebula in blue luminescence (BL) and extended red emission (ERE) are almost mutually exclusive. We also present the optical detection of the circum-binary disk of the RR in the light of the BL. The total intensities from the two optical band emissions (BL and ERE) when summed over the nebula are of comparable magnitude. Their spatial distributions with respect to the embedded illumination sources lead us to suggest that they may be attributed to different ionization stages of the same family of carriers. ",
        "introduction": "The optical spectrum of the Red Rectangle (RR) nebula consists of several components  of comparable intensities \\citep{schmidt80,wb90,vijh05}, and it is possible to use narrow-band imaging to isolate the morphological substructures where particular emissions are dominant. This is in contrast to most ordinary reflection nebulae, in which the scattered light component dominates \\citep{wb90}. The special status of the RR is a consequence of the optically-thick circumstellar disk surrounding the central illuminating source and seen edge-on, which prevents forward-directed scattered radiation from reaching the observer. Thus, it is the suppression of scattered radiation by at least one order of magnitude that makes blue luminescence (BL) and extended red emission (ERE) readily observable in the RR. The latter two radiation processes are most likely the result of isotropic emission processes, which are readily observable from nebular locations where the scattered light is preferentially directed away from the observer on account of the strongly forward-throwing scattering phase function of dust at optical wavelengths.  The spectral components of the optical light of the RR are: a)~Dust-scattered radiation: to first order, this spectrum is a somewhat bluer version of that of the more luminous star in the central binary system, HD~44179, a T$_\\mathrm{eff}\\sim$8250~K A-giant AGB star. b)~Blue luminescence (BL) band: this emission band with a peak near 380~nm (FWHM $\\sim$ 45~nm) has been attributed to fluorescence by neutral PAH molecules with 14 - 18 C~atoms. \\citep{vijh04, vijh05} c)~Extended red emission (ERE): this band peaks near 670~nm (FWHM $\\sim$ 180~nm)  and it is the primary cause for the unusual red color of the RR. Although observational constraints have indicated that the ERE carrier requires photons with E~$\\geq$~10.5 eV for their initiation, indicating a possible association with doubly ionized molecular species \\citep{witt06}, the carrier of the ERE is as yet unidentified \\citep{wv04}. d)~Sharp red emission features: this is a set of relatively narrow emission bands due to  as yet unidentified molecules. They appear in a wavelength range largely coincident with that of the ERE band \\citep{vanWinckel02}. It has also been  suggested that some of these bands are emission counterparts of corresponding  diffuse interstellar bands (DIBs) \\citep{scarrott92}. Spectroscopic studies have demonstrated that the intensities of the emission bands b),  c), and d) are comparable to those of the underlying scattered light continuum, allowing their relative distributions to be traced. This makes it possible to constrain the physical state(s) of the respective emitters, because the  almost-edge-on geometry of the bi-polar outflow allows one to associate specific  physical conditions with regard to density, radiation field, and ionization equilibrium with  specific nebular regions.  Our imaging study of the RR differs in significant aspects from previous efforts. Most earlier work was focused on sub-arcsecond resolution at near- and mid-infrared wavelengths with the aim to resolve the core of the RR \\citep{roddier95,cruzalbes96,mekarnia98,tuthill02,miyata04}, which resulted in the well-resolved imaging of the circumstellar disk and brightness distribution of the RR in the bright innermost region. The more extended brightness distribution of the RR was studied at high resolution by \\citet{cohen04}, using WFPC2 on HST. Their band-passes were limited by the filters available in WFPC2, with the F467M filter defining their shortest-wavelength band. Thus, they were unable to address the spatial distribution of the BL, which peaks below 400 nm in wavelength. They were, however, able to demonstrate the distinctly different spatial distribution of the dust-scattered radiation in the RR, which dominates the F467M band, and of the ERE, which was imaged in three different red bands (F588N, F622W, F631N). They also produced the best-resolved record of the ladder structure in the ERE crossing the outflow cones of the RR. Another important mapping effort is that of \\citet{breg93}, who imaged the extended brightness distribution of the RR in the two aromatic emission features at 3.3 \\micron\\ and 11.3 \\micron. They demonstrated that the 11.3 \\micron\\ brightness distribution followed the extended bipolar structure of the RR, while the 3.3 \\micron\\ isophotes were more concentrated near the center of the nebula. Our goal was to pursue the large-scale brightness distribution in the RR of specific optical band emissions, which we attempted to isolate with narrow-band filters and still narrower Fabry-Perot bandpasses, while the angular resolution was limited by seeing ($\\sim 1$ \\as). Our work includes, therefore, a first effort to image the BL in the RR. ",
        "conclusions": "\\begin{enumerate} \\item The morphology of the BL and ERE emissions in the RR nebula are almost mutually exclusive. The ERE is strongly enhanced along the walls and in the interior of the bipolar outflow cones of the RR nebula, while the BL is strongest in those regions that are located in the shadow of the circumstellar disk surrounding the central source, HD 44174.  \\item The presence of a circumstellar disk, viewed edge-on, divides the space around HD 44179 into two environments characterized by radiation fields of greatly different intensity and hardness. This leads us to suggest that the ERE and the BL arise from different ionization stages of the same family of carriers, such as PAH molecules. In case of PAH molecules, a large range of sizes and species are probably present. The BL preferentially detects small, neutral molecules while the mid-IR aromatic emission features arise from both neutral and ionised species. The ERE, if due to PAHs, arises from doubly ionized molecules with masses $>$ 200 amu.  \\item  Our observations at 3934~\\AA\\, near the peak of the BL band, show an intensity distribution elongated in E-W direction, with a maximum at the position of the circumstellar disk of HD 44179. If this structure is associated with the outer portions of the disk, its alignment differs from that of the inner disk observed in absorption at near-IR wavelengths and the disk seen in CO emission by 11 degrees.  \\item Our observations add to the growing evidence for greater chemical complexity of the RR nebula than previously proposed. \\end{enumerate}"
    },
    "0609/astro-ph0609299_arXiv.txt": {
        "abstract": "We present the first mid-infrared spectra of SMP~LMC\\,11 in the Large Magellanic Cloud. While this object resembles a planetary nebula in the optical, its infrared properties are more similar to an object in transition from the asymptotic giant branch to the planetary nebula phase.  A warm dust continuum dominates the infrared spectrum. The peak emission corresponds to a mean dust temperature of 330~K. The spectrum shows overlapping molecular absorption bands from 12 to 17\\,$\\mu$m corresponding to acetylene and polyacetylenic chains and benzene.  This is the first detection of C$_4$H$_2$, C$_6$H$_2$, C$_6$H$_6$ and other molecules in an extragalactic object.  The infrared spectrum of SMP~LMC~11 is similar in many ways to that of the pre-planetary nebula AFGL~618. The IRS spectrum shows little evidence of nitrogen-based molecules which are commonly seen in Galactic AGB stars. Polycyclic aromatic hydrocarbons are also absent from the spectrum.  The detection of the \\ion{[Ne}{2]}\\,12.8 $\\mu$m line in the infrared and other forbidden emission lines in the optical indicates that an ionized region is present. ",
        "introduction": "The transition from the asymptotic giant branch (AGB) to the planetary nebula (PN) phase remains an elusive stage in stellar evolution.  At the end of its lifetime on the AGB, a star loses mass at high rates, triggering the formation of molecules and dust in a thick circumstellar envelope \\citep[e.g.][]{Habing96,Blommaert:ISOlegacy}. Radiation pressure on the dust and coupling between the dust and gas drive the outflows, stripping the star and exposing its hot core.  As the effective temperature of the exposed core rises, it will ionize the nebular gas.  The detailed chemistry involved in this brief transition phase is poorly understood.  In carbon-rich stars, carbon is primarily incorporated first into carbon monoxide (CO) and then into acetylene (C$_2$H$_2$).  Acetylene and its derivatives are the precursors from which carbon-based compounds such as polycyclic aromatic hydrocarbons (PAHs) and soot are formed \\citep[][]{atb89,ff89}.  Here, we present the infrared spectrum of SMP~LMC~11, an object in the Large Magellanic Cloud (LMC) whose infrared spectral characteristics fall in the stage where most of this complex chemistry takes place.  SMP~LMC~11 was part of the \\citet{san} survey of PNe in the Magellanic Clouds (from which it takes its designation, SMP~LMC~11).  It was also included in the catalogue by \\citet{lei} of accurate positions of known PNe in the LMC.  The reported H$\\beta$ flux and extinction vary among papers in the literature. For instance, \\cite{lei2}, \\cite{woo} and \\cite{mea} report a log(H$\\beta$)=$-$13.14 (with H$\\beta$ in units of erg cm$^2$ s$^{-1}$), whereas \\citet{sha2} obtain log(H$\\beta$)=$-$13.94.  Similarly, differences in visual extinction amount to a factor of more than 3.  The presence of another source nearby in the IRAC images (see \\S2) may explain these discrepancies. \\citet{mor} presented optical spectroscopy of 97 faint PNe including SMP~LMC~11 and identified the forbidden emission lines of \\ion{[O}{3]}, \\ion{[N}{2]} and \\ion{[S}{2]}. Recently \\citet{sha2} observed it with the {\\it Hubble Space Telescope} ({\\it HST}) as part of their sample of objects in the LMC and SMC. They measured the strength of the \\ion{[O}{1]} line, along with the previously identified forbidden lines, and inferred a bipolar morphology with a size of 0.76\\arcsec$\\times$0.55\\arcsec.  In this letter we report the first mid-infrared spectrum of SMP LMC\\,11.  The spectrum reveals a wealth of molecular absorption bands and a few atomic emission lines.  To our knowledge, the only other object showing a similar infrared spectrum is the Galactic pre-planetary nebula\\footnote{We adopt the term ``pre-planetary nebula'' to avoid the confusing term ``protoplanetary nebula,'' which could imply an object evolving into a PN or a nebula with embedded protoplanets.}  (PPN) AFGL~618.  \\citet{buj} inferred that AFGL~618 became a PPN only 200\\,yr ago. \\cite{chi98} identified simple aliphatics (CH$_2$, CH$_3$) in its spectrum at 3.4 $\\mu$m. \\cite{cer01a} studied the spectrum of AFGL~618 from the Short-Wavelength Spectrometer \\citep[SWS,][]{deg}  on the {\\it Infrared Space Observatory} \\citep[{\\it ISO,}][]{kes}, and they describe the system as having a thick molecular envelope surrounding a B0 star and an ultra-compact H\\,II region.  They suggest that UV radiation from the central star and shocks in the high-velocity winds have significantly modified the chemistry.  \\cite{cer01a} detected C$_4$H$_2$, C$_6$H$_2$ and of C$_6$H$_6$ in a circumstellar environment for the first time.  In this letter we compare the observed features in SMP~LMC~11 to those of AFGL~618.  The next section describes the observations and the data reduction techniques.  \\S 3 gives the analysis of the observed spectral features and compares SMP~LMC~11 to AFGL~618.  \\S 4 discusses the evolutionary status of the object. ",
        "conclusions": "The infrared spectrum of SMP~LMC~11 reveals three components: a cool dust continuum, absorption bands from hydrocarbon molecules, and emission from at least one low-ionization forbidden line.  The mean temperature of the dust as inferred from the peak of the IR emission ($\\sim$330~K) falls closer to the most evolved AGB stars (T$\\gtrsim$300~K) than the youngest PNe (T$\\sim$150~K) \\citep{kwo}. The peak of the dust emission will shift to longer wavelengths as dust moves further away from the central star.  The {\\it Spitzer} spectrum of SMP~LMC~11 represents the first extragalactic detection of several molecules, including di-acetylene (C$_4$H$_2$), tri-acetylene (C$_4$H$_2$), benzene (C$_6$H$_6$), HC$_3$N, C$_2$H$_4$, and possibly propyne.  These molecules are the building blocks from which more complex hydrocarbons are produced. The absorption bands are due to the increase in thickness of the circumstellar shell, which is more dramatic at the very end of the AGB phase \\citep{gar}.  For free molecules, this absorption will not be a Gaussian or a Lorentzian, but will be a complex addition of many individual ro-vibrational lines with different strengths (which depend on the temperature), causing the strong absorption we see in the 12-17\\,$\\mu$m region (due to the P and R branches of the molecules present in that region).  The stronger absorption in this region compared to AFGL~618 indicates that SMP~LMC~11 has a larger molecular column density.  The spectrum of SMP~LMC~11 does not contain any PAH emission features. A mixture of PAHs and the PAH-precursor C$_2$H$_2$ are seen together in the low-resolution IRS spectrum of the post-AGB object MSX SMC\\,029 \\citep{kra}.  The absence of PAH emission in SMP~LMC~11 can be attributed to age or orientation effects.  \\cite{cer01a} propose that PAHs may not form until the end of the AGB phase, based on their observations of AFGL 618.  Alternatively, the dust in SMP~LMC~11 could be distributed in an edge-on torus while MSX SMC\\,029 is seen more pole-on or has a more spherically symmetric dust distribution.  The PAHs emit in photo-dissociation regions, and in an edge-on torus, we would not have a clear line of sight to the emitting region. The dust envelope is still too thick to allow a direct line of sight into the hotter central regions.  Despite the detection of HC$_3$N (and perhaps HC$_5$N), the spectrum of SMP~LMC~11 shows evidence for reduced abundances of nitrogen-based molecules such as HCN and NH$_3$ compared to Galactic objects like AFGL~618.  This agrees with the nitrogen abundance derived by \\cite{lei2} for SMP~LMC~11 ($\\log(N/H)+12=7.1$) which is one of the lowest in their sample of LMC PNe. This result fits with the finding of \\citet{mat} that the HCN bands that are prominent in the Galactic sample of carbon stars are significantly weaker or absent in the LMC sample.  The lower nitrogen abundance in the LMC compared to the Galaxy and higher efficiency of the third dredge-up at lower metallicities work together to decrease the nitrogen abundance relative to carbon.  The presence of an optically thick and relatively warm dust component between us and the central ionized region can explain not only the absence of PAH features, but also the absence of most ionized lines in the infrared.  We only see [Ne II] (and maybe \\ion{[Ne}{3]}), and this line is probably visible only in scattered light.  Scattering is more efficient in the optical than in the infrared, which would explain why more of the forbidden lines can be detected in the optical.  The dust also explains the presence of molecular absorption in the spectrum, since the molecules do not have a direct line of sight to the still shielded central ionized region.  It is also possible that some of the forbidden line emission is excited by high-velocity (v$>$100 km s$^{-1}$) shocks \\citep{lei2}.  All of these properties are more consistent an object which is in transition from the AGB to a PN than with a classic PN.  The bright H$\\beta$ luminosity and its large size, present some problems with this scenario of a transition object.  While the measured H$\\beta$ flux by \\cite{sha2} makes SMP~LMC~11 one of the faintest PNe in their LMC sample, its equivalent flux at 1 kpc (if we assume that it is at a distance of 50 kpc) is log(H$\\beta$)=$-$10.55, which is bright for Galactic standards.  Similarly, \\cite{sha2} state that it is one of the smallest bipolar nebulae in their LMC sample, but at a Galactic distance of 1 kpc it would be resolved. These characteristics resemble those of a typical PN. One could invoke then a massive molecular envelope surrounding the ionized region to explain the absorption features.  Still, the absence of PAHs and H$_2$ in the spectrum, and the high temperature of the dust (compared to PNe) would remain unexplained (unless perhaps invoking a complex geometry).   The rich chemistry seen in the IRS spectrum of SMP~LMC~11 makes it (together with AFGL~618, and MSX SMC\\,029) a key object in our understanding of the formation of hydrocarbons, deserving future study from the community. Particularly, CO measurements of the circumstellar AGB envelope may prove to be particularly useful."
    },
    "0609/astro-ph0609250_arXiv.txt": {
        "abstract": "We present the Boltzmann moment equation approach for the dynamics of stars (BEADS-2D), which is a finite-difference Eulerian numerical code designed for the modelling of {\\it anisotropic and non-axisymmetric} flat stellar disks. The BEADS-2D code solves the Boltzmann moment equations up to second order in the thin-disk approximation. This allows us to obtain the anisotropy of the velocity ellipsoid and the vertex deviation in the plane of the disk.   We study the time-dependent evolution of exponential stellar disks in the linear regime and {\\it beyond}. The disks are initially characterized by different values of the Toomre parameter $Q_{\\rm s}$ and are embedded in a dark matter halo, yielding a rotation curve composed of a rigid central part and a flat outer region. Starting from a near equilibrium state, several unstable modes develop in the disk. In the early linear phase, the very centre and the large scales are characterized by growing one-armed and bisymmetric positive density perturbations, respectively. This is in agreement with expectations from the swing amplification mechanism of short-wavelength trailing disturbances, propagating through the disk centre. In the late linear phase, the overall appearance is dominated by a two-armed spiral structure localized within the outer Lindblad resonance (OLR).  During the non-linear evolutionary phase, radial mass redistribution due to the gravitational torques of spiral arms produces an outflow of mass, which forms a ring at the OLR, and an inflow of mass, which forms a transient central bar. This process of mass redistribution is self-regulatory and it terminates when spiral arms diminish due to a shortage of matter. Finally, a compact central disk and a diffuse ring at the OLR are formed. An increase in $Q_{\\rm s}$ stabilizes the disks at $Q_{\\rm s}\\approx 3.15$, in agreement with the theoretical predictions.  Considerable vertex deviations are found in regions with strongly perturbed mass distributions, i.e.\\ near the spiral arms.  The vertex deviations are especially large at the convex edge of the spiral arms, whereas they are small at the concave edge.  The mean vertex deviations correlate well with the global Fourier amplitudes, reaching mean values of about 12$^\\circ$  in the saturation stage. Local values of the vertex deviation can reach up to almost 90$^\\circ$. Near the convex edge of the spiral arms, the ratio of radial to azimuthal components of the velocity ellipsoid can deviate considerably from the values predicted from the epicycle approximation. ",
        "introduction": "The majority of normal disk galaxies are characterized by non-axisymmetric structures like spirals or bars. These structural elements have been widely discussed in the literature as a result of gravitational instabilities which are connected to growing density waves or global instabilities of disks (e.g.\\ Binney \\& Tremaine 1987, afterwards BT87).  A first insight into the properties of galactic discs was provided by linear stability analysis. Structure growth in a stellar disk within the linear regime was studied via the linearised collisionless Boltzmann equation by Toomre (\\cite{toomre64}), Zang (\\cite{Zang}),  Pichon \\& Cannon (\\cite{Pichon}), Evans \\& Read (\\cite{Evans}), Jalali \\& Hunter (\\cite{Jalali}), and others. Toomre (\\cite{toomre64}) has deduced a simple stability criterion from the Jeans equations in case of stellar disks: a stellar disk is stable against axisymmetric perturbations if $Q_{\\rm s} \\equiv \\sigma_{rr} \\kappa / (3.36 G \\Sigma)$ exceeds a critical value $Q_\\mathrm{s,crit}=1$ ($\\sigma_{rr}$ is the radial velocity dispersion, $\\kappa$ is the epicycle frequency, $G$ is the gravitational constant, and $\\Sigma$ is the surface density of the disk). In case of non-axisymmetric perturbations the analysis becomes technically more difficult, but the Toomre parameter turned out to still be a good indicator for stability, if one increases its critical value (Bertin et al.\\ \\cite{bertin89}, Polyachenko et al.\\  \\cite{polyachenko97b}).   A disadvantage of linear stability analysis is its restriction to small perturbations, both in amplitude and wavelength. Once the perturbations exceed a critical level or when the initial perturbations are already large, numerical simulations are necessary. Hydrodynamical simulations revealed three qualitatively distinct phases for the growth of unstable disks (Laughlin et al.\\ \\cite{laughlin97}). Starting from small amplitudes, perturbations grow first linearly. When they reach a critical level, mode-mode coupling becomes important and the non-linear evolution starts. The self-interaction of the dominant mode might result in a rearrangement of the radial mass distribution ($m=0$ mode). The disk itself might become violently unstable, which is indicated by the perturbations reaching a ``macroscopic'' level. Finally, a quasi-equilibrium saturation stage can be reached when the growth of the instabilities is stopped due to shocks or feedback mechanisms heating up the disk dynamically. Otherwise, fragmentation continues until the disk is destroyed or transformed.   Though N-body simulations have been proven to be extremely useful for many studies of galactic dynamics, they are less suitable for studying the growth of very weak perturbations. The main reason is that N-body calculations of galaxies are usually performed with less particles than the actual number of stars in these systems. This causes an inherent artificial particle noise which cannot be neglected for typical particle numbers.  Alternatively, numerical hydrodynamics simulations became a primary tool for the analysis of growing spiral instabilities (e.g. Korchagin et al.\\ \\cite{korchagin00}, Orlova et al.\\ \\cite{orlova02}). First, they allow for a direct comparison with the results of a semi-analytic linear stability analysis, including predictions for the growth rates (Laughlin et al. \\cite{laughlin97,laughlin98}, Korchagin et al.\\ \\cite{korchagin00}), because the same set of basic equations is applied. Second, the standard grid methods used to solve the hydrodynamical equations allow to start with very small perturbations (basically limited by machine accuracy). Their evolution can be followed over many $e$-folding times until they reach the non-linear regime and beyond.  A major drawback of purely hydrodynamical simulations is that the dominant disk component is the stellar disk. The latter is characterized by an {\\it anisotropic velocity dispersion}, whereas the hydrodynamical equations deal with an {\\it isotropic pressure}. Secondly, one has to assume an equation of state connecting pressure, (surface) density and temperature (velocity dispersion). For instance, if one assumes a polytropic-like equation of state, i.e.\\ $P = C \\, \\Sigma^\\gamma$, the set of hydrodynamic equations is closed with the Euler equation describing the momentum transport (for an example of a stability analysis based on such assumptions see e.g.\\ Aoki et al.\\ \\cite{aoki79}). Since the temperature does not show up in the polytropic equation of state, no energy transport equation is required. Obviously, this is a convenient simplification which, however, might not be generally applicable. Moreover, the profile of the Toomre parameter is completely fixed by the surface density distribution.  A more complete ansatz for describing stellar disks is the collisionless Boltzmann equation (BT87) \\be \\label{eq_boltzmann} \\pdert{f} + \\bl{v} \\cdot \\pder{f}{\\bl{r}} + \\dot{\\bl{v}} \\cdot \\pder{f}{\\bl{v}} = 0, \\ee where $f(\\bl{r},\\bl{v},t)$ is the distribution function. The term $\\dot{\\bl{v}}$ includes accelerations like gravitational forces, pressure terms, etc.  Since a direct solution is already numerically very difficult, the problem is often reduced by taking the velocity moments of equation~(\\ref{eq_boltzmann}).  By this, one gets an infinite series of moment equations where the equations for moment $i$ include moments of the order $i+1$.  The knowledge of all velocity moments is equal to the knowledge of the distribution function itself. Practically, one has to terminate the set of equations by a closure condition. Setting all moments of third order to zero (zero-heat-flux condition) results in the Jeans equations which are well applicable to stellar systems with negligible two-body relaxation.  In that case, the second order moment equations which yield information on the velocity ellipsoid are included and stellar anisotropy can be considered.  An example are the one- and two-dimensional chemo-dynamical models by Theis et al.\\ (\\cite{theis92}) and Samland et al.\\ (\\cite{samland97}).  For the special case of galactic disks, Amendt \\& Cuddeford (\\cite{amendt91}) and Cuddeford \\& Amendt (\\cite{cuddeford91}) studied the higher order moments in detail. They found that for reasonable constraints to the distribution function the third order terms vanish to leading order in the plane of the disk. This corroborates the assumption of a zero-heat-flux. Extensions of the Boltzmann moment equations including collisional processes have also been used (e.g.\\ Larson \\cite{larson70}, Louis \\cite{louis90} or Giersz \\& Spurzem \\cite{giersz94} for the evolution of star clusters).     In this paper we present the Boltzmann equation approach for the dynamics of stars (BEADS-2D), which is a numerical code that solves the Boltzmann moment equations up to second order using the methods of finite-differences. The BEADS-2D code is applied to study the time-dependent evolution of flat, non-isotropic stellar disks in the linear regime and {\\it beyond}. In contrast to previous numerical studies, we include the {\\it non-axisymmetric} velocity dispersion terms into the Boltzmann moment equations. Therefore, our approach allows us to study the evolution of the anisotropy and the vertex deviation of the stellar component. We compare the stability properties of non-isotropic stellar disks with the predictions of linear stability analysis for non-isotropic disks by Polyachenko et al. (\\cite{polyachenko97b}). We provide a physical interpretation for growing instabilities in our model stellar disks.   In the following section we describe the Boltzmann moment equation approach. The basic principles of the BEADS-2D code are discussed in \\S~\\ref{sect_numericalmodelling}. The results of numerical modelling  are presented in \\S~\\ref{sect_results}. Finally, conclusions and a summary are given in \\S~\\ref{sect_summary}. ",
        "conclusions": "\\label{sect_summary}   We have developed the Boltzmann moment equation approach for the dynamics of stars (BEADS-2D), which is a finite-difference Eulerian code based on the Boltzmann moment equations up to second order. We adopt the zero-heat-flux approximation to close the system of moment equations. The BEADS-2D code is formulated in the polar coordinates $(r,\\phi)$, which makes it most suitable for the numerical modelling of flattened galactic stellar disks. By calculating the tensor of the velocity dispersions, we can follow the anisotropy and the vertex deviation of the stellar system. Compared to numerical codes assuming a polytropic equation of state (see e. g. Korchagin et al. \\cite{korchagin00}), the BEADS-2D code allows for a larger variety of simulations, because the mass and velocity dispersion profiles can be chosen independently. For the reader's convenience, we provide the full set of three-dimensional Boltzmann moment equations up to second order in Appendix~A . These equations are formulated in the cylindrical coordinates $(z,r,\\phi)$ using the usual zero-heat-flux approximation and are suitable for a global modelling of galactic disks without simplifying geometry assumptions.  As an example of the utility of the BEADS-2D code, we study the time-dependent evolution of exponential stellar disks characterized initially by different values of the Toomre parameter $Q_{\\rm s}$. The disks are embedded in a static dark matter halo, yielding a nearly flat rotation curve (except for the innermost several kiloparsecs characterized by rigid rotation). Specifically, we find the following.   (i) We confirm the results of linear stability analysis by Polyachenko et al. (\\cite{polyachenko97b}), who obtained a stability threshold value of $Q_{\\rm s} \\approx 3.15$ for stellar disks with a {\\it purely} flat rotation curve. A physical interpretation for the instability is swing amplification of short-wavelength trailing disturbances, which propagate through the disk centre and provide a feedback mechanism for the swing amplifier.  (ii) The character of the instability is distinct in the linear and nonlinear regimes, the tentative boundary between which we define by the global Fourier amplitude $C_2(t)=0.1$. As a reference model, we investigated a stellar disk with $Q_{\\rm s}=1.3$. In the early linear phase, instability demonstrates a typical lumpy pattern, which grows in amplitude with time and gradually transformes in to a global two-armed spiral in the late linear phase. This lumpy pattern results most probably from the interference between leading and trailing disturbances, propagating through the disk centre. A similar phenomenon is discussed in Binney \\& Tremaine (\\cite{binney87}}) in connection with the numerical simulations of Toomre (\\cite{toomre81}). In the nonlinear phase, the redistribution of mass and angular momentum by the gravitational torques of spiral arms comes into play. The gravitational torques produce an inflow of mass, which forms  a transient central bar inside corotation, and an outflow of mass, which forms a ring at the OLR. This process of mass redistribution continues until the two-armed spiral is weakened by the shortage of matter. Finally, a transient bar is transformed into the compact dense central disk and the spiral arms virtually disappear, leaving only the outer diffuse ring at the OLR.   (iii) The instability in disks with $Q_{\\rm s} \\la 2.5$ is mostly dominated by the symmetric $m=2$ mode, while the lopsided $m=1$ mode in disks  with $Q_{\\rm s}> 2.5$ may compete with or even prevail over the $m=2$ mode. This is due the fact that the $m=1$ mode has no inner Lindblad resonance, which makes it difficult to stabilize. On the other hand, the stabilizing effect of the inner Lindblad resonance on the $m=2$ mode is expected to grow along the sequence of increasing $Q_{\\rm s}$.   (iv) Stellar disks near a stability limit of $Q_{\\rm s}=3.15$ still  develop a clumpy structure.  However, both the growth timescales (comparable to the Hubble time) and the small amplitudes of positive density perturbations make such instabilities difficult to observe at best.  (v) Stellar disks are stabilized by one or more of the following procedures. a) An increase in the stellar velocity dispersions suppresses the growth of short-wavelength density perturbations and thus inhibits the swing amplifier.  b) Slower pattern speeds in hotter disk make it more likely for the inner Lindblad resonance to occur, thus providing the cutoff for the sing amplification feedback.    The presence of the non-diagonal velocity dispersion terms in the Boltzmann moment equations allows us to study the vertex deviation. In most regions of the stellar disk the velocity ellipsoid is well aligned with the radial axes of the polar grid. In agreement with analytical predictions by Kuijken \\& Tremaine (\\cite{KT}), considerable vertex deviations show up in regions with strongly perturbed mass distributions, i.e.\\ near the spiral arms.  The vertex deviations are large at the convex edge of the spiral arms, whereas they are small at the concave edge.  The mean vertex deviations correlate well with the global Fourier amplitudes. Near the convex edge of the spiral arms, the ratio of radial to azimuthal components of the velocity ellipsoid can deviate considerably from the values predicted from the epicycle approximation. A more detailed numerical study of the properties of velocity ellipsoids in spiral galaxies will be presented in a future paper."
    },
    "0609/astro-ph0609066_arXiv.txt": {
        "abstract": "Presently seven nearby radio-quiet isolated neutron stars discovered in ROSAT data and characterized by thermal X-ray spectra are known. They exhibit very similar properties and despite intensive searches their number remained constant since 2001 which led to their name ``The Magnificent Seven''. Five of the stars exhibit pulsations in their X-ray flux with periods in the range of 3.4 s to 11.4 s. XMM-Newton observations revealed broad absorption lines in the X-ray spectra which are interpreted as cyclotron resonance absorption lines by protons or heavy ions and / or atomic transitions shifted to X-ray energies by strong magnetic fields of the order of 10$^{13}$ G. New XMM-Newton observations indicate more complex X-ray spectra with multiple absorption lines. Pulse-phase spectroscopy of the best studied pulsars \\rxb\\ and \\rbs\\ reveals variations in derived emission temperature and absorption line depth with pulse phase. Moreover, \\rxb\\ shows long-term spectral changes which are interpreted as due to free precession of the neutron star. Modeling of the pulse profiles of \\rxb\\ and \\rbs\\ provides information about the surface temperature distribution of the neutron stars indicating hot polar caps which have different temperatures, different sizes and are probably not located in antipodal positions. ",
        "introduction": "\\label{intro} After the discovery of \\rxf\\ as an isolated neutron star in ROSAT data \\citep{1996Natur.379..233W,1997Natur.389..358W}, six further objects with very similar properties were found. Their X-ray emission is characterized by a soft, blackbody-like continuum which is little attenuated by photo-electric absorption by the interstellar medium, indicating small distances \\cite[of the order of a few 100 pc,][]{2006NS...Posselt}. For the brightest object, \\rxf, this was soon confirmed by the parallax measurement \\citep{2001ApJ...549..433W,2002ApJ...571..447K,2002ApJ...576L.145W}. Including new HST observations the most recent revision of the parallax yields a distance of 161$^{+18}_{-14}$ pc \\citep{2006NS...vanKerkwijk}. A preliminary value for the parallax of \\rxb\\ is presented by \\cite{2006NS...vanKerkwijk} which corresponds to a distance of 330$^{+170}_{-80}$ pc.  No indication for a hard, non-thermal X-ray emission component was found in the X-ray spectra. Compared to normal radio pulsars also no strong radio emission is detected: While for \\rxb\\ and \\rxc\\ deep radio observations failed to detect a radio counterpart \\citep{2003MNRAS.340L..43J,2003ApJ...590.1008K}, for \\rbs\\ and \\rbt\\ the detection of weak radio emission was claimed \\citep{2006ATel..798....1M,2006NS...Malofeev}. The constant X-ray flux on time scales of years, no obvious association with a supernova remnant and relatively high proper motion measurements for the three brightest objects suggest that we deal with a group of cooling neutron stars with ages of around 10$^6$ years. Recent reviews which summarize the observed properties of the seven thermally emitting INSs known today can be found in \\cite{2000PASP..112..297T}, \\cite{2001xase.conf..244M} and \\cite{2004AdSpR..33..638H,2005fysx.conf...39H}.  \\begin{table*} \\caption[]{X-ray and optical properties of the Magnificent Seven} \\centering \\label{tab-x-opt} \\begin{tabular}{lcccccc} \\hline\\noalign{\\smallskip} \\multicolumn{1}{l}{Object} & \\multicolumn{1}{c}{kT} & \\multicolumn{1}{c}{Period} & \\multicolumn{1}{c}{Amplitude} & \\multicolumn{1}{c}{Optical} & \\multicolumn{1}{c}{PM} & \\multicolumn{1}{c}{Ref.} \\\\  \\multicolumn{1}{l}{} & \\multicolumn{1}{c}{eV} & \\multicolumn{1}{c}{s} & \\multicolumn{1}{c}{\\%} & \\multicolumn{1}{c}{mag} & \\multicolumn{1}{c}{mas/year} & \\multicolumn{1}{c}{} \\\\[3pt] \\tableheadseprule\\noalign{\\smallskip} RX\\,J0420.0--5022   &  44    &  3.45 & 13     & B = 26.6\t\t&     & 1           \\\\ RX\\,J0720.4--3125   &  85-95 &  8.39 & 8-15   & B = 26.6\t\t&  97 & 2,3,4,5,6   \\\\ RX\\,J0806.4--4123   &  96    & 11.37 &  6     & B $>$ 24\t\t&     & 7,1         \\\\ RBS\\,1223$^{(a)}$    &  86    & 10.31 & 18     & m$_{\\rm 50ccd}$ = 28.6 &     & 8,9,10,11   \\\\ RX\\,J1605.3+3249     &  96    & 6.88? & ?      & B = 27.2\t\t& 145 & 12,13,14,15 \\\\ RX\\,J1856.5--3754   &  62    & $-$   & $<$1.3 & B = 25.2\t\t& 332 & 16,17,18,19 \\\\ RBS\\,1774$^{(b)}$    & 102    & 9.44  &  4     & B $>$ 26\t\t&     & 20,21,22    \\\\ \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} \\multicolumn{7}{l}{$^{(a)}$ = 1RXS\\,J130848.6+212708} \\\\ \\multicolumn{7}{l}{$^{(b)}$ = 1RXS\\,J214303.7+065419} \\\\ \\multicolumn{7}{l}{References:} \\\\ \\multicolumn{7}{l}{ (1)  \\cite{2004A&A...424..635H} (2)  \\cite{1997A&A...326..662H} (3)  \\cite{2001A&A...365L.302C} } \\\\ \\multicolumn{7}{l}{ (4)  \\cite{2004A&A...419.1077H} (5)  \\cite{2004A&A...415L..31D} (6)  \\cite{2003A&A...408..323M} } \\\\ \\multicolumn{7}{l}{ (7)  \\cite{2002A&A...391..571H} (8)  \\cite{1999A&A...341L..51S} (9)  \\cite{2002A&A...381...98H} } \\\\ \\multicolumn{7}{l}{ (10) \\cite{2002ApJ...579L..29K} (11) \\cite{2003A&A...403L..19H} (12) \\cite{1999A&A...351..177M} } \\\\ \\multicolumn{7}{l}{ (13) \\cite{2003ApJ...588L..33K} (14) \\cite{2004ApJ...608..432V} (15) \\cite{2005A&A...429..257M} } \\\\ \\multicolumn{7}{l}{ (16) \\cite{1997Natur.389..358W} (17) \\cite{2002ApJ...576L.145W} (18) \\cite{2003A&A...399.1109B} } \\\\ \\multicolumn{7}{l}{ (19) \\cite{2001A&A...378..986V} (20) \\cite{2001A&A...378L...5Z} (21) \\cite{2005ApJ...627..397Z} (22) \\cite{2006NS...Komarova} } \\\\ \\end{tabular} \\end{table*}  The discovery of the seven neutron stars (which are often called the Magnificent Seven, hereafter M7) with purely thermal X-ray spectra raised wide interest by theoreticians and observers as promising objects to learn about atmospheres and the internal structure of neutron stars \\citep[e.g.][]{1997ApJ...476L..47P}. The apparent absence of non-thermal processes which would hamper the analysis of the X-ray emission allows a direct view onto the stellar surface. Limits on the radius of the star \\rxf\\ were used to constrain the equation of state of neutron star matter \\citep{2002ApJ...564..981P,2004NuPhS.132..560T}. The ROSAT PSPC spectra with low energy resolution were consistent with Planckian energy distributions with blackbody temperatures kT in the range 40 -- 100 eV and little attenuation by interstellar absorption. More recent high resolution observations of the two brightest objects \\rxf\\ \\citep{2001A&A...379L..35B} and \\rxb\\ \\citep{2001A&A...365L.298P,2002nsps.conf..273P,2003ApJ...590.1008K} were performed using the low energy transmission grating (LETG) aboard Chandra and the reflection grating spectrometers (RGS) of XMM-New\\-ton. In particular the high signal to noise LETG spectrum from a 500 ks Chandra observation of \\rxf\\ did not reveal any significant deviation from a blackbody spectrum \\citep{2001A&A...379L..35B,2003A&A...399.1109B}. The statistical quality and energy band coverage of the RGS and LETG spectra of \\rxb\\ are however insufficient to detect subtle narrow features in the spectrum \\citep{2001A&A...365L.298P}.  Four of the M7 exhibit clear pulsations in their X-ray flux, indicating the spin period of the neutron star (\\rxa, \\rxb, \\rxc\\ and \\rbs) and there is evidence that also \\rbt\\ is a pulsar \\citep{2005ApJ...627..397Z}. A possible candidate period for \\rxe\\ needs future confirmation. A summary with blackbody temperatures derived from (phase-averaged) EPIC-pn spectra, spin periods, amplitudes for the flux modulation, optical brightness and measured proper motions is given in Table~\\ref{tab-x-opt} together with key references for each object. Despite extensive searches in ROSAT data \\citep[e.g. ][]{2003ApJ...598..458R,2005A&A...444...69C,2006AJ....131.1740A} no further neutron star with similar characteristic properties was found.  The first significant deviations from a blackbody spectrum were reported by \\cite{2003A&A...403L..19H} from EPIC spectra of \\rbs. A broad absorption feature near 300 eV was interpreted as proton cyclotron resonance absorption. Apart from this absorption line with an equivalent width of 160 eV the X-ray spectrum of \\rbs\\ is blackbody-like. \\cite{2003A&A...403L..19H} also suggested that changes in the soft part of the X-ray spectrum of the pulsar \\rxb\\ with pulse phase reported by \\cite{2001A&A...365L.302C} are caused by variable cyclotron absorption \\citep{2004A&A...419.1077H}. Cyclotron resonance absorption features in the 0.1--1.0 keV band are expected in spectra from strongly magnetized neutron stars with field strengths in the range of $10^{10} - 10^{11}$ G or $2 \\times 10^{13} - 2 \\times 10^{14}$ G if caused by electrons or protons, respectively \\citep[see e.g.][]{2001ApJ...560..384Z,2002nsps.conf..263Z}. Variation of the magnetic field strength over the neutron star surface leads to a broadening of the line \\citep{2004ApJ...607..420H}.  The relatively long spin periods in the range of 3.45~s $-$ 11.37~s (the majority of radio pulsars has periods less than 1~s) were the first indicator for strong magnetic fields. If the stars were born with millisecond spin periods, age estimates from neutron star cooling curves \\citep{2006NuPhA...Page} and proper motions \\cite[tracing the neutron star back to a likely birth place,][]{2006NS...Motch} of typically 10$^6$ years require B fields of the order of 10$^{13}$ G to decelerate the rotation of the stars to their current periods. This was recently confirmed by the accurate determination of the pulse period derivative $\\dot P$ of $0.698 \\times 10^{-13}$ s s$^{-1}$ and $1.120 \\times 10^{-13}$ s s$^{-1}$ for \\rxb\\ and \\rbs, respectively \\citep{2005ApJ...628L..45K,2005ApJ...635L..65K}. In the magnetic dipole braking model this yields characteristic ages of $1.9 \\times 10^{6}$ years and $1.5 \\times 10^{6}$ years and B field strengths of  $2.4 \\times 10^{13}$ G and $3.4 \\times 10^{13}$ G, respectively. Such strong fields supports the picture in which the broad absorption lines -- if due to cyclotron resonance -- at least for \\rxb\\ and \\rbs\\ are caused by protons.  \\begin{figure*} \\centering \\hbox{ \\resizebox{0.485\\hsize}{!}{\\includegraphics[clip=,angle=-90]{p015736_pn0_bb.ps}} \\resizebox{0.485\\hsize}{!}{\\includegraphics[clip=,angle=-90]{p015736_pn0_bbgauss.ps}} } \\hbox{ \\resizebox{0.485\\hsize}{!}{\\includegraphics[clip=,angle=-90]{p015736_pn0_bb2gauss.ps}} \\resizebox{0.485\\hsize}{!}{\\includegraphics[clip=,angle=-90]{p015736_pn0_bb3gauss.ps}} } \\caption{Modeling of the EPIC-pn spectra of \\rxe. From top left to bottom right the number of absorption lines increases from zero to three. As continuum a blackbody model with interstellar absorption is used.} \\label{fig-spectra} \\end{figure*} ",
        "conclusions": "The new results from optical and X-ray observations of the Magnificent Seven, thermally emitting isolated neutron stars, strongly support the model of cooling neutron stars with ages around a million years. Large proper motions measured for the three brightest objects make accretion from the interstellar medium to re-heat the neutron star very inefficient. Absorption features in the X-ray spectra of most of the Magnificent Seven and first measurements of the spin period derivative from the two pulsars \\rxb\\ and \\rbs\\ consistently point to magnetic field strengths of the order of 10$^{13}$ G to 10$^{14}$ G. This places them at the long spin period and high magnetic field end of the radio pulsar distribution. However, the $\\dot P$ measurements still indicate field strengths below those of the magnetars (Fig.~\\ref{fig-ppdot}). The discovery of a few radio pulsars with similar magnetic field strengths and long periods \\citep{2000ApJ...541..367C,2002MNRAS.335..275M,2003ApJ...591L.135M} shows that radio emission can still occur at inferred field strengths close to the ``quantum critical field\" $B_{cr}= m_e^2c^3/e\\hbar \\simeq 4.4\\times 10^{13}$ G. Therefore it remains unclear if the M7 exhibit no radio emission at all or if we do not detect them because their radio beam is very narrow due to their large light cylinder radius and therefore does not cross the Earth.  In a $\\sim$10$^{13}$ G magnetic field cyclotron resonance lines at soft X-ray energies are expected to be caused by protons or highly ionized atoms of heavy elements. The ratios of line strengths in consecutive harmonic lines scales with E$_{\\rm cyc}/(mc^2)$ \\citep{1980Ap&SS..73...33P} and for the involved high particle masses no harmonic lines should be detectable. There is now more and more evidence for additional absorption lines in the X-ray spectra of the Magnificent Seven. An explanation for the additional (or even all) lines can be atomic bound-bound or bound-free transitions. In high B fields atomic orbitals are distorted into a cylindrical shape and the electron energy levels are similar to Landau states, with binding energies of atoms strongly increased. For hydrogen in a magnetic field of the order of 10$^{13}$ G the strongest atomic transition is expected at energy E/eV $\\approx$ 75(1+0.13ln(B$_{13}$))+63B$_{13}$, with B$_{13}$ = B/10$^{13}$ G \\citep{2002nsps.conf..263Z}. For the line energies found in the spectra of thermal isolated neutron stars this requires similar field strengths to those derived assuming cyclotron absorption. Atomic line transitions are expected to be less prominent at higher temperatures because of a higher ionization degree \\citep{2002nsps.conf..263Z}. For a more detailed discussion of expected line features that can be produced in a hydrogen atmosphere see \\cite{2006NS...vanKerkwijk}. It is also not clear how important the contribution of heavier elements in the atmosphere is because of the strong gravitational stratification forces \\citep{2003astroph...Mori}. Further, the remarkable harmonic-like energy spacing of multiple lines needs to be understood. Studying their pulse phase dependence which can probably best be done for \\rbs\\ might shed light on this.  The modeling of pulse profiles in different X-ray energy bands is a powerful tool to unveil the surface temperature distribution of the neutron stars \\citep{2006MNRAS.366..727Z}. First studies of \\rxb\\ and \\rbs\\ yielded constraints on the polar cap geometry. It is remarkable that for both stars asymmetric configurations of the hot polar caps were inferred. In both cases different temperatures and sizes of the hot spots are found with the smaller spot being the hotter one. Probably, also in both stars the spots are not located in exactly antipodal positions. This could point to magnetic field configurations of an off-centered dipole or the involvement of higher-order multi-pole components. Studies of the effects of a strong magnetic field on the temperature distribution in the neutron star crust show that configurations which include dipolar poloidal and toroidal components can indeed reproduce the observed temperature distributions \\citep{2004A&A...426..267G,2005astroph...Geppert}.  \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics[clip=,angle=90]{p_pdot.ps}} \\caption{$P - \\dot P$ diagram for radio pulsars (black dots), magnetars (blue dots) and the Magnificent Seven. Red dots show the direct $P$ and $\\dot P$ measurements for \\rxb\\ and \\rbs, while the open diamonds mark the $\\dot P$ values expected from the magnetic field estimates assuming proton resonance as origin of the absorption lines. For \\rxb\\ and \\rbs\\ both field estimates are connected with a vertical line. Marked in green are radio pulsars which are also detected at high energies (incomplete for the millisecond pulsars).} \\label{fig-ppdot} \\end{figure}"
    },
    "0609/astro-ph0609585_arXiv.txt": {
        "abstract": "We report observations with the \\Chandra\\ X-ray Observatory of the single, cool, magnetic white dwarf GD~356. For consistent comparison with other X-ray observations of single white dwarfs, we also re-analyzed archival \\ROSAT\\ data for GD~356 (GJ~1205), G~99-47 (GR~290 = V1201~Ori), GD~90, G~195-19 (EG250 = GJ~339.1), and WD~2316+123 and archival \\Chandra\\ data for LHS~1038 (GJ~1004) and GD~358 (V777~Her). Our \\Chandra\\ observation detected no X rays from GD~356, setting the most restrictive upper limit to the X-ray luminosity from any cool white dwarf~---  $L_{X} < 6.0\\!\\times\\! 10^{25}$ erg~s$^{-1}$, at $99.7$\\% confidence, for a 1-keV thermal-bremsstrahlung spectrum. The corresponding limit to the electron density is $n_{0} < 4.4\\!\\times\\!10^{11}$ cm$^{-3}$. Our re-analysis of the archival data confirmed the non-detections reported by the original investigators. We discuss the implications of our and prior observations on models for coronal emission from white dwarfs. For magnetic white dwarfs, we emphasize the more stringent constraints imposed by cyclotron radiation. In addition, we describe (in an appendix) a statistical methodology for detecting a source and for constraining the strength of a source, which applies even when the number of source or background events is small. ",
        "introduction": "}  Several theorists \\citep[e.g., ][]{zhe84,ser90,tho95} have suggested that single, cool, magnetic white dwarfs might have coronae. Some observers \\citep{fon82,arn92,cav93,mus95,mus03} have previously searched for X radiation that might be emitted by hot gas above a white-dwarf photosphere. There were no persuasive detections, despite one false alarm~--- GR~290, in archival Einstein data \\citep{arn92}. We here report another upper limit, more stringent than the previous ones, for the white-dwarf star GD~356.  Three considerations motivated these X-ray searches: (a) the preponderance of magnetic white dwarfs among the X-ray emitting cataclysmic variables; (b) the possibility that some features in the optical spectra of magnetic white dwarfs might be cyclotron resonances \\citep{zhe84}; and (c) the feeling that coronal heating is not so well understood that constraints from other kinds of stars wouldn't be worthwhile. In the special case of GD~356, the presence of Balmer lines in emission imply at least a chromosphere, making more plausible the existence of a corona. Magnetic fields in most of the stars examined so far (including GD~356) are 1~MG or more, thought to be fossils from the white dwarfs' previous lives as Ap stars. More recent measurements of variable, weaker fields in some white dwarfs~--- the DBV GD~358 \\citep{win94} and the DA LHS~1038 \\citep{sch95}~--- suggest non-fossil magnetism. Calculations \\citep{tho95} indicate that cool white dwarfs with convective envelopes can support $\\alpha$--$\\omega$ dynamos capable of generating 100-kG fields.  Section~\\ref{s:coronal_activity} reviews recent theoretical considerations; Section~\\ref{s:previous_obs}, limits from previous searches. Next, Section~\\ref{s:chandra_obs} presents the results of our \\Chandra\\ observations of GD~356; Section~\\ref{s:reanalyze}, our re-analysis of prior \\ROSAT\\ and \\Chandra\\ observations of GD~356 and of other white dwarfs. Finally, Section~\\ref{s:discussion} discusses the implications of the new and re-analyzed results and emphasizes the importance of thermal cyclotron radiation in magnetic white dwarfs \\citep[cf.][and references therein]{zhe04}. ",
        "conclusions": "}  The objective of this investigation was to determine whether single, cool white dwarfs emit X radiation indicative of magnetic coronal activity. To assess the scientific implications of these null detections, we first (\\S \\ref{s:formation}) discuss issues related to the formation of a magnetic corona around a white dwarf. We then (\\S \\ref{s:cyclotron}) examine the rather severe constraints that cyclotron emission lines and radiative loss impose on a corona around a cool magnetic white dwarf, such as GD~356. Finally (\\S \\ref{s:conclude}), we summarize our results and conclusions about hypothesized hot coronae around magnetic white dwarfs.  \\subsection{Formation of a Corona\\label{s:formation}}  The existence of coronae around single, cool white dwarfs remains an unsettled issue: This, of course, motivated our search for X-ray evidence. Here we briefly address some questions relevant to the formation of a magnetic corona. \\begin{enumerate} \\item Is there a convection zone? Theoretical studies have indicated that convection can occur in white dwarfs for certain temperature ranges. Analyses of white-dwarf atmospheric elemental abundance also evidence convective activity. Thus, it is quite plausible that some white dwarfs possess a convection zone that generates acoustic waves. \\item How much acoustic wave energy is generated in the convection zone? Calculations \\citep{mus87,win94} have shown that the convection-generated acoustic flux could exceed $10^8$~erg~cm$^{-2}\\,$s$^{-1}$ and may reach $10^{10}$~erg~cm$^{-2}\\,$s$^{-1}$ for DA white dwarfs and $10^{11}$~erg~cm$^{-2}\\,$s$^{-1}$ for DB white dwarfs. \\item What fraction of the acoustic wave energy is transmitted to the white-dwarf surface? For white dwarf with a substantial magnetic field ($B \\approx 10^4$~G or higher), more than half of the wave energy generated in the convection zone can reach the white-dwarf atmosphere. If all the wave energy is converted into coronal X rays, the expected luminosity will be $\\approx 10^{27}\\!-\\! 10^{30}$~erg~s$^{-1}$. Our upper limit to the X-ray luminosity for GD~356 is an order of magnitude below $10^{27}$~erg~s$^{-1}$ over most of the range of assumed temperatures. This implies either that the putative convection zone generates less acoustic energy, that the efficiency of acoustic-wave transmission to the surface is smaller, or that a corona fails to form even if the acoustic energy reaches the white-dwarf surface. \\item Can a magnetic corona form, given sufficient energy provided by the acoustic flux? The null detection of coronal X rays from white dwarfs has led some \\citep [see e.g.,][]{mus05} to suggest that the emerging acoustic flux causes chromospheric activity rather than formation of a hot corona. An example of such chromospheric activity would be oscillations resulting from the stellar atmosphere to propagating acoustic waves in the presence of a temperature inversion \\citep{mus05}. Whether acoustic waves generate chromospheric activity remains unverified. Nevertheless, a number of systems exhibit chromospheric activity. The luminosity of H$\\alpha$ emission lines in GD~356 is $\\approx 1.8\\! \\times\\! 10^{27}$~erg~s$^{-1}$ \\citep[see][]{gre85}. Our \\Chandra\\ observation eliminates the possibility that irradiation of the atmosphere by coronal X rays powers the Balmer lines. Further, it is not clear that atmospheric oscillations can account for the luminosity of Balmer emission lines in GD~356. \\end{enumerate}  While coronal X radiation from cool ($T_{\\rm eff} < 10,000$~K) magnetic white dwarfs remains undetected, X radiation from hot optically thin thermal plasmas appears to occur from very hot ($1.2 \\times 10^5$~K) white dwarf KPD 0005+5106. Now the question is this: Without a magnetic corona, what supports the hot plasma envelope in the strong gravitational field of the white dwarf? Radiation driven envelopes~--- predicted for hot white dwarfs \\citep{bes90,zhe96}~--- could provide this support. Provided the white dwarf has a strong magnetic field, cyclotron-resonance radiation pressure from photospheric radiation could drive the wind, with an estimated mass-loss rate $\\approx 2\\! \\times\\! 10^{10}$~g~s$^{-1}$. Thus, a hot magnetic white dwarf could emit X rays from an optically thin, thermal plasma in a radiation-driven outflow.  \\subsection{Cyclotron Radiation\\label{s:cyclotron}}  The \\Chandra\\ observation sets stringent constraints on a supposed hot corona above the white-dwarf atmosphere. For GD~356, the X-ray luminosity $L_{X} < 6.0\\!\\times\\! 10^{25}$~erg s$^{-1}$ and electron number density $n_{0} < 4.4\\!\\times\\! 10^{11}~{\\rm cm}^{-3}$ at the base of a corona of (exponential) scale height of $2.0\\!\\times\\! 10^7~{\\rm cm}$ for a 1-keV plasma. However, for a magnetic white dwarf, \\citet{zhe04} note that electron-cyclotron emission lines and radiative losses even more severely constrain the parameters of a hypothesized hot corona.  For the 15-MG field of GD356, we computed the radiative transfer in a hot corona, including cyclotron and bremsstrahlung emissivities and opacities due to cyclotron radiation \\citep{cha89} and Faraday rotation and mixing. Figure~\\ref{f:cyclotron} displays the emergent spectrum for white-dwarf parameters appropriate to GD~356 (\\S \\ref{s:chandra_obs}) and coronal density and temperatures given in the caption. For a density comparable to the upper limit set by the \\Chandra\\ observation, the hypothesized hot corona would emit extremely strong thermal cyclotron lines in the near-infrared band~--- namely, around the second and third harmonics of the cyclotron frequency ($\\nu_{B} = 42$~THz, $\\lambda_{B} = 7.1$~$\\mu$m). The JHK photometry \\citep{skr06} shows no indication of such an excess, indicating that a hot corona does not exist or has a density much lower than the upper limit set by the X-ray observations. Alternatively, the magnetic field could be somewhat weaker than 15~MG, which would shift the strong third harmonic ($\\nu_{3} = 126$~THz, $\\lambda_{3} = 2.4$~$\\mu$m) to a lower frequency, longward of the K$_{s}$ band. Nevertheless, our calculation confirms the conclusion of \\citet{zhe04}: Infrared--visible spectrophotometry is potentially a powerful probe of any hot corona around a magnetic white dwarf.  In addition to producing potentially detectable emission lines, thermal cyclotron radiation would be the dominant cooling mechanism in a hot corona around a magnetic white dwarf. Figure~\\ref{f:power} displays the radiated thermal cyclotron and bremsstrahlung luminosity for a supposed isothermal corona around GD~356 ($B = 15$~MG), as a function of density for various electron temperatures between 0.125~keV and 2.0~keV. As the plot clearly shows, cyclotron cooling dramatically exceeds bremsstrahlung cooling for a hot, tenuous plasma above a magnetic white dwarf. Indeed, for GD~356, the coronal thermal cyclotron luminosity rivals the photospheric luminosity unless the electron density is very much less than the upper limit set by the X-ray observation.  Strong cyclotron cooling above a magnetic white dwarf imposes severe demands upon any coronal heating mechanism. As Figure~\\ref{f:power} demonstrates, even the weak requirement \\citep{zhe04} that the coronal thermal cyclotron luminosity not exceed the photospheric luminosity limits the electron density to $n_{0} < 3\\!\\times\\! 10^{9}~{\\rm cm}^{-3}$ for $kT = 1$~keV. If an acoustic flux as large as $10^{10}$~erg~cm$^{-2}$ s$^{-1}$ \\citep{mus87} efficiently heats a corona around GD~356, this heating rate can balance cyclotron cooling for $n_{0} \\approx 3\\!\\times\\! 10^{5}~{\\rm cm}^{-3}$ for $kT = 1$~keV. This is about 6-orders-of-magnitude less than the density limit set by X-ray non-detection!  For an electron density $n_{0} < 3\\!\\times\\! 10^{6}~{\\rm cm}^{-3}$~--- approximately independent of temperature~--- a corona above the white-dwarf photosphere would be transparent to cyclotron radiation in the 15-MG field of GD~356. However, when transparent, the cyclotron cooling time~--- again approximately independent of temperature~--- is only 2~$\\mu$s in this magnetic field. For electron densities and temperatures of interest, the mean time between collisions~--- about $(1.4\\!\\times\\! 10^{9}~{\\rm s~cm}^{-3})\\, (kT/(1~{\\rm keV}))^{1.5}\\, {n_{e}}^{-1}$~--- is much longer than this cyclotron cooling time. Thus, the plasma would not be in local thermodynamic equilibrium \\citep[LTE,][and references therein]{zhe04}, unless some collisionless process~--- e.g., scattering by Alfven waves~--- intervenes to transfer energy to the electrons' transverse degrees of freedom, on this very short timescale.  \\subsection{Summary\\label{s:conclude}}  In summary, our \\Chandra\\ observation of the single, cool white dwarf GD~356 limits the luminosity and density of a hypothesized hot corona to $L_{X} < 6.0\\!\\times\\! 10^{25}$ erg~s$^{-1}$ and $n_{0} < 4.4\\!\\times\\!10^{11}$ cm$^{-3}$ ($99.7$\\% confidence), for a 1-keV thermal-bremsstrahlung spectrum. We also re-analyzed archival \\ROSAT\\ data for this white dwarf, G~99-47, GD~90, G~195-19, and WD~2316+123, as well as archival \\Chandra\\ data for LHS~1038 and GD~358, using a statistical methodology (described in the Appendix) better suited to the low-count observations. Our upper limits are reasonably consistent with those of the original authors and lie somewhat above the limits we set for GD~356. As an aside to these searches for X-ray emission from cool white dwarfs, we have listed all \\Chandra-detected sources in the LHS-1038, GD-356, and GD-358 fields on the ACIS-S3 CCD.  We have shown \\citep[see also][]{zhe04} that for a magnetic white dwarf (such as GD~356), the non-detection of infrared--visible cyclotron emission lines can more severely constrain the parameters for a hot corona than does an X-ray non-detection. Furthermore, strong cyclotron cooling places extreme demands upon any coronal-heating mechanism. Indeed, our preliminary theoretical analysis suggests that cyclotron cooling around a magnetic white dwarf renders problematic the formation and maintenance of a hot corona. However, we require a more thorough study to prove this conclusively."
    },
    "0609/astro-ph0609120_arXiv.txt": {
        "abstract": "{ We consider an inhomogeneous but spherically symmetric Lemaitre-Tolman-Bondi model to demonstrate that spatial variations of the expansion rate can have a significant effect on the cosmological supernova observations. A model with no dark energy but a local Hubble parameter about $15$ \\% larger than its global value fits the supernova data better than the homogeneous model with the cosmological constant. The goodness of the fit is not sensitive to inhomogeneities in the present-day matter density, and our best fit model has $\\Omega_M(r) \\sim 0.3$, in agreement with galaxy surveys. We also compute the averaged expansion rate, defined by the Buchert equations, of the best fit model and show explicitly that there is no average acceleration.}  \\preprint{HIP-2006-39/TH} ",
        "introduction": "The simplest homogeneous and isotropic cosmological model within the framework of general relativity is parameterized by two numbers: the Hubble constant $H_0$ and the density parameter $\\Omega_M$, interpreted respectively as the average expansion rate and the average density of non-relativistic matter in the present universe. This model was generally considered a good description of the universe on the largest scales until the high redshift supernova observations in the late 90's \\cite{Riess:1998cb}. With the latest data from supernovae \\cite{Riess:2004nr,Astier:2005qq}, galaxy distributions \\cite{Eisenstein:2005su} and anisotropies of the cosmic microwave background \\cite{Spergel:2006hy} the simplest homogenous and isotropic model would now lead to a highly contradictory picture of the universe, as the best fit values of this model for the average matter density demonstrate:  \\begin{itemize} \\item Cosmic microwave background: $\\Omega_M \\sim 1$ \\item Galaxy surveys: $\\Omega_M \\sim 0.3$ \\item Type Ia supernovae: $\\Omega_M \\sim 0 $ \\end{itemize} The natural conclusion is that at least one of the assumptions of the model must be false.   As is well known, the problem has conventionally been remedied by introducing the cosmological constant $\\Lambda$ or vacuum energy $\\Omega_\\Lambda$, giving rise to an accelerated expansion of the universe. Indeed, the analysis of the cosmological data with both the vacuum energy and matter components included yields a consistent picture of the universe, known as the concordance $\\Lambda$CDM model, with the following best fit values for the density parameters: \\begin{itemize} \\item Cosmic microwave background: $\\Omega_M +\\Omega_\\Lambda \\sim 1$ \\item Galaxy surveys: $\\Omega_M \\sim 0.3$ \\item Type Ia supernovae\\footnote{In the supernova data analysis spatial flatness ($ \\Leftrightarrow  \\Omega_M +\\Omega_\\Lambda = 1$) has been assumed. However, these values are within the $1\\sigma$ error of the best fit values without the flatness constraint \\cite{Riess:2004nr}: $\\Omega_M \\sim 0.5 $ and $\\Omega_\\Lambda \\sim 1$. }: $\\Omega_M \\sim 0.3$ and $\\Omega_\\Lambda \\sim 0.7$ \\end{itemize} Moreover, the data seems to require $\\Omega_\\Lambda > 0$ at a high confidence level \\cite{Riess:2004nr,Astier:2005qq}.   Although the cosmological concordance $\\Lambda$CDM model fits all the observations well, it is plagued by theoretical problems \\cite{Copeland:2006wr}. We do not have a theory that would explain, not to mention predict, the required value of the cosmological constant. Various dark energy models, which have been studied intensively in order to provide a dynamical explanation for the cosmological constant, are not compelling from the particle physics point of view and often require fine-tuning \\cite{Copeland:2006wr,Straumann:2006tv}. Modifications of the general theory of relativity on cosmological scales appear to suffer from analogous problems; in fact, it has recently been argued that the cosmological constant seems to be essentially the only modification that fits all the cosmological data \\cite{Koivisto:2006ie}.   As a rejection of either matter domination or Einstein gravity leads to trouble, it is well motivated to study the validity of the third main assumption, the perfect homogeneity. Undoubtedly, at least on small scales, large inhomogeneities exist. Although the potential cosmological consequences of the inhomogeneities were recognized already around the same time when the homogeneous and isotropic models of the universe were first studied, their impact on the global dynamics of the universe is still unknown (see e.g.\\ \\cite{Krasinski}). Indeed, the role of the inhomogeneities has often been debated whenever there have been some ambiguities with cosmological observations, for example, in the 90's when in the determination of the age of the universe there was a discrepancy between the implications of cosmological and astronomical observations \\cite{Futamase,Russ:1996km}.   Most recently, inhomogeneities have been invoked as the culprit for the apparent acceleration of the expansion of the universe\\footnote{Inhomogeneities as an alternative to dark energy were first discussed in \\cite{Pascual-Sanchez:1999zr}.}, in particular by virtue of their so-called backreaction on the metric (for a recent discussion on the issues involved and a comprehensive list of references, see \\cite{Rasanen:2006kp}). The problems in these approaches have been mainly of technical nature, as reducing the symmetries of the metric rapidly complicates the calculations. A lesser difficulty arises from the fact that different inhomogeneities can lead to identical observations \\cite{Mustapha:1998jb}, so that not even ideal observations of light could determine the inhomogeneities uniquely. As a consequence, there have been two conceptually distinct ways to approach the inhomogeneities: the first one examines the effect of the inhomogeneities on the expansion of the universe, whereas the second tries to determine their impact on the observations directly.   In the first approach an effective description of the universe is constructed by averaging out the spatial degrees of freedom, i.e.\\ the inhomogeneities \\cite{Rasanen:2006kp,Ellis:2005uz,Buchert:1999er,Coley:2005ei}. As a result, one obtains averaged, effective Einstein equations which, in addition to terms found in the usual homogeneous case, include new terms that represent the effect of the inhomogeneities, called backreaction in this context. It has been argued that the backreaction might account for the accelerated expansion \\cite{Rasanen:2006kp,Buchert:1999er,Kai:2006ws}.   However, since we can only observe the redshift and energy flux of light arriving from a given source, not the expansion rate and the matter density of the universe nor their averages, one may wonder how the actual observables are related to the averaged equations. To wit, since we do not observe the expansion of the universe directly, its acceleration is also an indirect conclusion, arising from the fact that \\textit{in the perfectly homogeneous cosmological models} dark energy is required for a good fit. Consequently, there is no a priori reason to assume that the acceleration would be needed in the more general inhomogeneous models of the universe. Moreover, it is well recognized that in general it is not correct to integrate out constrained degrees of freedom as if they were independent. Indeed, the fact that we can make cosmological observations only along our past light cone makes the observable universe a constrained system.   The second approach avoids these problems by studying the effect of the inhomogeneities directly on the observable light \\cite{Mustapha:1998jb,Celerier:1999hp}. This would be virtually impossible in the presence of generic inhomogeneities but can be done in some simpler models, such as the spherically symmetric Lemaitre-Tolman-Bondi (LTB) model \\cite{Lemaitre:1933qe,Tolman:1934za,Bondi:1947av}. This model works best in describing smooth inhomogeneities at scales of $100$ ${\\rm{Mpc}}$ and larger; the spherical symmetry prevents to use it as a model for the random small scale lumpiness caused by galaxies, as noted in \\cite{Biswas:2006ub}. Indeed, the LTB model has been used to study the effect of the smooth inhomogeneities on the cosmological observations by several authors \\cite{Biswas:2006ub,Iguchi:2001sq,Godlowski:2004gh,Alnes:2005rw,Bolejko:2005fp,Vanderveld:2006rb,Garfinkle:2006sb,Chung:2006xh,Mustapha:1997xb,Brouzakis:2006dj}; a common conception is that these models inevitably contradict the observed homogeneity of the large scale galaxy distribution.   Although spherical symmetry is probably an unrealistic assumption for the entire universe, the LTB model can be regarded as describing observations that have been averaged\\footnote{This is naturally only a crude approximation; a more correct way would be to first solve the observables in a nonspherical model and only then average out their angular dependence.} over the celestial sphere and it is therefore useful at least on two counts. First, it serves as a simple testing ground for the effect of the inhomogeneities on the cosmological observations. Second, as the fits can be performed unambiguously, it can be used to study the connection of the backreaction driven effective acceleration to the observations; both of these points will be examined in this work.   As mentioned in the beginning, two independent parameters ($H_0$ and $\\Omega_M$) are needed to uniquely define the homogeneous matter dominated universe. In the presence of inhomogeneities, the values of these two quantities are needed at every spatial point; that is, the inhomogeneous dust models are defined by two functions of the spatial coordinates: $H_0(x^i)$ and $\\Omega_M(x^i)$. As a consequence, there are inhomogeneities of two physically different kind: inhomogeneities in the matter distribution, and inhomogeneities in the expansion rate. Although their dynamics are coupled via the Einstein equation, as boundary conditions they are independent. This opens up the possibility for a universe with inhomogeneous expansion but homogeneous present-day matter distribution; a model of this kind could potentially fit the supernova data as well as the galaxy surveys without invoking dark energy.   The paper is organized as follows. Sect. \\ref{ltb} discusses the general properties of the spherically symmetric, inhomogeneous LTB model. In Sect. \\ref{models} we fit the LTB model with inhomogeneous expansion but homogeneous present-day matter distribution to the supernova observations. For completeness, we consider both the cases of pure dust and dust plus the cosmological constant. The initial conditions of these models are discussed in Sects. \\ref{initialcond} and \\ref{agesect}. In Sect. \\ref{averagesec} we evaluate the expansion rate and shear in the appropriate averaged (the so-called Buchert \\cite{Buchert:1999er}) equations and demonstrate that an accelerated average expansion is not needed to fit the supernova data. Finally, Sect. \\ref{conclusions} contains our conclusions. ",
        "conclusions": "We have studied the effect of inhomogeneities on the cosmological supernova observations in the spherically symmetric matter dominated LTB universe. Our main conclusions are as follows:  \\begin{enumerate} \\item The inhomogeneous matter dominated LTB model fits the supernova data better than the homogeneous FRW model with the cosmological constant.  \\item Inhomogeneities of the expansion rate seem to have much bigger effect on the goodness of the fit than the inhomogeneities of the matter distribution.  \\item The model giving the best fit has perfectly uniform present-day matter density with $\\Omega_M \\sim 0.3$, in agreement with the galaxy surveys, but its Hubble parameter varies from the local value of $\\sim 70$ ${\\rm{km/s/Mpc}}$ to the value of $\\sim 60$ ${\\rm{km/s/Mpc}}$ over the distance scale of $\\sim500$ ${\\rm{Mpc}}$ from us.  \\item Neither local nor average acceleration is needed for a good fit to the supernova data.  \\item At least in the LTB universe, the averaged Buchert equations give a notably different prediction for the observable luminosity-redshift relation than the exact Einstein equations. \\end{enumerate}  The first of the above listed results is neither new nor surprising, as it was proved in \\cite{Mustapha:1998jb} that any isotropic set of observations can be explained by appropriate inhomogeneities, and later demonstrated with examples by several authors. Instead, the second result is a new one, as is the model that explains both the observations of supernovae and the observed homogeneity in the large scale distribution of galaxies without invoking dark energy. Moreover, we found in Sect. \\ref{initialcond} that within the dynamics of the Einstein equation, the universe may evolve to this kind of a configuration without any substantial amount of fine-tuning in the initial conditions. In addition, as the supernovae and galaxy surveys constrain the boundary condition functions $H_0(r)$ and $\\Omega_M (r)$ only for the small values of the radial coordinate $r$, the tails of these functions can still be freely chosen to give a good fit for the CMB power spectrum as well. In fact, Alnes, Amarzguioui and Gr{\\o}n have realized this in practice by fitting the LTB model to the first peak of the CMB \\cite{Alnes:2005rw}. It has also been shown that even a homogeneous model can fit the CMB data without dark energy \\cite{Blanchard:2003du}.   Our another focus was on the averaged dynamics of the dust universe, given by the Buchert equations that illustratively demonstrate the potential cosmological significance of the inhomogeneities. In Sect. \\ref{averagesec}, we showed that none of our dust models that fit the supernova observations have accelerated average expansion. On the other hand, one can easily construct examples that have average acceleration but are ruled out by observations, whereas dust models that would both have the average acceleration and fit the cosmological data are still to be presented. Moreover, as discussed at the end of Sect. \\ref{averagesec}, it seems that the Buchert equations cannot be used to calculate the observable luminosity-redshift relation. Although explicitly demonstrated only in the LTB universe, we expect that it might be the case in the more general models as well; after all, the Buchert equations are obtained by integrating over a spatial hypersurface whereas the observables are calculated by integrating over the past light cone. Altogether, averaging over a spatial hypersurface is perhaps not practical if one wants to make comparison with the observations. Instead, probably a better alternative is to relate the inhomogeneities directly to the observations, as done e.g.\\ in this paper and, if needed, take the averages only of the relevant observables, such as the energy flux and redshift of light.   Of course, some sort of smoothing is implicitly assumed also in the LTB model. Therefore, the study of observations in the LTB universe does not answer the interesting question about the potential cosmological consequences of the small scale lumpiness caused by galaxies. Their effect could be significant, possibly depending also on how smoothly dark matter, perhaps the dominant energy component of the universe, is distributed in the universe."
    },
    "0609/astro-ph0609646_arXiv.txt": {
        "abstract": "{We present the results of the first laboratory experiment of  checkerboard shaped pupil binary mask coronagraphs using visible light, in the context of the R\\&D activities for future mid-infrared space missions such as the 3.5\\,m SPICA telescope. } {The primary aim of this work is to demonstrate the coronagraphic performance of checkerboard masks down to a $10^{-6}$ peak-to-peak contrast, which is required to detect self-luminous extra-solar planets in the mid-infrared region. } {Two masks, consisting of aluminum films on a glass substrates, were manufactured using nano-fabrication techniques with electron beam lithography: one mask was optimized for a pupil with a 30\\% central obstruction and the other was for a pupil without obstruction. The theoretical contrast for both masks was  $10^{-7}$ and no adaptive optics system was employed. } { For both masks, the observed point spread functions were quite consistent with the theoretical ones. The average contrast measured within the dark regions was $2.7 {\\times} 10^{-7}$ and $1.1 {\\times} 10^{-7}$. } {The coronagraphic performance significantly outperformed the $10^{-6}$ requirement and almost reached the theoretical limit determined by the mask designs. We discuss the potential application of checkerboard masks for mid-infrared coronagraphy, and conclude that binary masks are promising for future high-contrast space telescopes. } ",
        "introduction": "Direct detection and spectroscopy of extra-solar planets is very important to the understanding of how planetary systems were born, how they evolve, and ultimately to find biological signatures on these planets. After the first report unambiguously confirming the existence of an extrasolar planet by Mayor \\& Queloz (1995), more than 200 extra-solar planets have been found to date, mostly by detailed observation of the Doppler shift of the spectrum of the central star (\\cite{Mayor1995}), or by the variability of the luminosity of the central star due to the transit of the planet (\\cite{Henry2000}; \\cite{Charbonneau2000}). However, such indirect observation does not usually provide the spectral features, luminosity and other important properties of the planets themselves, so the direct observation of extra-solar planets is of primary importance. However, the enormous contrast in luminosity between the central star and the planet, and the associated diffraction patterns from optical instruments requires some coronagraphic techniques to suppress, or at least attenuate them. Typically, the contrast at visible light wavelengths is $\\sim\\,10^{10}$ but is reduced to $\\sim\\,10^6$ in the mid-infrared region (\\cite{Burrows2004}), which is one of the most important advantages of infrared coronagraphy for the observation of extra-solar planets.  The Space Infrared telescope for Cosmology and Astrophysics (SPICA) is the next generation mission for infrared astronomy led by Japan (\\cite{Nakagawa2004}). The SPICA telescope uses on-axis Ritchey-Chretien optics with a 3.5\\,m diameter monolithic primary mirror. The whole telescope will be cooled to 4.5\\,K and infrared observations will be made at wavelengths within the 5--200\\,$\\mu$m range. A coronagraphic instrument is currently being considered for the SPICA mission (\\cite{Enya2006}; \\cite{Abe2006}). The primary target of this instrument is the direct detection and observation of extra-solar Jovian outer planets (typically beyond 5\\,AUs). The baseline requirements for the SPICA coronagraph are as follows: $10^{-6}$ peak-to-peak contrast, smallest possible inner-working-angle (IWA), and a core wavelength within the range  $5$--$20$\\,$\\mu$m.    Of the various current coronagraphic methods, coronagraphs using binary shaped  pupil  masks have some remarkable advantages (e.g., Jacquinot \\& Roizen-Dossier 1969; Vanderbei, Spergel \\& Kasdin 2003; Kasdin, Vanderbei \\& Spergel 2003). They are relatively simple in their implementation, essentially achromatic systems, and robust against pointing errors of the telescope (\\cite{green2004}). Therefore, binary masks have been studied as the baseline for the SPICA coronagraph development (\\cite{Tanaka2006}; \\cite{Enya2006}). Our aim is to demonstrate precise mask fabrication, and to check the performance of the masks with visible light in the laboratory before proceeding further.  In this paper, we present the results of a laboratory demonstration of the checkerboard mask coronagraphs in visible light. The target contrast to be demonstrated in this work was set to $10^{-6}$, which is required to detect self-luminous extra-solar planets in the mid-infrared. It was shown that the observed contrast went beyond this goal, very close to the theoretical contrast of $10^{-7}$. Expansion toward a mid-infrared demonstration test-bench is also discussed. ",
        "conclusions": "We report the results from a laboratory experiment using a binary checkerboard-type pupil shaped mask coronagraph. Although this kind of mask presents several limitations (large IWA and low throughput), its manufacture and implementation are very robust. Two masks, optimized for either obstructed pupil or not, were fabricated and tested. Both masks achieved a higher performance than the $10^{-6}$  requirement: $2.3\\,{\\times}\\,10^{-7}$ and $1.1\\,{\\times}\\,10^{-7}$. Therefore, we conclude that a binary checkerboard mask is a very attractive solution for the future MIR coronagraph of the SPICA space telescope, for which this study was carried out. Other coronagraphic solutions, more complex to implement are also considered, but the checkerboard mask can be regarded as the safest option."
    },
    "0609/astro-ph0609470_arXiv.txt": {
        "abstract": "We discuss  how tidal interaction between the Large Magellanic Cloud (LMC), the Small Magellanic Cloud (SMC), and the Galaxy triggers galaxy-wide star formation in the Clouds for the last $\\sim 0.2$ Gyr based on our chemodynamical simulations on the Clouds. Our simulations demonstrate that the tidal interaction induces the formation of asymmetric spiral arms with high gas densities and consequently  triggers star formation within the arms in the LMC. Star formation rate in the present LMC is significantly enhanced just above the eastern edge of the LMC's stellar bar owing to the tidal interaction. The location of the enhanced star formation is very similar to the observed location of  30 Doradus, which suggests that the formation of 30 Doradus is closely associated with the last Magellanic collision about 0.2  Gyr ago. The tidal interaction can dramatically compress gas initially within the outer part of the SMC so that new stars can be formed from the gas to become intergalactic young stars in the inter-Cloud region (e.g., the Magellanic Bridge). The metallicity distribution function of the newly formed stars in the Magellanic Bridge has a peak of [Fe/H] $\\sim$ $-0.8$, which is significantly lower than the stellar metallicity of the SMC. ",
        "introduction": "The Magellanic system composed of the LMC and the SMC is believed to be an interacting one where star formation histories of the Clouds have been strongly influenced by dynamical and hydrodynamical  effects of galaxy  interaction (\\cite[Westerlund 1997]{Westerlund97}). It is however unclear how galaxy interaction between the Clouds and the Galaxy triggers star formation in the gas disks of the Clouds. Recent observations on spatial distributions of HI (\\cite[Staveley-Smith \\etal\\ 2003]{Staveley-Smith03}), molecular gas (\\cite[Fukui \\etal\\ 1999]{Fukui99}), and young stars (\\cite[Grebel \\&  Brandner 1998]{Grebel98}) have provided vital information on galaxy-wide triggering mechanisms of star formation in the Clouds. By comparing numerical simulations of the Magellanic system with these observations, we here discuss (1) how the tidal interaction changes the spatial distribution of high-density gaseous regions where new stars can be formed in the LMC, (2) whether the formation of 30 Doradus is triggered by the interaction, and (3) how the interaction triggers star formation in the Magellanic Bridge (MB).   \\begin{figure} \\centerline{ \\scalebox{0.5}{% \\includegraphics{bekki.fig1.eps}% } } \\caption{Time evolution of distance between the LMC and the SMC (magenta), the LMC and the Galaxy (red), and the SMC and the Galaxy (blue) for the last 0.8 Gyr. Note that the LMC-SMC distance becomes minimum (8kpc) about 0.2 Gyr ago.} \\end{figure} ",
        "conclusions": "The present study has suggested that tidal interaction between the Clouds and the Galaxy is closely associated with the formation of 30 Doradus, the southern molecular ridge of the LMC,  and inter-Cloud stars with low metallicities. The observed asymmetric and clumpy distributions of young stars in the LMC are  demonstrated to be due to the last Magellanic interaction about 0.2 Gyr ago, which forms asymmetric spiral arms with high-density gas. The synchronized burst of star formation in the Clouds about 0.2 Gyr ago (by tidal interaction) can be proved by the observed age distributions of young star clusters in the Clouds (e.g., \\cite[Girardi \\etal\\ 1995]{Girardi95}). The simulated distributions of star-forming regions will be compared with the latest results of the Spitzer observations on young stellar objects (YSOs) in the Clouds."
    },
    "0609/astro-ph0609193_arXiv.txt": {
        "abstract": "{ Traditional binned statistics such as $\\chi^2$ suffer from information loss and arbitrariness of the binning procedure, which is especially important at low count rates as encountered in the \\XMM~Extended Survey of the Taurus Molecular Cloud (XEST). We point out that the underlying statistical quantity (the log likelihood $L$) does not require any binning beyond the one implied by instrumental readout channels, and we propose to use it for low-count data. The performance of $L$ in the model classification and point estimation problems is explored by Monte-Carlo simulations of {\\it Chandra} and \\XMM~X-ray spectra, and is compared to the performances of the binned Poisson statistic ($C$), Pearson's $\\chi^2$ and Neyman's $\\chi^2_N$, the Kolmogorov-Smirnov, and Kuiper's statistics. It is found that the unbinned log likelihood $L$ performs best with regard to the expected chi-square distance between true and estimated spectra, the chance of a successful identification among discrete candidate models, the area under the receiver-operator curve of reduced (two-model) binary classification problems, and generally also with regard to the mean square errors of individual spectrum parameters. The $\\chi^2$ ($\\chi^2_{\\rm N}$) statistics should only be used if more than 10 (15) predicted counts per bin are available. From the practical point of view, the computational cost of evaluating $L$ is smaller than for any of the alternative methods if the forward model is specified in terms of a Poisson intensity and normalization is a free parameter. The maximum-$L$ method is applied to 14 XEST observations, and confidence regions are discussed. The unbinned results are compared to binned XSPEC results, and found to generally agree, with exceptions explained by instability under re-binning and by background fine structures. In particular, HO Tau is found by the unbinned method to be rather cool ($kT$ $\\sim$ 0.2 keV), which may be a sign of shock emission. The maximum-$L$ method has no lower limit on the available counts, and allows to treat weak sources which are beyond the means of binned methods. ",
        "introduction": "Energy histograms are often used as spectrum estimators. While this procedure has the advantage of simplicity, the grouping of counts into a histogram is also associated with information loss. This becomes especially important if only few counts are available, so that the spectral fine structures are sparsely sampled by the observed counts. Such a situation arises in the \\XMM~Extended Survey of the Taurus Molecular Cloud (XEST; G\\\"udel et al. \\cite{guedel06a}), thus prompting the search for alternative unbinned methods.  Another motivation for histogram formation is to use $\\chi^2$ as a simple and simply coded measure of agreement between theory and observation. The role of $\\chi^2$ is thus to provide a plausibility ordering of alternative models, and to assess their absolute credibility. Since the histograms follow a multinomial distribution, which becomes gaussian only in the limit of infinite sample size, the use of $\\chi^2$ represents an approximation and corrections must be applied for small $n$ (e.g., Kendall \\& Stuart \\cite{kendall58}, Wachter \\cite{wachter79}, Nousek \\cite{nousek89}, Mighell \\cite{mighell99}, Arzner \\cite{arzner04}). Alternatively, one may use the binned Poisson ($C$) statistics. There exists, however, a simpler solution.  Namely, the relevant statistical quantity, the likelihood function, can be defined for an unlimited instrumental resolution without any reference to binning. Using this unbinned expression avoids arbitrariness of histogram formation, and thus a potential source of discrepant results. It has been successfully applied to ROSAT (Boese \\& Doebereiner \\cite{boese01}) and EGRET (Digel \\cite{digel00}) observations. Related approaches invoke piecewise-constant intensity models (Scargle \\cite{scargle98}; see Stelzer et al. \\cite{stelzer06} for an application to the XEST) or transformations to uniform null hypotheses (Kinoshita \\cite{kinoshita02}), and applications in other fields than astronomy include particle physics (Baker \\& Cousin \\cite{baker84}) and medicine (Miller et al. \\cite{miller02}). While the $C$ statistic is often used in astronomical applications (e.g., Dolphin \\cite{dolphin02} for star formation statistics; Babu and Feigelson \\cite{babu96} for a general overview), the unbinned (exact) Poisson likelihood is less often used. Alternative unbinned methods such as the Kolmogorov-Smirnov or Kuiper's statistics have the advantage of being (asymptotically) model-independent, which, however, also entails sub-optimal performance if the modeling was correct.  In this article we revisit the issue by Monte-Carlo simulations of \\XMM~and {\\it Chandra} CCD spectra of the Taurus Molecular Cloud (TMC). By simulating counts from a known spectrum, and applying different statistics in order to find the best-fit model, we assess the performances of the various statistics, and thereby gain a more differentiated picture. The present study represents an extension of the work of Nousek (\\cite{nousek89}) and Wachter et al. (\\cite{wachter79}) to more complex spectral models and a broader range of statistics.  The article is organized as follows. Section 2 introduces the models for sources and background spectra. Section 3 discusses the counting statistics. Section 4 describes measures af agreement between models, and between models and data, which can be used to assess the performance of various statistics. Section 5 describes the Monte-Carlo simulations and their numerical results. Real-data applications from the TMC are presented in Section 6, where also the issue of confidence regions is briefly touched, and the unbinned estimates are compared to the binned estimates from the XSPEC software. Section 7 contains a summary and conclusions. ",
        "conclusions": "Summary and conclusions}  We have used Monte-Carlo simulations to assess the performance of the unbinned (exact) Poisson likelihood ($L$), binned Poisson likelihood ($C$), Pearson's $\\chi^2$ and Neyman's $\\chi^2_{\\rm N}$, Kolmogorov-Smirnov ($D$), and Kuiper's ($V$) statistics in the model classification and point estimation problems of low-count \\XMM~and {\\it Chandra}/ACIS spectra parameterized by temperature ($kT$), hydrogen density ($N_H$) and normalization ($\\N$). By `unbinned' we mean that no grouping of instrumental energy channels is involved, so that the full readout resolution is used. The $L$ statistic equals the $C$ statistic on the finest possible (instrumental) binning. In all cases, the source normalization $\\N_{\\rm src}$ was taken from its maximum-likelihood estimator, while the shape parameters ($kT, N_H$) were taken such as to optimize the different statistics.  The outcome of the simulations can be summarized as follows:  \\begin{itemize}  \\item The unbinned Poisson likelihood performs best with regard to the following measures of performance: (i) the probability of a successful identification in a search over discrete $d_{\\chi^2}$-ordered candidate models, (ii) the expected chi-square distance between continuously indexed true and estimated models, (iii) the area under the operator-receiver curves of reduced binary (two-model) classifier problems, and (iv) and generally also with respect to the mean square errors of individual parameter errors.  \\item Under the $L$ statistic, two models can on average be distinguished if their chi square distance exceeds $\\sim$3. Similar statements hold for the other statistics ($C,\\chi^2,\\chi^2_{\\rm N},D,V$) and correspondingly larger distances (up to $\\sim$5). The chi square distance is thus empirically found to be a useful parameter-free measure of discrepancy between Poisson intensities.   \\item The $\\chi^2$ statistics should not be used unless more than 10 counts per bin are available, and $\\chi^2_{\\rm N}$ should not be used for less than 15 counts.  \\end{itemize}  We therefore argue that the unbinned Poisson likelihood $L$ is beneficial in cases with fewer counts than instrumental readout channels. Using $L$, the narrow spectral lines and steep gradients are taken into account at their exact location, unaffected by the bin size. Even more important, any instability of results under (arbitrary) choices of bin offsets and -sizes is avoided. In our view, this is a major benefit since it eliminates user-dependent sources of discrepant results.  The $L$ statistics has been applied to XEST data, and demonstrated to operate under real observational conditions. The major selection criteria were spectral simplicity (so that the 3-parameter model applies), low count rate, and comparably well-known background. The parameter space was rigorously explored to avoid potential difficulties with iterative optimization. The maximum-$L$ results usually agree with the minimum-$\\chi^2$ and minimum-$\\chi^2_{\\rm N}$ values within 68\\% confidence, but are not equal. There are, though, two cases (out of 14) where the unbinned and binned estimates disagree: HO Tau and GN Tau. Such disagreement, by itself, is a helpful indicator for potential problems in the forward modeling.  In particular, HO Tau is found by the $L$ method to be rather cool ($kT \\sim$ 0.2 keV). Such a cool X-ray-emitting plasma would be unusual for a T Tauri star, the majority of which have average temperatures of $> 10$~MK. In the present survey dominant plasma with similarly low temperature is only found in the low temperature component of the spectrum of the jet-driving sources DG Tau A, GV Tau A, and DP Tau, which may be due to shocks in the jet (G\\\"udel et al. \\cite{guedel05}, G\\\"udel et al \\cite{guedel06b}). However, DG Tau A also has significant plasma at very high temperatures ($\\sim 50$~MK) and HO Tau has no known jet. A dominant low-temperature corona has been found on the nearest classical T Tauri star, TW Hya, which is thus far unique in this respect. Apparent high densities in this cool plasma suggest that collisionally-ionized shocks in columns of material accreting onto the star may generate the X-ray emission (Kastner et al. \\cite{kastner02}; Stelzer and Schmitt 2004). TW Hya is an unusually old classical T Tauri star, approximately 10 Myr old. Using the Siess et al. (2000) stellar evolution models, HO Tau appears also to be a relatively old classical T Tauri star, at 8--9~Myr. Its mass accretion rate of $1.3 \\times 10^{-9}$~M$_{\\odot}$\\,yr$^{-1}$ (White and Ghez 2001) is several times higher than that of TW~Hya. Therefore it is plausible that HO~Tau may be analogous to TW~Hya, in which case the bulk of its X-ray emission may be generated by accretion shocks.  On the other hand, we can check whether the derived $N_H$ agrees with expectations from measured optical and near-infrared extinctions $A_V$ and $A_J$. For general interstellar matter, the conversions are $N_H \\approx 2\\times 10^{21}~A_V$~cm$^{-2}$ and $N_H \\approx 5.6\\times 10^{21}~A_J$~cm$^{-2}$ (see Vuong et al. 2002 and references therein). The optical extinction for HO Tau is $A_V = 1.11-1.13$~mag (Kenyon \\& Hartmann \\cite{kenyon95}, White \\& Ghez \\cite{white01}), and the near-IR extinction is $A_J = 0.32-0.46$~mag (Kenyon \\& Hartmann \\cite{kenyon95}, Brice\\~no et al. \\cite{briceno02}), thus implying $N_H = (1.8-2.6)\\times 10^{21}$~cm$^{-2}$ under the assumption of standard gas-to-dust ratios. These estimates are in better agreement with the ``hot'' solution (Fig. 15) although the error ranges are large. We also note that significant deviations from the standard interstellar relation between extinction and X-ray absorption are well known (e.g., in the jet-driving stars mentioned above). Such deviations may point to an ``anomalous'' gas-to-dust ratio for example as a consequence of dust evaporation.  We also note that both the binned and unbinned methods provide low $kT$ and high $N_H$. The discrepancy is in $L_X$ (see Fig. \\ref{comparison_fig}) which is a consequence of large uncertainties in the spectral integration because most of the soft emission from the cool plasma is strongly absorbed. Indeed, if the temperature estimate is changed from 0.14 keV (unbinned) to 0.28 keV (binned) then $L_X$ drops by about a factor ten. The $L_X$ of the unbinned method is indeed rather high given a stellar bolometric luminosity of this system of $L_* = 0.17~L_{\\odot} = 6.5\\times 10^{32}$~erg~s$^{-1}$. Adopting the usual X-ray saturation level of $L_X/L_* = 10^{-3}$ (e.g., Vilhu \\& Rucinski \\cite{vilhu83}), we would expect a maximum $L_X = 6.5\\times 10^{29}$~erg~s$^{-1}$, more in line with the binned solution. The low count-rate of HO Tau unfortunately precludes a high-resolution X-ray spectroscopic study to study density-sensitive triplet emission-lines from He-like ions of O, Ne and Mg, which suggest high densities in the TW Hya plasma. As for GN Tau, $A_J$ = 1.17~mag (Luhman \\cite{luhman04}), hence we expect $N_H = 6.6\\times 10^{21}$~cm$^{-2}$, very close to the result from the unbinned method.  Although the Monte Carlo simulations suggest that the maximum-$L$ estimates are closer to the true values when averaged over a (hypothetical) ensemble of similar observations, a few words of caution are in order. The present simulations have assumed that the true spectrum is contained in the set of candidate spectra. If this is not the case, or if the Poisson process assumption is violated, then the use of the most binding $L$ statistics also introduces the most severe misinterpretations. A major cause of flawed models stems from the background estimation. We have assumed here that this estimation is perfect, and have not considered background errors (e.g., Conrad et al. \\cite{conrad03}), nor the more complicated problem of estimating the background spectrum jointly with the source spectrum. An imprecise background model may be the cause for the discrepancy between binned and unbinned estimates found in GN Tau, where the most relevant low-energy background (0.3 to 1 keV) contains 200 counts only, implying about 30\\% statistical background uncertainty at resolution $\\Delta E \\sim$0.04 keV. Furthermore, we have focused here on the choice of likelihood functions for the point estimation and model classification problems. Accordingly, our treatment of parameterization issues and confidence regions was rather crude. In particular, we did not attempt to introduce any (Bayesian) a priori information other than implicit in the choice of the forward models.  Finally, it should be pointed out that the unbinned approach is in principle not restricted to the one-dimensional spectrum problem considered in this article, and that the energy tags might be replaced by tuples of photon energy, arrival time, and position on the detector. However, the corresponding forward models would involve many more degrees of freedom, to the point where a simple maximum-likelihood principle might no longer be sufficient to uniquely determine the solution. Additional information (such as a Bayesian prior) would then be required in order to regularize the problem. Also, the more elaborated forward models would be more vulnerable to systematic errors (i.e., from CCD boundaries). In view of the (sparse) data we shall not pursue these issues.  \\begin{appendix}"
    },
    "0609/astro-ph0609700_arXiv.txt": {
        "abstract": "We present metallicities and ages for 52 red giants in the remote Galactic dwarf spheroidal (dSph) galaxy  Leo\\,II.  These stars cover the entire surface area of Leo\\,II and are radial velocity members.  We obtained medium-resolution multi-fiber spectroscopy with the FLAMES multi-object spectrograph as part of a Large Programme with the Very Large Telescope at the European Southern Observatory, Chile. The metallicities were determined based on the well-established near-infrared \\ion{Ca}{2} triplet technique. This allowed us to achieve a mean random error of 0.16\\,dex on the metallicities, while other systematic effects, such as unknown variations in the dSph's [Ca/Fe]-ratio, may introduce a further source of uncertainty of the order of 0.1\\,dex. The resulting metallicity distribution is asymmetric and peaks at [Fe/H] $= -1.74$\\,dex on the Carretta \\& Gratton scale.  The full range in metallicities extends from $-2.4$ to $-1.08$\\,dex. As in other dSph galaxies, no extremely metal-poor red giants were found.  We compare Leo\\,II's observed metallicity distribution with model predictions for several other Galactic dSphs from the literature.  Leo\\,II clearly exhibits a lack of more metal poor stars, in analogy to the classical G-dwarf problem, which may indicate a comparable ``K-giant problem''. Moreover, its evolution appears to have been affected by galactic winds.  We use our inferred metallicities as an input parameter for isochrone fits to Sloan Digital Sky Survey photometry of our target stars and derive approximate ages.  The resulting age-metallicity distribution covers the full age range from 2 to about 15 Gyr on our adopted isochrone scale. During the first $\\sim 7$ Gyr relative to the oldest stars the metallicity of Leo\\,II appears to have remained almost constant, centering on the mean metallicity of this galaxy. The almost constant metallicity at higher ages and a slight drop by about 0.3 dex thereafter may be indicative of rejuvenation by low metallicity gas. Overall, the age-metallicity relation appears to support the formation of Leo\\,II from pre-enriched gas. Evidence for enrichment is seen during the recent 2 to 4 Gyr. Our findings support earlier derived photometric findings of Leo\\,II as a galaxy with a prominent old population and dominant intermediate-age populations.   We do not see a significant indication of a radial metallicity gradient nor age gradient in our current data. ",
        "introduction": "The majority of the galaxies in the Local Group are dwarf spheroidal (dSph) galaxies, the least massive, least luminous type of galaxies known.  They are mainly found within less than 300 kpc of massive galaxies like the Milky Way or M31 (see, e.g., Grebel, Gallagher, \\& Harbeck 2003; van den Bergh 1999a).  The origin and nature of these seemingly dark-matter-dominated dwarf galaxies is still poorly understood and is the subject of intense research.  DSphs are possibly the most dark-matter-dominated objects in the Universe and may share a common halo mass scale of $4 \\cdot 10^7$ M$_{\\odot}$ (e.g., Gallagher \\& Wyse 1994; Mateo et al. 1998; Odenkirchen et al.\\ 2001; Kleyna et al.\\ 2001; Klessen, Grebel, \\& Harbeck 2003; Wilkinson et al.\\ 2004, 2006a).  This makes it particularly interesting to understand their evolutionary histories.  The dSph census in the Local Group remains incomplete owing to the intrinsically low surface brightnesses of these gas-deficient galaxies.  The last three years alone saw the discovery of 13 new Local Group dSph candidates (Zucker et al.\\ 2004, Zucker et al.\\ 2006a, 2006b, 2006c; Willman et al.\\ 2005; Belokurov et al.\\ 2006a, 2006b, Martin et al. 2006), seven of which are new Milky Way companions.  This increases the number of known Galactic dSph satellites to 15, omitting accreted companions and additional objects whose nature remains yet to be confirmed (Grebel, Gallagher, \\& Harbeck 2006).  More than 75\\% of {\\em all} Milky Way companions are found within 150 kpc; only three additional dSphs believed to be Galactic satellites are found at larger distances. Leo\\,II, the dSph studied in the current paper, is one of the latter.  Leo\\,II was discovered by Harrington \\& Wilson (1950).  A number of mainly photometric studies were conducted in the subsequent decades. Earlier studies based on resolved color-magnitude diagrams (CMDs) placed Leo\\,II at a distance of $\\sim 230$ kpc (e.g., Demers \\& Harris 1983; Bellazzini, Gennari, \\& Ferraro 2005), $\\sim 215$ kpc (Demers \\& Irwin 1993; Lee 1995), and $\\sim 205$ kpc (Mighell \\& Rich 1996). These studies used different values for the (low) amount of Galactic foreground reddening, different mean metallicities, and different methods to derive the distance of Leo\\,II, e.g., the apparent luminosity of the horizontal branch or of the tip of the red giant branch (TRGB).  In the current study, we adopt the TRGB distance of $233 \\pm 15$ kpc obtained by Bellazzini et al.\\ (2005), which excludes the contribution of luminous asymptotic giant branch (AGB) stars and also shows consistency with the horizontal branch (HB) locus.  If this distance is representative of Leo\\,II's distance from the Milky Way during most of its existence, it should have been little affected by tidal stripping and other environmental effects exerted by the Galaxy.  Leo\\,II is a dSph galaxy with prominent old and intermediate-age populations.  The presence of old populations was first established by Swope (1967), who reported the detection of a red HB and of RR Lyrae variables.  Deeper CCD exposures by Demers \\& Irwin (1993) revealed the existence of a less well populated blue HB.  The most comprehensive variability study is that of Siegel \\& Majewski (2000), which presents 148 RR Lyrae stars and four anomalous Cepheids.  The RR Lyrae stars indicate a dominant metallicity of $-1.9$ dex with a metal-poor tail out to $-2.3$ dex.  The presence of intermediate-age populations was first inferred from the detection of luminous carbon stars in Leo\\,II (Aaronson, Olszewski, \\& Hodge 1983; Aaronson \\& Mould 1985, Azzopardi et al.\\ 1985).  Five out of the seven C stars and C star candidates probably belong to an intermediate-age population judging from their position above the TRGB, although the colors and luminosities of these stars suggest that they may be 3--4 Gyr older than the C stars detected in Leo\\,I (see Lee 1995 for details).  Moreover, a number of luminous AGB stars and a red clump were detected in Leo\\,II, which demonstrate the presence of an intermediate-age population beyond any doubt (Lee 1995; Mighell \\& Rich 1996).  Hence the stellar age tracers show that Leo\\,II experienced an extended star formation (SF) history.  The SF history of Leo\\,II was quantified using deep Hubble Space Telescope (HST) imaging extending below the old main-sequence turnoff of this galaxy and theoretical isochrones.  Mighell \\& Rich (1996) found a mean age of $9\\pm1$ Gyr with SF lasting for more than 7 Gyr.  They conclude that about half of the stars in Leo\\,II formed between 8 and 12 Gyr ago.  The oldest stars in Leo\\,II are as old as the oldest Galactic globular clusters, a common feature for all Galactic dSphs and other nearby galaxies (Grebel \\& Gallagher 2004).  Hernandez, Gilmore, \\& Valls-Gabaud (2000) suggest that the peak SF activity occurred $\\sim 8$ Gyr ago and that SF ended $\\sim 6$ Gyr ago.  They excised the HB populations from their CMD analysis, and their SF history does not reproduce the old populations of Leo\\,II.  Dolphin's (2002) CMD-fitting algorithm reveals an extended SF history starting some 15 Gyr ago and extending until 2 to 4 Gyr ago, although very little activity occurred within that age bin.  The main SF activity takes place in the 8 to 15 Gyr bin according to Dolphin (2002).  All three photometric analyses agree in finding an extended SF history for Leo\\,II that lasted many Gyr, and that the main SF activity occurred at relatively high ages of 8 Gyr or more.  In a recent,  shallower ground-based study Bellazzini et al.\\ (2005) discuss the detection of red giant branch bumps in Leo\\,II.  The luminosity of the main bump indicates that the majority of the stars in Leo\\,II is $\\sim 4$ Gyr younger than regular Galactic globular clusters of comparable low metallicity to the dominant stellar population in Leo\\,II. They also demonstrate that the red HB and red clump stars in Leo\\,II are more centrally concentrated than the ancient HB population, akin to the population gradients found in many dSphs (e.g., Stetson, Hesser, \\& Smecker-Hane 1998; Hurley-Keller, Mateo, \\& Grebel 1999; Da Costa et al.\\ 2000; Harbeck et al.\\ 2001).  Interestingly, in all three deep HST CMD studies the isochrone fits require a relatively high metallicity of $\\sim -1.6$ dex or even $\\sim -1.1$ dex.  Indeed, when comparing isochrones for old populations to globular clusters with metallicity determinations on the traditionally used Zinn \\& West (1984) scale, the isochrones favor too high a metallicity for metal-poor populations (Grebel 1997, 1999).  As discussed by Dolphin (2002), this can be in part circumvented by using the Carretta \\& Gratton (1997) scale for globular clusters instead (see also discussion in Bellazzini et al.\\ 2005).  Nevertheless, the age-metallicity degeneracy that plagues purely photometric derivations of SF histories from resolved stellar populations remains.  The age-metallicity degeneracy can be resolved through spectroscopic metallicities.  The first spectroscopic metallicity determination for Leo\\,II was carried out by Suntzeff et al.\\ (1986), who measured three red giants in this galaxy and found [Fe/H] $\\sim -1.9$ dex.  Their metallicity scale corresponds to the Zinn \\& West (1984) scale.  More recently, Bosler et al.\\ (2004) spectrometered 41 red giants in Leo\\,II and found a mean metallicity of $-1.57$ dex with a spread ranging from $-1.26 >$ [Fe/H] $> -2.32$ on the Carretta \\& Gratton (1997) scale.  In our current paper, we present spectroscopic metallicity measurements for 52 red giants.  \\textsection 2 introduces the data and briefly describes the reduction strategy. In \\textsection 3 the calibrations of the measured CaT equivalent widths (EWs) onto a reference metallicity scale are presented and compared to existing data sets.  The resulting metallicity distribution function (MDF) is analyzed in \\textsection 4 and \\textsection 5 discusses this MDF in the context of simple models of chemical evolution. In \\textsection 6 we turn to the derivation of stellar ages and to the age-metallicity relation of Leo\\,II.  Finally, in \\textsection 7 we summarize our findings. ",
        "conclusions": "After Leo\\,I, Leo\\,II is the most remote dSph companion of the Milky Way. Its prevailing intermediate-age stellar population as well as its significant population of old stars make it an ideal testbench to study  star formation that extended over many billions of years. Existing photometric and spectroscopic data as well as our own results suggest that SF activity may have ceased only 2 Gyr ago (Dolphin 2002; Bosler et al.\\ 2004).  In the framework of a VLT Large Programme tailored to investigate kinematical and chemical evolutionary aspects of Galactic dSphs we have determined CaT metallicities for 52 red giants in this system, thereby doubling the largest previously published data sets (Bosler et al.\\ 2004; B06). Our targets cover the entire area of the galaxy, thus allowing us to efficiently trace potential spatial metallicity variations.  By employing the well-established CaT technique we showed that Leo\\,II's MDF peaks at $-$1.74\\,dex (CG scale), in agreement with previous spectroscopic and photometric estimates (e.g., Bellazzini et al.\\ 2005 and references therein). We infer a range of metallicities of $\\sim 1.3$ dex.  We emphasize that Leo\\,II does not contain any extremely metal-poor stars of comparably low metallicities as found in the Galactic halo field stars. The lowest metallicities inferred in our study are similar to those of metal-poor Galactic globular clusters. In the compilation  of Venn et al. (2004), 7\\% of the (kinematically selected) halo stars lie in the regime of [Fe/H]$\\le-2.5$ and only 1\\% lie below $-3$\\,dex. If these numbers were to be considered also representative of the dSphs' MDFs, we would expect to find about three to four stars more metal poor than $-2.5$\\,dex in Leo\\,II. A similar lack of extremely metal-poor stars has also been observed in other Galactic dSphs (e.g., Shetrone et al.\\ 2001; Pont et al.\\ 2004; Bosler et al.\\ 2004; Tolstoy et al.\\ 2004; Koch et al.\\ 2006a,b).  While these studies all tried to cover the full color range of the RGB to ensure that stars of very low or very high metallicities, if present, would not be missed, the aforementioned deficiency might still be blamed on small number statistics.  For want of adequately sophisticated chemical evolution models of the Leo\\,II dSph  itself, we compared the observed MDF to the model predictions for other Galactic dSphs from LM04. These models include the effects of supernovae of type Ia and II and galactic winds.  None of these models really fits Leo\\,II's MDF well, which may not be surprising given that each Galactic dSph has experienced its own unique evolutionary history (e.g., Grebel 1997, 1999).  LM04's model for the Carina dSph approximates the MDF of Leo\\,II best, and although Car experienced strongly episodic star formation its age range is comparable to that of Leo\\,II.  The range in [Fe/H] covered by our MDF, combined with the fact that there is no obvious trend of metallicity with the color of the red giants indicate that there is a broad range of ages present in Leo\\,II, a fact already known from earlier photometric CMD analyses. By means of isochrone fitting using our spectroscopic metallicities as an additional constraint, we could show that the ages span the full range  from 2 to 15\\,Gyr. There is a considerable fraction of ``younger'' and more metal-rich stars with ages of a few Gyr in our data, but one cannot exclude the possibility that these may in fact be AGB contaminants.  Our AMR suggests constant star formation across our entire age range, although we note that our measurement accuracy and the age resolution would not allow us to detect any potential gaps in Leo\\,II's SF history. The metallicity remained fairly constant for most of the active periods of Leo\\,II, indicating early star formation from pre-enriched gas and possibly continued rejuvenation by little-enriched gas until some 5 Gyr ago.  However, we caution that this qualitative interpretation is subject to the many uncertainties affecting the derivation of ages.  There appears to be no radial population gradient from the spectroscopic point of view, at least not in metallicity and at best weakly indicated in our inferred age distribution.  Bellazzini et al.\\ (2005) show that Leo\\,II follows the common trend of containing more centrally concentrated intermediate-age populations as compared to the old component. In the case of Leo\\,II the main driver of the population gradient appears to be age, since we have demonstrated that the AMR shows little evidence for metallicity variations (i.e., enrichment) across most of the age range."
    },
    "0609/astro-ph0609536_arXiv.txt": {
        "abstract": "GRO~J1655--40 is a galactic microquasar, i.e. a short-period binary with relativistic radio-jets and where the companion is a black-hole. A little after its discovery in 1994 by $BATSE$, a distance of 3.2$\\pm$0.2~kpc has been estimated, and at which the radio jets appeared superluminal. Since that time, this distance of GRO~J1655--40 has been discussed in many studies, often strengthening the value of 3.2~kpc, and used in numerous models. However, recently, \\citet{Foellmi-etal-2006b-astroph} used new VLT-UVES and {\\it published} photometric data to  show that GRO~J1655--40 must be closer than 1.7~kpc and that the accuracy of the main distance estimators for this source can be questioned. Still, the details on how really the distance of 3.2~kpc has been build was not fully clear. It is the purpose of this article to show that while the {\\it upper} limit to the distance is rather firm, the lower limit is not. We draw some conclusions about the new understanding we have of GRO~J1655--40, and finally present a new and promising method that can be used to determine the distance of GRO~J1655--40, and maybe many faint and embedded stars. ",
        "introduction": "Microquasars are short-period stellar binaries in which one of the component is a compact object, i.e. a neutron star or a black-hole. The presence of such object is causing (sometimes persistent) strong radio jets and X-ray flares. In some cases, the relativistic jets appear to be superluminal, i.e. having an apparent speed larger than the speed of light. Microquasars are galactic laboratories of high-energy phenomena, and they must be seen as part of a large paradigm where AGNs, microquasars and possibly gamma-ray bursts (GRBs) share similar physics \\citep{Mirabel-2004}. The main advantages of microquasars over their large-scale parents, is their short timescales, allowing much more detailed dynamical studies.  GRO~J1655--40 (a.k.a. Nova Sco 94) has been discovered as a Soft X-ray Transient (SXT) on July 27, 1994 with \\emph{BATSE} on board the \\emph{Compton Gamma Ray Observatory} \\citep{Zhang-etal-1994}. Strong evidence that the compact object in GRO~J1655--40 ~is a black hole was presented by \\citet{Bailyn-etal-1995b} \\citep[see also][]{Balucinska-Church-2001}. They obtained through spectroscopy a mass function of the secondary (i.e. the absolute minimum mass of the compact object) of $f$ = 3.16$\\pm$0.15 $M_{\\odot}$; that has been later updated by \\citet{Orosz-Bailyn-1997} to $f$ = 3.24$\\pm$0.09 $M_{\\odot}$. This value is above the theoretical maximum mass of a neutron star: $\\sim$3 $M_{\\odot}$ \\citep{Arnett-Bowers-1977} \\citep[but see also][ for a discussion on the dependance of this maximum mass with the equation of state]{Nauenberg-Chapline-1973,Prakash-etal-1988,vanKerkwijk-etal-1995}. Later, \\citet{Shahbaz-etal-1999} revised down the mass function below this limit ($f$ = 2.73$\\pm$0.09 $M_{\\odot}$), but published the masses of the compact object and the secondary star: 5.5--7.9 $M_{\\odot}$ and 1.7--3.3 $M_{\\odot}$ respectively thanks to the inclination angle, clearly confirming the presence of a black-hole. The jets of GRO~J1655--40 were observed in radio, and given a distance of about 3 kpc, they appeared to be superluminal \\citep[see e.g.][]{Rees-1966}: 1.5$\\pm$0.4~$c$ \\citep{Tingay-etal-1995}, 1.05~$c$ \\citep{Hjellming-Rupen-1995}, where $c$ is the speed of light.  All studies until the most recent ones \\citep[e.g.][]{Barret-etal-1996,Regos-etal-1998,vanderHooft-etal-1998,Phillips-etal-1999,Shahbaz-etal-1999,Kuulkers-etal-2000,Soria-etal-2000,Greene-etal-2001,Combi-etal-2001,Buxton-Vennes-2001,Yamaoka-etal-2001,Gierlinski-etal-2001,Kubota-etal-2001,Remillard-etal-2002,Kong-etal-2002,Kobayashi-etal-2003,Stevens-etal-2003,Willems-etal-2005,Brocksopp-etal-2006} usually quote the distance of 3.2$\\pm$0.2~kpc determined by \\citet{Hjellming-Rupen-1995}, although \\citet{Mirabel-etal-2002} pointed out that a distance of about 1~kpc cannot be ruled out with the current data. At the time of writing, the latest discovery about GRO~J1655--40 is the fact that magnetic fields are the only explanation for the launch of jets in this object by \\citet{Miller-etal-2006}, who also use the distance of 3.2 kpc. The purpose of this paper is to complement the work of \\citet{Foellmi-etal-2006b-astroph}, who have shown that the distance of GRO~J1655--40 must be smaller than 1.7~kpc, but focusing on how the distance of 3.2~kpc has been determined, and to show that this value is far from being firm. Other distance methods will be discussed in a forthcoming paper. ",
        "conclusions": "It is important to stress that the exact and accurate distance of GRO~J1655--40 is still an open question. Before the direct measurement by the european satellite \\emph{GAIA}, it should be possible to test the distance value using methods that have never been used so far. In particular, we want to mention a new (and exploratory) possibility using bisectors.  First, \\citet{Gray-2005} has shown that the bluemost point of single-line bisector is an indicator of the luminosity for late-type stars. This very interesting relation is weakened by the fact that high-quality single-line bisectors require a high spectroscopic resolution, and a very high Signal-to-Noise ratio (above, say, 300). However, the recent study by \\citet{Dall-etal-2006} has shown that the bisector of the cross-correlation function (CCF) can be used as much the same way as single-line bisectors. The combination of these two results could lead to a first-order \"direct\" method of estimating the luminosity of the secondary star in GRO~J1655--40. Given the much larger brightness of GRO~J1655--40 in the near infrared (NIR, $m_K \\sim 13$) than in the optical ($m_V \\sim 17$), a CCF bisector could be obtained in a reasonable amount of time with the new near-IR echelle spectrograph on the VLT: CRIRES, which is being commissioned at the time of writing.  Such investigation is already underway. This new method might also proved to be useful for many other faint and embedded stars in star-forming regions."
    },
    "0609/astro-ph0609309_arXiv.txt": {
        "abstract": "Deep X-ray imaging spectroscopy of the bright pulsar wind nebula 3C58 confirms the existence of an embedded thermal X-ray shell surrounding the pulsar \\psr. Radially resolved spectra obtained with the \\xmm\\ telescope are well-characterized by a power-law model with the addition of a soft thermal emission component in varying proportions. These fits reproduce the well-studied increase in the spectral index with radius attributed to synchrotron burn-off of high energy electrons. Most interestingly, a radially resolved thermal component is shown to map out a shell-like structure $\\approx 6^{\\prime}$ in diameter. The presence of a strong emission line corresponding to the \\ion{Ne}{9} He-like transition requires an overabundance of $\\sim 3 \\times [{\\rm Ne/Ne}_{\\sun}]$ in the Raymond-Smith plasma model. The best-fit temperature $kT \\sim 0.23$ keV is essentially independent of radius for the derived column density of $N_{\\rm H} = (4.2\\pm 0.1) \\times 10^{21}\\ {\\rm cm}^{-2}$. Our result suggests that thermal shells can be obscured in the early evolution of a supernova remnant by non-thermal pulsar wind nebulae emission; the luminosity of the 3C58 shell is more than an order of magnitude below the upper limit on a similar shell in the Crab Nebula. We find the shell centroid to be offset from the pulsar location. If this neutron star has a velocity similar to that of the Crab pulsar, we derive an age of 3700 yr and a velocity vector aligned with the long axis of the PWN. The shell parameters and pulsar offset add to the accumulating evidence that 3C58 is not the remnant of the supernova of CE 1181. ",
        "introduction": "A long standing puzzle in supernova physics is the apparent lack of an associated thermal shell surrounding a few young pulsars with bright relativistic wind-powered nebulae (PWNe). Though relatively rare, with fewer than ten examples known, the lack of a supernova remnant shell is at odds with our current understanding of supernova evolution and neutron star formation. The canonical examples, the Crab and 3C58, are both putative historical remnants less than 1000 yrs old.  To date, no evidence has been found for thermal emission associated with the Crab Nebula pulsar despite repeated searches (Seward et al. 2006 and references therein). However, recent \\xmm\\ and \\chandra\\ observations have identified thermal emission within the 3C58 remnant (Bocchino et al. 2001; Slane et al. 2004).  The availability of a large quantity of XMM archival data on this source derived from calibration observations allows us to construct very sensitive images and spectra to explore the evidence for a surrounding SNR shell with unprecedented sensitivity.  The morphology of 3C58 is well-documented in both the X-ray and radio bands (Slane et al. 2004 and Bietenholz et al. 2006 and references therein). It is center-filled with a 66~msec pulsar near the center, axisymmetric lobes, and an elaborate network of wisps and filaments. Failure to find a radio shell around 3C58 led Reynolds \\& Aller (1985) to conclude that there is no evidence for its interaction with an external medium.  Coupled with the low velocities of the optical filaments (Fesen 1983) compared to the mean expansion velocity required for the remnant to reach its current extent in 820 yr, these results led to suggestions that the remnant was considerably older. Over the past twenty years, further evidence has accumulated that the remnant's age is inconsistent with an origin in SN1181. The images and spatially resolved spectra we present below add incrementially to the case and provides a testable prediction concerning the pulsar velocity which could support or refute the historical association.  \\begin{deluxetable*}{lccccc} \\tabletypesize{\\small} \\tablewidth{0pt} \\tablecaption{\\xmm\\ Observation log for 3C58$^a$} \\tablehead{ \\colhead{ObsID} & \\colhead{Date} & \\colhead{Pointing}& \\colhead{pointing} & \\colhead{Exposure}       & \\colhead{Exposure} \\\\ \\colhead{     } & \\colhead{}   & \\colhead{R.A}   & \\colhead{Decl.}    & \\colhead{EPIC-MOS$^b$} & \\colhead{EPIC-pn}\\\\ \\colhead{     } & \\colhead{(UT)} & \\colhead{(J2000)} & \\colhead{(J2000)}  & \\colhead{(ks)}           & \\colhead{(ks)}     \\\\ } \\startdata 153752201\\sp & 2002-09-11~04:29:35 & $2^{\\rm h} 4^{\\rm m}43.6^{\\rm s}$ & $+64^{\\circ}51^{\\prime}13{\\farcs}2$ & 36777 & 15579 \\\\ 153751801\\sp & 2002-09-11~13:09:58 & $2^{\\rm h} 5^{\\rm m}03.0^{\\rm s}$ & $+64^{\\circ}47^{\\prime}38{\\farcs}6$ & 39561 & 16328 \\\\ 153752501\\sp & 2002-09-11~20:03:41 & $2^{\\rm h} 5^{\\rm m}23.4^{\\rm s}$ & $+64^{\\circ}43^{\\prime}52{\\farcs}3$ & 35195 & 14721 \\\\ 153752401\\sp & 2002-09-12~03:20:44 & $2^{\\rm h} 6^{\\rm m}32.4^{\\rm s}$ & $+64^{\\circ}48^{\\prime}06{\\farcs}8$ & 26184 & \\dots \\\\ 153752101    & 2002-09-13~04:22:11 & $2^{\\rm h} 5^{\\rm m}38.0^{\\rm s}$ & $+64^{\\circ}49^{\\prime}40{\\farcs}0$ & 17363 &  3022 \\\\ 153751701    & 2002-09-13~10:57:39 & $2^{\\rm h} 5^{\\rm m}18.9^{\\rm s}$ & $+64^{\\circ}53^{\\prime}23{\\farcs}7$ & 13515 &  6042 \\\\ 153751901\\sp & 2002-09-13~17:51:22 & $2^{\\rm h} 5^{\\rm m}13.1^{\\rm s}$ & $+64^{\\circ}51^{\\prime}41{\\farcs}8$ & 34114 & 14569 \\\\ 153752001\\sp & 2002-09-13~23:45:05 & $2^{\\rm h} 5^{\\rm m}52.6^{\\rm s}$ & $+64^{\\circ}55^{\\prime}27{\\farcs}7$ & 48631 & 21038 \\\\ \\enddata \\tablenotetext{a}{\\footnotesize All observations were obtained in pointed imaging mode with the thin transmission window in the focal plane; the EMOS and EPN data were collected with the {\\tt PrimeFullWindow} and {\\tt PrimeLargeWindow} submodes, respectively} \\tablenotetext{b}{\\footnotesize  EPIC exposure times are after filtering; the EPIC-MOS exposure times are for the sum of the two MOS cameras.} \\label{ta:log} \\end{deluxetable*}   In \\S2 we describe the available observations and discuss their analysis. Section 3 presents the detailed spatially resolved spectroscopy which allows us to detect clear evidence for a shell of thermal emission, overabundant in Neon, and extending to the edge of the synchrotron nebula; we interpret this result in \\S4, concluding that the pulsar is offset from the center of the symmetrically expanding shell implying a velocity aligned with the symmetry axis of the remnant. If the pulsar has a Crab-like velocity, this offset provides another argument against an association with SN1181. ",
        "conclusions": "In Figure~6, we illustrate the relative contributions of the thermal and non-thermal components as a function of radius. The thermal shell is clearly present. Note, however, that this thermal emission represents less than~4\\% of the observed (absorbed) flux of the remnant in the $0.5-10$~keV band, explaining why it has proven so elusive to date. Figure~7 displays the radial dependence of the shell temperature and Neon abundance in the remnant shell -- both are flat. In contrast, the power law spectral index $\\Gamma$ monotonically increases from $\\Gamma = 1.9$ at the remnant center to $\\Gamma = 2.8$ at the edge (and $\\Gamma = 2.9$ in the eastern lobe). This latter result is generally consistent with earlier work (Torii et al 2000; Slane et al. 2002) although, owing to the much larger number of photons collected and the spectral-spatial decomposition we have performed, the measurements extend to a greater radius and differ in detail, remaining flat between $2.2^{\\prime}$ and $3.2^{\\prime}$.  The widths of the annuli provide largely independent measurements of the spectrum at each radius given the size of the PSF. Some contamination is present from the inner, brighter rings to the outer ones, however. Since the dominant power law is flatter in the central regions, photons scattered to the outer regions bias the power law measured there. We have determined the degree of contamination from power law photons in each region by using the PSF binned in $0.5^{\\prime}$ annuli and the geometrical intersection of those annuli with each extraction annulus (the minimal dependence of the PSF on photon energy allows us to ignore this complicating factor). The corrections to the measured power law slopes reported in Table~2 are modest, albeit outside the formal errors. In the central $0.5^{\\prime}$ circle, the power law spectral index falls from 2.02 to 2.00, while in the first four annuli beyond this, the indices changes as follows: 2.29 becomes 2.38, 2.45 becomes 2.50, 2.55 becomes 2.58, and 2.73 becomes 2.78.  Thus, the rise in spectral slope is slightly steeper than shown in Figure~7.  PSF contamination effects are negligible for the thermal component since the count rate is so low.  The parameters we derive for the shell are in excellent agreement with those inferred in the early work of Bocchino et al. (2001) and with those found from a portion of the shell in the deep Chandra observation of 3C58 (Slane et al. 2004): the temperature and Neon abundance are identical at $kT=0.23$~keV and~3.2 solar, respectively. While these previous authors have cited the Neon overabundance as evidence that we are seeing emission from the ejecta, a growing body of work (Cunha et al. 2006 and references therein) suggests that the solar abundance of Neon has been systematically underestimated by a factor of 2.5~to~3. In that our fits do not require the statistically marginal enhancement of a factor of two in the Mg abundance reported by Slane et al. (2004), we argue that the provenance of the thermal emission remains an open question.  The total observed X-ray luminosity we derive for the shell component is $5.9 \\times 10^{32} d_{3.2}^2$~ergs s$^{-1}$, within the range found in Bocchino et al. (2001). This is significantly below the upper limit derived by Seward et al.  (2006) for the missing Crab Nebula shell. The total mass is $\\sim 0.77 M_{\\odot}$ (for a mean molecular weight of 0.6 which may not be appropriate if the bulk of the radiating material is ejecta).  Evidence has been accumulating for some time that, despite the apparently robust conclusions drawn from the historical records that 3C58 is coincident with SN1181 (Stephenson and Green 1999), the remnant properties are inconsistent with such a young age.  Bietenholz (2006) provides a comprehensive summary of the arguments against an association, as well as a list of alternative scenarios. Most of the evidence suggests a remnant age of $\\sim 3000-5000$ years (e.g., Chevalier 2005). Our spatially resolved image of the shell adds further to the case for an age $>820$~yrs.  Using the X-ray emitting mass we have derived and the analytic models of PWN evolution from Chevalier (2005) we can set limits on the apparent age of the remnant. From Eq.~27 of Chevalier (2005), we can combine the observed supernova remnant radius of $2.8 \\ d_{3.2}$~pc and the pulsar timing-derived parameter $\\dot E = 2.7 \\times 10^{37}$ erg s$^{-1}$ to obtain the age $t_3$ in units of 1000~yrs,  $$t_3 \\approx 1.8 \\ M_{ej}^{0.4} \\ E_{51}^{-0.2}$$  \\noindent where $M_{ej}$ is the ejected mass in units of solar mass $M_{\\odot}$. If we make the extreme assumption that all of the observed $0.77 M_{\\odot}$ of X-ray-emitting gas is ejecta, and that it represents some fraction $\\leq 1$ of the total ejected mass of the supernova, we find $t_3 \\simgt 1.62 \\ E_{51}^{-0.2}$.  Thus, the pulsar must be at least twice the age of SN1181 for reasonable explosion energies and, for a more plausible ejected mass of $\\sim 5 M_{\\odot}$, the age is 3400~yrs. Alternatively, if we assume all of the observed gas has been swept up by the expanding PWN, from Chevalier's (2005) eq. 28, we find $t = 5100$~yrs (cf. the pulsar's characteristic age of 5390~yrs).  The pulsar's location with respect to the circular shell of thermal X-ray emission provides a potential test that could resolve the age controversy.  The location of the pulsar is marked with a crosshair in Figure 8. It is apparent that this is not the center of the shell; a circle best matching the observed thermal emission is centered $27^{\\prime\\prime} \\pm 5^{\\prime\\prime}$ west (and slightly north) of the pulsar's location. For an age of 821 years, a distance of 3.2~kpc, and spherical expansion (an assumption supported by the fact we see no thermal emission in the eastern elongation of the remnant), this requires a two-dimensional pulsar velocity of $500d_{3.2}$~km s$^{-1}$. While not impossible for a young neutron star, it is worth noting that if the object has the same tranverse velocity as the Crab pulsar, the remnant age would be $3750d_{3.2}$~yrs, consistent with other estimates. We further note that the implied trajectory for the pulsar aligns to within a few degrees of the long axis of the nebula, behavior commonly seen among PWNe. If Chandra survives for another decade, it will be possible to resolve this issue directly, as the implied proper motion of $0.5^{\\prime\\prime}$ for the high-velocity scenario will be directly measurable.  \\bigskip"
    },
    "0609/astro-ph0609623_arXiv.txt": {
        "abstract": "The extended X-ray emission from 'the Cap' region located \\timeform{11'} (11.6~kpc) above the disk of the starburst galaxy M82 has been observed with Suzaku and XMM-Newton. Owing to the good energy resolution and the large collecting area of the XIS on Suzaku, combined with similar properties of the EPIC instrument on XMM-Newton, we have clearly detected K-shell emission lines from O~VII, O~VIII, Ne~X, Mg~XI, Mg~XII and the Fe-L complex. Two optically-thin thermal plasma components are required to fit the observed X-ray spectra. We have determined the metal abundances of O, Ne, Mg, Si and Fe in this region for the first time. Their metal abundance ratios agree well with those of metal-poor stars and the model prediction of metals synthesized by type-II supernovae, but they are not consistent with the metallicities of type-Ia supernovae. This result is support for the idea that the origin of the metals in the Cap is type-II supernovae explosions occurring in the starburst regions in the M82 galaxy. We discuss the possible contribution from sputtered dust grains to the metals in the Cap. An emission line consistent with the C~VI transition of $n=4$ to $1$ at 0.459~keV is marginally detected, although it is not statistically significant at the 99\\% confidence level; the presence of this line would suggest charge-exchange processes in the Cap. ",
        "introduction": "% Starbursts inject a large amount of kinetic and thermal energy into intergalactic space via outflows enriched with heavy metals synthesized in the massive stars that result in supernova explosions. The metal-enriched ejecta mix with the intergalactic medium and now form part of the warm-hot intergalactic medium (WHIM) and intra-cluster medium (ICM). Using ASCA results, \\citet{FukazawaPASJ98_Si_Fe_ICM_ASCA}, \\citet{FinoguenovApJ00_ASCA_Cl_Metal} and \\citet{FinoguenovApJ01_ASCA_PoorCl_Metal} showed the importance of the metal enrichment of the ICM by type-II supernova activity in starbursts at early epochs. The entropy excess in the ICM above the self-similar prediction implies the importance of non-gravitational heating. One of the plausible heat sources is that from the superwinds which result from starbursts \\citep{PonmanNAT99_ICM_Entropy, LloydDaviesMN00_ICM_Entropy}. Thus, starburst activity is a key phenomenon when trying to understand the metal and thermal evolution of galaxies and clusters of galaxies. In spite of the fact that we cannot see the initial starbursts directly, the observation of hot plasma currently being produced by the nearby starburst galaxies can reveal the basic processes and give us insight into the same processes at high redshift.  Since M82 is a prototypical starburst galaxy and one of the brightest starburst galaxies in the local universe, it has been studied at almost all wavelengths, ranging from low frequency radio up to TeV $\\gamma$-rays. Essentially all X-ray astronomical satellites have observed M82. The extended X-ray structure basically consists of a multi-phase thermal plasma extending from the nuclear region toward the halo, mainly along the minor axis of the galaxy \\citep{WatsonApJ84_M82_HRI, Fabbiano88ApJ_M82_N253_Ein, StricklandAA97_M82_ROS}. The point sources in M82 include the hyper luminous X-ray source, M82 X-1, dominating X-ray luminosity above $\\sim 3$~keV \\citep{Matsu99PASJ_M82AGN, Ptak99ApJL_M82AGN, Matsumoto2001_M82_CXO, KaaretMN01_M82_CXO_HRC}.  The spectral and spatial structure of the thermal plasma component of M82 has been vigorously studied with CCDs or IGSPCs onboard ASCA, BeppoSAX, Chandra and XMM-Newton \\citep{Moran_Apj97_M82, Ptak1997_ASCA_M82, Tsuru1997_ASCA_M82, Cappi_AA99_SAX_M82_N253, WeaverApJ00_M82_N253_SBGII, GriffithsSci00_M82_ACIS, StevensMNL03_M82_EPIC, RanalliPH05_M82_XMM_GasAbundance}. The results from the observation with RGS were reported by \\citet{ReadMN02_M82_RGS} and \\citet{OrigliaApJ04_M82_Star_Gas_Abundance}. The X-ray spectra are well represented by a model consisting of multiple thermal plasma components (\\KT$\\sim 0.3-1.5$~keV) with absorption due to the cool matter in the M82 galaxy. The measurements of absolute metal abundances differ by a factor of several to ten among the measurements with ASCA, XMM-Newton/EPIC and RGS \\citep{Tsuru1997_ASCA_M82, ReadMN02_M82_RGS, StevensMNL03_M82_EPIC, OrigliaApJ04_M82_Star_Gas_Abundance}. The abundances of O and Fe are lower than those of Si and S by a factor of several, a common result from the ASCA and XMM-Newton/RGS observations. Any combination of the metal production of type-Ia and type-II supernovae cannot explain the abundance ratios among the metals. Recently, \\citet{RanalliPH05_M82_XMM_GasAbundance} have reported that the metal abundances and their ratios significantly change from the nucleus toward the halo. Additionally, the Fe abundance of $\\sim 0.3$ solar was obtained for the high temperature diffuse plasma in the core of M82 \\citep{GriffithsSci00_M82_ACIS}. Thus, the metal abundances of the X-ray plasmas in the M82 galaxy itself are still under debate.  Searching for very extended X-ray emission far beyond the halo of the M82 galaxy, \\citet{Tsuru_PASJ90_GingaM82halo} have claimed evidence for X-ray emission extending for 100~kpc around M82 from scanning observations with the Ginga satellite. Extended emission named 'the Cap' in the \\HALPHA\\ and soft X-ray bands was discovered at approximately $11'$, corresponding to 11.6~kpc at the distance of 3.63~Mpc, to the north of the center of M82 \\citep{Devine1999_M82_NalphaCap, Lehnert1999_ROSAT_cap}. Comparing the \\HALPHA\\ and the X-ray morphology, \\citet{Lehnert1999_ROSAT_cap} have revealed a close relationship between the two emissions. This X-ray emission has been confirmed with XMM-Newton by \\citet{StevensMNL03_M82_EPIC}. \\citet{Lehnert1999_ROSAT_cap} explained its X-ray emission as the result of shock heating associated with an encounter between the starburst-driven galactic superwind and a large photoionized cloud in the halo of M82. \\citet{Hoopes2005_GALEX_M82_N253} detected UV emission at the Cap with GALEX and suggested that the most likely emission mechanism is scattering of stellar continuum from the starburst in M82 by dust in the Cap. The X-ray spectra of the Cap region obtained with ROSAT and XMM-Newton can be fitted with a thermal plasma model with a temperature of \\KT$=0.80\\pm 0.17$~keV or $0.65^{+0.69}_{-0.62}$~keV \\citep{Lehnert1999_ROSAT_cap, StevensMNL03_M82_EPIC}. The metal abundance of this plasma is key information in unraveling the origin and fate of the ejecta. However, no successful measurement has been made so far due to the limited spectral resolution and/or statistics.  Thus, we made a 109~ks observation of the Cap with the Suzaku XIS. The XIS has good energy resolution, especially in the low energy band from 0.3 to1.0~keV, and low detector background. These properties have enabled us to obtain the best quality X-ray spectrum in the Cap region and to elucidate the nature of the plasma, and especially to measure the metal abundances. We also analyzed the archival data of the XMM-Newton observation of M82, in which the Cap was also observed in the same field of view. Throughout this paper, we adopt a distance of $D=3.63$~Mpc to M82 \\citep{Freedman1994_HST_M81_Dist}. The solar metal abundances adopted in this paper are shown in table~\\ref{tab:SolarAbundance}.  \\begin{table} \\caption{The numbers of atoms per a H atom for the solar metal abundances adopted in this paper \\citep{AndersGeCoA89_MeteoSolarAbundance}. } \\label{tab:SolarAbundance} \\begin{center} \\begin{tabular}{cccccc} \\hline\\hline atom & number &atom & number & atom  & number \\\\ \\hline He & $0.0977$             & C  & $3.63\\times 10^{-4}$ & N  & $1.12\\times 10^{-4}$ \\\\ O  & $8.51\\times 10^{-4}$ & Ne & $1.23\\times 10^{-4}$ & Mg & $3.80\\times 10^{-5}$ \\\\ Si & $3.55\\times 10^{-5}$ & S  & $1.62\\times 10^{-5}$ & Ar & $3.63\\times 10^{-6}$ \\\\ Ca & $2.99\\times 10^{-6}$ & Fe & $4.68\\times 10^{-5}$ & Ni & $1.78\\times 10^{-6}$ \\\\ \\hline \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": "% \\subsection{Metal Abundance Ratios} Using the XIS and EPIC data, we demonstrated that the X-ray plasma of the Cap region consists of two thermal components. The high temperature component has \\KT$=0.63\\pm 0.02$~keV, which agrees with the previously reported ones of \\KT$=0.80\\pm 0.17$~keV and \\KT$=0.65^{+0.04}_{-0.03}$~keV \\citep{Lehnert1999_ROSAT_cap, StevensMNL03_M82_EPIC}. The emission integral of this component is also consistent with other observations. The existence of the O~VII emission line and the low temperature component with \\KT$=0.2$~keV is a new finding of this paper. It would reflect the nature of the plasma in the Cap region. However, we do not go into details in this paper.  The XIS image shows the existence of  diffuse emission in the connecting region between M82 and the Cap (figure~\\ref{fig:Suzaku_BI_Image}), as already pointed out by \\citet{StevensMNL03_M82_EPIC}. This may be expected from a large number of sequential supernovae, mostly type-II as a result of starburst activity occurring in M82. Thus, one would expect that the X-ray plasma in the Cap originates from the type-II supernovae in M82, a hypothesis which can be tested with the pattern of relative ratios of the metal abundances.  Thanks to the large effective areas and good energy resolution of the XIS and EPIC in the soft X-ray band, we successfully derived the metal abundances of O, Ne, Mg, Si and Fe in the X-ray plasma of the Cap for the first time. The abundances of O, Ne, Mg and Si are $1\\sim 2$ solar, while that of Fe is about 0.5 solar, which is generally consistent with the idea of the plasma originating from type-II supernovae. The emissivities of the lines depend on the plasma temperature, determined by the shape of the Fe L line complex, suggesting that the abundances obtained might be strongly coupled with the Fe abundance. Therefore, we have checked the reliability of the results. Figure~\\ref{fig:Fe_O_Mg_Ne_Si_ConfCont} shows the confidence contours for the measured abundances of  O, Ne, Mg and Si against Fe. This shows that the metal abundance ratios of O, Ne and Mg to Fe are higher than unity in units of the solar abundance, at a confidence level of 99\\%. An abundance ratio of unity between Si and Fe is permitted at a confidence level of 99\\%, but still rejected at the level of 90\\%. Thus, it is safely concluded that the Fe abundance is significantly lower than the others.  We have compared the metal abundance ratios of the Cap with those synthesized by type-II and type-Ia supernovae (figure~\\ref{fig:MetalAbundances}). We adopted the W7 model in \\citet{Nomoto1984_SN1a_Metal} and \\citet{Thielemann1986_SN1a_Metal} for the metal masses synthesized by type-Ia supernovae. The model plot for type-Ia supernovae normalized with the Si abundance at the solar value shows clear disagreement with our observational result. This result is easily expected since O, Ne of Mg are not synthesized by type-Ia supernovae, while the metal abundances actually observed are much higher than that of Fe.  For type-II supernovae, we used the results of $m_u=50$~\\UNITSOLARMASS\\ by \\citet{Tsujimoto1995_SN2_Metal}, in which the averages of the synthesized masses are 1.8, 0.23, 0.12, 0.12, 0.041, 0.0080 and 0.0091~\\UNITSOLARMASS\\ for O, Ne, Mg, Si, S, Ar and Fe, respectively. We plot the metal abundance ratios synthesized by type-II supernovae, normalized with the Ne abundance at the solar value, in figure~\\ref{fig:MetalAbundances}. Comparison with our results shows that metal synthesis by type-II supernovae agrees well with the observed metal abundance ratios in the Cap.  Another good (experimental) sample for metal synthesis by type-II supernovae comes from observations of metal-poor stars; these observations are free from the ambiguities resulting from the different supernova models. The chemical composition of these old stars reflects that of the gas from which they formed in the early universe, where type-II supernovae dominated type-Ia. Table~6 in \\citet{ClementiniMN99_Ab_MetalPoorStar} shows the average metal abundance ratios in metal-poor stars obtained with Hipparcos and ground-based observations. After the adjustment of the definition of the solar abundance, the metal abundance ratios relative to Fe in metal-poor stars are $3.1\\sim 4.4$, $3.5$ and $2.4\\sim 3.0$ for O, Mg and Si, respectively. Figure~\\ref{fig:MetalAbundances} shows that the metal abundance ratios of the metal-poor stars and the Cap agree well with each other. Note that the absolute metal abundances of the metal-poor stars in the figure are meaningless since they are normalized so that the abundance of Fe is $0.4$ solar.  Thus, the results of the comparison with models of type-II supernovae and metal-poor stars are consistent with the idea that the X-ray plasma in the Cap originates from type-II supernovae. On the other hand, type-Ia supernovae are not a major contributor to the metals in the Cap. Thus, this naturally leads to the suggestion that most of the metals ejected by M82 are synthesized by type-II supernovae occurring through its starburst activity.  \\begin{figure} \\begin{center} \\FigureFile(150mm,120mm){20060710_MetalContour_v0_ph.eps} \\end{center} \\caption{The confidence contours for the O, Ne, Mg and Si abundances relative to the Fe abundance, obtained from the fits using the 2vMEKAL model. The contours are given at confidence levels of 68\\%, 90\\% and 99\\%. The dashed lines show the solar abundance ratio between the metals. } \\label{fig:Fe_O_Mg_Ne_Si_ConfCont} \\end{figure}    \\subsection{Comparison with the Mass of the Dust in the Cap} \\citet{Hoopes2005_GALEX_M82_N253} detected UV emission at the Cap with GALEX. They suggest that the most likely UV emission mechanism is scattering of stellar continuum from the starburst in M82 by dust in the Cap. They also argue that the dust may have been pushed out of M82 by the starburst, or stripped from either M82 or M81 by the tidal interaction between the two galaxies. The former picture is supported by the SCUBA observation showing the outflowing dust from M82 \\citep{AltonAA99_SBG_DustOutFlow}. The Spitzer result by \\citet{EngelbrachtApJL06_M82_MIR} preferentially supports the latter assumption.  Taking the composition of dust grains into account, Si and/or Fe with significant masses could exist in the form of dust in the Cap. Since the mass of the dust is not reported by \\citet{Hoopes2005_GALEX_M82_N253}, it is impossible to make direct comparison between the metal masses in the X-ray plasmas and those in the dust at the moment. An upper limit on the \\HI\\ column density of $2.7\\times 10^{19}$~\\UNITNH\\ is given by \\citet{Yun1993_M82_HI} and \\citet{Yun1994_M81group_HI}. Adopting a volume of $3.7\\times 3.7\\times 0.9{\\rm kpc}^3$ for the emission region of the Cap following \\citet{Lehnert1999_ROSAT_cap}, the upper limit on the mass of the \\HI\\ cloud is $\\sim 7\\times 10^5$~\\UNITSOLARMASS. Assuming a gas-to-dust ratio of 100, the upper limit on the dust mass is $\\sim 7\\times 10^3$~\\UNITSOLARMASS.  Dust grains are sputtered in an X-ray plasma with a time scale of $t_{\\rm sp} \\sim 10^8\\left(a/0.1{\\rm \\mu m}\\right)\\left(n/10^{-3} ~{\\rm cm^{-3}}\\right)^{-1}{\\rm yr}$ \\citep{YamadaPASJ05_Dust_ICM}, where $a$ and $n$ are the size of the dust grains and the density of the plasma, respectively, where $f_{\\rm x}$ is the filling factor. Estimating $n\\sim 5\\times 10^{-3}~{\\rm cm^{-3}}\\cdot f_{\\rm x}^{-0.5}$, the sputtering time scale is $t_{\\rm sp}\\sim 2\\times 10^7{\\rm yr} \\cdot f_{\\rm x}^{0.5}$. This is comparable with the travel time scale of $\\sim 1.5\\times 10^7{\\rm yr}$ from M82 to the Cap at the shock speed of $\\sim 740~{\\rm km\\ s^{-1}}$, where this speed is required to produce the plasma with \\KT$\\sim 0.63$~keV, following the estimation in \\citet{Lehnert1999_ROSAT_cap}. This suggests that some fraction of the dust grains have been possibly sputtered and mixed into the X-ray plasmas.  The masses of Si and Fe in the plasma phase obtained with our observation are $1.4\\times 10^3\\cdot f_{\\rm x}^{0.5}$~\\UNITSOLARMASS\\ and $1.8\\times 10^3\\cdot f_{\\rm x}^{0.5}$~\\UNITSOLARMASS, respectively, where $f_{\\rm x}$ is the filling factor. The sum of these masses is comparable with the upper limit on the mass of the dust. Thus, it would seem possible that the dust could contribute a significant part of the metals in the X-ray plasmas. In other words, dust depletion due to sputtering could have an impact on the measured abundances from our X-ray results. Quantitative observation of the dust is crucial to understanding the ejecta from starburst activity and its fate. Further analyses and observations of the dust in the UV and submillimeter bands are very important.   \\subsection{Line Emission through The Charge-Exchange Process} \\citet{LallementAA04_ChargeExchange} has pointed out that the charge-exchange process could be important in X-ray line emission from the Cap, where the ionized superwind from M82 can be assumed to collide with cool ambient gas located at the Cap, as suggested by \\citet{Lehnert1999_ROSAT_cap}. The charge-exchange process contributes only emission lines to the X-ray spectrum. However, we have detected continuum emission in the observed X-ray spectra of the Cap. This observational result suggests that charge-exchange is not the major process in the Cap, although a minor contribution cannot be denied.  According to the charge-exchange emission model for comets by \\citet{KrasnopolskySSRv04_Comet_Xray_EUV}, the $n=4$ to $1$ transition line is enhanced for C~VI at 0.459~keV. During the charge-exchange process, the electron from the donor neutral atom is first trapped at high levels, typically $n=4$ for carbon. The Suzaku and XMM-Newton spectra given in figures~\\ref{fig:SuzakuSimpleBGDsubtSpec}, \\ref{fig:SuzakuMethod2_CAP_NTE_fit} and \\ref{fig:XMMSpec} show a hint of an emission line around  0.46~keV. Thus, we added a narrow gaussian line around 0.459~keV to the 2vMEKAL model and examined its confidence. Figure~\\ref{fig:CVIconfCont} shows the confidence contour map for the narrow gaussian line. A local minimum is found at an energy consistent with 0.459~keV, which suggests marginal detection of the C~VI emission line although it is not statistically significant at the 99\\% confidence level.  We have estimated possible contributions to emission lines, especially O, due to the charge-exchange process. Following the model of \\citet{Lehnert1999_ROSAT_cap}, we again assume a box with dimensions of $3.7\\times 3.7\\times 0.9{\\rm kpc}^3$ for the Cap. An upper limit of $2\\times 10^{-3}{\\rm cm}^{-3}$ has been obtained for the density of the \\HI\\ cloud at the Cap \\citep{Yun1993_M82_HI, Yun1994_M81group_HI}. Thus, the upper limit on the column density of the \\HI\\ cloud along the direction of the superwind from M82 is estimated to be $8\\times 10^{18}\\ {\\rm cm}^{-2}$. Since the cross-section of the charge-exchange between an O ion and a H atom is $\\sim 10^{-15}-10^{-14}\\ {\\rm cm}^2$ \\citep{WegmannPSS98_CometXray_ChargeExchange}, it is expected that an O ion in the superwind  would encounter a \\HI\\ cloud thick enough for the charge-exchange process.  In this paper, we will not discuss the details of the process by which the superwind and O ions collide with the \\HI\\ cloud. Nevertheless, we can estimate an upper limit on the photon flux of O emission lines assuming an extreme case in which all O ions suffer the charge-exchange process with \\HI\\ and emit O K emission lines. Assuming a density of $\\sim 1\\times 10^{-3}\\ {\\rm cm^{-3}}$, the velocity of the superwind of $\\sim 740~{\\rm km\\ s^{-1}}$ and solar abundance of O in the plasma, the upper limit on the photon flux is estimated to be $5\\times 10^{-6}$~\\UNITPFLUX. The observed unabsorbed photon flux of O~VII is $\\sim 6\\times 10^{-6}$~\\UNITPFLUX (tables~\\ref{tab:SimultFit}). We note that the same logic gives an upper limit of $2\\times 10^{-6}$~\\UNITPFLUX\\ on the photon flux of C due to the charge-exchange process, and this is consistent with the upper limit on the one possible C~VI emission line at 0.459~keV. Thus, the upper limit on the photon flux of the O K emission lines due to charge-exchange is at the same order of the one actually observed. This suggests that the charge-exchange process may make a significant contribution to the observed O K emission lines.  We thus expect that the charge-exchange process could be one of the key aspects of the understanding of the physical processes and the measurement of the metal abundances in the Cap, and hence the ejecta from starburst galaxies and its cosmic evolution. More extensive studies of the possible charge-exchange are definitely necessary. This could not be achieved with the limited spectral resolution of X-ray CCDs. The realization of the non-dispersion type of spectrometer with ultra high spectral resolution, i.e. the $\\mu$-calorimeters onboard the NeXT and DIOS missions is essential \\citep{KuniedaSPIE06_NeXT, OhashiSPIE06_DIOS}.  \\begin{figure} \\begin{center} \\FigureFile(120mm,120mm){0730SX2ondoga_CVIcont_4c_ph.eps} \\end{center} \\caption{The confidence contour map of the C~VI emission line energy at 0.459~keV and its normalization obtained from the fitting with 2vMEKAL $+$ 1GAUS. The normalization of the line is defined as an unabsorbed photon flux, where the normalization unit is ${\\rm ph\\ cm^{-2}\\ s^{-1}}$. The contours are given in confidence levels of 68\\%, 90\\% and 99\\%. } \\label{fig:CVIconfCont} \\end{figure}"
    },
    "0609/astro-ph0609415_arXiv.txt": {
        "abstract": "We report on the properties of a sample of ultraviolet luminous galaxies (UVLGs) selected by matching the {\\it Galaxy Evolution Explorer (GALEX)} All-sky Imaging and Medium Imaging Surveys with the Sloan Digital Sky Survey Third Data Release. The overlap between these two surveys is roughly 450 square degrees. Out of 25362 galaxies (with SDSS spectroscopy) between $0.0<z<0.3$ detected by \\galex, there are 215 galaxies with $L>2\\times10^{10}$~L$_{\\odot}$ at 1530~\\AA\\ (observed wavelength). The properties of this population are well correlated with ultraviolet surface brightness. We find that the galaxies with low UV surface brightness are primarily large spiral systems with a mixture of old and young stellar populations, while the high surface brightness galaxies consist primarily of compact starburst systems, with an approximate boundary at a surface brightness of $I_{1530}=10^8$~L$_{\\odot}$~kpc$^{-2}$. The large galaxies appear to be the high-luminosity tail of the galaxy star formation function, and owe their large luminosity to their large surface area.  In terms of the behavior of surface brightness with luminosity, size with luminosity, the mass-metallicity relation, and other parameters, the compact UVLGs clearly depart from the trends established by the full sample of galaxies. The subset of compact UVLGs with the highest surface brightness ($I_{1530}>10^9$~L$_{\\odot}$~kpc$^{-2}$; ``supercompact UVLGs'') have characteristics that are remarkably similar to Lyman Break Galaxies at higher redshift. They are much more luminous (and thus have much higher star formation rates) than typical local ultraviolet-bright starburst galaxies and blue compact dwarf galaxies. They have metallicities that are systematically lower than normal galaxies of the same stellar mass, indicating that they are less chemically evolved. In all these respects, they are the best local analogs for Lyman Break Galaxies. ",
        "introduction": "%  Over the past decade enormous progress has been made toward mapping the cosmological history of star formation in the universe ({\\it e.g.,} Madau et al. 1996; Giavalisco et al. 2004). This has mainly been accomplished using large samples of high-redshift galaxies selected by their rest-frame ultraviolet (UV) colors ({\\it e.g.,} Steidel et al 1996, 2003; Dickinson et al. 2004). These surveys indicate that the global star formation rate of the universe has been in decline since $z\\sim1-2$, and was generally constant at higher redshift out to at least $z=6$ \\citep{g04}.  This picture still contains some uncertainty resulting from several factors. The star formation rate density at low redshift ($z=0-1$) has been determined through different techniques (${\\it e.g.,}$ \\ha\\ luminosity) than those used for higher redshift galaxies (${\\it e.g.,}$ rest-frame ultraviolet luminosity). These techniques are affected differently by extinction and radiative transfer effects, and they fundamentally probe star formation over different time scales.  One way around this problem is to obtain rest-frame ultraviolet (UV) measurements for a large sample of galaxies at low redshift, enabling the measurement of star formation rates using the same techniques that are used at higher redshift. This requires a UV telescope in space with a large field of view, something that has not been available until the {\\it Galaxy Evolution Explorer (GALEX)} mission \\citep{cm05}. \\galex\\ is obtaining UV fluxes for more than $\\sim10^7$ galaxies in the redshift range of $0<z<2$. Initial results on the UV luminosity density show strong evolution from $z=2$ to $z=0$, with the strongest evolution occurring in the most UV luminous galaxies \\citep{ds05,arn05}. The fraction of galaxies with $L_{1530}>0.2$L$_{*, z=3}$ fell by a factor of 30 from $z=1$ to $z=0$ (using L$_{*, z=3}=6\\times10^{10}$ L$_{\\odot}$; Steidel et al. 1999).  These UV luminous galaxies at high redshift are more commonly called Lyman Break Galaxies (LBGs; Steidel et al. 1999). These high-redshift galaxies are so named because they are identified by the effects of the Lyman break on their broadband colors \\citep{sh93}. LBGs are UV bright galaxies undergoing intense star formation with low to moderate stellar masses (log $M_*=9.5$ to $11.0$ M$_{\\odot}$), and are candidates for the precursors of present-day elliptical galaxies (see {\\it e.g,} Giavalisco 2002). LBGs are common at $z>2$, and they are clearly important as the sites of a significant fraction of all the star formation in the universe. Since strong evolution has made objects like LBGs extremely rare in the local universe, all of the information on this important galaxy population has come from very distant samples, which are inherently difficult to study. Thus, there has been little detailed information available on the processes driving the evolution of star formation in the population of LBGs.  Using local UV-bright starbursts as local analogs to LBGs has contributed significantly toward understanding these objects \\citep{h98,m99}. However, local starbursts differ from LBGs in important ways. Local starbursts are usually dwarf galaxies or small (sub-kpc) regions in the nuclei of larger galaxies, while LBGs have typical sizes of a few kpc \\citep{ferg04}. Luminous local starbursts are usually very dusty systems in which only a small fraction of the UV light escapes, while LBGs with similar bolometric luminosities (star formation rates) typically contain modest amounts of dust (e.g. Reddy et al. 2006; Erb et al. 2006b).  Given these differences, it is not clear that the conditions in local starbursts or the triggers of star formation are identical to those in LBGs, and so there is a need for better LBG analogs in the local universe.  Since LBGs are found in part by their large UV luminosity, LBG analogs in the local universe should also be UV luminous. The large area UV sky surveys being carried out by \\galex\\ provide an ideal data set for finding rare UV luminous galaxies in the local universe. Heckman et al. (2005; hereafter Paper I) described the properties of the most UV luminous galaxies (UVLGs) in the local universe based on cross-matching the initial \\galex\\ surveys with the Sloan Digital Sky Survey (SDSS) first data release (DR1; Abazajian et al. 2003). The UVLGs were composed of two basic types of galaxies: Large UVLGs, which are characterized by lower UV surface brightness and high mass, and compact UVLGs, which have higher UV surface brightness and lower mass. Many of the compact UVLGs have properties very similar to LBGs.  Although this sample was very illuminating, several questions remain. The extent of the similarity between the compact UVLGs and LBGs is a crucial question. More generally, it is not known whether these galaxies are truly a distinct population of objects in an earlier phase of evolution, {\\it i.e.,} remnants of the epoch of galaxy formation, or whether they are simply the high end of the UV luminosity function. Many of these questions could be better addressed if there were more such galaxies available for study, so we present an analysis of a larger sample of UVLGs, based on more recent \\galex\\ and SDSS data. ",
        "conclusions": "%  We have used \\galex\\ and the SDSS to identify and study the most UV luminous galaxies in the local universe. Our main results are as follows:  \\begin{enumerate}  \\item{ The most UV luminous galaxies in the local universe comprise a diverse group, with properties that are well correlated with UV surface brightness. Although there is not a sharp transition, we can use a surface brightness boundary of $I_{1530}<10^8$~L$_{\\odot}$~kpc$^{-2}$ to divide the UVLG sample into two groups: large and compact UVLGs.}  \\item{ The large UVLGs are massive, metal-rich disk galaxies with UV surface brightnesses only slightly larger than typical star forming galaxies. They are similar to normal galaxies in most respects, but are very luminous primarily because of their large size.}  \\item{ The compact UVLGs are low-mass, relatively metal-poor systems which often have a disturbed or interacting morphology. The high UV surface brightness and high specific star formation rate in these compact UVLGs indicates that intense star formation is ongoing.}  \\item{ It is possible to isolate a sample of local LBG-analogs with UV surface brightness criterion of $I_{1530}>10^9$~L$_{\\odot}$~kpc$^{-2}$. The galaxies in the resulting sample have many properties in common with LBGs, including luminosity, mass, star formation rate, specific star formation rate, extinction, and metallicity.}  \\item{ Compact UVLGs stand out from the trends established by the full DR3/GR1 sample in that they have much smaller sizes and higher surface brightness than would be predicted for galaxies of their mass or luminosity. They have metallicities that are generally lower by a factor of two to three compared to normal galaxies of the same mass. These properties suggest that they are a distinct population of objects, perhaps at a different phase of evolution than the bulk of the galaxies in the local universe.}  \\item{ The high UV luminosity and implied high star formation rate of compact UVLGs distinguishes them from any previously studied local galaxy population, including local UV-bright starbursts and blue compact dwarf galaxies.}  \\end{enumerate}"
    },
    "0609/astro-ph0609765_arXiv.txt": {
        "abstract": " ",
        "introduction": "The dominant source of phase errors in interferometric radio observations at frequencies above 5\\,GHz are fluctuations in the tropospheric water vapour content along the line of sight of the individual interferometer elements. The effect of the fluctuations scales linearly with frequency, making observations at frequencies of tens of GHz (at wavelengths less than about one cm) particularly prone to atmospheric phase changes. For a radio interferometer operating at frequencies of up to 115\\,GHz such as the Australia Telescope Compact Array (ATCA) it is therefore desirable to identify the periods within a year that are suitable for observations at these high frequencies. A well-tested means of characterizing the phase stability above a radio observatory are two-element interferometers observing unmodulated beacons of geostationary satellites (e.g. \\citealt{Radford1996, Hiriart2002}).  \\subsection{Atmospheric seeing in the radio regime}  A excellent introduction to radio seeing can be found in \\citealt{Thompson2001}. It is generally assumed that the water vapour distribution in the troposphere can be described as follows. Water vapour does not mix well with air because the air temperature is close to the condensation point of water. Its distribution can be approximated as nested parcels of air with a variety of water vapour densities. Kinetic energy is transferred from larger parcels to smaller and smaller parcels until it is dissipated by viscous friction, a situation known as Kolmogorov turbulence. A sketch of the situation can be found in Figure~II in \\cite{Masson1994}. The rms phase fluctuations ($\\sigma_\\phi$), measured over a long period, are expected to exhibit a power law relationship with baseline length, $b$:  \\begin{equation} \\sigma_\\phi^2 \\propto b^{\\beta}. \\label{eq:bk} \\end{equation}  Generally the turbulence is considered to be confined to a specific layer.  When the layer is much thinner than the interferometer spacing, two-dimensional turbulence theory is applicable, which suggests a power law exponent of $\\beta/2 = 1/3$ for the rms phase. When the layer is deep compared with the interferometer spacing, three-dimensional theory is appropriate, and a power law exponent of $\\beta/2 = 5/6$ is expected.  Furthermore, it is assumed that the mixture of air and water vapour is ``frozen'' in the atmosphere and is transported relative to the ground at a particular speed, $v_{\\rm s}$. This wind speed aloft couples the temporal and spatial behaviour of the phase screen, and $v_{\\rm s}$ therefore is called the phase screen speed. ",
        "conclusions": "We have presented a description of the ATCA's seeing monitor along with an analysis of 21 months of data. The main conclusions of this analysis are that:  \\begin{itemize}  \\item the seeing conditions on summer nights are not substantially different from winter days;  \\item the turbulence seems to be concentrated quite near the ground; and  \\item the turbulence appears to be in transition between being described by 2D and 3D models. This suggests the turbulent layer is comparable to the seeing monitor baseline length.  \\end{itemize}  The latter two points build a consistent picture of the turbulence being substantially confined to a layer in the lower $\\sim50 - 100$\\,m of the atmosphere where the interaction with the ground is likely to be strongest."
    },
    "0609/astro-ph0609286_arXiv.txt": {
        "abstract": "We present an updated investigation of the relation between large scale disk circular velocity, $v_c$, and bulge velocity dispersion, $\\sigma_c$. New bulge velocity dispersions are measured for a sample of 11 low surface brightness (LSB) and 7 high surface brightness (HSB) spiral galaxies for which $v_c$ is known from published optical or HI rotation curves.  We find that, while LSB galaxies appear to define the upper envelope of the region occupied by HSB galaxies (having relatively larger $v_c$ for any given $\\sigma_c$), the distinction between LSB and HSB galaxies in the \\vs~plane becomes less pronounced for $\\sigma_c \\la 80$ km s$^{-1}$. We conclude that either the scatter of the \\vs~ relation is a function of $v_c$ (and hence galaxy mass) or that the character of the \\vs~relation changes at $v_c \\sim 80$ \\kms. Some inplications of our findings are discussed. ",
        "introduction": "The existence of supermassive black holes (SBHs) in galactic nuclei has been accepted since the detections in the nearby spiral NGC 4258 \\citep{miyoshi} and in our own galaxy \\citep{schodel,ghez}. Compelling cases now exist in three dozen additional galaxies (see Ferrarese \\& Ford 2005 for a review), making it possible to investigate the existence of scaling relations linking the SBH mass, $M_{BH}$, to the overall properties of the host.  Indeed, $M_{BH}$ correlates with the blue luminosity $L_B$ of the host bulge \\citep{kormendy,marconi}, with the central light concentration \\citep{graham}, and with the bulge velocity dispersion $\\sigma_c$ \\citep{laura2,geb}.  Because of its tight scatter (0.34 dex in $M_{BH}$), the $M_{BH}-\\sigma_c$ relation is thought to allow us to peer into the mechanisms controlling the joint formation/evolution of SBHs and galaxies.  However, cosmological simulations show that a possibly more fundamental relation should be expected between $M_{BH}$ and the total gravitational mass $M_{tot}$ of the galaxy, or the mass $M_{DM}$ of the dark matter (DM) halo \\citep[e.g.][]{tiziana,kawakatu05,wyithe05,adams,monaco,haehnelt1,cattaneo,silk,haehnelt2,loeb}. Indirect observational evidence for such a link has been found in the form of a correlation between the circular velocity $v_c$ of spirals' disks (obtained by either HI or deep optical observations) and the $\\sigma_c$ of their bulge component (Ferrarese 2002; Baes et al. 2003; Pizzella et al. 2005):  \\begin{equation} \\log{v_c}=(0.84\\, \\pm\\, 0.09)\\log{\\sigma_c} + (0.55\\, \\pm\\, 0.29) \\end{equation}  Further exploring the $v_c-\\sigma_c$ relation is of interest because, by linking galactic components on vastly different scales, the relation is a reflection of the interplay between disks and spheroids during galaxy formation/evolution.  Furthermore, if $v_c$ and $\\sigma_c$ are used as surrogates for $M_{DM}$ and $M_{BH}$ respectively, equation (1) implies a strong causal relation between SBHs and DM haloes.  Pizzella et al. (2005) report evidence of a separation in the $v_c-\\sigma_c$ plane between High and Low Surface Brightness spiral galaxies (HSB and LSB respectively), the latter having distinctly larger $v_c$ at a given $\\sigma_c$ compared to the former. Based on this, the authors argue against the importance of baryonic collapse in shaping the density profiles of DM haloes in LSBs, and suggest that for a given DM halo mass, LSBs might host less massive SBHs than HSBs. The behaviour of HSBs themselves is somewhat controversial: while Pizzella et al. (2005) find that HSB in the range $50 \\la \\sigma_c \\la 350$ km\\, s$^{-1}$ (corresponding to $150 \\la v_c \\la 550$ km\\, s$^{-1}$) follow a linear relation, Ferrarese (2002) and Baes et al. (2003) claim that the relation is mildly non-linear for $\\sigma_c \\ga 80$\\, km\\, s$^{-1}$ ($v_c \\ga 140$\\, km\\, s$^{-1}$),  and breaks down altogether at lower $\\sigma_c$, possibly revealing the inability of the least massive haloes to form central SBHs \\citep{shankar06,haehnelt1,silk,loeb}  The goal of this paper is to investigate the $v_c-\\sigma_c$ relation for a combination of LSB and HSB galaxies, with $v_c$ in the critical range $110-220$ km\\, s$^{-1}$. The data reduction and analysis are discussed in \\S2 and \\S3 respectively. Discussion and conclusions can be found in \\S4.  \\begin{figure} \\includegraphics[angle=0,width=8.5cm]{chi2.eps} \\caption{Spectra for four of the program galaxies are shown in black. The top two panels show spectra obtained at the VLT, while the bottom two panels show Calar Alto spectra. To illustrate the quality of the data, for each telescope we chose to show the two spectra with the lowest (first and third panel from the top) and highest (second and fourth panel from the top) signal-to-noise ratio. The red lines show the best $\\chi^2$ fits to each spectrum, overplotted to the real data in the spectral region used to perform the fits. The narrow emission features present in three of the spectra are artifacts resulting from imperfect sky subtraction} \\end{figure} \\label{ppxf} ",
        "conclusions": "The $v_c - \\sigma_c$ relation for HSB galaxies is plotted in the left panel of Fig.~\\ref{vs}. The 11 LSB galaxies studied in this paper populate the region between $40 \\leq \\sigma_c \\leq 80$ \\kms, where previously only three galaxies could be found. Furthermore, in the same region, we more than double the number of HSB galaxies, from three to eight. Pizzella et al. (2005) conclude that, relative to the \\vs~relation defined by HSB galaxies, LSB galaxies appear to have larger circular velocities, $v_c$, at any given value of $\\sigma_c$. This conclusion is reaffermed by our data.  However, with the addition of the present sample, it has now become more evident that there is some overlap in the locations occupied by LSB and HSB galaxies in the \\vs~plane. In other words, LSB galaxies appear to occupy the {\\it upper envelope} of the strip defined by HSB galaxies. Furthermore, at low velocity dispersions $30 \\leq \\sigma_c \\leq 80$ \\kms, the \\vs~relation for both LSB and HSB galaxies appears to be shifted towards larger circular velocities relative to the relation defined by the (HSB) galaxies with larger velocity dispersion.  A least-square fit, taking into account the errors on both quantities \\citep[FITEXY;][]{press92}, to the sample of HSBs with $\\sigma_c \\geq 80$ \\kms~gives:  \\begin{equation} \\log{v_c}=(0.68 \\pm 0.04)\\log{\\left(\\frac{\\sigma_c}{129 {\\rm~km~s^{-1}}}\\right)+(2.33 \\pm 0.01)}, \\end{equation}  \\noindent with $\\chi^2_r=2.58$. For $30 \\leq \\sigma_c \\leq 80$ \\kms, on the other hand,  \\begin{equation} \\log{v_c}=(1.33 \\pm 0.86)\\log{\\left(\\frac{\\sigma_c}{63 {\\rm~km~s^{-1}}}\\right)+(2.21 \\pm 0.03)}, \\end{equation}  \\noindent with $\\chi^2_r=0.23$ for HSB galaxies, and  \\begin{equation} \\log{v_c}=(1.91 \\pm 0.68)\\log{\\left(\\frac{\\sigma_c}{58 {\\rm~km~s^{-1}}}\\right)+(2.20 \\pm 0.02)}, \\end{equation}  \\noindent with $\\chi^2_r=1.55$ for LSB galaxies.  Although the uncertainties are large, equations (3) and (4) are consistent with each other well within 1$\\sigma$, but inconsistent with equation (2) (see also right panel of Figure~\\ref{vs}).  Two interpretations are open at this point. The samples studied to date are far from complete or homogeneous, and it is conceivable that, as more data are collected, more galaxies will fill the gap between the LSB galaxies, which are mainly found in the low-velocity dispersion region, and the extrapolation to low $\\sigma_c$ of equation (2), defined mainly by HSB galaxies. This interpretation allows for the possibility that the log-log \\vs~relation, including both HSB and LSB galaxies, is linear, but that its {\\it scatter} increases going from large to small velocity dispersions. Alternatively, it is possible that the log-log \\vs~relation is not linear, and that indeed its character changes dramatically below $v_c \\sim 200$ km s$^{-1}$, as originally suggested by Ferrarese (2002) and Baes et al. (2003). In either case, it seems unavoidable to conclude that the efficiency of bulge (and, perhaps, SBH) formation, is regulated by intrinsic parameters in addition to the depth of the global potential well.  The implications of these findings for the central SBHs cannot be easily quantified. The $M_{BH} - \\sigma_c$ relation is not characterised for LSBs, for which no SBH detection has been attempted. It is also not defined below $\\sigma_c \\leq 80$ km s$^{-1}$, which corresponds to SBHs too small to be detected given current space or ground-based instrumentation. We reiterate the conclusion of Ferrarese et al. (2002) and Baes et al. (2003) that the behaviour of the $v_c - \\sigma_c$ relation at $\\sigma_c \\la 80$ km s$^{-1}$ might reflect the inability of these galaxies to form a central SBH, as argued on theoretical grounds by, for instance, \\cite{shankar06,haehnelt1,silk,loeb}.  Numerical simulations \\citep{tiziana,robertson} predict that the $M_{BH}-\\sigma_c$ relation originates from the feedback between the central SBH and the progenitor of the hot stellar component at the early stages of the formation of both; therefore, in low-mass, late-type HSB and LSB galaxies internal factors might lead to the suppression of such feedback, creating a dynamical relation between disk, bulge and SBH which is starkly different from that displayed by more massive systems.  \\begin{table*} {\\scriptsize \\begin{tabular}{lcclrrrlrcrc} \\hline Name & R.A. & Dec. &  $v_c$ & $\\sigma_c$ & $r_e$ & $v_{rad}$ &  Morph & $m_I$ & Incl. & Exp.Time & S/N \\\\ & (J2000) & (J2000) & (km s$^{-1}$) & (km s$^{-1}$) & (arcsec) & (km s$^{-1}$) & & (mag) & (degrees) & (seconds) & (@ 8500 \\AA)\\\\ \\hline \\multicolumn{12}{c}{}\\\\ \\multicolumn{12}{c}{LSB galaxies}\\\\ \\multicolumn{12}{c}{}\\\\ \\hline ESO215-G39 & 11h17m04s & $-$49$^{\\circ}$12$'$05$''$ & 162$\\pm$3 & 71$\\pm$11\\phantom{0} & 4.5 & 4335 & SABc & 12.01 & 50 & 500  & 20.03\\\\ ESO268-G44 & 12h48m42s & $-$45$^{\\circ}$00$'$29$''$ & 190$\\pm$9 & 70$\\pm$9\\phantom{0} &  3.5 & 3477 & Sb & 12.22 & 62 & 450  & 25.68\\\\ ESO322-G19 & 12h29m07s & $-$40$^{\\circ}$40$'$24$''$ & 141$\\pm$9 & 49$\\pm$10\\phantom{0} & 0 & 3100 & SBc & 12.64 & 79 & 900  & 22.19\\\\ ESO323-G42 & 12h55m01s & $-$40$^{\\circ}$58$'$20$''$ & 158$\\pm$4 & 54$\\pm$9\\phantom{0} & 4 & 4203 & Sc & 11.53 & 69 & 450  & 30.32\\\\ ESO323-G73 & 13h04m02s & $-$38$^{\\circ}$11$'$57$''$ & 163$\\pm$7 & 70$\\pm$7\\phantom{0} & 2 & 4929 & Sbc & 12.44 & 48 & 450  & 35.29\\\\ ESO374-G03 & 09h51m57s & $-$33$^{\\circ}$04$'$25$''$ & 146$\\pm$3 & 40$\\pm$6\\phantom{0} & 6 & 2931 & SABc & 11.51 & 71 & 600  & 21.80\\\\ ESO382-G06 & 13h05m32s & $-$32$^{\\circ}$57$'$38$''$ & 171$\\pm$4 & 54$\\pm$8\\phantom{0} & 2 & 4809 & Sab & 13.20 & 54 & 600  & 21.01\\\\ ESO444-G21 & 13h23m31s & $-$30$^{\\circ}$06$'$52$''$ & 116$\\pm$21 & 53$\\pm$9\\phantom{0} & 4.5 & 4265 & Sc & 12.86 & 84 & 1900  & 20.46\\\\ ESO444-G47 & 13h28m25s & $-$31$^{\\circ}$51$'$28$''$ & 154$\\pm$6 & 61$\\pm$14\\phantom{0} & 0 & 4389 & SBc & 12.82 & 71 & 900  & 22.01\\\\ ESO509-G91 & 13h40m03s & $-$25$^{\\circ}$28$'$28$''$ & 203$\\pm$9 & 62$\\pm$11\\phantom{0} & 6.5 & 5113 & SBc & 12.70 & 79 & 900  & 20.91\\\\ ESO376-G10 &10h42m00s & $-$36$^{\\circ}$56$'$07$''$ & 220$\\pm$11 & 54$\\pm$12\\phantom{0} & 7 & 3184 & SBd & 11.00 & 76 & 600  & 20.27\\\\ \\hline \\multicolumn{12}{c}{}\\\\ \\multicolumn{12}{c}{HSB galaxies}\\\\ \\multicolumn{12}{c}{}\\\\ \\hline NGC 3953 & 11h53m49s& +52$^{\\circ}$19$'$36$''$ & 223$\\pm$5& 115$\\pm$9 & 66 & 1052 & Sbc & 9.03 & 63 & 2500  & 60.01 \\\\ NGC 3877 & 11h46m08s& +47$^{\\circ}$29$'$40$''$ & 167$\\pm$11& 84$\\pm$8 & 53.5 & 895 & Sc & 9.91 & 83 & 3600  & 66.08 \\\\ NGC 4088 & 12h05m34s& +50$^{\\circ}$32$'$21$''$ & 173$\\pm$14& 62$\\pm$7 & 76 & 757 & SABc & 9.51 & 71 & 3600  & 58.13 \\\\ NGC 3949 & 11h53m41s& +47$^{\\circ}$51$'$32$''$ & 164$\\pm$7& 63$\\pm$5 & 36 & 800 & Sbc & 10.23 & 56 & 4000  & 57.15 \\\\ NGC 4157 & 12h11m04s& +50$^{\\circ}$29$'$05$''$ & 185$\\pm$10& 71$\\pm$6 & 35 & 774 & SABb & 9.93 & 90 & 4000  & 63.30 \\\\ NGC 3769 & 11h37m44s& +47$^{\\circ}$53$'$35$''$ & 122$\\pm$8& 80$\\pm$7 & 29 & 737 & Sb & 11.08 & 78 & 8000  & 66.25\\\\ NGC 3726 & 11h33m21s& +47$^{\\circ}$01$'$45$''$ & 162$\\pm$9& 63$\\pm$10 & 84 & 866 & Sc & 9.50 & 49 & 6600 & 38.44\\\\ \\hline \\end{tabular}} \\caption{Column 1 gives the galaxy's name, while columns 2 \\& 3 give the J2000 coordinates. The asymptotic circular velocity $v_c$ is tabulated in column 4, it is taken directly from Verheijen (2001) for the HSB galaxies, and measured from the published rotation curves for the LSB galaxies.  Central velocity dispersion ($\\sigma_c$) at $r_e/8$ and the effective radius of the galaxy can be found in columns 5 and 6.  The heliocentric radial velocity, Hubble classification, total $I$ band magnitude, inclination can be found in columns 7 to 10 and are taken from Verheijen (2001) for the HSB galaxies and Palunas \\& Williams (2000) for the LSB galaxies and the online NED database. Finally, the total exposure time and S/N/pixel at 8500\\AA\\ are given in columns 11 and 12.\\label{list}} \\end{table*}"
    },
    "0609/astro-ph0609079_arXiv.txt": {
        "abstract": "\\noindent We examine the correlation between globular cluster (GC) color and magnitude using HST/ACS imaging for a sample of 79 early-type galaxies ($-21.7<M_B<-15.2$ mag) with accurate surface-brightness fluctuation distances from the ACS Virgo Cluster Survey. Using the \\kmm mixture modeling algorithm, we find a highly significant correlation, $\\gamma_z \\equiv \\frac{d(g-z)}{dz} = -0.037 \\pm 0.004$, between color and magnitude for the subpopulation of blue GCs in the co-added GC color-magnitude diagram of the three brightest Virgo cluster galaxies (M49, M87 and M60). The sense of the correlation is such that brighter GCs are redder than their fainter counterparts. For the single GC systems of M87 and M60, we find similar correlations; M49 does not appear to show a significant trend. There is no correlation between $(g-z)$ and $M_z$ for GCs belonging to the red subpopulation.  The correlation $\\gamma_g\\equiv \\frac{d(g-z)}{dg}$ for the blue subpopulation is much weaker than $\\gamma_z$. Using Monte Carlo simulations, we attribute this finding to the fact that the blue subpopulation in $M_g$ extends to higher luminosities than does the red subpopulation, which biases the \\kmm fit results. The highly significant correlation between color and $M_z$, however, is a real effect: this conclusion is supported by biweight fits to the same color distributions. We identify two environmental dependencies which influence the derived color-magnitude relation: (1) the slope of the color-magnitude relation decreases in significance with decreasing galaxy luminosity, although it remains detectable over the full luminosity range of our sample; and (2) the slope is stronger for GC populations located at smaller galactocentric distances. These characteristics suggest that the observed trend is, at least partially, shaped by external agents. We examine several physical mechanisms that might give rise to the observed color-magnitude relation including: (1) presence of contaminants like super-clusters, stripped galactic nuclei, or ultra-compact dwarfs; (2) accretion of GCs from low-mass galaxies;  (3) stochastic effects; (4) the capture of field stars by individual GCs; and (5) GC self-enrichment. Although none of these scenarios offers a fully satisfactory explanation of the observations, we conclude that self-enrichment and field-star capture, or a combination of these processes, offer the most promising means of explaining our observations. ",
        "introduction": "It is by now an accepted paradigm that the formation of globular clusters (GCs) is closely linked to the formation of their host galaxies (e.g. Searle \\& Zinn~1978, Ashman \\& Zepf~1992, Kissler-Patig~1997, Hilker {\\it et al.}~1999b, Forbes \\etal~1997, C\\^{o}t\\'{e} {\\it et al.}~1998, Beasley \\etal~2002). In recent years, observational data for extragalactic GC systems have made a huge leap forward in quality, thanks mainly to high-resolution imaging from the Hubble Space telescope, and to wide field imaging and multi-object spectroscopy from 8m-class ground based telescopes (see West \\etal~2004 and references therein).  As a rule, studies of extragalactic GCs have tended to focus on correlations between the global photometric properties of the GC systems (GCSs) and those of their host galaxy. For instance, Peng \\etal~\\cite{Peng05}, examined the color distributions for GCs belonging to 100 galaxies observed in the ACS Virgo Cluster Survey (ACSVCS; C\\^ot\\'e \\etal~2004; see also below), finding nearly all of these galaxies ($-22<M_B<-15$ mag) to possess bimodal, or at least asymmetric, GC color distributions. Consistent with previous findings (e.g., Gebhardt \\& Kissler-Patig~1999; Kundu \\etal~2001; Larsen \\etal~2001), the mean GC color was found to correlate with galaxy luminosity in the same sense as the well-known color-magnitude relation for early type galaxies (e.g. Bower, Lucey \\& Ellis~1992; Karick \\etal~2003; Hilker \\etal~2003; Ferrarese \\etal~2006). A corresponding GC color - galaxy color relation was also found by Peng \\etal~\\cite{Peng05}, confirming previous findings (e.g., Burgarella \\etal~2001; Larsen \\etal~2001; Lotz \\etal~2004).  At the same time, analyses of the color magnitude diagrams for the GCSs of early-type galaxies revealed some tentative evidence for a dependence of {\\it individual} GC color on luminosity.  Ostrov \\etal~\\cite{Ostrov98} found from Washington photometry of NGC 1399 that the color peaks of the GC color distribution seem to merge at the highest luminosities. Dirsch et al.~\\cite{Dirsch03} confirmed this finding for the same galaxy based on an expanded sample of GCs. There have also been hints of such a behavior in other galaxies such as NGC 5128 (Harris \\etal~2004) and M87 (Whitmore \\etal~1995; Larsen \\etal~2001).  With the improved imaging capabilities of HST provided by the Advanced Camera for Surveys (ACS; Ford \\etal~1998), it has become possible to investigate this issue in more detail.  Harris \\etal~\\cite{Harris06} used ACS imaging to investigate the distribution of GCs in the $(I, B-I)$ color magnitude plane for eight brightest cluster galaxies in the range $1800 \\lesssim cz \\lesssim 3200$ km s$^{-1}$. For the four galaxies with the largest number of GCs --- and also the joint sample of eight GC systems --- they found a significant trend towards redder colors for increasing luminosity for the subpopulation of blue GCs. This finding led them to suggest a mass-metallicity relation of the form $Z \\propto M^{0.55}$ for this subpopulation.  Curiously, no color-magnitude relation was found for the red GCs.  Strader \\etal~\\cite{Strade05} used the publicly available data from the ACSVCS to look for color-magnitude trends in the three brightest Virgo Cluster galaxies M49, M87 and M60. For the blue GC populations of M87 and M60, they found a significant trend of redder color with brighter GC $z$-band luminosity similar to that found by Harris \\etal~\\cite{Harris06}. They too interpreted this trend as evidence for a mass-metallicity relation among the blue (metal-poor) GCs and suggested self-enrichment as a possible explanation. No significant color-magnitude trend was found for M49, the brightest member of the Virgo Cluster.  Both the Harris \\etal~\\cite{Harris06} and Strader \\etal~\\cite{Strade05} studies used heteroscedastic mixture models to fit the colors of an assumed double-Gaussian color distribution to the observed GC color-magnitude diagrams (CMDs) as a function of magnitude. Harris {\\it et~al.} use {\\kmm} (e.g. Ashman, Bird \\& Zepf~1994), while Strader {\\it et~al.} use {\\tt Nmix} (Richardson \\& Green~1997).  Two important issues that were not addressed by these studies are the subject of this paper: first, we will investigate the color-magnitude correlations not only in the red (i.e., $z$), but also in the blue band (i.e., $g$); second, we will investigate how the blue peak slope varies as a function of host galaxy luminosity and galactocentric distance. Note that the Harris \\etal~\\cite{Harris06} study targeted only brightest cluster galaxies, so it spanned a relatively limited range in galaxy luminosity.  We finally point out that there is recent evidence for the existence of a GC color-magnitude relation also in spiral galaxies. Based on ACS imaging, Spitler {\\it et~al.}~\\cite{Spitle06} found a significant color-magnitude trend for the Sombrero galaxy (NGC 4594) implying $Z \\propto M^{0.30}$, which is a somewhat weaker trend than suggested by Harris \\etal~\\cite{Harris06} and Strader \\etal~\\cite{Strade05} for giant ellipticals.  The aim of this paper is to investigate the color-magnitude trend of GCs over the entire host galaxy luminosity range covered by the ACSVCS. Special emphasis will be put on environmental dependencies of this trend and checks of the fitting technique. The paper is structured as follows: In \\S\\ref{gcsel} we outline how GC candidates are selected in the ACSVCS. In \\S\\ref{realvssim} we compare \\kmm fits to simulated CMDs with fits to observed CMDs. In \\S\\ref{environment} we investigate the environmental dependence of the color-magnitude trend. In \\S\\ref{discussion} we examine a variety of physical effects that may create these trends. We finish this paper with the summary and conclusions in \\S\\ref{conclusions}. ",
        "conclusions": "In this paper we have analysed the color-magnitude relation for the GC subpopulations in early-type galaxies belonging to the Virgo Cluster. Our sample consists of 79 galaxies, spanning the range $-21.7<M_B<-15.2$~mag, drawn from the ACS Virgo Cluster Survey (C\\^ot\\'e \\etal 2004).  Our principal findings can be summarized as follows:  \\begin{enumerate} \\item Based on \\kmm fits, we find a highly significant correlation $\\gamma_z \\equiv \\frac{d(g-z)}{dz} = -0.037 \\pm 0.004$ for the blue GCs in the co-added CMD of the three brightest Virgo Cluster galaxies (M49, M87, M60). The sense of the correlation is such that brightest GCs in this subpopulation are redder than their faint counterparts. For M87 and M60 alone, we find comparable correlations; by contrast, the blue GCs in M49 do not appear to follow such a trend.  In no galaxy do we find a significant correlation between $(g-z)$ and $M_z$ for the GCs associated with the red subpopulation.  \\item Based on \\kmm fits, we find the slope $\\gamma_g \\equiv \\frac{d(g-z)}{dg}$ within the blue subpopulation to be much weaker than $\\gamma_z$. In turn, we find a mild {\\it positive} correlation $\\gamma_g$ for the red subpopulation, in the opposite sense to what is found for the blue subpopulation.  \\item We test \\kmm on artificial CMDs with, and without, implemented slopes for the blue subpopulation, and no slope for the red population. \\kmm correctly reproduces the implemented slopes when using the CMD $(g-z)$ vs. $M_z$. For $(g-z)$ vs. $M_g$, \\kmm gives a biased estimate for the color of the blue population and yields a positive slope for the red subpopulation, although no slope had been implemented. We trace this bias to the fact that in $M_g$ -- unlike in $M_z$ -- the luminosity function of blue GCs is populated to significantly brighter luminosities than for the red GCs. We conclude that for the real data, the strong slope $\\gamma_z$ is {\\it not} an artifact of the fitting method, while the weaker slope $\\gamma_g$ is an artificial feature of the \\kmm fit.  \\item We derive correlations independent of \\kmm by using a biweight estimator for GCs divided on the basis of a luminosity independent color.  The results confirm the significant slope found by \\kmm in $M_z$ and reveal no significant discrepancies between the slope measured using $M_z$ and $M_g$.  \\item The slope measured from the blue subpopulation decreases with host galaxy luminosity. However, even among the faintest galaxies ($-18.4<M_B<-15.2$~mag) the slope is nonzero.  \\item The color-magnitude relation exhibited by the blue GCs is stronger for clusters with smaller projected galactocentric distances. This difference is driven mainly by GCs with $M_z<-10$ mag.  \\end{enumerate}  We examine a number of physical mechanisms that might give rise to a color-magnitude relation with these characteristics: (1) contamination by super-star-clusters, stripped galactic nuclei, or ultra-compact dwarf galaxies; (2) accretion of GCs from low-mass galaxies; (3) stochastic effects; (4) capture of galactic field stars by GCs; and (5) self-enrichment of individual GCs.  Neither contaminants nor accretion of low-mass galaxies appear likely to fully explain our observations. Self-enrichment in more massive GCs may be able to account for the observed color-magnitude relation, although more work is needed to determine if the dependence on galactocentric radius can be explained within the context of the self-enrichment scenario. Although the capture of field stars {\\it does} lead to a dependence on galactocentric radius that is at least qualitatively in agreement with our observations, this mechanism does not appear efficient enough to explain the overall size of the observed trend.  While none of these scenarios can fully account for the full set of observations in a straightforward manner, self-enrichment and capture of field stars seem the most promising approaches. There is, of course, no {\\it a priori} reason that a single mechanism is responsible for the observed color-magnitude relation, and conceptually, the combined effects of self-enrichment and field star capture may be able to explain the observations: i.e., field star capture could act as a second order effect, producing environmental variations of a global mass-metallicity relation induced by self-enrichment.  This would alleviate the discrepancy between real and required field star capture rate estimated in this work. Dedicated N-body simulations going beyond these simple estimates will be required to accurately address the efficiency of field star capture in different environments.  \\label{conclusions}"
    },
    "0609/astro-ph0609779.txt": {
        "abstract": "The existence of extrasolar planets with short orbital periods suggests that planetary migration induced by tidal interaction with the protoplanetary disk is important. Cores and terrestrial planets may undergo migration as they form. In this paper we investigate the evolution of a population of cores with initial masses in the range 0.1--1~earth mass embedded in a disk.  Mutual interactions lead to orbit crossing and mergers, so that the cores grow during their evolution.  Interaction with the disk leads to orbital migration, which results in the cores capturing each other in mean motion resonances.  As the cores migrate inside the disk inner edge, scatterings and mergers of planets on unstable orbits together with orbital circularization causes strict commensurability to be lost. Near commensurability however is usually maintained.  All the simulations end with a population of typically between two and five planets, with masses depending on the initial mass.  These results indicate that if hot super--Earths or Neptunes form by mergers of inwardly migrating cores, then such planets are most likely not isolated.  We would expect to always find at least one, more likely a few, companions on close and often near--commensurable orbits.  To test this hypothesis, it would be of interest to look for planets of a few to about 10~earth masses in systems where hot super--Earths or Neptunes have already been found. ",
        "introduction": "\\label{intro}  The recent announcements of the detection of extrasolar planets of a few earth masses (M$_\\oplus$) (OGLE--05--390L~b, 5.4~M$_\\oplus$, Beaulieu et al. 2006; Gliese~876~d, 7.3~M$_\\oplus$, Rivera et al. 2005) gives support to the initial solid core accumulation model for planet formation.  Although it has been proposed that such ``super--Earths'' could have formed through gravitational instability in a gaseous protoplanetary disk leading to a giant planet which subsequently lost its gaseous envelope through the action of external UV radiation (Boss 2006), a formation mechanism through the accumulation of planetesimals seems more natural.  In addition to these super--Earths, nine planets with a mass comparable to that of Uranus or Neptune (in the range 10.5 to 18.5~M$_\\oplus$) have been reported.  OGLE--05--390L~b is at a distance of 2.1~astronomical unit (au) from the central star and was detected through microlensing.  Gliese~876~d, detected through radial velocity measurements, is at 0.02~au from its parent star.  Among the nine  planets with masses similar to that of Neptune, four are within 0.1~au from the central star.  Gliese~876 is a M--type star.  The temperature at 0.02~au from the star is therefore low enough for heavy elements to condense.  Thus it is possible to consider that Gliese~876~d formed {\\em in situ} by accumulation of heavy material that spiraled in with the gas through the circumstellar disk. However, the existence of close orbiting giant planets, the so--called 'hot Jupiters' and 'hot Neptunes', has been taken as an indication of the operation of large scale migration induced by the interaction of recently formed protoplanets with protoplanetary disks (e.g., Papaloizou~\\& Terquem~2006 and references therein). Such migration processes may have operated during the formation of lower mass planets as well.  In particular, planets in the earth mass range are expected to undergo type~I migration (e.g., Ward~1997).  Accordingly, here we envisage a scenario in which cores assemble further away from the central star, and migrate inwards due to tidal interaction with the disk.  The disk is supposed to be truncated at some inner edge so that the planets do not fall onto the star.  Mutual perturbations of the cores leads to orbit crossing, collisions and possibly mergers during the migration phase, resulting in the formation of a smaller number of more massive planets on short period orbits.  Such a scenario has been considered by Brunini \\& Cionco (2005), who calculated by means of $N$--body simulations the evolution of 100~protoplanets of 0.5~M$_\\oplus$ together with 200~planetesimals of 0.1~M$_\\oplus$ subject to their mutual interaction and the tidal interaction with the disk.  Focusing on the mass and semi--major axis of the final largest solid core, they concluded that Neptune--like planets on short orbits should be common.  Here we focus on the final configuration of the multi--planet systems we obtain.  Evolution of an ensemble of cores in a disk almost always leads to a system of a few planets, with masses that depend on the total initial mass, on short orbits with mean motions that frequently exhibit near commensurabilities and, for long enough tidal circularization times, apsidal lines that are locked together. Starting with a population of 10 to 25 planets of 0.1 or 1~M$_\\oplus$, we end up with typically between two and five planets with masses of a few tenths of an earth mass or a few earth masses, depending on the total initial mass, inside the disk inner edge.  Interaction with the central star leads to tidal circularization of the orbits which, together with possible close scatterings and final mergers, tends to disrupt mean--motion resonances that are established during the migration phase.  The system, however, often remains in a configuration in which the orbital periods are close to commensurability.  Apsidal locking of the orbits, if established during migration, is often maintained through the action of these processes.  The plan of the paper is as follows.  In section~\\ref{sec:secular}, we review studies of the evolution of migrating resonant planets.  We also show that tidal circularization by the central star results in the disruption of strict commensurability, although the mean motions may remain nearly commensurable and apsidal lines remain locked.  In section~\\ref{sec:model}, we describe the model we use to follow the evolution of a population of planets, and describe our initial conditions.  To calculate the mutual interactions between the cores, we use a $N$--body code.  A dissipative force is included to model the tidal interaction with the disk, which leads to orbital decay and eccentricity and inclination damping.  Relativistic effects and tidal interaction with the central star are also included, as they affect the eccentricity of the orbits close to the star.  We also incorporate the possibility of corotation torques acting in the edge region. The potential importance of these in reversing type~I migration has been indicated by Masset et al.~(2006). We discuss the effects of such torques on eccentric orbits.  In section~\\ref{sec:results}, we describe our results.  As mentioned above, all the runs end with a few planets on close orbits inside the inner cavity that frequently exhibit near commensurability and apsidal line locking.  We study the effect of varying the circularization timescale and of a hypothetical reversal of the torque near the disk inner edge due to corotation or other effects.  Finally, in section~\\ref{sec:disc}, we summarize and discuss our results. ",
        "conclusions": "\\label{sec:disc}  We have calculated the evolution of a population of cores/planets with masses in the range 0.1--1~M$_\\oplus$ embedded in a disk.  They evolve due to gravitational interaction with the central star, mutual gravitational interactions, tidal interaction with the disk and the star.  Mutual interactions lead to orbit crossing and mergers, so that the cores grow during their evolution.  Interaction with the disk leads to orbital migration.  As cores with different masses and at different locations migrate with different rates, they capture each other in mean motion resonances. Such captures enable planets to migrate inside the cavity interior to the disk inner edge.  As they approach closer to the central star, for small enough cavities their orbits are circularized through tidal dissipation and strict commensurabilities are lost. That process may be also aided by scatterings and mergers of planets on unstable orbits that occur interior to the disk inner edge.  Near--commensurability however may be maintained.  Note too that if apsidal locking is established during migration, it can be preserved through the operation of these processes.  All the simulations end with a population of typically between two and five planets, with masses depending on the initial mass.  Note that the disk tidal torque may be reversed near the disk inner edge.  When this is the case,  although it is possible that some planets can still penetrate inside the cavity due to scatterings and/or weakening of the torques because of a finite eccentricity,  some planets may be left in the disk, just beyond the inner edge, until that disperses.  The qualitative results do not depend on the detail of the initial conditions.  As long as a population of cores is able to migrate inwards at different rates, the system evolves toward a family of a few planets which are almost always on near--commensurate orbits.  The orbital migration and eccentricity damping timescales we have adopted in this paper have been derived for type~I migration in inviscid disks.  Note that type~I migration has been shown to follow a random walk in a turbulent disk (Nelson \\& Papaloizou 2004).  The studies done in the present paper assume that the cores can migrate down to the disk inner edge, in regions where the gas is ionized and magnetic turbulence can develop (Fromang et al. 2002).  These studies would of course not apply if there were no systematic inwards migration of the cores/protoplanets of the type we consider here.  As of today, there is no indication that type~I migration in a turbulent disk has a systematic trend (Nelson 2005), but this cannot be ruled out either.  It is likely that type~I migration does depend at least to some extent on the torque exerted by the material that corotates with the planet, which so far has not been taken into account in the simulations, which lack the required resolution.  If there is a systematic trend, as long as the time averaged migration rate is inwards, by the averaging principle and as confirmed in test studies, short term fluctuations do not qualitatively change the results we have presented in this paper.  It has very recently been suggested (Paardekooper \\& Mellema 2006) that, because of effects arising from radiation trapping, type~I migration in an optically thick laminar disk could be outward.  This result, should it be confirmed for both laminar and turbulent disks, suggests that cores would not migrate inwards in the disk inner parts as long as the dust opacity is high enough there.  However, after the dust settles and agglomerates to form cores, the opacity decreases and migration could then resume.  It is indeed an observational fact that the disk inner parts (whithin a few au) become optically thin before the rest of the disk is depleted.  The calculations done in this paper show that if hot super--Earths or Neptunes form by mergers of inwardly migrating cores, then such planets are most likely not isolated.  We would expect to always find at least one, more likely a few, companions on close and often near--resonant orbits.  To test this hypothesis, it would be of interest to look for planets of a few to $\\sim 10$~earth masses in systems where hot super--Earths or Neptunes have already been found and there is no destabilizing influence of a giant planet close by.  It has been speculated that the cores of giant planets could form in a way similar to that investigated here, by accumulation of cores in the disk inner parts (e.g., Papaloizou \\& Terquem 1999).  The calculations presented in this paper suggest that to assemble a massive core in the inner disk, significantly more mass in smaller cores may be needed to begin with.  Indeed, most of the simulations end with at least three planets, and not with a single planet containing all the initial mass."
    },
    "0609/astro-ph0609590.txt": {
        "abstract": "We investigate the structure of the core surrounding the recently identified deeply embedded young stellar object Barnard 1c which has an unusual polarization pattern as traced in submillimeter dust emission.  Barnard 1c lies within the Perseus molecular cloud at a distance of 250 pc.  It is a deeply embedded core of 2.4 $M_\\odot$ (Kirk et al.) and a luminosity of $4\\pm 2 L_\\odot$.  Observations (and resolutions) of \\co\\ (9.2\\arcsec\\ $\\times$ 5.9\\arcsec), \\thirteenco\\ and \\ceighteeno\\ (14.3\\arcsec\\ $\\times$ 6.7\\arcsec), \\hcop\\ (7.6\\arcsec\\ $\\times$ 5.8\\arcsec), and \\ntwohp\\ (5.9\\arcsec\\ $\\times$ 4.6\\arcsec) were obtained with the Berkeley-Illinois-Maryland Association array, together with the continuum at 3.3 mm (6.4\\arcsec\\ $\\times$ 4.9\\arcsec) and 2.7 mm (9.5\\arcsec\\ $\\times$ 6.3\\arcsec). The field of view of the BIMA array antennas at 3 mm is 2.1\\arcmin. Single-dish measurements of \\ntwohp\\ and \\hcop\\ with FCRAO reveal the larger scale emission in these lines with resolutions of 57.5\\arcsec\\ and 60.5\\arcsec, respectively.  The \\CO\\ and \\HCOP\\ emission traces the outflow extending over the full field of view, which coincides in detail with the S-shaped jet recently found in Spitzer IRAC imaging. The \\NTWOHP\\ emission, which anticorrelates spatially with the \\ceighteeno\\ emission, originates from a rotating envelope with effective radius $\\sim 2400$ AU and mass $2.1 - 2.9 \\ M_\\odot$, as derived from the 3.3 mm continuum emission.  \\NTWOHP\\ emission is absent from a 600~AU diameter region around the young star, offset from the continuum peak.  The remaining \\NTWOHP\\ emission may lie in a coherent torus of dense material. With its outflow and rotating envelope, B1c closely resembles the previously studied object L483-mm, and we conclude that it is a protostar in an early stage of evolution, i.e., Class 0 or in transition between Class 0 and Class I.  We hypothesize that heating by the outflow and star has desorbed CO from grains which has destroyed \\NTWOHP\\ in the inner region and surmise that the presence of grains without ice mantles in this warm inner region can explain the unusual polarization signature from B1c. ",
        "introduction": "Class 0 sources represent the youngest phase of low-mass star formation.  They are characterized by higher infall rates than more evolved sources based on outflow activity \\citep{bon96,ww01}, an absence of optical emission, and a high ratio of submillimeter to bolometric luminosity.  Observations of the internal structure of Class 0 objects rely on interferometers because the sources remain deeply embedded within their parent molecular clouds.  To understand the kinematics of the collapse process requires observation of multiple molecular lines because the chemistry within these dense, cold cores is complex.  Carbon-bearing species are observed to be strongly depleted within the core interiors due to freeze-out onto grains. Nitrogen-bearing species were thought to deplete much more slowly than carbon-bearing species because of the low binding energy of the N$_2$ molecule, requiring long periods of time or very high densities to show depletion \\citep{ber97,ber02}.  However, recent laboratory results indicate that CO binding energies are comparable to those of N$_2$.  These studies also indicate that the degree of mixing of N$_2$ and CO ices impacts the desorption timescales \\citep{obe05,bis06}.  Nonetheless, it is clear that \\NTWOHP\\ (and its daughter product NH$_3$) show a strong rise in abundance as its main destroyer, \\CO, depletes from the gas, making \\NTWOHP\\ an effective tracer of internal core kinematics \\citep[e.g.,][]{aik05}.  When \\CO\\ is present however, reactions between C$^+$ or \\CO\\ lead to the eventual destruction of the two standard nitrogen-bearing tracers: NH$_3$ and \\NTWOHP.  Recent observations of several Class 0 sources have shown an absence of \\NTWOHP\\ emission at the source position on scales within several hundred AU of the protostar \\citep[i.e.,][]{jor04b}.  Typically, \\NTWOHP\\ emission is depressed at the core center and has been interpreted as destruction of \\NTWOHP\\ due to evaporation of CO and its isotopes from dust grains \\citep[e.g., L483-mm (hereafter L483),][]{jor04} or depletion of \\NTWOHP\\ at high densities \\citep[e.g., IRAM 04191+1522 (hereafter IRAM 04191),][]{bel02}.  An absence of \\NTWOHP\\ cannot be purely an indicator of age; L483 is thought to be in transition between Class 0 and Class I \\citep{taf00} while IRAM 04191 is estimated to be among the youngest Class 0 sources known \\citep{and99}.  Interestingly, in both cases, the \\NTWOHP\\ emission morphology is double peaked around the source center with a signature indicating rotation, and an anticorrelation is noted between \\NTWOHP\\ emission and the emission from tracers dominated by the outflow (\\CO\\ and \\HCOP). However, in L483, \\CEIGHTEENO\\ is centrally peaked on the source indicating that \\CO\\ has been evaporated from dust grains at the core center.  The combination of \\NTWOHP\\ and \\CEIGHTEENO\\ morphology could thus discriminate between whether an absence of \\NTWOHP\\ is due to depletion onto dust grains or the destruction of \\NTWOHP\\ in the presence of \\CO.  In this paper, we present observations of the internal structure of the protostellar core Barnard 1c (B1c) from the Berkeley-Illinois-Maryland Association (BIMA) array and the Five Colleges Astronomical Radio Observatory (FCRAO) 14 m telescope.  B1c was discovered during 850 \\micron\\ polarimetry mapping in Barnard 1 by \\citet{mw02}.  Recent IRAC imaging from Spitzer reveals that this source is highly reddened and deeply embedded in the B1 cloud with an extensive, highly collimated outflow \\citep{jor06}. The presence of central cavities within young protostellar cores is of particular relevance to B1c because a heated central region could explain why B1c has a unique signature in polarized emission.  Its polarization pattern suggests that the polarized intensity rises to the center of the core \\citep{mw02}, rather than flattening out as seen in other cores. When compared to the total intensity, a flat distribution in polarized intensity produces a declining ratio toward the peaks of cores.  All other low-mass starless and star-forming cores observed in polarized dust emission have a so-called ``polarization hole'' at high intensities, thought to arise from changes in magnetic field geometry or dust grain physics \\citep{mat05}.  One of the favored explanations for the polarization holes is that the grains at core centers are ineffective polarizers (due to changing grain physics).  Heating removes the outer grain mantles and could increase the polarization efficiency of the grains within cavities \\citep{whi01}.  Barnard 1 is part of the Perseus molecular cloud complex \\citep{bc86a} which is one of the closest star-forming regions to the Sun.  Its distance is the subject of some debate, with estimates ranging from 200 pc \\citep[based on extinction studies,][]{cer90} to 330-350 pc \\citep[based on the Per OB2 association,][]{dez99,bb64,her83}.  It has been suggested that the complex may be comprised of two clouds at 200 and 300 pc \\citep{cer85}, that there is a distance gradient along its length \\citep{sar79}, or that it is comprised of independent clouds at varying distances \\citep{rid05}.  To facilitate comparison with recent works on the Perseus cloud \\citep[i.e.,][]{eno06,jor06}, we adopt a distance of $250 \\pm 50$ pc to the Barnard 1 cloud, as determined from recent estimates of extinction \\citep{cer03} and measurements of parallax in members of IC 348 \\citep{beli02}.  This paper presents high resolution interferometric data from several molecular species and continuum emission at 2.7 and 3.3 mm.  The main objective is to determine whether B1c exhibits a central cavity which could help explain its unique polarization properties.  The observations and data reduction techniques are described in $\\S$ \\ref{obs}.  We present the continuum results and derive the core mass in $\\S$ \\ref{dustres}.  The molecular line data are presented in $\\S$ \\ref{moldatares}.  We discuss these data in $\\S$ \\ref{disc}.  Our findings are summarized in $\\S$ \\ref{sum}. ",
        "conclusions": "\\label{disc}  \\subsection{Column Density and Chemistry}  \\subsubsection{Column Density of \\NTWOHP} \\label{n2hpmass}  In order to estimate the \\NTWOHP\\ column density, we require estimates of the FWHM of the lines and estimates of the line opacity. Figure \\ref{deltav_grid} shows the distribution of line widths and opacities derived from the \\ntwohp\\ spectra. We derived the \\NTWOHP\\ column density using a simple curve of growth determined by assuming LTE excitation conditions at \\tex\\ $ = 12$ K, a line width of 0.75 \\kms, and the RADEX escape probability code\\footnote{www.strw.leidenuniv.nl$/ \\sim$moldata} \\citep{sho05}. This simple approach takes into account line opacity, but a more detailed envelope model with proper density and temperature slope will be required to get more accurate estimates (Matthews et al. in preparation).  Figure \\ref{n2hpcolumn} shows the column density distribution of \\ntwohp\\ emission in B1c.  Since the LTE analysis predicts opacities on the order of unity while the spectra and hfs component fitting suggest significantly higher opacities, it is likely that the excitation is not in LTE (i.e., \\tex\\ $< 12$ K). However, Figure \\ref{n2hpcolumn} is illustrative as it shows the same depression in \\NTWOHP\\ emission seen in the zeroth moment map of Figure \\ref{uniform}.  \\begin{figure} \\plotone{f16.ps} \\caption{Column density distribution of \\ntwohp\\ emission in B1c.  Contours are sum of the \\ntwohp\\ emission over all channels with signal exceeding 3 times  the rms level. } \\label{n2hpcolumn} \\end{figure}  \\subsubsection{Column Density of \\CEIGHTEENO} \\label{disc-c18o}  To derive the density at the core center, we estimate the column density at the peak of \\ceighteeno\\ emission.  We can derive the column density from the expression \\begin{eqnarray} \\lefteqn{N_{\\rm C^{18}O} = \\frac{3.34 \\times 10^{14}}{\\nu \\ \\mu^2 \\ (1 - \\exp(- h\\nu/kT_{\\rm ex}))} } \\nonumber \\\\ && {} \\times \\ \\frac{\\int{T_{\\rm MB} dV}}{\\exp(-Jh\\nu/2kT_{\\rm ex})} \\ \\frac{\\tau}{1 - \\exp(-\\tau)} \\label{c18ocolumn} \\end{eqnarray}  \\noindent where $\\nu$ is the frequency in GHz (109.78217 GHz), $\\mu$ is the dipole moment in Debye (0.11 Db), \\tex\\ is the excitation temperature, and $J$ is the lower rotational level for the transition. We assume $T_{\\rm kin} =$ \\tex\\ which is constant for all rotational levels, LTE is valid for the \\CEIGHTEENO\\ gas and that the optical depth effects are completely accounted for by the factor $\\tau / (1 - e^{- \\tau})$.  We evaluate the integral over the line width from the peak of the moment zero map, integrated over a linewidth of 1 \\kms. We estimate $\\tau$ from the peak temperature of the line via the expression \\begin{equation} T_{\\rm peak} \\approx \\frac{h\\nu}{k} \\ \\frac{1 - \\exp(-\\tau)}{\\exp(h\\nu/kT_{\\rm kin} - 1)} \\label{tpeak} \\end{equation}  \\noindent where $T_{\\rm kin}$ is the kinetic temperature. From the spectrum at the peak position, we evaluate $\\tau$ for $T_{\\rm peak} = 2$ K and $T_{\\rm kin} \\approx T_{\\rm ex} = 12$ K.  We find $\\tau = 0.244$, which is consistent with the \\CEIGHTEENO\\ emission being optically thin.  The strength of the emission at the peak is 1.16 \\jybeam\\ \\kms.  Using the conversion factor for our BIMA array synthesized beam at 109 GHz (1.12 K/\\jybeam), $\\int T_{\\rm MB} dV$ is 1.3 K \\kms.  Substitution of these values in equation (\\ref{c18ocolumn}) yields a column density of $1.03 \\times 10^{15}$ cm$^{-2}$ (the uncertainty, based largely on the flux calibration, is $\\sim 30$\\%).  Utilizing the abundance ratio [C$^{18}$O]/[H$_2$] $= 1.7 \\times 10^{-7}$ \\citep{fre82} yields an H$_2$ column density of $1.75 \\times 10^{22}$ cm$^{-2}$.  The effective radius of the beam in the \\CEIGHTEENO\\ observation is $1.8 \\times 10^{16}$ cm.  Assuming the core is as deep as it is wide, the density at the peak is $\\sim 9.5 \\times 10^5$ cm$^{-3}$.  This value is three times the estimated central density of $3.2 \\times 10^5$ cm$^{-3}$ from the analysis of dust emission maps by \\cite{kir06}, but a magnitude less than the estimate based on 3.3 mm continuum data in $\\S$ \\ref{disc-mass}. The difference is partly attributable to the different spatial filtering between the observations. The 2.7 mm data, with more comparable $(u,v)$ coverage to the \\CEIGHTEENO\\ data, yields densities of $(3.3 \\pm 1.0) \\times 10^6$ cm$^{-3}$, closer to the value derived above.  Aside from the uncertainties introduced by the different degrees of spatial filtering, the discrepency between the H$_2$ column density measured with dust and \\CEIGHTEENO\\ most likely arises due to depletion from standard ISM abundances, leading to underestimates of $N({\\rm H}_2)$.  Although the grains must be heated in B1c due to the centrally peaked distribution of \\CEIGHTEENO, some of the gas could remain frozen onto the dust grains.  Assuming the dust represents the total column of the inner envelope and the gas is depleted, then the [C$^{18}$O]/[H$_2$] ratio is not $1.7 \\times 10^{-7}$ but instead $1.3 \\times 10^{-8}$ (from 2.7 mm dust continuum), a factor of 13 less.  This value is similar to the values measured in Class 0 sources by \\cite{jor02}.  \\subsection{A Depression in \\NTWOHP\\ Emission near the Core Center}  As demonstrated by the integrated intensity and moment maps, the \\NTWOHP\\ emission is diminished near the core's center.  We are able to produce a more coherent picture of the central ``cavity'' in \\NTWOHP\\ emission by reweighting the \\NTWOHP\\ data presented in Figure \\ref{comb_n2hp} to uniform when creating the maps. The downweighting of shorter baselines degrades the S/N of the individual channels significantly, but the moment map produced retains high signal by utilizing all seven hyperfine components. Figure \\ref{uniform} shows the moment zero map produced from the uniformly weighted BIMA array data combined with FCRAO data. A high threshold (3$\\sigma$) was placed on inclusion of individual channels into the moment map to avoid channels between components.  The cross marks the position of the continuum source, and the triangle represents the pointing center of the JCMT observations presented in \\citet{mw02}.  The ring depicts the 14\\arcsec\\ beam of the JCMT.  The scale of the evaporated cavity is such that it would not have been resolved by the JCMT.  The scale of the cavity is 2.7\\arcsec\\ $\\times$ 2.15\\arcsec, at a position angle of 7\\degr.  This corresponds to 675 $\\times$ 540 AU at a distance of 250 pc.  \\begin{figure} \\plotone{f17.ps} \\caption{Moment zero map of the \\NTWOHP\\ emission from B1c with uniform weighting.  The new resolution is 3.09\\arcsec\\ $\\times$ 2.52\\arcsec, almost a factor of 2 improved over the data in $\\S$ \\ref{moments}. The cavity can now be seen as an absence of \\NTWOHP\\ emission offset to the west of the continuum peak, marked by a cross. The triangle marks the position of the source as determined from the SCUBA data; the ring depicts the JCMT beam of 14\\arcsec.  Contours are the same as in Figure \\ref{n2hpcolumn}.} \\label{uniform} \\end{figure}  The depression of \\NTWOHP\\ emission in the inner envelope and along the outflow path is likely due to the destruction of\u00a0\\ \\NTWOHP\\ by molecules that are released from grains, most notably CO and H$_2$O. Chemical models of an evolving core with a central luminosity source (i.e., Class 0) suggest that by 10$^5$ years a cavity of $\\sim 700$ AU will be seen in the\u00a0\\ \\NTWOHP\\ abundance and emission \\citep{lbe04}. This is the direct result of stellar heating providing enough energy to evaporate CO from grains within the core center.  This model may not be directly applicable to B1c due to differences in masses and/or luminosity, but it provides a rough estimate of the\u00a0\\ time that has passed since stellar ignition.  In addition, observations and models suggest that\u00a0 N$_2$H$^+$ will also be destroyed within the outflow \\citep{bac01,bnm98}, and we see a strong anticorrelation between the position of the outflow and the presence of \\NTWOHP\\ emission.  A central depression in \\NTWOHP\\ emission has also been detected in the transition Class 0/I source L483 by \\citet{jor04} and in the low luminosity object IRAM 04191 by \\citet{bel04}.  \\citet{lee05} estimate the inner edge of the IRAM 04191 cavity has a mean radius of 1400 AU. This is a factor of four times larger than the cavity detected in B1c. In a simple model where cavities grow over time, this suggests that B1c could be even younger than IRAM 04191's estimated age of $3 \\times 10^4$ yr since the onset of accretion \\citep{and99}.  The gradients and line widths measured in \\ntwohp\\ are strikingly similar in these two cores, indicating strong dynamical similarities, although we stress that IRAM 04191 is one of the new Class of very low luminosity objects \\citep[VeLLOs,][]{dif06,you04} while B1c is a relatively luminous ($4 \\pm 2 \\ L_\\odot$) Class 0 object (Matthews et al.\\ 2006, in prep).  We can estimate a kinematic age for this source based on the outflow data.  The projected linear extent of the emission coupled with the distance from the source yields estimates of the source age.  Using the brightest peak of the blue lobe detected in \\co\\ and assuming an inclination angle of 45$^\\circ$, the dynamical age is $3.7 \\times 10^4$ yr, comparable to that of the source IRAM 04191 in Taurus ($3 \\times 10^4$ years).  To produce an age on the order of $10^5$ years, the outflow inclination angle must be $\\sim 15$\\degr\\ (i.e., the outflow must lie almost along the line of sight).  The absence of \\NTWOHP\\ in IRAM 04191 is attributed to depletion within the cold core interior \\citep{bel02}.  The strong centrally peaked detection of \\ceighteeno\\ emission from B1c (see Figure \\ref{co_moments}) shows that depletion is not significant in B1c, and the anticorrelation between \\NTWOHP\\ with the core center and outflow lobes suggests that \\NTWOHP\\ is absent due to destruction by CO.  In addition, the strong outflow detected with Spitzer (Figure \\ref{spitzer_co}) is highly collimated, suggestive of a young source, but its spatial extent suggests that the collapse of the driving source may not have been as recent as that in IRAM 04191.  These characteristics are similar to those of L483 which is thought to be in transition to a Class I object \\citep{taf00}.   \\citet{lee05} note that the centers of the molecular outflow and envelope surrounding IRAM 04191 are offset from the the continuum source by approximately 560 AU. They suggest that the two cavities seen to the south of that source could be due to a binary in which one companion is too young to have any appreciable dust emission. There is also evidence for binarity within B1c, particularly in the behavior of its outflow ($\\S$ \\ref{res-outflow}), which shows evidence of precession.  The cavity center is also offset from the continuum source position.  There is as yet no detection of a second dust peak nor a second outflow in B1c.  \\subsection{Rotational Support}  The observed global rotation velocities within the B1 cloud are insufficient to support the cloud against collapse by a factor of $\\sim 8$ \\citep{bach90}.  The ages of embedded but optically visible objects LkH$\\alpha$ 327 and LkH$\\alpha$ 328 are between 4-6 $\\times 10^6$ yr \\citep{ck79}.  Based on this, \\citet{bach90} conclude that another mechanism must be providing substantial support to the B1 cloud.  Within B1c, the two peaks of the rotating torus can be used to derive the rotational energy relative to the $V_{LSR}$ at the core center of B1c itself.  The inner edges of the torus are offset from the position of the continuum peak.  For a sphere, the rotational energy is given by \\begin{equation} E_{\\rm rot} = \\frac{2}{5} \\ \\rho \\ V \\ v_r^2 \\label{rotenergy} \\end{equation}  \\noindent where $\\rho$ is the mean density, $V$ is the volume, and $v_r$ is the radial velocity relative to the center.  Based on estimates of the central density with \\CEIGHTEENO\\ ($\\S$ \\ref{disc-c18o}) and dust emission \\citep{kir06}, we adopt a mean density of $5 \\times 10^5$ cm$^{-3}$ for B1c.  The radial velocity of the two peaks relative each other is $\\sim 0.06$ \\kms\\ (as derived from Figure \\ref{spectra}).  Thus, we estimate the radial velocity relative to the core center to be a relatively small shift of $\\sim 0.03$ \\kms.  The mean offset of the peaks from the core's velocity center is $\\sim 4$ \\arcsec, or 1000 AU (for $d=250$ pc), encompassing a volume of $1.4 \\times 10^{49}$ cm$^3$.  Thus, based on equation (\\ref{rotenergy}), the rotational energy of the core is $\\sim 1.2 \\times 10^{39}$ ergs.  \\citet{mw02} derived the energetics of the B1c core based on estimates of the magnetic field dispersion and found rough equipartition between the gravitational, kinetic and magnetic energies (all $\\sim 10^{48}$ ergs).  The rotational energy is thus likely negligible to the support of the core compared to the other means of support available through nonthermal gas motions and magnetic field support.  The radius of the torus ($\\sim 1000$ AU) detected in B1c leads naturally to a discussion of a ``pseudo-disk'' \\citep{gal93} around the protostar within this core.  Such a disk is magnetically supported and grows with time.  \\citet{mw02} estimated the magnetic field strength in the low density regime around the B1 cores to be on the order of 30 $\\mu$G.  While the field strength within the B1c core may be higher, there is no Zeeman splitting measurement of sufficient spatial resolution to isolate the field strength of the individual cores within B1.  Assuming the presence of a pseudo-disk as described by \\citet{gal93}, we can use their equation (24) to estimate the age of the disk around B1c.  For a field strength of 30 $\\mu$G and a sound speed of 0.35 \\kms, a disk of radius 1000 AU has an age of $1.8 \\times 10^5$ yr, consistent with the expected age of a Class 0 protostar."
    },
    "0609/astro-ph0609235_arXiv.txt": {
        "abstract": "We have obtained new mid-infrared observations of 65 Class I/Flat-Spectrum (F.S.) objects in the Perseus, Taurus, Chamaeleon I and II, $\\rho$ Ophiuchi, and Serpens dark clouds. These objects represent a subset of the young stellar objects (YSOs) from our previous near-infrared multiplicity surveys. We detected 45/48 (94\\%) of the single sources, 16/16 (100\\%) of the primary components, and  12/16 (75\\%) of the secondary/triple components of the binary/multiple objects surveyed. One target, IRS 34, a 0$\\farcs$31 separation F.S. binary, remains unresolved at near-infrared wavelengths.  The composite spectral energy distributions (SEDs) for all of our sample YSOs are either Class I or F.S., and, in 15/16 multiple systems, at least one of the individual components displays a Class I or F.S. spectral index.  However, the occurrence of mixed pairings, such as F.S. with Class I, F.S. with Class II, and, in one case, an F.S. with a Class III (Cha I T33B), is suprisingly frequent. Such behavior is not consistent with that of multiple systems among  T Tauri stars (TTS), where the companion of a classical TTS also tends to be a classical TTS, although other mixed pairings have been previously observed among Class II YSOs. Based on an analysis of the spectral indices of the individual binary components, there appears to be a higher proportion of mixed Class I/Flat-Spectrum systems (65-80\\%) than that of mixed Classical/Weak Lined T Tauri systems (25-40\\%), demonstrating that the envelopes of Class I/Flat-Spectrum systems are rapidly evolving during this evolutionary phase. In general, the individual binary/multiple components suffer very similar extinctions, A$_{v}$, suggesting that most of the line-of-sight material is either in the foreground of the molecular cloud or circumbinary. We report the discovery of a steep spectral index secondary companion to ISO-Cha I 97, detected for the first time via our mid-infrared observations.  In our previous near-infrared imaging survey of binary/multiple Class I and F.S.YSOs, ISO-Cha I 97 appeared to be single. With a spectral index of $\\alpha$ $\\geq$ $+$3.9, the secondary component of this system is a member of a rare class of very steep spectral index YSOs, those with $\\alpha$ $>$ $+$3. Only three such objects have previously been reported, all of which are either Class 0 or Class I. ",
        "introduction": "We have known for many years that the majority of solar-to-late-type field stars are binaries or multiples (Abt \\& Levy 1976; Duquennoy \\& Mayor 1991; Fischer \\& Marcy 1992), but only in the past decade have significant numbers of pre-main-sequence (PMS) stars been surveyed for multiplicity. Near-infrared surveys of nearby, young dark cloud complexes have shown that young, low-mass TTSs have binary fractions which are greater than or equal to that of the field (e.g., Ghez, Neugebauer, \\& Matthews 1993; Mathieu 1994; Simon et al. 1995; Ghez et al. 1997; Barsony, Koresko, \\& Matthews 2003), and appear to be coeval, that is, at the same evolutionary age (Hartigan, Strom, \\& Strom 1994; Brandner \\& Zinnecker 1997). Observations at millimeter continuum wavelengths also find multiple sources at the origins of extended molecular outflows and optical jets (Looney, Mundy, \\& Welch 1997). Thus, the formation of binary and multiple systems appears to be the rule, rather than the exception, in star-forming regions.  Recently, multiplicity surveys have been extended to include even younger, self-embedded YSOs, those with Class I and flat-spectrum SEDs (Lada 1987; Adams, Lada, \\& Shu 1987, 1988). The first survey for multiplicity among such objects was conducted by Looney, Mundy, \\& Welch (2000) with the BIMA millimeter interferometer at 2.7 mm. Of the eight Class 0 and Class I YSOs surveyed, all were members of binary systems or small groups. Not long thereafter, Reipurth (2000) analyzed the multiplicity of 14 young (ages $\\leq$ 10$^{5}$ yr) sources which drive giant Herbig-Haro flows. Between 79\\% and 86\\% (depending on the range of separations considered) of these sources had at least one companion and, of these, half were higher order multiple systems.  These fractions are even larger than those found among either the PMS TTS or field star populations.  In order to determine the multiplicity properties of embedded Class I and flat-spectrum sources with reasonable statistical confidence, which the previous surveys do not provide, Haisch et al. (2002, 2004) and Duch\\^{e}ne et al. (2004) conducted the first near-infrared imaging surveys of nearby molecular clouds which permitted observations at resolutions $\\leq$ 1$^\\prime$$^\\prime$, and therefore enabled detections of companions as close as 100 - 300 AU, depending on the distance to the clouds. Merging the three surveys into one, these authors derived a ``restricted\" companion star fraction ($\\Delta$K $\\leq$ 4 magnitudes and a separation range of 300 - 1400 AU) of 16\\% $\\pm$ 3\\%, with all clouds presenting fully consistent fractions. This frequency is in excellent agreement with that obtained for TTSs in the same star-forming regions, and is approximately twice as high as that observed for late-type field dwarfs.  While near-infrared data are very efficient in identifying binary/multiple components, such observations are not at a long enough wavelength to determine the evolutionary state of the objects via their spectral indices (as defined in Section 3). However, since contamination from photospheric emission is minimal at mid-infrared wavelengths, observations at 10 $\\mu$m, in conjunction with near-infrared $JHKL$ data, can be used to ascertain the spectral index, and thus the evolutionary state of the individual binary/multiple YSOs. The 10 $\\mu$m radiation from each component does not originate from the star itself, but from a ``photosphere\" of surrounding dust heated to several hundred degrees. Radiative transfer models of Class I YSOs have shown that the 10 $\\mu$m photosphere is located about 1 AU from a low-mass protostar (e.g., Kenyon, Calvet, \\& Hartmann 1993; Chick \\& Cassen 1997). Thus, the individual components should be point sources in the mid-infrared. Indeed, recent 10 $\\mu$m observations of Class I objects have often found them to be unresolved (e.g., Barsony, Ressler, \\& Marsh 2005). Finally, no current space-based observatory can match the angular resolution achievable from the ground at mid-infrared wavelengths, which is ideal for studying very young binaries.  Mid-infrared studies of small samples of Class I/flat-spectrum binaries have already been published (Liu et al. 1996; Girart et al. 2003; Ciardi et al. 2003). In this paper, we present the results of a mid-infrared imaging survey of 65 single and binary/multiple Class I and flat-spectrum YSOs from our previous near-IR surveys (Haisch et al. 2002; 2004) in the $\\rho$ Ophiuchus ($d$ = 125 pc; Knude \\& H{\\o}g 1998), Serpens ($d$ = 310 pc; de Lara, Chavarria-K., \\& L\\'{o}pez-Molina 1991), Taurus ($d$ = 140 pc; Kenyon \\& Hartmann 1995), Perseus ($d$ = 320 pc; Herbig 1998), and Chamaeleon I and II ($d$ = 160 pc and $d$ = 178 pc; Whittet et al. 1997) star-forming regions, respectively. Following the near-infrared multiplicity survey of Haisch et al. (2004), we impose a lower limit for detectable separations of 100 -- 140 AU for Taurus, Chamaeleon, and $\\rho$ Ophiuchi and $\\sim$ 300 AU for Perseus and Serpens, which was set by the seeing in the near-infrared. We have also imposed an upper limit to the separations of 2000 AU in order to avoid including sources which are not gravitationally bound systems (e.g., Reipurth \\& Zinnecker 1993; Simon et al. 1995). In addition, sensitivity calculations from our near-infrared data indicate that we can detect a $K$ = 4 magnitude difference between the primary and companion at a separation of 1\\arcsec \\hspace*{0.05in}at the 5$\\sigma$ confidence level. Thus, we restrict our analysis to component magnitude differences $\\Delta$$K$ $\\leq$ 4 mag.  We discuss our mid-infrared observations and data reduction procedures in $\\S$2. In $\\S$3, we present the results of our survey, and discuss the results in $\\S$4. We summarize our primary results in $\\S$5. ",
        "conclusions": "We have obtained new mid-infrared observations of 64 Class I/flat-spectrum objects in the Perseus, Taurus, Chamaeleon I and II, $\\rho$ Ophiuchi, and Serpens dark clouds. These objects represent a subset of the YSOs from our previous near-infrared multiplicity surveys (Haisch et al. 2002, 2004). We detected 45/48 (94\\%) of the single sources, 16/16 (100\\%) of the primary components, and 12/16 (75\\%) of the secondary/triple components of the binary/multiple objects surveyed. The tight mid-IR binary, IRS34, remains unresolved at near-IR wavelengths.  While the composite SEDs for all of our sample YSOs are either Class I or Flat-Spectrum, individual source components may display Class II, or, in one case, Class III, spectral indices.  The SED classes of the primary and secondary components are frequently different. For example, a Class I object may be found to be paired with a Flat-Spectrum source, or a Flat-Spectrum source paired with a Class II YSO. Such behavior is not consistent with what one typically finds for TTSs, where the companion of a classical TTS also tends to be a classical TTS (Prato \\& Simon 1997; Duch\\'{e}ne et al. 1999). Mixed pairings, however, have been previously observed among Class II YSOs (Ressler \\& Barsony 2001; McCabe et al. 2006).  Based on an analysis of the spectral indices of the individual binary components, there appears to be a higher proportion of mixed Class I/Flat-Spectrum systems (65-80\\%) than that of mixed Classical/Weak Lined T Tauri systems (25-40\\%). This demonstrates that the envelopes of Class I/Flat-Spectrum systems are rapidly evolving during this evolutionary phase, although they may still be coeval.  In general, the individual binary/multiple components suffer very similar extinctions, A$_{v}$, suggesting that most of the line-of-sight material is either in the foreground of the molecular cloud or circumbinary. However, the GY 244/WL 5 pair, whose A$_{v}$ values differ by 11 magnitudes, is a notable exception.  If the projected separation of this pair ($\\sim$ 1100 AU), equivalent to its physical separation, then this system could easily be gravitationally bound. However, it is also possible that GY 244 and WL 5 represent a chance projection.  ISO-Cha I 97 was detected as a single star in our near-infrared imaging survey of binary/multiple Class I and Flat-Spectrum YSOs (Haisch et al. 2004). Our mid-infrared observations have revealed that this source is in fact binary. Combining our $K$-band sensitivity limit from HGBS04 with our 10 $\\mu$m flux yields a lower limit to the spectral index of the secondary component of ISO-Cha I 97 of $\\alpha$ $\\geq$ $+$3.9. This very steep spectral index places the secondary of ISO-Cha I 97 in a class of YSO that has heretofore been rarely known, i.e., those with $\\alpha$ $>$ $+$3. Further very steep spectrum YSOs are expected to be discovered with the Spitzer Space Telescope."
    },
    "0609/astro-ph0609003_arXiv.txt": {
        "abstract": "{We study the constraints on the inflationary parameter space derived from the 3 year WMAP dataset using ``slow roll  reconstruction'', using the SDSS galaxy power spectrum to gain further leverage where appropriate.  This approach inserts the inflationary slow roll  parameters directly into a Monte Carlo Markov chain estimate of the  cosmological parameters, and uses the inflationary flow hierarchy to compute the parameters' scale-dependence.   We work with the first three parameters ($\\epsilon$, $\\eta$ and $\\xi$) and pay close attention to the possibility that the 3 year WMAP  dataset contains evidence for a ``running'' spectral index, which is dominated by the $\\xi$ term.  Mirroring the WMAP team's analysis we find that the permitted distribution of $\\xi$ is broad, and centered away from zero. However, when we require that inflationary parameters yield at least 30 additional e-folds of inflation after the largest observable scales leave the horizon, the bounds on $\\xi$ tighten dramatically.  We make use of the absence of an explicit pivot scale in the slow roll reconstruction formalism to determine the dependence of the computed parameter distributions on the pivot. We show that the choice of pivot has a significant effect on the inferred constraints on the inflationary variables, and the spectral index and running derived from them. Finally, we argue that the next round of cosmological data can be expected to place very stringent constraints on the region of parameter space open to single field models of slow roll inflation. } \\preprint{}  \\begin{document} ",
        "introduction": "The primordial perturbation spectrum poses two distinct problems  for cosmology. Observationally, we wish to extract the  spectrum's amplitude and scale dependence from data, while theoretically we seek to understand the origin of  the perturbations.   After the release of the 3-year WMAP dataset \\cite{Hinshaw:2006ia,Page:2006hz,Spergel:2006hy} there is considerable optimism that we will soon  measure the departure from scale  invariance in the spectrum at significance levels greater than   $3 \\sigma$.\\footnote{WMAP is an acronym for the {\\sl Wilkinson Microwave Anisotropy Probe}. Here WMAPI refers to the first year's data from the satellite, and WMAPII refers to the second WMAP data release, which is drawn from the first three years' data.}     On the theoretical front, inflation \\cite{Guth:1980zm,Albrecht:1982wi,Linde:1981mu}  remains the leading theoretical paradigm for understanding  the very early universe and  the origin of the perturbations.  In general terms, the latest round of cosmological data has confirmed the broad predictions of inflation -- the universe is very close to being flat, it has a roughly scale-invariant perturbation spectrum and $\\langle TE \\rangle$, the  cross correlation between the  E mode of the polarization and the temperature, has the form predicted by inflation.  Beyond this rosy picture, however, lies a number of qualifications. Firstly, the observed form of  $\\langle TE \\rangle$ is not a unique prediction of inflation, but is also a generic feature  of models where the perturbations are laid down during a collapsing phase prior to a bounce, such as the ekpyrotic scenario \\cite{Khoury:2001wf}.\\footnote{The $\\langle TE \\rangle$ signal  is the same for any mechanism that induces adiabatic perturbations that are correlated on scales much larger than the post-recombination horizon volume.  ``Bouncing'' models do this by laying down perturbations during the contracting phase, when the comoving horizon size is shrinking. }   Secondly, inflation can be implemented in a vast number of different ways, and each mechanism produces a potentially distinctive perturbation spectrum. Consequently, the inflationary paradigm does not make an unambiguous prediction for the detailed form of the perturbation spectrum.  Some authors have posited that inflation has a ``default'' perturbation spectrum \\cite{Boyle:2005ug}.  The underlying arguments leading to this position rely on simplicity and naturalness, rather than a watertight ``no go'' theorem.  Consequently, inflation itself will not be ruled out if the perturbation spectrum turns out to differ from this simple expectation, but it is certainly true that many well-known models   would be excluded.  Looking at data, it is always tempting to focus on possible anomalies -- parameters that differ noticeably from their ``expected'' values, but for which the discrepancy falls well short of a $3\\sigma$ detection. Since the release of the original WMAPI dataset \\cite{Bennett:2003bz,Spergel:2003cb,Peiris:2003ff}, three aspects of the power spectrum have received considerable attention: the ``low''  values of $C_\\ell$ at small $\\ell$, the glitches around the first Doppler peak and at $\\ell \\sim 40$, and the apparent evidence for a running, or scale-dependent, spectral index that arose when the WMAPI data was combined with other probes of the power spectrum at shorter scales.   The status of all of these anomalies has evolved since the WMAPI release. The quadrupole is still $\\sim 2 \\sigma$ lower than the LCDM best-fit value, but the octupole has moved closer to the LCDM value, thanks to an optimal extraction of the spectrum.  The glitch near the first Doppler peak is not present in the WMAPII dataset with its smaller errorbars, while the $\\ell \\sim 40$ discrepancy is still seen \\cite{Cline:2006db,Covi:2006ci}.  Finally, the WMAPII dataset contains a preference for a running spectral index  when viewed alone, in combination with other CMB data, and/or large scale structure information such as 2dFGRS \\cite{Cole:2005sx} and the Sloan Digital Sky Survey \\cite{Tegmark:2003uf}.  However, a reanalysis by Seljak {\\em et al.}, which includes SN1a data and information from the Lyman-$\\alpha$ forest, puts a tight constraint on any running \\cite{Seljak:2006bg} --- this result seems to hinge on the inclusion of the Lyman-$\\alpha$ data. Further analysis of a possible running index is provided by \\cite{Feng:2006ui}.  In this paper, we focus on the running of the spectral index, and the constraints it places on the inflationary parameter space.  Conventionally, one specifies the spectrum at a pivot point in terms of the spectral index $n_s$ and the running $\\frac{dn_s}{d\\ln{k}}$.  Inflation predicts $n_s$ and $\\frac{dn_s}{d\\ln{k}}$ in terms of the slow roll parameters, which are functions of the potential and its derivatives.   In \\cite{Easther:2006tv}  we showed that if the running is as large as the central value extracted from the WMAPII dataset,  no {\\em simple\\/} model of inflation can simultaneously provide both this running and a sufficient amount of inflation.\\footnote{Which we take to be a minimum of 30 e-foldings, and by  ``simple'' we mean a single inflaton, minimally coupled to gravity, whose potential can be well specified in terms of its first three derivatives.}  Currently, there are simple models of inflation within the errorbars, but if the constraints from the data continue to tighten around the median value of the running, they will be ruled out. We arrived at this result by considering the Hubble Slow Roll hierarchy, otherwise known as the inflationary flow equations \\cite{kinney:2002,easther/kinney:2003,liddle:2003,Peiris:2006ug}.  These codify the scale dependence of the slow roll parameters: given the values of the slow roll parameters at  a single moment during the inflationary era, we can compute their values at all other times.  A further virtue of this approach is that when the infinite slow roll hierarchy is truncated at finite order, the truncation is preserved by the evolution equations. Using the slow roll hierarchy means that we effectively dispense with the pivot point, other than when comparing our results to treatments using the standard $n_s$ and $\\frac{dn_s}{d\\ln{k}}$ parameterization of the spectrum.  As we  showed in \\cite{Peiris:2006ug}, we can include the slow roll parameters $\\{\\epsilon,\\eta,\\xi,\\cdots\\}$ in the cosmological parameter estimation process, and entirely avoid the use of the spectral index  and its running. Further, the slow roll hierarchy can be solved to produce the underlying inflationary potential, leading us to dub this approach ``slow roll reconstruction.''\\footnote{The overall amplitude of the spectrum  is fixed by an arbitrary constant of integration, which can only be measured directly if the tensor spectrum is observed.}    Of course, this presupposes that the perturbations were produced by inflation, so slow roll reconstruction does not replace fits to $n_s$ and $\\frac{dn_s}{d\\ln{k}}$, which are purely empirical characterizations of the spectrum and contains no strong theoretical priors.\\footnote{The index and running are effectively the first two terms in a  Taylor series expansion, so their use does make the tacit assumptions that the spectrum is a smooth, differentiable function of the comoving wavenumber $k$.} However, if we do adopt an inflationary prior, fitting to the slow roll parameters directly and using the flow equations to obtain their scale dependence ensures that we gain the maximum possible leverage from imposing this constraint.  It is well known that the scalar and tensor perturbations obey a consistency condition at lowest order in slow roll,\\footnote{In practice, almost all estimates of the cosmological parameters assume that the amplitude and slope of the tensor spectrum have the correlation predicted by minimal models of inflation. Since the tensor spectrum is merely bounded from above,  rather than observed directly,  this constraint makes a Monte Carlo Markov Chain  analysis more tractable without excluding viable models.} but there are further, more subtle correlations at second order and beyond which tie together the running and  the  scalar spectral index \\cite{Cortes:2006ap}.  By using slow roll reconstruction we automatically include these correlations in our analysis.  Moreover, this removes a potential source of ambiguity inherent in the introduction of a scale dependent spectral index -- which is how to best parametrize the scale dependence?  While $\\frac{dn_s}{d\\ln{k}}$ is arguably the simplest measure of the scale dependence, it is by no means unique. In particular, by adopting it we are implicitly assuming that while $n_s$ varies with $k$, $\\frac{dn_s}{d\\ln{k}}$ does not.  By specifying the spectrum in terms of the slow roll parameters we ensure our spectrum has the scale dependence predicted by inflation, and the primary limitation is the highest order  slow roll parameter we include in our fit. While this is obviously a choice that will be guided by the quality of the data, it means that we can describe the scale dependence in a relatively unambiguous way.\\footnote{This issue mirrors the concern when the dark energy equation of state is allowed to vary with time, as there is no unique way of specifying that time dependence \\cite{Corasaniti:2002vg,Linder:2005ne}.}   Finally, since we obtain the overall form of the potential as a by-product of this analysis, we can include constraints based on the total number of e-folds of inflation.  The closest analogue to slow roll reconstruction is the approach of Leach and collaborators using the Horizon Flow Functions \\cite{Leach:2002ar,Leach:2002dw,Leach:2003us}, who compute the spectral index and running as functions of the slow roll parameters. As in slow roll reconstruction, one has an explicit inflationary prior, and the slow roll parameters appear directly in the parameter set.  The advantage of this approach is that it can be implemented more simply in existing codes, since it translates between the slow roll formalism and  $n_s$ and $\\frac{dn_s}{d\\ln{k}}$, whereas in order to implement slow roll reconstruction we have to numerically solve the differential equation that maps the field evolution into the comoving wavenumber $k$.\\footnote{In practice this is not computationally expensive, especially when compared to the numerical evaluation of the CMB spectrum and likelihood codes.} However, slow roll reconstruction can be more easily extended to include higher order terms in slow roll, and copes naturally with heterogeneous datasets that span a wide range of wavelengths -- this is not impossible in the formalism of \\cite{Leach:2002ar,Leach:2002dw,Leach:2003us}, but would require  more algebraic labor. Secondly, slow roll reconstruction can naturally include constraints based on the total duration of inflation. Finally, and perhaps most importantly, because \\cite{Leach:2002ar,Leach:2002dw,Leach:2003us} use $n_s$ and $\\frac{dn_s}{d\\ln{k}}$, they are effectively Taylor expanding the power spectrum round a fixed pivot scale, $k_0$.  Conversely,  we can easily post-process our chains to present parameter constraints  at an arbitrary pivot.   In \\cite{Finelli:2006fi}, the formalism of  \\cite{Leach:2002ar,Leach:2002dw,Leach:2003us}  is applied to the WMAPII dataset to produce constraints on the slow roll parameters, and this work is directly comparable to the results we present below.   Slow roll reconstruction also draws on the stochastic Monte Carlo reconstruction algorithm, described in \\cite{easther/kinney:2003}.  This procedure does not yield unambiguous estimates of the slow roll variables, but it permits estimates of the impact of higher order terms  in the slow roll expansion  \\cite{Peiris:2006ug} and has been applied to cosmological data in \\cite{Peiris:2003ff,Kinney:2003uw,Malquarti:2003ia,Kinney:2006qm}; a similar approach can also be found in \\cite{Hansen:2001eu,Caprini:2002jy}. An alternative approach of computing exact power spectra for specific potentials has been applied to the WMAPII data by \\cite{Martin:2006rs}.   Our specific goal here is to extend our analysis in \\cite{Peiris:2006ug} to include the third  slow-roll parameter, $\\xi$, in a Monte Carlo Markov Chain  [MCMC] analysis.  This parameter dominates  the running of the spectrum, so it must be incorporated in any discussion of a running spectral index, especially in the light of the hints in the WMAPII data that the spectrum has a non-trivial scale-dependence.  We will also include constraints based on $N$, the number of e-folds of inflation remaining after the fiducial scale has left the horizon.       We consider three different sets of chains:  WMAPII + $N>30$, WMAPII + SDSS + $N>30$, and WMAPII + SDSS with no $N$ constraint. While we put a cut on $N$, we do not take the further step of correlating $N$ with the inflationary energy scale.  A low $N$ generally implies that inflation occurs well below the GUT scale but making this statement precise requires introducing assumptions about the physics of reheating, which we wish to avoid.\\footnote{For instance, if the post-inflationary  universe has an epoch in which  $a(t) \\sim t$ -- as might happen if the end of inflation is followed  by the formation of a ``frustrated'' network of cosmic strings --  $N$ and the inflationary energy scale are largely uncorrelated (see e.g. \\cite{Burgess:2005sb}). This is admittedly an extreme example,  and we may explore constraints arising from the connection between $N$ and the inflationary scale  in future work.}  Roughly speaking, however, imposing $N>30$ is equivalent to setting the inflationary scale at or above the TeV scale.   This ensures that inflationary energy scales are not probed by precision electroweak measurements, and allows  baryogenesis to be driven by electroweak physics.  Given that the GUT scale is often regarded as the ``natural'' inflationary scale, assuming that inflation happens above the TeV scale is a very mild theoretical prior. However, demanding $N>30$ leads to very tight bounds on $\\xi$. Translating our results into bounds on $\\frac{dn_s}{d\\ln{k}}$ yields a tighter constraint than that of Seljak {\\em et al.} \\cite{Seljak:2006bg} when we add the $N>30$ prior.  When we drop the constraint on $N$the central value of $\\xi$ is unexpectedly large, and $\\xi = 0$ is apparently excluded at the $2\\sigma$ level by the WMAPII+SDSS dataset. When we compute $N$ for these chains we find a distribution peaked at 10-15 e-folds, consistent with   \\cite{Easther:2006tv}, but in conflict with successful cosmological inflation.     We then turn to the pivot dependence of the bounds on the inflationary and spectral parameters, and show that in the absence of a prior on $N$, choosing a different pivot scale (comparing constraints at $0.002 \\mbox{ Mpc}^{-1}$ to those at $0.02 \\mbox{ Mpc}^{-1}$, in this case) somewhat tightens the bounds on $\\xi$ and $\\frac{dn_s}{d\\ln{k}}$, while shifting $n_s$ significantly to the red. In the case with a ``sufficient e-folds\" prior, the pivot dependence of the constraints is not strong. We stress that these results must be interpreted with caution, and make no claims to have found strong evidence for a non-zero value of $\\xi$, as we have not yet made use of the full cosmological dataset. However, these results help us understand the significance of the weak preference for a negative running seen in the WMAPII dataset, and give us cause for optimism that future data will allow us to tightly constrain the inflationary parameter space.    This paper is arranged as follows. In Section 2 we briefly review the slow roll hierarchy, and its use in Monte Carlo Markov Chain analyses. Section 3 contains the constraints on the inflationary parameter space derived from our chains. In Section 4 we give our interpretation and analysis, including a discussion of the pivot-dependence of our constraints on the inflationary parameter space. ",
        "conclusions": "}     \\TABLE[!t]{ \\begin{tabular}{||c|c|c|c||} \\hline Parameter           & WMAPII+SDSS   \t   &  WMAPII             &  WMAPII+SDSS \\\\ & (no $N$ restriction) & ($N>30$) \t         & ($N>30$) \\\\ \\hline $\\epsilon$          & $<0.013$ (95\\% CL)    & $<0.033$  (95\\% CL)  &  $<0.015$ (95\\% CL)\\\\ \\hline $\\eta$      \t    & $0.029 \\pm 0.021$    & $0.024^{+0.029}_{-0.028}$  &  $0.0 \\pm 0.014$ \\\\ \\hline $\\xi$               & $0.029 \\pm 0.013$    & $0.0 \\pm 0.003$    &  $0.0 \\pm 0.002$ \\\\ \\hline $\\ln[10^{10} A_s]$  & $3.17 \\pm 0.06$      & $3.03 \\pm 0.10$    &  $3.17 \\pm 0.06$ \\\\ \\hline \\end{tabular} \\caption{Constraints on HSR parameters defined at $k=0.002$ Mpc$^{-1}$. Constraints are at the 68\\% confidence level unless otherwise noted.}\\label{table:hsr_constraints}}   \\TABLE[!t]{ \\begin{tabular}{||c|c|c|c||} \\hline Parameter           & WMAPII+SDSS   \t   &  WMAPII             &  WMAPII+SDSS \\\\ & (no $N$ restriction) & ($N>30$) \t         & ($N>30$) \\\\ \\hline $n_s$               & $1.08 \\pm 0.05$    & $0.99 \\pm 0.02$    &  $0.98 \\pm 0.02$ \\\\ \\hline $d n_s/d\\ln k$      & $-0.057^{+0.027}_{-0.026}$   & $0.004 \\pm 0.006$   &  $0.0^{+0.004}_{-0.003}$ \\\\ \\hline $r$                 & $<0.22$ (95\\% CL)   & $<0.56$ (95\\% CL)  &  $<0.24$ (95\\% CL) \\\\ \\hline $n_t$               & $>-0.026$ (95\\% CL)  & $>-0.067$ (95\\% CL)  &  $>-0.030$ (95\\% CL) \\\\ \\hline \\end{tabular} \\caption{Constraints on derived primordial parameters defined at $k=0.002$ Mpc$^{-1}$. Constraints are at the 68\\% confidence level unless otherwise noted.}\\label{table:derived_constraints}}      In this paper we have used slow roll reconstruction to  extract bounds on the inflationary potential from observational data. These bounds make strong use of an inflationary prior, and thus apply to cosmological models where the density perturbation spectrum is laid down by  single, slow rolling, minimally coupled inflaton field.   Our results align with the broad picture supplied by previous analyses of the WMAPII dataset: there is noticeable but by no means compelling evidence for a significant scale dependence (running) in the power spectrum when one considers WMAPII and SDSS on their own. However, adding a prior constraint that rules out models with an unacceptably small number of e-folds significantly tightens the bounds on any possible running.  The combination of WMAPII and SDSS on their own are compatible with a broad class of perturbation spectra, but this range shrinks substantially when further constraints are added to the estimation process.  In particular, we consider inflationary models where at least 30 e-folds of inflation occur after CMB scales have left their horizon.  Similar results are seen in conventional analyses of the spectrum when Lyman-$\\alpha$ forest data is included \\cite{Seljak:2006bg}.   We further analyze the impact on the choice of pivot on the quoted ranges for the spectral parameters.  Slow roll reconstruction provides an estimate for the inflationary potential over its entire range, thanks to the closed nature of the truncated slow roll hierarchy.  Consequently, we actually have access to the underlying power spectrum and inflationary potential, although these quantities are fully specified by the truncated slow roll hierarchy -- which here consists of three independent parameters. However, we can use the flow equations to post-process our chains to give the constraints on the standard spectral parameters at an arbitrary pivot. We make use of this feature, and study the change in constraints on the inflationary parameter space when changing $k_0$ from  $0.002 \\mbox{ Mpc}^{-1}$ to $0.02 \\mbox{ Mpc}^{-1}$. Without a prior on the number of e-folds $N$, the constraints on $n_s$ and $\\frac{dn_s}{d\\ln{k}}$ are found to be significantly dependent on the pivot scale. The large running seen without this prior cannot be removed by shifting the pivot -- changing $k_0$ shifts $n_s$ well to the red, so the hints of an overall scale dependence cannot be expunged simply by changing $k_0$.  One lesson from this analysis is that pivot-dependent estimates of power-law observables, while valid, need to be interpreted with great care. For example, in the case of the constraints presented at $k=$ 0.02 Mpc$^{-1}$ for WMAPII+SDSS without any prior on $N$, the fact that this pivot scale is an inflection point ``hides'' a tendency to run from blue on large scales to red on small scales as you move away from this pivot. That is, when there is a strong running, any slice in $k$ of the primordial power spectra can look quite different in $n_s$, $r$, $\\frac{d n_s}{d\\ln{k}}$ space. In our view, it is a better aid to interpretation to present constraints over the entire $k$ range implied by one's parameterization, as in Figure \\ref{fig:pk} -- in principle,  this is a simple thing to do in the MCMC methodology for {\\sl any} parameterization of $P(k)$ -- and then compare with specific inflationary models over this entire range of scales. What does matter is whether the constraints at a wide range of $k$ either encompass or exclude the prediction from that specific potential. For absolute completeness, the ``band'' of constraints from data on $P(k)$ should be compared with a ``band'' for the model under consideration, encompassing the theoretical uncertainty in both the model parameters and the reheating process  \\cite{Kinney:2005in}.   Where they are directly comparable, our results broadly agree with  those of Finelli {\\em et al.\\/} \\cite{Finelli:2006fi}.    However,  the authors of \\cite{Finelli:2006fi} use a different formulation of the slow roll expansion from the one adopted here, which allows their $\\epsilon_3$ to be of order unity, and they find they must pay careful attention to their prior for this parameter.\\footnote{In the formalism of \\cite{Finelli:2006fi}, $\\epsilon_3$ only appears with a small factor, so while this parameter need not be small, its overall contribution is still typically sub-dominant.} Conversely, our ``third'' parameter, $\\xi$ needs no such special handling, and we simply take all our slow roll parameters to lie within $[-1,1]$, which is far larger than the range allowed by the data. Secondly, our running ($\\alpha$) is effectively scale-dependent, since $\\xi$ is a function of scale, but  \\cite{Finelli:2006fi} specify their running at a fixed pivot, and it is not a function of scale.  Finally, by using slow roll reconstruction we are able to include constraints on the remaining number of e-folds $N$, a possibility which is not considered in the analysis of  \\cite{Finelli:2006fi}.    The tension between the running and the number of e-folds confirms our theoretical analysis in \\cite{Easther:2006tv}.  Algebraically, the running index is dominated by $\\xi$, since $\\epsilon$ and $\\eta$ are both required by the data to be relatively small, and these terms only contribute quadratically to $\\frac{dn_s}{d\\ln{k}}$. When this parameter becomes large and positive the flow equations ensure that $\\epsilon$ and $\\eta$ will grow quickly. In this case $\\epsilon$ becomes equal to unity in $\\sim 15$ e-folds, and inflation terminates far too quickly for it to solve the usual cosmological problems.  In this analysis we have post-processed our chains to obtain likelihood contours for $N$ in the $(\\xi,N)$ plane, and we see that $N$ does in fact peak around 10 e-folds. Conversely, if we impose a prior on $N$, the permitted range of $\\frac{dn_s}{d\\ln{k}}$ shrinks considerably.  At large values, $\\xi$ has a long degeneracy because of the lack of constraining power in the current dataset. This is possibly a marginalization artifact, as we found that a large running is correlated with values of $\\Omega_m$ beyond the range seen in the standard concordance cosmology. If ``simple'' models of inflation are a good description of the universe, the rising quality of astrophysical data will eventually break this degeneracy, and the running will fall somewhere within the range found when the $N>30$ prior is applied here. Conversely if the data contracts around the large negative median value of the running, we would learn that inflation is non-minimal in some way, as discussed in \\cite{Easther:2006tv}. In either case, slow roll reconstruction will be able to put tighter constraints on $N$ as the constraints on $\\xi$ tighten.  In a subsequent paper we plan to assess the ability slow roll reconstruction to put bounds on the inflationary parameter space with the data from different proposed and projected experiments.  In addition, we will need to include $\\lambda_3$, the fourth slow roll parameter, in these calculations, to be sure than any conclusions we reach are not a function of a premature truncation of the slow roll hierarchy. On the other hand, if this latter analysis shows that $\\lambda_3$ is necessary to describe the data four independent parameters would be needed to describe the inflationary potential, which may be an indication that inflation is non-minimal.  Finally we stress that while we find evidence that a significant breaking of scale-invariance in the primordial power spectrum is consistent with the 3-year WMAP dataset in the presence of an inflationary prior, there is no compelling evidence for a running spectral index. However, the analysis here gives us cause for optimism that this question can be settled in the near future, and that slow roll reconstruction provides an elegant and powerful framework for analyzing the cosmological constraints  on slow roll inflation."
    },
    "0609/astro-ph0609529_arXiv.txt": {
        "abstract": "We describe deep radio imaging at 1.4\\,GHz of the 1.3 square degree Subaru/\\textit{XMM-Newton\\/} Deep Field (SXDF), made with the Very Large Array in B and C configurations. We present a radio map of the entire field, and a catalogue of 505 sources covering 0.8 square degrees to a peak flux density limit of 100\\,$\\umu$Jy. Robust optical identifications are provided for 90\\,per cent of the sources, and suggested IDs are presented for all but 14 (of which 7 are optically blank, and 7 are close to bright contaminating objects). We show that the optical properties of the radio sources do not change with flux density, suggesting that AGNs continue to contribute significantly at faint flux densities. We test this assertion by cross-correlating our radio catalogue with the X-ray source catalogue and conclude that radio-quiet AGNs become a significant population at flux densities below 300\\,$\\umu$Jy, and may dominate the population responsible for the flattening of the radio source counts if a significant fraction of them are Compton-thick. ",
        "introduction": "Surveys of extragalactic radio sources find two broad populations of object: active galactic nuclei (AGNs), and star-forming galaxies. The two processes which power these objects (accretion onto supermassive black holes, and star formation) are believed to be intimately linked by a mechanism or mechanisms collectively known as `AGN-driven feedback' (e.g., Croton et al.\\ 2006; Bower et al.\\ 2006) which produces the observed tight correlation between the masses of the central black hole and the stellar bulge (e.g., Ferrarese \\& Merritt 2000; Gebhardt et al.\\ 2000). Radio surveys therefore provide the opportunity to study the cosmic evolution of both these important processes.  Spectroscopically-complete samples of radio sources exist at equivalent 1.4-GHz flux densities of $\\sim100$\\,mJy (e.g., Willott et al.\\ 2002), and these samples are dominated by objects with powerful radio-emitting jets which carry a kinetic power comparable to their optical/X-ray photoionizing luminosity (e.g., Rawlings \\& Saunders 1991). These objects are commonly referred to as `radio-loud' AGNs, and have radio luminosities and morphologies which place them above the Fanaroff \\& Riley (1974) break ($L_{\\rm178MHz} \\sim 2 \\times 10^{25}\\rm\\,W\\,Hz^{-1}\\,sr^{-1}$) and they are classified as FR Class~II objects. At increasingly fainter flux densities, the observed source counts follow a power-law to $\\sim1$\\,mJy (e.g., Windhorst et al.\\ 1993) and the spectroscopic completeness of samples selected at these flux density limits decreases, in part because an increasing fraction of the objects have spectra which show little sign of optical nuclear activity (e.g., Waddington et al.\\ 2001), apparently because their supermassive black holes are accreting at a much lower rate. While these are still genuine radio-loud AGNs due to their ratio of optical and radio energy outputs, the radio luminosities place them below the Fanaroff--Riley break as FR\\,I sources.  The cosmic evolution of these two classes of radio-loud AGNs appears to differ.  While it has been known for many years that the population of the most powerful (FR\\,II) extragalactic radio sources undergoes very strong positive evolution (e.g., Dunlop \\& Peacock 1990), the less powerful (FR\\,I) sources evolve less strongly, possibly with a constant comoving space density (e.g., Clewley \\& Jarvis 2004). By modelling the evolution of the radio luminosity functions for these sources, it can be shown that their contribution to the total radio source counts should continue as a power law to even fainter flux densities (e.g., Jarvis \\& Rawlings 2004). However, below 1\\,mJy the observed counts flatten (when normalized to Euclidean), and this is attributed to the appearance of a new population of radio sources which do not contribute significantly at higher flux densities, but which start to dominate at $S_{1.4} \\la 300\\,\\umu$Jy.  This new population is usually believed to be star-forming galaxies (Condon 1984; Windhorst et al.\\ 1985), and the faint source counts have been successfully modelled on the basis of this interpretation (e.g., Seymour et al.\\ 2004). Recently, however, Jarvis \\& Rawlings (2004) have suggested that the source counts may flatten due to the `radio-quiet' AGN population, by which we mean supermassive black holes with a high accretion rate but a low ratio of radio to optical/X-ray luminosity. For quasars, where the non-stellar optical radiation is seen directly, the distinction between radio-loud and radio-quiet can also be made on the basis of the ratio of radio to optical fluxes (e.g., Kellermann et al.\\ 1989).  Since radio-quiet AGNs are known to be far more prevalent than their radio-loud counterparts, this population could indeed be important.  The fact that we still do not know the nature of the population(s) which dominate at flux densities $\\la 300\\,\\umu$Jy can be attributed to two factors. First, most of the effort to follow up radio source samples spectroscopically to high completeness has been devoted to samples selected at either bright ($\\ga 100$\\,mJy; e.g., Willott et al.\\ 2002) or very faint flux densities ($\\la 50\\,\\umu$Jy) from multi-wavelength studies of deep (and hence small-area) fields (e.g., Richards et al.\\ 1998). As a result, there has been very little study of the sources in the range $100\\,\\umu{\\rm Jy} \\la S_{1.4} \\la 500\\,\\umu\\rm Jy$, where the new population(s) begin to appear. Even when studies have focused on this flux density range (e.g., Hopkins et al.\\ 2003), the optical counterparts span such a large range in apparent magnitude that it is difficult to obtain IDs and subsequent photometry and spectroscopy of all members of a sample. Results have therefore been equivocal and even contradictory.  Windhorst et al.\\ (1985) identified such sources with faint blue galaxies whose colours suggested that they were undergoing significant star formation. Spectroscopy by Benn et al.\\ (1993) supported this interpretation, but later work by Gruppioni, Mignoli \\& Zamorani (1999) indicated that early-type galaxies were the dominant population. They attributed this difference to their deeper limit for optical identifications.  The key to solving this issue is therefore to increase the identification fraction of faint radio surveys as close to 100\\,per cent as possible by combining deep multi-colour imaging and moderate resolution (a few arcseconds, so as not to resolve out galactic-scale starbursts and hence bias a sample in favour of active galaxies) radio mapping over a degree-size field, so that there will be a statistically useful number of such objects. The Subaru/\\textit{XMM-Newton\\/} Deep Field (SXDF; Sekiguchi et al.\\ 2001, 2006) is ideal for this task, as it is home to extremely deep optical and near-infrared imaging (Furusawa et al.\\ 2006; Lawrence et al.\\ 2006). Deep (50--100\\,ksec) X-ray exposures with \\textit{XMM-Newton\\/} (Ueda et al.\\ 2006) have also been made in the SXDF, enabling the active galactic nucleus (AGN) content of the radio source population to be studied. We have therefore undertaken deep radio observations of this field to enable the $\\sim100\\,\\umu$Jy radio source population to be identified and studied with greater completeness than has previously been possible.  The format of this paper is as follows. In Section~2 we describe the radio observations and reduction methods, and in Section~3 we explain how we have determined and corrected for incompleteness in the catalogue. In Section~4 we describe how optical identifications were made, and report on our analysis of the faint radio source population. In Section~5 we summarize our main results. Radio maps of sources which are either significantly extended or have an uncertain optical identification are presented in an Appendix. All magnitudes quoted in this paper are calibrated on the AB scale. ",
        "conclusions": "We have presented a radio map covering 1.3 square degrees in the region of the Subaru/\\textit{XMM-Newton\\/} Deep Field, and described the production of a catalogue covering 0.81 square degrees containing all sources with a peak flux density greater than 100\\,$\\umu$Jy\\,beam$^{-1}$. We have made reliable identifications for 90\\,per cent of the sources, which is a significantly higher fraction than other surveys covering a similar area. We have demonstrated that there is no significant difference between the optical colours of radio sources selected at different flux densities, and have combined this with X-ray data to show that an important fraction of the faint ($100\\,\\umu{\\rm Jy} \\leq S_{1.4} < 300\\,\\umu\\rm Jy$) radio source population must be radio-quiet AGNs, including Type~2 objects which contribute to the Cosmic X-ray Background."
    },
    "0609/gr-qc0609044_arXiv.txt": {
        "abstract": "This paper is devoted to the computation of compact binaries composed of one black hole and one neutron star. The objects are assumed to be on exact circular orbits. Standard 3+1 decomposition of Einstein equations is performed and the conformal flatness approximation is used. The obtained system of elliptic equations is solved by means of multi-domain spectral methods. Results are compared with previous work both in the high mass ratio limit and for one neutron star with very low compactness parameter. The accuracy of the present code is shown to be greater than with previous codes. Moreover, for the first time, some sequences containing one neutron star of realistic compactness are presented and discussed. ",
        "introduction": "\\label{s:intro} Motivated by the various gravitational wave detectors coming online \\cite{ligo,virgo}, numerical simulations of binary compact objects have been extensively investigated in the last years. With progress made for the evolution of both binary neutron stars (BNS) \\cite{ShibaT06} and binary black holes (BBH) \\cite{Preto06,BakerCCKM06,CampaLMZ06}, it seems timely to turn to the last interesting type of binary: a system composed of one black hole and one neutron star (BHNS). The evolution synthesis codes indicate that the detection rate of BHNS with LIGO/Virgo will be as high (if not higher) than for BNS systems (see Table 6 of \\cite{BelczKB01}). Simulations of coalescing binary systems usually proceed in two steps. First, one needs to produce initial data that verify the Einstein constraint equations and that are as physically relevant as possible. Then, the initial configurations are evolved forward in time. Those steps involve rather different techniques and are equally challenging.  Most of the initial data for coalescing binaries rely on the quasiequilibrium hypothesis that assumes that the objects are on exact closed circular orbits. This is of course only an approximation because no closed orbits can exist for those systems in general relativity. However, this description should be rather accurate, at least for large separations. The quasiequilibrium approximation has been applied to BNS \\cite{BaumgCSST98a,BaumgCSST98b,UsuiUE00,UryuE00,GourgGTMB01} and BBH systems \\cite{PfeifTC00,GourgGB02,GrandGB02,CaudiCGP06}. In the BNS case, the evolution of such initial data have exhibited circular-like trajectories \\cite{ShibaT06}. The long term evolution of the data in the BBH case is still underway. The aim of this paper is to apply the same technique to a BHNS, without assuming extreme mass ratios like in \\cite{TanigBFS05}. On the contrary, a moderate mass ratio of 5 is used. This value is chosen mainly for comparisons with previous work \\cite{TanigBFS06} but is sufficient to demonstrate the ability of the code to handle binaries with close masses. Let us mention that the simulations of the formation of compact binaries indicate a slightly lower mass ratio to be more probable \\cite{BulikGB04}. The present work shares some properties with \\cite{TanigBFS06} even if some details are different. Moreover, one of the main shortcomings of \\cite{TanigBFS06} is the fact that only a NS of unrealistic compactness, as small as $0.0879$, is considered. This implies a mass close to $0.7 \\, M_\\odot$ for most of the available EOS, which is much smaller than the canonical value of $1.2-1.4 \\, M_\\odot$. On the contrary, in this work and for the first time, a BHNS with a realistic NS is computed. Indeed, the most compact star presented here has a compactness of $0.15$ which makes its mass in the range $1.2-1.5 \\, M_\\odot$ depending on the EOS. ",
        "conclusions": ""
    },
    "0609/astro-ph0609832_arXiv.txt": {
        "abstract": "\\suzaku observed the molecular cloud MBM12 and a blank field less than $3^{\\circ}$ away to separate the local and distant components of the diffuse soft X-ray background.  Towards MBM12, a local (D$<\\sim275$ pc) O{\\sc vii} emission line was clearly detected with an intensity of 3.5 ph cm$^{-2}$\\,s$^{-1}$\\,sr$^{-1}$\\,(or line units, LU), and the O{\\sc viii} flux was $<0.34$ LU.  The origin of this O{\\sc vii} emission could be hot gas in the Local Hot Bubble (LHB), charge exchange between oxygen ions in the solar wind (SWCX) and geocoronal or interplanetary material, or a combination of the two.  If entirely from the LHB, the emission could be explained by a region with 100 pc radius, an electron density of 0.0087 cm$^{-3}$, and a temperature of $1.2\\times10^6$\\,K.  However, the implied temperature and emission measure would predict 1/4 keV emission in excess of observations. There is no evidence in the X-ray light curve or solar wind data for a significant contribution from geocoronal SWCX.  However, the larger spatial extent of interplanetary SWCX washes out the rapid time variations of contributions from this source.  In any case, the observed O{\\sc vii} flux represents an upper limit to both the LHB emission and interplanetary SWCX in this direction, and both are thought to be nearly isotropic at low to intermediate latitudes.  The off-cloud observation was performed immediately following the on-cloud.  The net off-cloud O{\\sc vii} and O{\\sc viii} intensities were (respectively) $2.34\\pm0.33$\\ and $0.77\\pm0.16$\\,LU, after subtracting the on-cloud foreground emission.  Assuming the LHB and SWCX components did not change, these increases can be attributed to more distant Galactic disk, halo, or extragalactic emission.  If the distant O{\\sc vii} and O{\\sc viii} emission is from a thermal plasma in collisional equilibrium beyond the Galactic disk, a temperature of $(2.1\\pm0.1)\\times10^6$\\,K with an emission measure of $(4\\pm0.6)\\times10^{-3}$cm$^{-6}$pc is inferred. ",
        "introduction": "The soft ($< 2$\\,keV) diffuse X-ray background was a relatively early discovery of X-ray astronomy (see review by \\citet{TB77}).  Unlike the diffuse hard X-ray ($> 2$\\,keV) background, whose isotropy demonstrated it was dominated by extragalactic sources, the origin of soft component was and is more uncertain.  While at high Galactic latitudes extragalactic emission contributes to the observed flux, at lower latitudes the emission must be local to the Galaxy since at energies of 3/4 keV, absorption is significant (one optical depth is only N$_{\\rm H} = 2\\times10^{21}$\\,cm$^{-2}$).  Despite this, both the Wisconsin M band \\citep{McCammon83} and the ROSAT 3/4 keV \\citep{Snowden95} surveys showed surprisingly little latitude dependence away from the Galactic bulge.  It is now known that at high latitudes, where N$_{\\rm H}$\\ is generally less than $10^{21}$\\,cm$^{-2}$, $\\sim 40$\\% of the 3/4 keV emission is due to AGN, and from the XQC sounding rocket flight we know that at least 42\\% of the high-latitude flux must be due to oxygen emission lines coming from $z < 0.01$\\ \\citep{XQC02} although it is not known if these are within the Galaxy or in the halo.  The situation in the Galactic plane is more confusing.  Dwarf M stars must contribute some of the 3/4 keV emission \\citep{KS01}, but at least 50\\% of the emission is of unknown origin \\citep{MS90}.  O{\\sc vii}\\ and O{\\sc viii}\\ contribute most of the line emission in the 3/4 keV band, although the fraction in lines versus continuum remains uncertain.  Oxygen's dominance is due both to its large cosmic abundance (compared to other metals), and its strong emission lines at 0.57 keV (the O{\\sc vii}\\ triplet from $n=2\\rightarrow 1$) and 0.65 keV (the O{\\sc viii}\\ Ly$\\alpha$\\ transition).  Absorption limits the observed in-plane 3/4 keV emission to regions within 1-2 kpc (assuming $\\langle n \\rangle \\sim 1$\\,cm$^{-3}$\\ in the Galaxy).  Some of the 3/4 keV emission must be from the same source as the 1/4 keV X-ray background, which has been attributed to a combination of emission from a ``Local Hot Bubble'' (LHB) \\citep{Snowden90} and charge exchange from solar wind ions (hereafter SWCX) \\citep{Cox98, Lallement04}.  In both cases, the X-rays must be truly ``local'', originating within $\\sim 100$\\,pc in the former case or within the Solar System in the latter.  To distinguish between the ``local'' 3/4 keV X-ray emission and the more distant Galactic and halo components, we used \\suzaku to observe the nearby molecular cloud MBM12 (Lynds 1457), along with a nearby ``blank-sky''position not occulted by the cloud.  Earlier observations of MBM12 obtained with ROSAT \\citep{SnowdenMBM12}(SMV93) and \\chandra \\citep{Smith05} had low spectral resolution (ROSAT) or were strongly affected by high background due to solar flares (\\chandra).  Our hope with this observation was to use \\suzaku's low background, good spectral resolution, and sizable effective area $\\times$\\ solid angle product to measure the components of both the local and more distant contributors to the 3/4 keV emission.  As described in \\citet{Smith05}, the true distance to MBM12 is uncertain, with estimates ranging from $60\\pm30$\\,pc to $275\\pm65$\\,pc.  Our goal, however, is only to use MBM12 as a curtain that separates local components such as the LHB and SWCX from more distant components such as a hot halo or extragalactic emission.  It seems unlikely there is a significant component to the soft X-ray emission that is beyond the LHB but in front of MBM12, so the distance uncertainty is not particularly important to this analysis.  The total column density due to the MBM12 cloud is also somewhat uncertain.  We follow \\citet{Smith05}, who argued for N$_{\\rm H} = 4\\times10^{21}$\\,cm$^{-2}$\\ as a reasonable value for the densest region of MBM12.  Although the \\suzaku XIS1 has more than four times the field of view of the \\chandra ACIS-S3 ($17.8'\\times17.8'$\\ versus $8.5'\\times8.5'$), this is of similar size to the densest part of MBM12 and so we expect a similar total column density.  For the O{\\sc vii}\\ and O{\\sc viii}\\ lines most relevant to our analysis the optical depths for this column density are 3.5 and 2.4, respectively.  In addition, with a nearby off-cloud measurement, we will be able to directly determine the distant contribution and estimate its effect on the on-cloud emission.  ROSAT observations were only able to put a $2\\sigma$\\ upper limit of 270 counts s$^{-1}$\\,sr$^{-1}$\\ on the 3/4 keV emission seen towards MBM12 (SMV93)\\footnote{We present all surface brightnesses in units of steradians, and note that 1 sr = $1.18\\times10^{7}$\\,arcmin$^{2}$}. SMV93 fit a ``standard'' $10^6$\\,K collisional ionization equilibrium (CIE) LHB model \\citep{RS77} assuming a pathlength of $\\sim65$\\,pc, and found a good match to the observed 1/4 keV emission with an emission measure of $0.0024$\\,cm$^{-6}$\\,pc.  This model generates only $\\sim 47$ counts s$^{-1}$\\,sr$^{-1}$\\ in the 3/4 keV band, primarily due to $\\sim 0.28$\\,ph\\,cm$^{-2}$\\,s$^{-1}$\\,sr$^{-1}$ (hereafter LU, for ``line units'') generated by the O{\\sc vii}\\ triplet (based on the ATOMDB v1.3.1 atomic database).  However, the ROSAT PSPC had little spectral resolution in this band and could not separate the O{\\sc vii}\\ and O{\\sc viii}\\ lines from the background continuum, and possible Fe L line emission.  \\citet{Smith05} used the \\chandra ACIS instrument to redo the SMV93 observations with higher spectral and spatial resolution.  The results were affected by a large solar flare during part of the observation, which likely led to increased emission from SWCX (see \\citet{SCK04}\\ for more details on SWCX).  \\citet{Smith05} detected strong O{\\sc vii}\\ and O{\\sc viii}\\ emission lines with surface brightnesses of $1.92^{+0.61}_{-0.60}$\\ and $2.35^{+0.59}_{-0.43}$\\,LU respectively, much larger than the prediction from SMV93.  The O{\\sc viii}\\ emission itself was also unexpected, as \\citet{Smith05}\\ showed that it cannot come from any of the standard LHB models, either equilibrium or strongly recombining.  They suggested that the observed O{\\sc viii} emission was from the SWCX, although this could not be proven. ",
        "conclusions": "\\subsection{On-Cloud}  By observing both MBM12 and a nearby ``off-cloud'' field in quick succession with a low-background imaging X-ray spectrometer, we can reliably measure both the foreground and distant emission.  Our on-cloud result indicates a rather high value for the ``local'' diffuse O{\\sc vii} surface brightness of about 3.5 LU.  This emission is almost certainly generated in front of MBM12 since the cloud is optically thick, with transmission $<3\\%$ at O{\\sc vii}.  Assuming the O{\\sc vii} surface brightness behind MBM12 is the same as the value given in Table~\\ref{tab:mbm12_distant}, the background contribution to the on-cloud emission is $<0.06$\\,LU.  As first suggested by \\citet{Cox98}, the most likely sources for the foreground emission are either the LHB, SWCX, or both.  Although larger than expected for standard LHB models, the local O{\\sc vii} surface brightness is still lower than most measured values towards high-latitude sightlines.  \\citet{XQC02} found $4.8\\pm0.8$\\,LU of O{\\sc vii} towards a 1 sr region at high latitude, while the 2$\\sigma$\\ upper limit set toward MBM12 by \\citet{SnowdenMBM12} with ROSAT is equivalent to 7.1 LU at O{\\sc vii}.  ASCA observations towards high-latitude sightlines found a surface brightness of $2.3\\pm0.3$\\,LU in O{\\sc vii}, although this was model dependent \\citep{Gendreau95}.  \\citet{Smith05} did find a lower value for the local O{\\sc vii} surface brightness ($1.92^{+0.61}_{-0.60}$\\,LU, statistical errors only) when observing MBM12 with \\chandra.  However, \\chandra's particle background is both larger and more uncertain than \\suzaku's. In addition, a large solar flare occurred during the \\chandra observation, and \\citet{Smith05} noted that this increased the uncertainty.  \\citet{Smith05} also measured an O{\\sc viii} surface brightness of $2.35^{+0.59}_{-0.43}$\\,LU, which they attributed largely to SWCX particle background (along with an unknown portion of the O{\\sc vii} emission).  We find a marginal detection of $0.24\\pm0.1$\\,LU of O{\\sc viii} towards MBM12.  It seems likely that including the poorly-understood systematic errors on the \\chandra O{\\sc vii} result would bring it into agreement with the \\suzaku measurement, although we cannot say with precision how much of either is due to SWCX and not the LHB.  In contrast, the O{\\sc viii} seen with \\chandra is almost certainly due to SWCX \\citep{Smith05}.  It is difficult to generate the observed O{\\sc vii} surface brightness from a solar abundance LHB model consistent with ROSAT and other observations, as a model that generates the observed O{\\sc vii} emission will predict too much 1/4 keV band emission.  As described in \\citet{Smith05}, the dominant source of O{\\sc vii} flux in a thermal plasma is collisional excitation of O$^{+6}$\\,ions, unless the plasma is substantially out of equilibrium.  We are, however, reasonably certain the LHB is near equilibrium.  The strongly recombining model proposed by \\citet{BS94} was ruled out by \\citet{Smith05}.  They showed that the 2$\\sigma$\\ upper limits to the O{\\sc vii} and O{\\sc viii} surface brightnesses from the \\chandra MBM12 observation, combined with upper limits on the O{\\sc vi} surface brightness from FUSE \\citep{Shelton03} and the ROSAT All-sky surface 1/4 keV band brightness \\citep{Snowden95} excluded this type of model.  Conversely, \\citet{EC93} considered models of the LHB arising from a strongly ionizing plasma, created by a recent supernova.  They concluded that generating the observed soft X-ray emission required an unrealistically high ambient ISM pressure (p/k $>3\\times10^4$\\,cm$^{-3}$K).  In addition, these types of models typically predict significant amounts of high-velocity O{\\sc vi}, which have not been observed \\citep{Shelton02}.  We therefore consider models in collisional equalibrium.  A CIE plasma at $10^6$\\,K, the ``typical'' temperature for the LHB (SMV93), has an O{\\sc vii} triplet emissivity of $\\Lambda = 5.7\\times10^{-16}$\\,ph cm$^3$\\,s$^{-1}$, assuming the oxygen abundance relative to hydrogen is $8.51\\times10^{-4}$ \\citep{SmithAPEC}.  This rises rapidly with temperature, reaching $1.5\\times10^{-15}$\\,ph cm$^3$\\,s$^{-1}$ at $1.2\\times10^6$\\,K, and peaking (for T $=2\\times10^6$\\,K) at $6.4\\times10^{-15}$\\,ph cm$^3$\\,s$^{-1}$.  However, the $2\\sigma$\\ upper limit on the O{\\sc viii}/O{\\sc vii} ratio of 13\\% puts an upper limit on the equilibrium temperature of $1.7\\times10^6$\\,K (where $\\Lambda = 4.8\\times10^{-15}$\\,ph cm$^3$\\,s$^{-1}$).  Using these values, and assuming a constant density and temperature throughout the LHB, we can express the total surface brightness of O{\\sc vii} as: \\begin{equation} L_{S}(\\hbox{O{\\sc vii}}) = {{1}\\over{4\\pi}} R_{\\hbox{\\small LHB}} {{n_e^2}\\over{1.2}} \\Lambda_{\\hbox{\\small {O{\\sc vii}}}} \\label{eq:sbsimp} \\end{equation} where L$_S$\\ is in LU, $n_e$\\ is the electron density and R$_{\\hbox{\\small LHB}}$\\ is the bubble radius.  This assumes that hydrogen and helium are fully ionized, so $n_e \\approx 1.2 n_H$. Taking our lower value of 2.3 LU of O{\\sc vii}, Equation~(\\ref{eq:sbsimp}) gives $n_e^2 R_{\\hbox{\\small LHB}} = 0.020$\\,cm$^{-6}$pc at $10^6$\\,K, or 0.0075\\,cm$^{-6}$pc at $1.2\\times10^6$\\,K.  We can obtain a lower limit of 0.0023\\,cm$^{-6}$pc on this value using our upper limit of $T=1.7\\times10^6$\\,K.  This final value is similar to the emission measure found by SMV93 (0.0024\\,cm$^{-6}$pc), although at a significantly higher temperature.  Assuming R$_{\\hbox{\\small LHB}} = 100$\\,pc, we require electron densities of 0.014, 0.0087, or 0.0048 cm$^{-3}$, and a pressure of $p/k =3.0, 2.2,$\\ or $1.7\\times10^4$\\,cm$^{-3}$K at $T = 10^6$\\,K, $1.2\\times10^6$\\,K, or $1.7\\times10^6$\\,K, respectively.  Interestingly, \\citet{Cox05} found that the midplane pressure required to support the various layers of the Galaxy ({\\it e.g.}\\ cold and warm H{\\sc i}, diffuse H{\\sc ii}, etc) is $2.2\\times10^4$\\,cm$^{-3}$K, in agreement with our value at $T = 1.2\\times10^6$\\,K.  In addition, \\citet{Snowden00} used X-ray shadows (such as those created by MBM12) seen in the ROSAT All-sky Survey observations at 1/4 keV to measure the temperature of the local diffuse soft X-ray component.  They also found a best-fit temperature of $1.2\\times10^6$\\,K, although this is based on the \\citet{RS77}(RS77) plasma code.  In particular, using this temperature and the pressure with the 1993 update to the RS77 plasma code (using solar abundances) would predict $\\sim 3\\times$\\ the observed 1/4 keV band surface brightness seen by ROSAT.  This could perhaps be explained if the Si, Fe, and other high-Z elements that create the 1/4 keV band emission were depleted relative to oxygen in the LHB; more modelling is needed to test this hypothesis.  Despite the suggestive agreement in temperature and pressure described above, there are issues in other wavebands.  \\citet{Hurwitz05} has placed a 95\\% upper limit on the emission measure of a local $10^6$\\,K component of 0.0004 cm$^{-6}$pc, based on CHIPS observations of diffuse EUV iron lines and assuming a solar abundance for iron.  Even if iron is fully depleted, they still find a 95\\% upper limit of $\\sim 0.005$\\,cm$^{-6}$pc for any CIE model with T$<1.6\\times10^6$\\,K, based on the non-detection of O{\\sc v}\\ and O{\\sc vi}\\ lines near 171-173\\AA.  The fully-depleted \\citet{Hurwitz05} upper limit disagrees with our value of 0.0075 cm$^{-6}$\\,pc by a factor of at least 50\\%.  In addition, the solar-abundance CHIPS limit strongly disagrees with the value found by SMV93 and from the ROSAT All-Sky Survey in the 1/4 keV band ($\\sim 0.0018 - 0.0058$\\,cm$^{-6}$pc) which also assumes solar abundances \\citep{Snowden98}.  \\citet{Hurwitz05} noted these discrepancies and suggested that some depleted abundance pattern might exist that brings the X-ray and EUV observations into agreement. Nonetheless, as it stands the fully-depleted CHIPS limits suggest that at least a third of the O{\\sc vii} we detect is not from the LHB.  One possibility is that this emission comes from SWCX.  More analysis of the solar wind data will be needed to determine if the observations were truly done during a period of relative quiescence; for example, the absolute O$^{+7}$\\ and O$^{+8}$\\ fluxes can be derived from ACE data with additional effort.  While Figure~\\ref{fig:sw_params}, based on the automatically processed ACE data, does not show any signs of increased oxygen flux, more data are needed to confirm this.  \\subsection{``Distant'' emission\\label{ss:offcloud}}  Figure~\\ref{fig:sw_params} shows that the solar wind conditions were similar during both observations, and the LHB intensity is not expected to change over an angle of less than $3^\\circ$.  If both oxygen lines are from an unabsorbed plasma in CIE, the O{\\sc viii}/O{\\sc vii} ratio ($0.33\\pm0.08$) implies $T = (2.2^{+0.1}_{-0.2})\\times10^6$\\,K.  At this temperature, the predicted emission measure is $(1.9\\pm0.3)\\times10^{-3}$\\,cm$^{-6}$pc using ATOMDB v1.3.1 emissivities \\citep{SmithAPEC}.  If, as is more likely, the plasma is behind the Galactic hydrogen layer ($N_{\\rm H} = 8.7\\times10^{20}$\\,cm$^{-2}$), then the unabsorbed O{\\sc viii}/O{\\sc vii} ratio would be $0.26\\pm0.06$.  In this case, $T = 2.1\\pm0.1\\times10^6$\\,K and the emission measure is $(4.0\\pm0.6)\\times10^{-3}$\\,cm$^{-6}$pc.  In either case, our results consistent with previous measurements of distant hot halo gas. However, our result does not touch on the question of whether the halo has one \\citep{Pietz98} or two \\citep{KS00} dominant temperatures; further \\suzaku observations will be necessary to address this question.  The line at 0.876 keV is a mystery, although we stress it is at best a $ 3\\sigma$\\ detection.  Between 0.7-1.3 keV, the strongest emission lines in a collisional plasma are typically from neon or iron.  The closest strong neon line to 0.876 keV is the Ne{\\sc ix} forbidden line, but this would require a 5\\% gain error in the XIS1.  The oxygen lines at lower energies show $<2\\%$\\ gain shift, and the calibration lines at higher energies (see Table~\\ref{tab:bglines}) have less than 1\\% gain shift.  The strongest iron lines near this energy are from $2p^43d\\rightarrow2p^5$\\ transitions in Fe {\\sc xviii}, but any identification with an Fe line is problematic since many other lines of Fe{\\sc xviii}\\ (such as the $2p^43s\\rightarrow2p^5$\\ line at 0.775 keV) would also be expected.  In particular, the 2$\\sigma$\\ upper limit on the Fe{\\sc xvii} 0.826 keV line of 0.19 LU strongly limits any Fe line identification for the line at 0.876 keV.  It is possible it is an as-yet unidentified weak instrumental line, although this raises the question of why it is not present in the on-cloud data.  Intriguingly, the Lyman limit for O{\\sc viii}\\ is 0.8704 keV, so it is possible that this is not a line, but rather a recombination edge resulting from cool electrons interacting with O$^{+8}$\\ ions.  If so, we wonder at the origin of the O$^{+8}$\\ ions--are they local to the Solar system due to a sudden change in the solar wind during the off-cloud observation, or from a distant recombining plasma?  As more data from ACE and \\suzaku becomes available, we may be able to answer this question."
    },
    "0609/astro-ph0609373_arXiv.txt": {
        "abstract": "The Absolute Radiometer for Cosmology, Astrophysics, and Diffuse Emission (ARCADE) is a balloon-borne instrument designed to measure the temperature of the cosmic microwave background at centimeter wavelengths. ARCADE searches for deviations from a blackbody spectrum resulting from energy releases in the early universe. Long-wavelength distortions in the CMB spectrum are expected in all viable cosmological models. Detecting these distortions or showing that they do not exist is an important step for understanding the early universe. We describe the ARCADE instrument design, current status, and future plans. ",
        "introduction": "The cosmic microwave background has proven to be an invaluable probe of physical conditions in the early universe. Temperature anisotropy in the CMB records the density and gravitational potential at an epoch when the fluctuations leading to the present-day large scale structure were still in the linear regime. The CMB spectrum, by contrast, provides information on the {\\em energetics}, including the critical period after recombination when the currently-observed large-scale structures first formed. Precise measurements of the CMB spectrum provide important information on the early universe.  The frequency spectrum of the cosmic microwave background (CMB) carries a history of energy transfer between the evolving matter and radiation fields in the early universe. Energetic events (particle decay, star formation) heat the diffuse matter which then cools via interactions with the background radiation, distorting the radiation spectrum away from a blackbody. The amplitude and shape of the resulting distortion depend on the magnitude and redshift of the energy transfer \\cite{zeldovich/sunyaev:1969, burigana/etal:1991, burigana/etal:1995}.  \\begin{figure}[t] \\centerline{ \\psfig{file=arcade_irvine_fig_1.ps,width=2.8in,angle=90}} \\caption{Precise measurements of the CMB thermodynamic temperature. The dotted line represents a blackbody spectrum. The gray bands show the ARCADE frequencies. The CMB spectrum is poorly constrained at centimeter or longer wavelengths.} \\label{spectrum_fig} \\end{figure}  Measurements across the peak of the CMB spectrum limit deviations from a blackbody to less than 50 parts per million \\cite{mather/etal:1990, gush/etal:1990, fixsen/mather:2002}. Direct observational limits at longer wavelengths, however, are weak: distortions as large as 5\\% could exist at wavelengths of several centimeters or longer without violating existing observations (Figure \\ref{spectrum_fig}). Plausible physical processes are expected to generate observable distortions without violating limits established at shorter wavelengths. Free-free emission from reionization and subsequent structure formation adds photons to the diffuse background, creating a spectral distortion $ \\Delta T_{\\rm ff} = T_{\\gamma} Y_{\\rm ff} / {x^2}$, where $T_{\\gamma}$ is the undistorted photon temperature, $x$ is the dimensionless frequency $h \\nu / k T_{\\gamma}$, \\begin{equation} Y_{\\rm ff} = \\int^z_0 ~\\frac{ k[T_e(z) - T_{\\gamma}(z)] }{ T_e(z) } \\frac{ 8 \\pi e^6 h^2 n_e^2 g } { 3 m_e (kT_{\\gamma})^3 \\sqrt{6\\pi m_e k T_e} } \\frac{dt}{dz^\\prime} dz^\\prime, \\label{Yff_definition} \\end{equation} is the optical depth to free-free emission, and g is the Gaunt factor \\cite{bartlett/stebbins:1991}. The distorted CMB spectrum is characterized by a quadratic rise in temperature at long wavelengths. The amplitude of the free-free signal depends on the column $\\int n_e^2$ of ionized gas and thus on the redshift $z_r$ at which the first collapsed objects formed. Such a distortion {\\it must} exist, with predicted amplitude of a few mK at frequency 3 GHz \\cite{haiman/loeb:1997, gnedin/ostriker:1997, oh:1999}. Detection of the free-free distortion would place important constraints on the era of luminous object formation and the extent of clumping in galactic halos.  The decay of massive particles or other relics produced near the Big Bang will also distort the CMB spectrum. Energy released at an early epoch to either charged particles or photons will heat free electrons, which then cool via Compton scattering from the colder CMB photons. For energy released at redshift $z < 10^4$ the gas is optically thin, resulting in a uniform decrement $ \\Delta T_{\\rm RJ} = T_{\\gamma} (1 - 2y) $ in the Rayleigh-Jeans part of the spectrum where there are too few photons, and an exponential rise in temperature in the Wien region with too many photons. The magnitude of the distortion is related to the total energy transfer $\\Delta {\\rm E} / {\\rm E} = 4y$ where \\begin{equation} y = \\int^z_0 ~\\frac{ k[T_e(z) - T_{\\gamma}(z)] }{ m_e c^2} \\sigma_T n_e(z) c \\frac{dt}{dz^\\prime} dz^\\prime, \\label{compton_y_definition} \\end{equation} is the dimensionless integral of the electron pressure $n_e k T_e$ along the line of sight, $m_e$, $n_e$ and $T_e$ are the electron mass, spatial density, and temperature, $T_{\\gamma}$ is the photon temperature, $k$ is Boltzmann's constant, $z$ is redshift, and $\\sigma_T$ denotes the Thomson cross section \\cite{sunyaev/zeldovich:1970}. Energy transfer at higher redshift $10^4 < z < 10^7$ approaches the equilibrium Bose-Einstein distribution, characterized by the dimensionless chemical potential $ \\mu_0 = 1.4 \\Delta {\\rm E} / {\\rm E}. $ Free-free emission thermalizes the spectrum at long wavelengths. Including this effect, the chemical potential becomes frequency-dependent, \\begin{equation} \\mu(x) = \\mu_0 \\exp(- \\frac{2x_b}{x}), \\label{mu_vs_freq_eq} \\end{equation} where $x_b$ is the transition frequency at which Compton scattering of photons to higher frequencies is balanced by free-free creation of new photons. The resulting spectrum has a sharp drop in brightness temperature at centimeter wavelengths \\cite{burigana/etal:1991}.   \\begin{figure}[b] \\centerline{ \\psfig{file=arcade_irvine_fig_2.ps,width=2.6in,angle=90}} \\caption{ Current 95\\% confidence upper limits to distorted CMB spectra, along with the predicted signal from reionization and structure formation. The FIRAS data and ARCADE 1 mK anticipated error box are also shown. ARCADE has both the sensitivity and control of systematic errors to measure the expected signal from the high-redshift universe.} \\label{distortion_fig} \\end{figure}  A chemical potential distortion is a primary signature for the decay of relics from GUT and Planck-era physics. Such relics are expected to exist. Growing evidence points to the existence of ``dark matter'' observed only through gravitational effects, with density $\\Omega_m = 0.22^{+0.01}_{-0.02}$ compared to the critical density \\cite{spergel/etal:2006}. A leading candidate for dark matter composed of weakly interacting massive particles is the lightest supersymmetric particle $\\chi$, the neutralino \\cite{jungman/etal:1996}. Neutralinos annihilate primarily into quark-antiquark pairs, which in turn cascade to photons, electrons, positrons, and neutrinos. The annihilation rate $ \\Gamma = n_\\chi n_{\\bar \\chi} \\langle \\sigma v \\rangle, $ where $\\langle \\sigma v \\rangle ~\\sim 6 \\times 10^{-26}$ cm$^3$ s$^{-1}$ at freeze out and the neutralino space density $n_\\chi$ varies as $z^3$ from the overall expansion of the universe. Gamma radiation from neutralino annihilation in the Galactic dark matter halo has been proposed to explain the excess flux observed by the EGRET mission \\cite{fornengo/etal:2004,valle:2004}. A chemical potential distortion to the CMB spectrum is an inevitable result of neutralino annihilation: annihilations sufficient to produce the EGRET excess would necessarily produce a chemical potential distortion. Measurements of the CMB spectrum can thus provide important constraints on the mass, cross section, and decay modes of dark matter candidates.  Figure \\ref{distortion_fig} shows current upper limits to spectral distortions at long wavelengths. The CMB spectrum is poorly constrained at centimeter or longer wavelengths, where observable signals from reionization or relic decay are expected to exist. New measurements with mK accuracy can provide important checks for both cosmology and high-energy physics. The Absolute Radiometer for Cosmology, Astrophysics, and Diffuse Emission (ARCADE) is a balloon-borne instrument to provide precisely such new data. ",
        "conclusions": ""
    },
    "0609/astro-ph0609659_arXiv.txt": {
        "abstract": "IZw18 has been recurrently claimed to be a young galaxy, but stars of increasingly older ages are found every time deeper magnitude levels are reached with high-resolution photometry: from the original few Myrs to, possibly, several Gyrs. We summarize the history of IZw18's age and an HST project which will allow us to derive both its distance and age. ",
        "introduction": "The Blue Compact Dwarf galaxy IZw18 is the most metal poor star-forming galaxy ever discovered, with a metallicity between 1/30 and 1/50 of solar, depending on the adopted solar value (e.g., \\cite[Searle \\& Sargent 1972]{}, \\cite[Izotov \\& Thuan 2004]{}). IZw18 is not only extremely metal poor, but also very blue and full of gas, three properties that originally suggested it to be forming now its first stars and make it the system in the local Universe most similar to primeval galaxies. Soon after its discovery, the basic question became whether IZw18 is so metal poor because a) it has started so recently to form stars that they haven't had time to significantly pollute the medium, or b) because its star formation (SF) activity, although occurring over a long period of time, has proceeded at a rate too low for an efficient chemical enrichment, or c) because strong galactic winds removing from the system most of the metals have accompanied a {\\it normal} SF activity. More than thirty years later, its age and distance are still subjects of hot debates.  The safest way to estimate the age of nearby galaxies is from the Colour-Magnitude Diagram (CMD) of their resolved stellar populations, but for IZw18, at 10--20 Mpc, this is feasible only with HST and yet challenging. Our novel starts with the advent of HST. ",
        "conclusions": ""
    },
    "0609/astro-ph0609145_arXiv.txt": {
        "abstract": "We present Giant Metrewave Radio Telescope \\ion{H}{i} observations of an eXtremely Metal Deficient galaxy SBS~1129+576. SBS~1129+576 has a weighted mean oxygen abundance of 12~+~$\\log$~(O/H)~=~7.41~$\\pm$~0.07, or 1/18 of the solar value. Our \\ion{H}{i} observations show that the galaxy is strongly interacting with a companion (projected separation $\\sim$27~kpc) galaxy, SBS~1129+577. \\ion{H}{i} emission from a third, smaller galaxy, SDSS~J113227.68+572142.3, is also present in the data cube. We study the \\ion{H}{i} morphology and kinematics of this small group at angular resolutions ranging from $\\sim$40 to 8~arcsec. The low-resolution map shows a bridge of emission connecting the two larger galaxies and a large one-armed spiral distortion of the disc of SBS~1129+577. We measure \\ion{H}{i} masses of $\\sim$4.2~$\\times$~10$^{8}$, $\\sim$2.7$~\\times$~10$^{9}$, and $\\sim$2.1~$\\times$~10$^{8}$~M$_{\\odot}$ for SBS~1129+576, SBS~1129+577 and the gas in the bridge, respectively. Assuming that most of the bridge gas originally came from SBS~1129+576, approximately one-third of its original gas mass has been stripped off. The third smaller galaxy has an \\ion{H}{i} mass of (M$_{\\rm HI}$~$\\sim$1.1$\\times$10$^{7}$ M$_{\\odot}$) and does not show any sign of interaction with the other two galaxies. The higher-resolution maps show that SBS~1129+577 has a central bar and a ring surrounding the bar; there is also a hint of an integral-shaped warp in SBS~1129+576.  All these features are very likely to have been induced by the tidal interaction. In both SBS~1129+576 and SBS~1129+577, there is, in general, a good correspondence between regions with high \\ion{H}{i} column density and those with ongoing star formation. The two brightest \\ion{H}{ii} regions in SBS~1129+576 have (inclination-corrected) gas column densities of $\\sim$1.6$\\times$10$^{21}$ and $\\sim$1.8$\\times$10$^{21}$ atoms~cm$^{-2}$, respectively. The inclination-corrected \\ion{H}{i} column density near the \\ion{H}{ii} regions in SBS~1129+577 is generally above $\\sim$2.0$\\times$10$^{21}$ atoms~cm$^{-2}$. These values are close to the threshold density for star formation observed in other blue compact galaxies. In contrast to SBS~1129+576 and SBS~1129+577 which are very gas-rich, the third member of this group, SDSS~J113227.68+572142.3, is gas-poor. ",
        "introduction": "\\label{sec:intro}  A very small fraction of low-mass, gas-rich galaxies have very low observed interstellar medium (ISM) metallicities, 12~+~$\\log$~(O/H)~$\\leq$~7.65. About 50 such galaxies are now known (see e.g. the review by \\cite{KO2000}) and are termed here as eXtremely Metal Deficient (XMD) galaxies. A few XMD galaxies, for example, I~Zw~18 and SBS~ 0335$-$052 E and W, where there is no evidence for an old stellar population, are considered as good candidates for young galaxies in the local Universe. In general, however, old stellar populations have been found in most (of the as of yet small numbers) of well-studied XMD galaxies. Despite the presence of an older stellar population, the very low ISM metallicities of XMDs indicate that they are relatively unevolved systems. Further evidence for their relatively unevolved status comes from the fact that they are very gas rich; neutral hydrogen is often the dominant baryonic component in these galaxies. The study of the neutral hydrogen content, distribution and kinematics is hence important in understanding this small, but interesting subclass of dwarf galaxies. As the best local analogues of young low-mass galaxies in the high-redshift Universe, detailed studies of the connections between the gas distribution, kinematics and star formation in these galaxies could provide useful input into models of galaxy formation and evolution. A systematic, \\ion{H}{i} survey of 22 XMD galaxies was conducted at the Nan\\c {c}ay Radio Telescope (NRT) by \\cite{MP05}. This survey increased the number of XMD galaxies with available integrated \\ion{H}{i} parameters by almost a factor of 2 (earlier surveys were done by e.g. \\cite{thuan99} 1999). Follow-up interferometric observations of a subsample of the galaxies observed at the NRT are being done using the Giant Metrewave Radio Telescope (GMRT). SBS~1129+576, the subject of this paper, is the first object from this subsample.  SBS~1129+576 is a very metal-poor (with weighted mean oxygen abundance 12~+~$\\log$~(O/H)~=~7.41~$\\pm$~0.07; \\cite{guseva03} 2003) nearby blue compact galaxy (BCG). It was discovered as part of the Second Byurakan Survey (SBS) (\\cite{markarian83}; \\cite{lipovetsky88}). \\cite{Pustilnik01} measure $M_{\\rm B} \\sim$16.50 mag for this galaxy. \\cite{guseva03} (2003) present $V$- and $I$-band CCD images as well as spectroscopic data for this galaxy. The CCD images show the galaxy to contain a chain of compact \\ion{H}{ii} regions within an elongated low surface brightness (LSB) envelope. The \\ion{H}{ii} regions and the LSB component both are bluer than typical for dwarf irregular or even blue compact dwarf (BCD) galaxies. The blue colour and the very low metallicity of SBS~1129+576 both are indicative of a relatively young age for the dominant stellar population. However, as is generally the case with very low metallicity galaxies, the spectroscopic and photometric data do not allow one to rule out models in which the galaxy contains stellar populations as old as 10~Gyr, albeit with the majority of the current luminosity coming from a much younger ($\\sim$100~Myr) population.  \\cite{thuan99} (1999) presented NRT \\ion{H}{i} data for SBS~1129+576. The measured heliocentric velocity is 1566~$\\pm$~2 \\kms, which corresponds to a distance of 26.3 Mpc (after correction for Virgocentric infall). \\cite{thuan99} (1999) noted the possibility of confusion with a nearby dwarf irregular galaxy, SBS~1129+577 (MCG+10-17-010). \\cite{pustilnik01} (2001a) suggested that interaction with this companion probably triggered starburst in SBS~1129+576, whereas \\cite{petrosian02} noted that SBS~1129+576 has a peculiar disturbed morphology and classified it as a 'merger'.  We present here GMRT \\ion{H}{i} 21-cm observations of emission from the SBS~1129+576/577 pair. The rest of this paper is divided as follows. In Section~\\ref{sec:obs}, the observations and data reduction are described, while the results of the observations are described in Section~\\ref{sec:res}. In Section~\\ref{sec:dis}, we compare SBS~1129+576 with the other very metal-poor galaxies, and Section~\\ref{sec:summary} provides a summary of our results. ",
        "conclusions": "\\label{sec:dis}  \\begin{figure} \\includegraphics[width=7.25cm,angle=270]{vfield57730k.eps} \\caption{ The velocity field of SBS~1129+577 at a resolution of 8~$\\times$~7~arcsec$^{-2}$ overlaid on the integrated \\ion{H}{i} emission (grey-scale over range of \\ion{H}{i} column density from 0 to 6.6~$\\times$~10$^{21}$~atoms~cm$^{-2}$). The velocity contours are regularly spaced at an interval of 6.6~\\kms\\ and range from 1528~\\kms\\ to 1607~\\kms.} \\label{fig:highresmom1} \\end{figure}  In Table~\\ref{tab:main}, we summarise the main observational and derived parameters of SBS~1129+576 and its two companion galaxies. These two companion galaxies are situated $\\sim$27~kpc north (SBS~1129+577) and south-east (SDSS~J113227.68+572142.3) of the XMD galaxy, respectively. The total projected size of this compact triplet is $\\sim$43~kpc and the range of their radial velocities is $\\sim$70~\\kms. The oxygen abundance in SBS~1129+577 is 12~+~$\\log$~(O/H)~=~7.94~$\\pm$~0.08 (based on Special Astrophysical Observatory (SAO) 6-m telescope observations; \\cite{Pustilnik06} 2006). This value is typical of BCGs with similar values of luminosity and suggests that it is an old galaxy. In contrast, SBS~1129+576 has 12~+~$\\log$~(O/H)~=~7.41~$\\pm$~0.07 (\\cite{guseva03} 2003); there are only about a dozen or so known BCGs with metallicites comparable or lower than this. The third galaxy in the group, SDSS~J113227.68+572142.3, has $g$ and $r$ magnitudes (as given in Sloan Digital Sky Survey (SDSS) Data Release 3 (DR3) data base; \\cite{DR3}) of 16.95 and 16.81~mag, respectively. Using the formulae in \\cite{smith02}, this translates to $B$~=~17.19~mag, M$_{\\rm B}$~=~--14.8~mag. Similarly, translating the SDSS $g$ and $r$ magnitudes for SBS~1129+576 and SBS~1129+577 to $B$ band, yields $B_{\\rm tot}$ of 16.44 and 15.50~mag, respectively. Combining these with the \\ion{H}{i} masses derived in this paper gives M$_{\\rm HI}$/L$_{\\rm B}$ of 2.1 and 3.8 (in solar units) for SBS~1129+576 and SBS~1129+577, respectively. Note that, in the value of M$_{\\rm HI}$/L$_{\\rm B}$ for SBS~1129+576 we included the mass of gas in the bridge, assuming that it consists mainly of the gas pulled out of the less-massive galaxy. These values of M$_{\\rm HI}$/L$_{\\rm B}$ are among the largest known for starburst galaxies and confirm that both these galaxies are gas-rich. In contrast, SDSS J113227.68+572142.3 is rather gas-poor, with M$_{\\rm HI}$/L$_{\\rm B}$ = 0.08. The origin and evolution of such a system of quite different types of dwarf galaxies is probably worth a separate deeper study.  \\subsection{Star formation in SBS~1129+576, SBS~1129+577 and SDSS~J113227.68+572142.3}  As mentioned in Section~\\ref{sec:intro}, \\cite{guseva03} (2003) presented detailed models of the star formation history of SBS~1129+576. Most of the current star formation is concentrated in two bright knots of emission (called regions 'a' and 'b' in \\cite{guseva03} 2003). These two bright knots can also be seen in the broad-band optical image in Fig.~\\ref{fig:highres}, where it can also be seen that they coincide with peaks in the \\ion{H}{i} column density distribution. From the measured metallicity, the H$\\alpha$ equivalent widths (EWs), and also assuming a standard Salpeter initial mass function (IMF), \\cite{guseva03} (2003) estimate that these star-forming regions are very young, with ages $\\lesssim$10~Myr. Like in many other BCD galaxies, the current star formation activity in SBS~1129+576 appears to be occurring only above a critical \\ion{H}{i} threshold density. The observed peak \\ion{H}{i} column densities corresponding to \\ion{H}{ii} regions 'a' and 'b' is $\\sim$5.9~$\\times$~10$^{21}$ and $\\sim$6.7~$\\times$~10$^{21}$~atoms~cm$^{-2}$, respectively, at an angular resolution of 8~arcsec which corresponds to a linear resolution of 1.02~kpc at the assumed distance of 26.3~Mpc. These observed peak \\ion{H}{i} column densities correspond to deprojected values of the total gas density (i.e. including a 10 per cent contribution from He, we note that this is a slight overestimate of the He abundance in metal-poor galaxies) of $\\sim$1.6~$\\times$~10$^{21}$ and $\\sim$1.8~$\\times$~10$^{21}$~atoms~cm$^{-2}$, similar to the values seen for other BCDs and dwarf irregular galaxies (e.g. \\cite{skillman87}; \\cite{taylor94}).  \\begin{table*} \\centering \\caption{\\label{tab:main} Main observational and derived parameters of the studied dwarf galaxies} \\begin{tabular}{lrrr} \\\\ \\hline Parameter               & SBS~1129+576     & SBS~1129+577     & SDSS~J113227+572142  \\\\ \\hline $\\alpha\\,$(J2000)       & 11$^{h}$32$^{m}$02$^{s}$.54 & 11$^{h}$32$^{m}$04$^{s}$.00  & 11$^{h}$32$^{m}$27$^{s}$.68\\\\ $\\delta\\,$(J2000) & +57\\degr22\\arcmin45\\arcsec.7 & +57\\degr26\\arcmin20\\arcsec.0 & +57\\degr21\\arcmin42\\arcsec.3 \\\\ $\\int f(v)dv$ (Jy~\\kms) & 3.9$^{a}$              & 16.7             & 0.07           \\\\ $V_{hel}^{\\,b}$ (\\kms)  & 1506~$\\pm$~3.3     &  1575~$\\pm$~3.3    & 1540:    \\\\ W$_{\\rm 50}$ (\\kms)     &   67~$\\pm$~3.3     &     85~$\\pm$~3.3   &  18:           \\\\ m$_{B}^{\\;\\;c}$         &   16.44          &     15.50        & 17.19          \\\\ M$_{B}^{\\;\\;d}$         &   $-$15.65       &     $-$16.59     & $-$14.8        \\\\ $M_{\\rm HI}$ (10$^8$M$_{\\odot}$)&  6.3$^a$        &     27.0         &  0.11          \\\\ $M_{\\rm HI}$/L$_{\\rm B}$$^e$   &  2.1$^a$         &     3.8          &  0.08          \\\\ 12~+~$\\log$~(O/H)          & 7.41~$\\pm$~0.07    & 7.94~$\\pm$~0.08    & --             \\\\ \\hline \\\\[-0.35cm] \\multicolumn{4}{l}{$^a$ \\ion{H}{i} flux/mass of the bridge is included; $^b$ Heliocentric \\ion{H}{i} velocity;} \\\\ \\multicolumn{4}{l}{$^c$ Translated from SDSS $g$ and $r$ magnitudes;} \\\\ \\multicolumn{4}{l}{$^d$ Absolute blue magnitudes corrected for the MW extinction and D=26.3 Mpc;} \\\\ \\multicolumn{4}{l}{$^e$ In solar units.} \\\\ \\end{tabular} \\end{table*}   The distribution of star-forming regions in SBS~1129+577 is more complicated. A long-slit spectroscopy of this galaxy was done by \\cite{Pustilnik06} (2006), using the SAO 6-m telescope. The slit was positioned at a position angle of 90\\degr\\ and passing through the brightest \\ion{H}{ii} region A1 (J113206.5+572620; $\\sim$20~arcsec to the east of the centre of the main body seen in continuum). The peak \\ion{H}{i} column density at this position is 2.8~$\\times$~10$^{21}$ atoms~cm$^{-2}$ at an angular resolution of 8~arcsec. Besides this region, the star formation is also seen to take place in three more regions along the slit, namely A2, A3 and A4. A2 is $\\sim$20~arcsec to the east of A1, while A3 is $\\sim$35~arcsec to the west of A1. The \\ion{H}{i} column densities are 3.1~$\\times$~10$^{21}$ and 2.4~$\\times$~10$^{21}$~atoms~cm$^{-2}$ at the positions of A2 and A3, respectively. A4, which is rather faint \\ion{H}{ii} region near the centre of the galaxy, has an \\ion{H}{i} peak with column density of 3.2~$\\times$~10$^{21}$~atoms~cm$^{-2}$. We note that there are other brighter knots visible in the continuum optical image which also coincide with peaks in \\ion{H}{i}. For example, there is a bright optical condensation (J11329.9+572630) in the ring which coincides with the highest peak (4.1~$\\times$~10$^{21}$~atoms~cm$^{-2}$) in the \\ion{H}{i} column density map. These too may be star-forming regions, but the confirmation by H$\\alpha$-imaging or spectroscopy is required. For the \\ion{H}{ii} regions identified above via the long-slit spectroscopy, based on their measured EW(H$\\beta$), we estimate ages of 4--8~Myr for the case of an instantaneous starburst model with the standard Salpeter IMF (\\cite{Leitherer} 1999). The inclination-corrected total gas column densities for regions A1, A2, A3, and A4 are 2.3~$\\times$~10$^{21}$, 2.6~$\\times$~10$^{21}$, 2.0~$\\times$~10$^{21}$, and 2.6~$\\times$~10$^{21}$, respectively, all of which are similar to the canonical threshold density.  As mentioned above, SDSS~J113227.68+572142.3 is gas-poor, with $M_{\\rm HI}$/L$_{\\rm B}$~=~0.08. It does not appear to be strongly interacting with the other two galaxies and, in particular, there is no evidence for the exchange of matter between it and the two brighter galaxies of the group. The SAO 6-m telescope spectrum (unpublished) for this object as well as its SDSS DR3 spectrum show an H$\\alpha$ emission line with EW~=~28~\\AA, a faint [\\ion{O}{iii}]$\\lambda$5007 line and Balmer absorption lines in blue. According to the Starburst-99 models of \\cite{Leitherer} (1999), for an instantaneous starburst with the metallicity $z$~=~0.0004 and a standard Salpeter IMF, all of these features are consistent with a star formation episode with an age of $\\sim$20~Myr. This age is larger but comparable with that of the starbursts in SBS~1129+576 and SBS~1129+577. For SDSS~J113227.68+572142.3, the peak \\ion{H}{i} column density (at a resolution of $\\sim$8~arcsec or 1~kpc) is $\\sim$3.9~$\\times$~10$^{20}$~atoms~cm$^{-2}$, corresponding to a deprojected total gas column density of $\\sim$4.0~$\\times$~10$^{20}$~atoms~cm$^{-2}$. This is a factor of $\\sim$2 below the nominal star formation threshold. We note that \\cite{begum} also found that star formation in very faint dwarf galaxies does occur at gas column densitites approximately two times lower than the nominal threshold density. The size of the galaxy is comparable to our beam size, even at the highest resolution (i.e. $\\sim$8~arcsec). It is thus very likely that our observed peak column density is a severe underestimate of the true column density. It is possible that the star formation in SDSS J113227.68+572142.3 was triggered by the accretion of a small amount of gas onto a gas-poor (dwarf elliptical ?) progenitor. Deeper \\ion{H}{i} observations may help to confirm or rule out this hypothesis.  \\subsection{Comparison with other XMD galaxies}  The SBS~1129+576/577 system is morphologically very similar to \\ion{H}{i}~1225+01 (\\cite{salzer92} 1992, \\cite{chengalur95}). Both systems consist of two components with a bridge between the two. In both systems, the southern component is almost edge-on, whereas the northern component is more face-on. The \\ion{H}{i} finger at the south-western edge of SBS~1129+576 is a countertail similar to that in the south-western component of \\ion{H}{i}~1225+01. The north-east clump of \\ion{H}{i}~1225+01 has a bar as SBS~1129+577 does have. Like in \\ion{H}{i}~1225+01, we assume that the \\ion{H}{i} in the bridge originates from the smaller galaxy in the pair. SBS~1129+576, like the south-western component in \\ion{H}{i}~1225+01, is on a prograde orbit. The \\ion{H}{i} masses of the less-massive components in both systems are similar (where we assume a distance of 15~Mpc to \\ion{H}{i}~1225+01). The ratio of the masses of the two components is, however, different, with a value of $\\sim$1:1.9 for \\ion{H}{i}~1225+01 and $\\sim$1:4.3 for SBS~1129+576/577. In computing these ratios, we assume that \\ion{H}{i} in the bridges originally belonged to less-massive components. The major difference between the systems SBS~1129+576/577 and \\ion{H}{i}~1225+01 is that no optical counterpart has been found in south-west clump of \\ion{H}{i}~1225+01 to a limiting surface  brightness of 27.0 $V$-mag~arcsec$^{-2}$ (\\cite{salzer91}).  It is interesting to compare the  environments of these systems. \\cite{salzer92} (1992) found that there is no dwarf galaxy within $\\sim$1~Mpc of \\ion{H}{i}~1225+01, while the nearest luminous galaxy is at $\\sim$1.2~Mpc from it (for an adopted distance of 15~Mpc). We re-examined the environment of \\ion{H}{i}~1225+01, using the data available in NED. We found that only one dwarf VCC~1208 (which is 124~arcmin north of \\ion{H}{i}~1225+01 and has $\\delta V$~=~7~\\kms) lies within the $\\sim$1~Mpc limit given by \\cite{salzer92} (1992). This galaxy lies at a projected distance of 540~kpc (for  D~=~15~Mpc). The region around SBS~1129+576 was surveyed as a part of SDSS. Since there are no galaxies with 1300~\\kms\\ $<$ V$_{\\rm hel}$ $<$ 1700~\\kms\\ within 30~arcmin (corresponding to the projected distance of 230~kpc), it seems to be an isolated group of dwarfs. There are several galaxies at the projected distances of 150--500 kpc with the velocities in the range $\\sim$1700--1800~\\kms\\ (including a rather luminous SB(s)c galaxy NGC~3683, with $M_{\\rm B}$ of $-$19.6~mag), but their radial velocities imply that all of them are more than $\\sim$2~Mpc away from the SBS~1129+576/577 system. There are four dwarf galaxies with radial velocities from 1450 to 1565~\\kms\\ at the projected distances of $\\sim$350--400 kpc. The nearest luminous galaxy with close radial velocity is NGC~3619 ($M_{\\rm B} = -$19.5~mag) at $\\sim$0.8~Mpc. We conclude that both systems are at a distance of $\\sim$1~Mpc from the nearest luminous galaxies and at several hundred kpc from the nearest dwarf galaxy.  \\par SBS~0335-052 (\\cite{pustilnik01} 2001a) is another interesting object to compare to the SBS~1129+576/577 system. Both the components, SBS~0335$-$052 E and W, at a projected distance of 22~kpc (\\cite{pustilnik01} 2001a) are candidate young galaxies (e.g. \\cite{izotov97}; \\cite{papaderos98}; Pustilnik, Pramskij \\& Kniazev 2004; \\cite{izotov05}). The \\ion{H}{i} envelope surrounding these components is similar to the one seen in SBS~1129+576/577 system. Its total gas mass is comparable to that of SBS~1129+577 (i.e. $\\sim$10$^{9}$~M$_{\\odot}$). The morphology and kinematics of SBS~0335-052 in \\ion{H}{i} and at other wavelengths suggest that the most plausible interpretation of this system is a merger. The progenitors are either extremely unevolved very gas-rich objects or protogalaxies. Their mutual tidal interaction has triggered the current bursts of star formation. The peak measured \\ion{H}{i} column densities are 1.0~$\\times$~10$^{21}$ and 0.7~$\\times$~10$^{21}$~atoms~cm$^{-2}$ in the west and east components, respectively, at a linear resolution of 5.4~$\\times$~3.9~kpc. This system, however,  seems to be not very isolated, lying at the outskirts of the loose group LGG 103. SBS~0335-052 is at a projected distance of $\\sim$150 kpc, and a velocity difference of $\\sim$120~\\kms\\ from the luminous spiral NGC~1376.  Similarly, the current starburst in the low-metallicity galaxy HS~0822+3542 appears to be triggered by interaction with its companion SAO~0822+3545 (\\cite{pustilnik03}; \\cite{chengalur06}; but see \\cite{corbin05} for an alternative interpretation). For I~Zw~18, an XMD system with the second-lowest known metallicity, there is an extended \\ion{H}{i} distribution with four distinct components (\\cite{vanzee98}), only two of which have optical counterparts.  Interpreting the observed configuration of SBS~1129+576 as the first close encounter (as is for the case of \\ion{H}{i}~1225+01), we would expect that the significant exchange of matter has not  yet taken place. This is consistent with the large difference of metallicities (O/H differ by a factor of $\\sim$3, or 0.5~dex). As noted by \\cite{pustilnik04}, in some of the probable XMD mergers, which are at more advanced stages, the differences in O/H in different components are significantly smaller. This is consistent with the gas exchange leading to a more homogeneous composition of ISM in the merging components towards the late stages of merger. In particular, for HS~0122+0743 (\\cite{pramskij03}; Pustilnik et al., in preparation), in which optical components are in contact, O/H differs by less than 0.1~dex. On the other hand, for SBS~0335--052 E and W, which seem to have their first close encounter the situation appears to be more complicated. While the galaxies are not expected to have exchanged much gas, their O/H differs by only 0.17 dex. This may be the result of the metal-enrichment by the first star formation episodes in both components leading to a typical value of Z~$\\sim$1/30~Z$_{\\odot}$ (\\cite{kunth86}). Summarising the comparison of SBS~1129+576 with several other XMD BCGs, we find that the tidal-interaction-triggered star formation is a common theme for all of them.    We summarise our conclusions about this \\ion{H}{i} study of the metal-poor galaxy SBS~1129+576 and its immediate surroundings as follows:  \\begin{enumerate}  \\item SBS~1129+576 is undergoing a strong interaction with the nearby dwarf galaxy SBS~1129+577. We assume that \\ion{H}{i} has been pulled out from it to form a bridge between the two galaxies. The mass of \\ion{H}{i} in the bridge is 2.1~$\\times$ 10$^{8}$~M$_{\\odot}$, roughly half the \\ion{H}{i} mass remaining in the disc of SBS~1129+576. The observed morphology of the system suggests that it is in a prograde orbit. The system is probably on its first encounter and no significant exchange of material has taken place yet. This is consistent with the very different ISM metallicities of the two systems. The enhanced star formation in these gas-rich galaxies ($M_{\\rm HI}$/L$_{\\rm B}$ of 2.1 and 3.8) has very likely been triggered by mutual tidal interaction.  \\item  The two \\ion{H}{i} peaks in SBS 1129+576 are at face-on surface densities of $\\sim$1.5~$\\times$~10$^{21}$ and $\\sim$1.7~$\\times$~10$^{21}$~atoms~cm$^{-2}$. These coincide with the positions of the two brightest \\ion{H}{ii} regions, respectively. The face-on \\ion{H}{i} column density in \\ion{H}{ii} regions identified in SBS~1129+577 using the slit spectroscopy is above $\\sim$1.8~$\\times$~10~$^{21}$~atoms~cm$^{-2}$. This is similar in value to the threshold density for star formation seen in other BCGs.  \\item  SBS~1129+577 has a central bar with a ring surrounding it, while the outer gas has been distorted into a strong one-armed spiral pattern. These are also likely to be consequences of the tidal interaction.  \\item   The velocity field of SBS~1129+577 shows that the disc is warped in the outer region, while in the inner region the gas orbits appear to be non-circular, probably as a response to the potential of the central bar. There is a hint of an integral-shaped warp in the inner parts of disc of SBS~1129+576.  \\item   The system has a third dwarf galaxy with \\ion{H}{i} mass of 1.1$\\times$10$^{7}$ M$_{\\odot}$. It has an $M_{\\rm HI}$/L$_{\\rm B}$ value of 0.08, that is, much lower than those for its more-massive companions. While the optical spectrum of this galaxy indicates a recent star-formation episode ($\\sim$20 Myr, comparable to that for starburst ages in SBS~1129+576 and SBS~1129+577), its low $M_{\\rm HI}$/L$_{\\rm B}$ suggests that it has had a very different evolutionary history. It is further unusual in having ongoing star formation, despite having a peak observed gas column density lower than the nominal threshold for star formation. We note, however, the true column density may be considerably higher than that derived from our comparatively low poor resolution maps.  \\item   The system SBS~1129+576/577 is similar in many aspects to two other well-known XMD systems. These are \\ion{H}{i}~1225+01 and SBS~0335--052 E and W, where the current star burst appears to have been triggered by tidal interaction.  \\end{enumerate}  {\\it Acknowledgment.} SAP acknowledges the partial support from the Russian federal program 'Non-stationary phenomena in astronomy'. The authors acknowledge the spectral and photometric information from the SDSS data base, used for this study. The Sloan Digital Sky Survey (SDSS) is a joint project of the University of Chicago, Fermilab, the Institute for Advanced Study, the Japan Participation Group, the Johns Hopkins University, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, Princeton University, the United States Naval Observatory, and the University of Washington. Apache Point Observatory, site of the SDSS telescopes, is operated by the Astrophysical Research Consortium (ARC). We thank the GMRT staff for making these observations possible. The GMRT is run by the National Centre for Radio Astrophysics of the Tata Institute of Fundamental Research."
    },
    "0609/astro-ph0609151_arXiv.txt": {
        "abstract": "A simple analytical model is used to calculate the X-ray heating of the IGM for a range of black hole masses. This process is efficient enough to decouple the spin temperature of the intergalactic medium from the cosmic microwave background (CMB) temperature and produce a differential brightness temperature of the order of $\\sim 5-20~\\mathrm{mK}$ out to distances as large as a few co-moving Mpc, depending on the redshift, black hole mass and lifetime. We explore the influence of two types of black holes, those with and without ionising UV radiation.  The results of the simple analytical model are compared to those of a full spherically symmetric radiative transfer code.  Two simple scenarios are proposed for the formation and evolution of black hole mass density in the Universe.  The first considers an intermediate mass black hole that form as an end-product of Population~III stars, whereas the second considers super-massive black holes that form directly through the collapse of massive halos with low spin parameter. These scenarios are shown not to violate any of the observational constraints, yet produce enough X-ray photons to decouple the spin-temperature from that of the CMB.  This is an important issue for future high redshift 21~cm observations. ",
        "introduction": "\\label{introduction}  One of the most startling findings made in the last few years is the discovery of super-massive black holes at redshifts $\\gsim 5.7$ with black hole masses of the the order of $10^9 M_\\odot$ (Fan \\etal 2003 and 2006). The origin and seeds of these black holes remain uncertain.  Currently, there are two main scenarios for creating such massive black holes. One is as the end-product of the first metal free stars (Population~III stars) that have formed through molecular hydrogen cooling (Abel, Bryan, Norman 2000, 2002, Bromm, Coppi \\& Larson 2002, Yoshida \\etal 2003). Given the low cooling rate provided by molecular hydrogen, the collapsing initial cloud is expected not to be able to fragment into small masses and thus produce very massive stars (for reviews, see Bromm \\& Larson 2004, Ciardi and Ferrara 2005).  These stars are expected to burn their fuel very quickly and to produce black holes with masses in the range $30 - 1000 M_\\odot$ (O'Shea and Norman 2006), with the exception of the mass range of $140 - 260 M_\\odot$ where the pair-instability supernovae leave no black hole remnants (Bond, Arnett \\& Carr 1984; Heger \\& Woosley 2002; Rakavy, Shaviv \\& Zinamon 1967).  Such objects grew their masses very efficiently by accretion up to $10^9 M_\\odot$ by $z\\approx 6$ (Volonteri \\& Rees 2005, Rhook \\& Haehnelt 2006).  The second avenue for producing even more massive black holes is through the collapse of very low angular momentum gas in rare dark-matter halos with virial temperatures above $10^4 \\mathrm{K}$ (see Shapiro 2004 for a recent review). Under such conditions, atomic cooling becomes efficient and black holes with masses $\\gg 10^3 M_\\odot$ can be formed (Bromm \\& Loeb 2003). Fragmentation of the initial gas into smaller mass objects due to efficient cooling can be prevented by trapping the Lyman-$\\alpha$ photons within the collapsing gas (Spaans \\& Silk 2006).  Notwithstanding the origin of these massive black holes, their impact on the intergalactic medium is expected to be dramatic in at least two ways. Firstly, these objects produce very intense ionising radiation with power-law behaviour that creates a different ionization pattern around them from that associated with thermal (i.e., stellar) sources. The ionization aspect of the miniquasar radiation has been explored by several authors (Madau, Meiksen \\& Rees 1997, Ricotti \\& Ostriker 2004a, 2004b, Madau \\etal 2004, Zaroubi \\& Silk 2005). Recently, however, it has been argued (Dijkstra, Haiman \\& Loeb 2004, Salvaterra, Haardt \\& Ferrara 2005) that miniquasars can not reionize the Universe as they will produce far more soft X-ray background radiation than currently observed (Moretti et al. 2003; So{\\l}tan 2003) and at the same time satisfy the WMAP 3$^{rd}$ year polarisation results (Page \\etal 2006 ; Spergel \\etal 2006) and the reionization constraints from the IGM temperature at redshift $\\approx 3$ (Theuns \\etal 2002a, 2002b, Schaye \\etal 2000). It should be noted, however, that the Dijkstra \\etal 2004 \\& Salvaterra \\etal\\ 2005 calculations have been carried out assuming specific black-hole mass density evolution histories and spectral energy distributions of UV/X-ray radiation emanating from the miniquasars.  Secondly, due to their x-ray radiation, even the intermediate mass black holes (IMBH) are very efficient in heating up their surroundings.  Nusser (2005) has pointed out that this heating facilitates observation of the redshifted 21 cm radiation in either emission or absorption by the neutral hydrogen in the high-redshift IGM. The observation of this radiation is controlled by the 21 cm spin temperature, $T_{spin}$, defined through the equation $n_1/n_0=3 \\exp(-T_\\ast/T_{spin})$. Here $n_1$ and $n_0$ are the number densities of electrons in the triplet and singlet states of the hyperfine levels, and $T_\\ast=0.0681~\\mathrm{K}$ is the temperature corresponding to the 21 cm wavelength. For the 21 cm radiation to be observed relative to the CMB background, it has to attain a different temperature and therefore must be decoupled from the CMB (Wouthuysen 1952; Field 1958, 1959; Hogan \\& Rees 1979). The decoupling is achieved through either Lyman$-\\alpha$ radiation or collisional excitations and heating. For the objects we are concerned with in this paper, \\ie\\ miniquasars, the collisional excitation and heating are much more important.  In general, throughout this paper, we will ignore the influence of Lyman-$\\alpha$ photons emitted by the quasar on $T_{spin}$. However, one should point out that collisional excitations due to x-ray photons results in a ``secondary'' Lyman$-alpha$ pumping which will dominate the spin temperature and CMB temperature decoupling in some regions around the miniquasar; this effect has been recently point out by Chuzhoy, Alvarzez \\& Shapiro (2006).  For recent papers that discuss X-ray heating, see Chen \\& Miralda-Escude (2006) and Pritchard \\& Furlanetto (2006).  Collisional decoupling of $T_{spin}$ from $T_{CMB}$ is caused by very energetic electrons released by the effect of the x-ray miniquasar radiation on the IGM. Shull \\& van Steenberg (1985) have estimated that more than a tenth of the energy of the incident photons is absorbed by the surrounding medium as heating (this fraction increases rapidly with the ionized fraction). The increase in the temperature is observable at radio frequencies in terms of the differential brightness temperature, $\\delta T_b$, which measures the 21 cm intensity relative to the CMB. A similar fraction of the absorbed energy also goes into collisional excitation, where this fraction decreases rapidly with the ionised fraction. These two processes, heating and excitation decouple the spin temperature from the CMB temperature and render the IGM observable through its 21 cm emission.   Recently, Kuhlen \\& Madau (2005) and Kuhlen, Madau \\& Montgomery (2006) have performed a detailed numerical study of the influence of 150$M_\\odot$ IMBH on its surroundings and calculated the gas, spin and brightness temperatures. They have shown that heating by 150$M_\\odot$ IMBH at $z=17.5$ can enhance the 21 cm emission from the warm neutral IGM. The filaments enhance the signal even further and may make the IGM visible in  future radio experiments (e.g., the LOFAR-Epoch of Reionization key science project\\footnote{For more details on the LOFAR radio telescope see http://www.lofar.org}).  In this paper, we adopt a complementary theoretical approach to the numerical one adopted by Kuhlen \\& Madau (2005). This allows us to explore the influence of power-law radiation fields from a range of black hole masses that are presumed to reside in the centres of primordial miniquasars. Furthermore, the effect of X-ray induced collisional excitations on the 21 cm spin temperature is included (Chuzhoy, Alvarez \\& Shapiro 2006) -- this effect is not taken into account in the Kuhlen \\& Madau (2005) work. We test two main classes of x-ray emitting miniquasars, those with UV ionizing radiation and those without.  We show that in both cases these miniquasars might play an important role in heating the IGM without necessarily ionizing it completely.  In addition, two simple scenarios for the formation of (mini)quasars as a function of redshift are presented. This is done using the extended Press-Schechter algorithm to predict the number density of forming black holes either with H$_2$ cooling or with atomic cooling. We also discuss the implications of these scenarios for the mass density of quasars at redshift~6, the soft X-ray background (SXRB)in the energy range $0.5-2~\\mathrm{keV}$ (Dijkstra, Haiman \\& Loeb 2004), the number of ionizing photons per baryon and, finally, the optical depth for Thomson scattering of CMB photons.   The paper is organised as follows: Section~\\ref{Theory} describes the theoretical methods used here and derives the ionization and kinetic temperature profiles around miniquasars without UV ionizing radiation. Section~\\ref{Tspin} calculates the spin and brightness temperature around the same quasars. In section~\\ref{IMBH} we show the ionization and heating profiles around quasars with UV ionizing radiation.  Section~\\ref{PS} present the two formation scenarios and their implications. The paper concludes with a summary section (\\S~\\ref{summary}). ",
        "conclusions": "This paper explores the feasibility of heating the IGM with quasars without violating the current observational constraints. Such heating is essential in order to be able to observe the 21 cm emission from neutral hydrogen, prior to and during the epoch of reionization. We have shown that miniquasars with moderate black hole masses can heat the surrounding IGM out to radii of a few comoving mega-parsecs.  In this paper, two Press-Schechter based black hole mass density evolution scenarios have been proposed, IMBH and SMBH. The first model assumes the black hole population is the end product of pop~III stars that leave behind black hole masses of the order of 10-100$M_\\odot$. The second model assumes direct formation of black holes as a result of the collapse of low angular momentum primordial halos.  For these two scenarios, we have explored three different quasar spectral templates: a power-law with ionization UV radiation, a power-law without ionising UV radiation and a Sazonov et al. (2004) type template.  With the exception of the models that have a $10\\%$ duty cycle, we have shown that the quasars are not able to fully ionise the IGM -- especially if one assumes the template that does not have ionising UV photons -- while the SXRB constraint is satisfied. We conclude that based on the mass evolution history shown here, there is enough mass in the quasars to heat up the IGM by redshift 15. For example, for quasars with a power-law index of $-1$ and no ionising UV radiation, quasars with black hole masses of $10^{3-4}M_\\odot$ can heat up the IGM over a $\\approx 0.1-1 \\mathrm{Mpc}$ comoving radius from the (mini-)quasar (see figure~\\ref{fig:dTb}). The models with $6\\%$ duty cycle reach such mass per comoving Mpc$^3$ at redshift larger than 10 for both scenarios.  Curiously, for the black hole mass density evolution, with $10\\%$ duty cycle and rigid disk model, proposed by Begelman \\etal (2006) we find that this scenario does not violate the SXRB observational constraint and produce about 10 ionising photons per baryon by $z=6$, normally thought to be enough to ionise the Universe. We also find that for a model in which the Universe ionizes suddenly at z=6, that this scenario predicts a Thomson scattering optical depth of $0.075$, consistent with the WMAP 3$^{rd}$ year results.  The main result presented in this paper is ``good news'' for the new generation of low frequency radio telescopes designed to probe the high redshift IGM through its 21~cm emission, such as LOFAR, MWA and PAST. It clearly shows that the quasar population could easily decouple the spin-temperature from that of the CMB.  However, since the spin temperatures achieved are not very high, this means that the brightness temperature will carry the signature not only of the ionized fraction and density fluctuations, but also of the variations in the spin temperature. This complicates the interpretation of the observed brightness temperature in terms of its link to the cosmological fields. Nevertheless, the high spin temperature bubbles are expected to overlap before those of the ionization, a factor that will mitigate this complication. Furthermore, one can turn this around and argue that these fluctuations will teach us more about the ionising sources than about cosmology. An extended tail in the spin temperature will be a clear signature of power-law radiation, {\\it i.e.}, quasars, while a short tail will be a clear signature of thermal radiation, {\\it i.e.}, stars.   \\label{summary}"
    },
    "0609/astro-ph0609198_arXiv.txt": {
        "abstract": "We use the results of the SCUBA Local Universe Galaxy Survey, a submillimetre survey of galaxies in the nearby Universe, to investigate the relationship between the far-infrared--submillimetre and radio emission of galaxies at both low and high redshift. At low redshift we show that the correlation between radio and far-infrared emission is much stronger than the correlation between radio and submillimetre emission, which is evidence that massive stars are the source of both the far-infrared and radio emission. At high redshift we show that the submillimetre sources detected by SCUBA are brighter sources of radio emission than are predicted from the properties of galaxies in the local Universe. We discuss possible reasons for the cosmic evolution of the relationship between radio and far-infrared emission. ",
        "introduction": "\\label{intro}  One of the most intriguing relationships in astronomy is that between the non-thermal radio and far-infrared (FIR) emission from galaxies. It is one of the strongest correlations in astronomy, with a tight correlation between FIR and radio luminosity over five decades of luminosity (e.g. Price \\& Duric 1992; Yun, Reddy \\& Condon 2001). The relationship holds for a remarkably wide variety of galaxy types, and is observed not only globally but also on $\\sim$kiloparsec scales within the disks of individual galaxies (e.g. Boulanger \\& P\\'erault 1988; Bicay \\& Helou 1990; Fitt et al. 1992; Murphy et al. 2006). Some authors also find evidence of the relationship at high redshifts (e.g. Garrett 2002; Appleton et al. 2004; Kov\\'acs et al. 2006). Yet, although the relationship has been known for over two decades (e.g. Condon 1992 and references therein), the cause of the relationship is still unclear.  The standard explanation of the relationship is that both the FIR and radio emission are caused by high-mass stars, with the stars both heating the dust, which then emits FIR emission, and producing the relativistic electrons responsible for synchrotron radiation when at the end of their lives they explode as supernovae (see e.g. Condon 1992 and references therein). Alternative explanations have been proposed. However, most workers in the field have assumed that the standard explanation (i.e. ongoing star-formation) is the correct one.  Nevertheless, despite its widespread acceptance, there is surprisingly little direct evidence for the standard explanation and it has one major problem. The problem is that the slope of the FIR--radio relation is roughly unity (Helou, Soifer \\& Rowan-Robinson 1985; Devereux \\& Eales 1989) whereas the synchrotron emissivity does not just depend on the number-density of relativistic electrons but also on the strength of the magnetic field. Although there are only a few measurements from Zeeman splitting of the magnetic fields in galaxies (Thompson et al. 2006 and references therein), it is possible to estimate the magnetic field in a galaxy using the `minimum energy' argument (Longair 1983), and estimates of the magnetic field made in this way can vary by as much as a factor of 100 from normal galaxies like our own to vigorously star-forming galaxies such as Arp 220 (Thompson et al. 2006). The radio spectrum of a galaxy roughly follows a power-law ($S \\propto \\nu^{-\\alpha}$) and the synchrotron emissivity is proportional to $B^{1.0+\\alpha}$. The radio emissivity in Arp 220 will therefore be much higher than in our own galaxy, not only because of the much greater production rate of relativistic electrons per cubic kiloparsec, but also because of the stronger magnetic field. However, it is still possible to explain the unity slope of the FIR--radio relation if the relativistic electrons lose all their energy within their galaxy, and thus an electron in Arp 220 loses its energy at a faster rate but ultimately only loses the same amount of energy as in a galaxy like our own.  This scenario is described by the `calorimeter' model (V\\\"olk 1989; Lisenfeld, V\\\"olk \\& Xu 1996), in which the dust-heating stellar UV photons and the relativistic electrons lose most of their energy within the galaxy. The calorimeter model will only give a unity slope for the FIR--radio relation if the ratio of inverse Compton to synchrotron losses is roughly the same in all galaxies, which leads to the prediction that the energy density of the interstellar radiation field and that of the magnetic field have a constant ratio for galaxies. Lisenfeld et al. (1996) used observations of 114 `normal' late-type galaxies to estimate the mean energy densities of the interstellar radiation field and, using the `minimum energy' argument, the magnetic field. They showed a linear correlation between the two, suggesting turbulent dynamo activity in galaxies as a possible explanation. Calorimeter-type models have also been presented by several other authors (e.g. Pohl 1994).  A more fundamental problem with the calorimeter model is that the radio spectra of star-forming galaxies do not exhibit the steepening that one would expect if the electrons lose all their energy within the galaxy in which they were formed. A recent interesting suggestion has been that the magnetic fields in galaxies are actually much higher than the estimates from the minimum energy argument, which would definitely lead to an electron losing all its energy before it diffuses out of its galaxy (Thompson et al. 2006). Thompson et al. explain the lack of clear curvature in the radio spectra as the result of ionization losses, which flatten the spectrum at low frequencies.  An alternative idea is that the FIR--radio relation holds for the `optically thin' scenario in which the majority of the stellar UV photons and relativistic electrons do escape from the galaxy disk (e.g. Helou \\& Bicay 1993). The FIR--radio correlation is then attributed to a constant ratio between the production rates of dust-heating UV photons and synchrotron-emitting relativistic electrons.  Niklas \\& Beck (1997) argue that observations do not support either the calorimeter or `optically-thin' model and suggest an alternative model in which the FIR--radio relation arises from a combination of correlations; that  between the volume density of cool (neutral) gas and the star formation rate and magnetic field strength, and that between the star formation rate and the FIR luminosity, assuming equipartion between the cosmic ray energy density and the magnetic field strength. In this scenario the cool neutral gas clouds play a central role. Such a model, unlike either the calorimeter (optically thick) or optically thin models, can explain the fact that the FIR--radio relation is not only observed globally but is also observed on kiloparsec scales within the disks of individual galaxies.  Alternative explanations have also been proposed by a number of other authors. Bettens et al. (1993) proposed that the cosmic ray flux, which as in the `standard explanation' produces the non-thermal radio emission, drives the evolution of molecular clouds; a higher flux of cosmic rays produces a higher ionization fraction in the cloud, which delays the cloud's collapse, and ultimately leads to a larger fraction of OB stars and more FIR emission from the surrounding dust. However, other authors (e.g. Niklas \\& Beck 1997, Murgia et al. 2005) have pointed out problems with this model; for example it fails to take into account the effect of the the interstellar magnetic field strength on the synchrotron luminosity.  Murgia et al. (2005) proposed a model in which the relation between CO, radio continuum and FIR emission results from hydrostatic pressure in galaxies, which can explain the correlation on both global and local scales. In this model the relationship between the CO and radio is explained by hydrostatic pressure acting as a regulating mechanism between the molecular gas fraction and the magnetic field strength, while the link to the FIR emission comes from models by Dopita et al. (2004) which connect FIR emission with interstellar pressure. This model thus avoids an explicit dependence on star formation.  Most of the various models proposed to explain the FIR--radio relationship require a link between the gas density and the magnetic field strength. One suggestion as to possible physical processes behind such a model is magnetohydrodynamic turbulence (Groves et al. 2003).  Whatever its cause, the strong correlation between FIR and radio emission in nearby galaxies provided submillimetre (submm) astronomers with a valuable tool for estimating the redshifts of the galaxies discovered in the first deep submm surveys (Smail, Ivison \\& Blain 1997; Hughes et al. 1998; Eales et al. 1999). These galaxies are very faint at optical wavelengths, which has made it very difficult to measure redshifts using conventional spectroscopy, and until recently very few spectroscopic redshifts existed (Chapman et al. 2005). Carilli \\& Yun (1999) first suggested that an alternative to optical spectroscopy would be to assume that the FIR--radio relation is the same at all redshifts, and then estimate redshifts from the ratio of submm to radio flux densities.  In this paper we explore the FIR-radio relationship at both high and low redshifts, using the results of the recently-completed SCUBA Local Universe Galaxy Survey (SLUGS). This is a submm survey of $\\sim 200$ galaxies in the local Universe, half drawn from a sample selected in the FIR (Dunne et al. 2000b; herein the IRS sample) and half drawn from an optically-selected (herein OS) sample (Vlahakis, Dunne \\& Eales 2005). Our submm measurements allow us to test a basic prediction of the standard theory. Because of the greater intensity of the interstellar radiation field in regions containing large numbers of OB stars, the dust in these regions will be hotter than the dust in the general interstellar medium. Therefore, the standard hypothesis makes the prediction that the correlation between FIR and radio emission will be tighter than the correlation between submm and radio emission. Previous authors have found evidence both in favour of and against this prediction by decomposing the FIR emission into warm and cold components (e.g. Hoernes, Berkhuijsen and Xu 1998; Pierini et al. 2003), but this previous work was based on FIR measurements alone rather than measurements at submm wavelengths. In Section~\\ref{fir-radio} we use our SLUGS data to test the prediction. In Section~\\ref{indicator} we then consider the FIR--radio relationship at high redshift. Because SLUGS is the only large sample of galaxies for which there are measured spectral energy distributions extending through the FIR and submm wavebands, we can predict the ratios of radio to submm emission one would see at high redshift {\\it if there is no cosmic evolution in the FIR--radio relationship.} Dunne, Clements \\& Eales (2000) performed a similar analysis using the IRS data alone. Using the spectroscopic redshifts from Chapman et al. (2005), we show that the sources detected in deep SCUBA surveys are generally brighter radio sources than one would predict from the FIR--radio relationship seen in the local Universe. ",
        "conclusions": "Using the results of the OS SLUGS sample we have investigated the relationships between FIR--submm and radio properties for `normal' local galaxies, compared to bright \\IRAS\\/ galaxies (using the IRS SLUGS sample from Dunne et al. 2000b). The results for the OS galaxies have then been used to assess the reliability of the Carilli \\& Yun (1999, 2000) radio--submm redshift estimator technique, for the first time using a sample containing significant fractions of cold dust.  We find the following results. \\begin{itemize} \\item{There is a very tight correlation between 1.4\\,GHz luminosity and 60-\\micron\\/ luminosity for the OS sample galaxies. There is much more scatter in the correlation between radio and 850\\mic luminosity, thus providing evidence that the standard explanation for the FIR--radio luminosity correlation is correct and that it is the population of young (OB) stars which both heats the dust and provides the source of radio emission. }  \\item{Using the SEDs of the 17 OS galaxies detected at both 850- and 450-\\mic we have shown that for a sample of `normal' galaxies there is much more scatter in the $\\alpha^{850}_{1.4}$-redshift relation than seen for the 25 most luminous \\IRAS\\/ galaxies of the IRS SLUGS sample, most likely since the OS objects contain large fractions of cold dust. From this we conclude that in order for the Carilli \\& Yun radio--submm redshift indicator method to be reliable as a redshift estimator for deep submm sources one would need measurements of the temperature of the dust in those objects.}  \\item{We have compared our OS $\\alpha^{850}_{1.4}$-redshift relations to the deep submm observations of Chapman et al. (2005) and demonstrate that the large majority of these deep SCUBA sources must have very different properties to our sample of low-redshift SLUGS galaxies. We find that the majority of the deep SCUBA galaxies are brighter at radio wavelengths than any of our predictions. We cannot exclude the possibility that this simply reflects that our low-redshift galaxies are not suitable templates for the high-redshift SCUBA galaxies, but we think it more likely that this shows the FIR--radio relation is evolving with redshift.} \\end{itemize}"
    },
    "0609/astro-ph0609684_arXiv.txt": {
        "abstract": "{} {To calculate transition rates from ground and excited states in neutral oxygen atoms due to electron collisions for non-LTE modelling of oxygen in late-type stellar atmospheres, thus enabling reliable interpretation of oxygen lines in stellar spectra.} {A 38-state $R$-matrix calculation in $LS$-coupling has been performed.  Basis orbitals from the literature (Thomas~et~al.) are adopted, and a large set of configurations are included to obtain good representations of the target wavefunctions.  Rate coefficients are calculated by averaging over a Maxwellian velocity distribution.  } {Estimates for the cross sections and rate coefficients are presented for transitions between the seven lowest $LS$ states of neutral oxygen.   The cross sections for excitation from the ground state compare well with existing experimental and recent theoretical results.   } {} ",
        "introduction": "\\label{sect:intro}  Oxygen is one of the most abundant and important elements in the universe.  Oxygen is widely accepted to be predominantly produced in core-collapse supernovae resulting from massive stars, and thus detailed study of the chemical evolution of oxygen allows us to understand star formation history and the initial mass function, as well as supernovae yields.  The main observational evidence regarding the evolution of oxygen in the Galaxy comes from the measurement of oxygen abundances in the atmospheres of cool stars.  The magnitude and trend of measured oxygen abundance with age or metallicity is not yet well constrained, in particular the results at low metallicity differ between different studies (e.g. Israelian~et~al.~\\cite{israelian98}, Boesgaard~et~al.~\\cite{boesgaard99}, Nissen et~al.~\\cite{nissen02}, Garcia-Perez et~al.~\\cite{garciaperez06}).  The use of different spectral diagnostics is most likely the reason for these discrepancies.  The use of the \\ion{O}{i} triplet at 777~nm is appealling as the lines are strong and unblended, even in metal-poor stars.  As they are in the visual region they are easily accessible with standard telescopes and spectrographs.  On the other hand, it is well known that these lines are not formed in conditions of local thermodynamic equilibrium (LTE) in cool star atmospheres (see Kiselman~\\cite{kiselman01} and references therein).  If the non-LTE formation of these lines could be understood, then they would be very useful diagnostics for measuring stellar oxygen abundances.  However, we also note that the lines originate from highly excited levels, and are thus quite temperature sensitive, and furthermore may be saturated at higher abundances.  The ability to solve the non-LTE problem is dependent on the availability of adequate atomic data for the radiative and collisional processes occuring to the oxygen atoms in the atmosphere.  The data for the collisional processes is the greatest stumbling block to reliable non-LTE modelling.  Collisions with electrons are particularly important, due to their large frequency (due to high velocity) and efficiency (collisions are not adiabatic).  Though collisions with hydrogen atoms are less frequent and expected to be less efficient (adiabatic) they may also be important due to sheer weight of number (typically $N_\\mathrm{H} \\ga 10^4 N_e$).  Though Kiselman~(\\cite{kiselman01}) identifies the collisional cross-sections due to hydrogen as the most pressing issue since there is very little information on these processes (but see recent paper by Krems~et~al.~\\cite{krems06}), the situation for electrons is still far from ideal.  Past studies of electron-impact excitation of neutral oxygen have predominantly focussed on excitation from the ground {\\ch configuration}.  However, for the purposes of the non-LTE problem in stellar atmospheres, data for excited states are also needed.  Previous work on the oxygen non-LTE problem have used either approximate general formulae (e.g.\\ Allende~Prieto~et~al.~\\cite{allende03}) or a range of data calculated by different methods (e.g.\\ Carlsson \\& Judge~\\cite{carlsson93}, Kiselman~\\cite{kiselman93}).  In some cases data (except the simple formulae) do not exist, such as for the coupling between $LS$ states of different spin, e.g. triplet--quintet system coupling.  Kiselman performed test calculations for  three stellar atmosphere models and found that intersystem collisional coupling was probably not important.  Recent work by Fabbian~et~al.~(private communication) confirms the results of Kiselman for the considered models, but finds that triplet-quintet intersystem coupling may be important in the case of metal-poor turnoff stars (e.g. $T_\\mathrm{eff}=6500$~K, $\\log g=4$, $\\mathrm{[Fe/H]}=-3$).  In any case, it is clear that a consistently computed set of reliable data for electron-impact excitation of oxygen is highly desirable for non-LTE studies of oxygen line formation in stellar atmospheres.  In this paper, we use the $R$-matrix approach to calculate electron-impact excitation cross sections for low energy collisions for transitions between the seven lowest $LS$ states of neutral oxygen, from which we calculate rate coefficients for cool star atmosphere temperatures.  We compare our results with experiment and those in literature, and discuss the differences. ",
        "conclusions": "The computed cross sections are plotted in Fig.~\\ref{fig:cross}.  Note, that in both Figs.~\\ref{fig:cross} and~\\ref{fig:comp}, some regions which were affected by pseudo-resonances at energies above the ionisation threshold were either smoothed or interpolated to remove any sharp transient behaviour which is unlikely to be physical.  The results of the 38-state and 19-state calculations are shown, both smoothed in the same manner.  Comparison of these results, along with the 33-state calculation, allows pseudo-resonances to be identified as features which change between the calculations thus are very likely to be spurious (due to incompleteness of the model), while features that remain unchanged are most likely real. In general, we may note that the near threshold behaviour, most important for the temperatures of interest, are not greatly affected.  However, it is clear that many of the cross sections are probably affected by pseudo-resonances particularly at {\\ch total} energies around 15-20~eV.  Note, the feature at around 14~eV in the $^1D$--$^1$S cross section is a real feature, it is a resonance in the $^2$S partial cross section (see e.g.\\ Thomas~et~al.~\\cite{thomas97}).  As expected, the singlet-quintet transitions have null cross-sections since in $LS$-coupling total electronic spin must be conserved.  In addition, the $^1S$$\\rightarrow$$^3S^o$ transition is forbidden as $S^e$$\\leftrightarrow$$S^o$ transitions are forbidden in general (Goddard~et~al.~\\cite{goddard71}).  \\begin{figure*} \\begin{center} \\resizebox{170mm}{!}{\\rotatebox{0}{\\includegraphics{fig1.eps}}} \\end{center} \\caption{The cross sections for electron-impact excitation as a function of initial electron energy between the seven lowest target states.  The full lines are the 38-state calculation, while the dotted lines are the 19-state calculation.  The cross sections are in units of $10^{-18}$~cm$^2$.} \\label{fig:cross} \\end{figure*}  Experimental results at low energy are available for some transitions from the ground state.  We compare our results with experiment in Fig.~\\ref{fig:comp}. In the case of the $^3P$$\\rightarrow$$^3S^o$ transition, we also plot the cross section from the 41-state $R$-matrix with pseudostates calculation tabulated in Johnson~et~al~(\\cite{johnson03a}) and described in Tayal~(\\cite{tayal02}).  In general the computed data compare quite well with the experimental data considering the experimental error bars.  For the $^3P$$\\rightarrow$$^1D$ transition, the experimental results suggest a peak in the cross section at an impact energy of about 6~eV.  A similar peak is seen in the computed cross section, though only roughly half the magnitude.  Compared to the cross section of Thomas~et~al., the peak is about of the same magnitude but at lower energy, with our cross section not showing the slow rate of increase in the cross section of that of Thomas~et~al., but rather a sharp increase at the threshold.  The cases of the $^3P$$\\rightarrow$$^1S$ and $^1D$$\\rightarrow$$^1S$ (see Fig.~\\ref{fig:cross}) transitions are quite similar, again rather than the slow rate of increase in the cross section just above threshold of Thomas~et~al., we find a sharp increase.  The cross section for the $^3P$$\\rightarrow$$^5S^o$ transition compares well with experiment at low energy, but is significantly larger than the experimental values at energies at and above 20~eV.  The cross sections also compare well with those of Zatsarinny \\& Tayal~(\\cite{zatsarinny02}), having the same magnitude and form near threshold.  The peak at about 18~eV seems to be the result of pseudo-resonances, as judged from the fact that the peak varied significantly between the 19-, 33- and 38-state calculations (see Fig.~\\ref{fig:cross}).  We note the peak is not present in the results of Zatsarinny \\& Tayal, which uses the non-orthogonal orbital method which should be more robust against such pseudo-resonances.  Similarly, the $^3P$$\\rightarrow$$^5P$ and $^3P$$\\rightarrow$$3p ^3P$ results seem to be affected by pseudo-resonances at {\\ch total energies} of about 18~eV.  However, near threhold the results are similar to those of Zatsarinny \\& Tayal.  The resonance transition $^3P$$\\rightarrow$$^3S^o$ is the best studied, with a significant number of experimental and theoretical studies.  As seen in Fig.~\\ref{fig:comp}, our result compares well with the experimental results, both in terms of magnitude and form of the cross sections.  The cross section compares well with the theoretical result of Tayal~(\\cite{tayal02}) at energies between 15 and 30~eV, which is also plotted.  Our cross sections are around 5--10\\% larger, as expected (see discussion in previous section) due to our neglect of a large part of the coupling to the continuum.  Regarding the near threshold behaviour, our results are very similar to those of Zatsarinny \\& Tayal~(\\cite{zatsarinny02}), the cross sections having very similar magnitude and form, including a sharp resonance at the threshold.   \\begin{figure*} \\begin{center} \\resizebox{\\hsize}{!}{\\rotatebox{0}{\\includegraphics{fig2.eps}}} \\end{center} \\caption{The collision cross sections as a function of initial electron energy compared with experiment. The full lines are the results of the present work.  Experimental and other theoretical results in each case are from: {\\bf (a)} Doering~(\\cite{doering92}), circles; Shyn \\& Sharp~(\\cite{shyn86a}, squares; in the latter case, the error bars are slightly offset for clarity. {\\bf (b)} Doering \\& Gulcicek~(\\cite{doering89a}), circles; Shyn~et~al.~(\\cite{shyn86a}), squares; in the latter case the error bars are slightly offset for clarity. {\\bf (c)} Doering \\& Gulcicek~(\\cite{doering89b}), circles. {\\bf (d)} Doering \\& Yang~(\\cite{doering01}), circles; Johnson~et~al~(\\cite{johnson03a}), small filled squares;  model-based results from experimental electron-impact-induced emission cross sections from Johnson~et~al~(\\cite{johnson03b}), open squares; 41-state calculation of Tayal~(\\cite{tayal02}), thick dotted line. {\\bf (e)} Gulcicek~et~al.~(\\cite{gulcicek88}), circles. {\\bf (f)} Gulcicek~et~al.~(\\cite{gulcicek88}), circles. } \\label{fig:comp} \\end{figure*}  Rate coefficients are computed by integrating a Maxwellian velocity distribution with the cross-sections, and are presented in Table~\\ref{tab:ratecoeffs}.  The calculations are done assuming $LS$ coupling, and as such no fine structure is included.  Fine structure of certain levels of interest are often included in non-LTE modelling, and one may wish to somehow redistribute the rate coefficent between fine structure levels.  It is often the practice to set fine structure transition rates to large values such that the fine structure levels have populations according to their statistical weights (e.g. Kiselman~\\cite{kiselman93}).  As such, it is irrelevant how the rates are distributed so long as the total rate is conserved.  However, we note that even neglecting relativistic effects in the Hamiltonian, a simple redistribution of the rate coefficients is not possible as a transformation from $LS$ to $J_aj$ coupling (where $J_a$ and $j$ are the initial total angular momenta of the atom and scattered electron respectively) must be done for each partial wave and scattered electron angular momentum (e.g.\\ Saraph~\\cite{saraph70}).   \\begin{table*} \\begin{center} \\caption{Rate coefficients $\\langle \\sigma v \\rangle$, in units of cm$^3$~s$^{-1}$, for electron-impact excitation of neutral oxygen, for selected temperatures in the {\\ch range $T=1000$--$50000$~K}.  Note, $a(b) = a \\times 10^b$.} \\label{tab:ratecoeffs} \\begin{tabular}{lrrrrrr} \\hline \\hline Initial & \\multicolumn{6}{c}{Final state} \\\\ state  & 2p$^4$ $^1D  $&2p$^4$ $^1S  $&    3s $^5S^o$&    3s $^3S^o$&    3p $^5P  $&     3p $^3P  $             \\\\ \\hline & \\multicolumn{6}{c}{\\underline{ 1000~K}} \\\\ 2p$^4$ $^3P  $&   $5.72(-21)$&   $5.13(-31)$&   $1.03(-55)$&   $3.40(-57)$&   $6.35(-64)$&    $5.48(-65)$             \\\\ 2p$^4$ $^1D  $&              &   $1.91(-20)$&         $ 0.$&   $3.34(-49)$&         $ 0.$&    $6.34(-56)$             \\\\ 2p$^4$ $^1S  $&              &              &         $ 0.$&         $ 0.$&         $ 0.$&    $8.85(-45)$             \\\\ 3s $^5S^o$&              &              &              &   $9.98(-10)$&   $1.11(-15)$&    $1.98(-17)$             \\\\ 3s $^3S^o$&              &              &              &              &   $9.91(-14)$&    $3.44(-15)$             \\\\ 3p $^5P  $&              &              &              &              &              &    $2.56(-09)$             \\\\ & \\multicolumn{6}{c}{\\underline{ 3000~K}} \\\\ 2p$^4$ $^3P  $&   $2.98(-13)$&   $4.28(-17)$&   $9.33(-25)$&   $3.46(-25)$&   $9.03(-28)$&    $4.89(-28)$             \\\\ 2p$^4$ $^1D  $&              &   $4.17(-13)$&         $ 0.$&   $6.87(-24)$&         $ 0.$&    $9.66(-26)$             \\\\ 2p$^4$ $^1S  $&              &              &         $ 0.$&         $ 0.$&         $ 0.$&    $4.53(-22)$             \\\\ 3s $^5S^o$&              &              &              &   $2.36(-08)$&   $6.23(-10)$&    $4.43(-11)$             \\\\ 3s $^3S^o$&              &              &              &              &   $1.65(-09)$&    $9.26(-10)$             \\\\ 3p $^5P  $&              &              &              &              &              &    $3.51(-08)$             \\\\ & \\multicolumn{6}{c}{\\underline{ 5000~K}} \\\\ 2p$^4$ $^3P  $&   $1.38(-11)$&   $3.07(-14)$&   $1.30(-18)$&   $7.53(-19)$&   $1.53(-20)$&    $1.26(-20)$             \\\\ 2p$^4$ $^1D  $&              &   $1.24(-11)$&         $ 0.$&   $6.11(-19)$&         $ 0.$&    $9.93(-20)$             \\\\ 2p$^4$ $^1S  $&              &              &         $ 0.$&         $ 0.$&         $ 0.$&    $1.41(-17)$             \\\\ 3s $^5S^o$&              &              &              &   $3.89(-08)$&   $1.02(-08)$&    $8.34(-10)$             \\\\ 3s $^3S^o$&              &              &              &              &   $1.09(-08)$&    $1.29(-08)$             \\\\ 3p $^5P  $&              &              &              &              &              &    $5.83(-08)$             \\\\ & \\multicolumn{6}{c}{\\underline{ 8000~K}} \\\\ 2p$^4$ $^3P  $&   $1.24(-10)$&   $1.34(-12)$&   $3.42(-15)$&   $2.66(-15)$&   $1.72(-16)$&    $1.88(-16)$             \\\\ 2p$^4$ $^1D  $&              &   $8.25(-11)$&         $ 0.$&   $3.75(-16)$&         $ 0.$&    $2.63(-16)$             \\\\ 2p$^4$ $^1S  $&              &              &         $ 0.$&         $ 0.$&         $ 0.$&    $7.98(-15)$             \\\\ 3s $^5S^o$&              &              &              &   $4.57(-08)$&   $4.97(-08)$&    $3.95(-09)$             \\\\ 3s $^3S^o$&              &              &              &              &   $2.86(-08)$&    $5.65(-08)$             \\\\ 3p $^5P  $&              &              &              &              &              &    $6.98(-08)$             \\\\ & \\multicolumn{6}{c}{\\underline{12000~K}} \\\\ 2p$^4$ $^3P  $&   $4.31(-10)$&   $1.10(-11)$&   $2.53(-13)$&   $2.47(-13)$&   $2.99(-14)$&    $3.96(-14)$             \\\\ 2p$^4$ $^1D  $&              &   $2.31(-10)$&         $ 0.$&   $1.50(-14)$&         $ 0.$&    $2.18(-14)$             \\\\ 2p$^4$ $^1S  $&              &              &         $ 0.$&         $ 0.$&         $ 0.$&    $3.30(-13)$             \\\\ 3s $^5S^o$&              &              &              &   $4.48(-08)$&   $1.20(-07)$&    $8.36(-09)$             \\\\ 3s $^3S^o$&              &              &              &              &   $4.38(-08)$&    $1.26(-07)$             \\\\ 3p $^5P  $&              &              &              &              &              &    $6.89(-08)$             \\\\ & \\multicolumn{6}{c}{\\underline{20000~K}} \\\\ 2p$^4$ $^3P  $&   $1.18(-09)$&   $5.84(-11)$&   $7.37(-12)$&   $9.56(-12)$&   $1.91(-12)$&    $3.00(-12)$             \\\\ 2p$^4$ $^1D  $&              &   $5.15(-10)$&         $ 0.$&   $3.24(-13)$&         $ 0.$&    $8.02(-13)$             \\\\ 2p$^4$ $^1S  $&              &              &         $ 0.$&         $ 0.$&         $ 0.$&    $6.05(-12)$             \\\\ 3s $^5S^o$&              &              &              &   $3.74(-08)$&   $2.38(-07)$&    $1.30(-08)$             \\\\ 3s $^3S^o$&              &              &              &              &   $5.25(-08)$&    $2.31(-07)$             \\\\ 3p $^5P  $&              &              &              &              &              &    $5.81(-08)$             \\\\ & \\multicolumn{6}{c}{\\underline{50000~K}} \\\\ 2p$^4$ $^3P  $&   $2.71(-09)$&   $2.35(-10)$&   $1.32(-10)$&   $3.04(-10)$&   $8.95(-11)$&    $1.77(-10)$             \\\\ 2p$^4$ $^1D  $&              &   $1.05(-09)$&         $ 0.$&   $5.85(-12)$&         $ 0.$&    $2.57(-11)$             \\\\ 2p$^4$ $^1S  $&              &              &         $ 0.$&         $ 0.$&         $ 0.$&    $6.50(-11)$             \\\\ 3s $^5S^o$&              &              &              &   $2.03(-08)$&   $4.12(-07)$&    $1.26(-08)$             \\\\ 3s $^3S^o$&              &              &              &              &   $4.02(-08)$&    $3.62(-07)$             \\\\ 3p $^5P  $&              &              &              &              &              &    $3.23(-08)$             \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}  It is of interest to compare these results, at least in broad terms, with those that have been used in past non-LTE calculations.  {\\ch Below we make comparisons at $T=5000$~K, which is representative of the typical temperatures of interest in cool stellar atmospheres.}  First we make a comparison with the predictions of approximate yet general and simple formulae used by Allende Prieto~et~al.~(\\cite{allende03}); this paper may be consulted for the formulae.  For allowed radiative transitions they used a formula due to Mihalas~(\\cite{mihalas72}), which has been derived from the approximate formula of van Regemorter~(\\cite{vanregemorter62}) adopting an effective Gaunt coefficient of 0.25.  For forbidden transitions an effective collision strength of $\\Upsilon_{ij}=0.05$ (for excitation) has been arbitrarily adopted.  The collision rate coefficient is computed from the effective collision strength via what Allende Prieto~et~al.\\ have referred to as the Eissner-Seaton formula (Eissner \\& Seaton~\\cite{eissner74}). We calculated rate coefficients at 5000~K, and these along with the ratio of these values to our values are presented in Table~\\ref{tab:allende}.  We see for allowed transitions, the ratios range from about 0.6 to 3.0.  As mentioned, in $LS$-coupling transitions between singlet and quintet states, and the $^1S \\rightarrow ^3S^o$ transition, are not collisionally allowed and thus we cannot make a comparison here.  For the remaining forbidden radiative transitions, the ratios range between roughly 0.003 and 60 indicating much poorer agreement in general.  Based on the evidence here, it appears that the approximate formula based on van Regemorter's formula gives at best order of magnitude estimates, while the simple approximation based on an arbitrarily chosen constant $\\Upsilon_{ij}$, is considerably less reliable.  We also compared with the data in the model atom prepared by Carlsson \\& Judge~(\\cite{carlsson93}) at similar temperatures.  The first notable aspect of the comparison is that for several transitions no collisional coupling was included due to lack of data.  For the transitions $^3P$$\\rightarrow$$^1D$, $^1S$ and $^1D$$\\rightarrow$$^1S$ they used data from a compilation by Mendoza~(\\cite{mendoza81}), which in fact originate from Berrington \\& Burke~(\\cite{berrington81}).  At 5000~K, the corresponding excitation rate coefficients are $1.82\\times10^{-11}$, $1.27\\times10^{-14}$ and $1.04\\times10^{-11}$~cm~s$^{-1}$ respectively. The ratios with respect to our values are roughly 1.3, 0.4 and 0.8 respectively.  For the $^3P$$\\rightarrow$$^5S^o$, $^3S^o$ transitions, they calculated collision rates from the cross sections of Rountree~(\\cite{rountree77}).  The cross sections of Rountree are significantly larger at the threshold than modern calculations, including our results, resulting in their rates being a factor of about 2.5 greater than ours.  For the $^5S^o$$\\rightarrow$$^5P$ and $^3S^o$$\\rightarrow$$^3P$ transitions, Carlsson \\& Judge computed rates using the impact parameter method of Seaton~(\\cite{seaton62}).  The rates are about a factor of five and two greater than ours, respectively.  \\begin{table*} \\begin{center} \\caption{Rate coefficients $\\langle \\sigma v \\rangle$ in units of cm$^3$~s$^{-1}$ for electron-impact excitation of neutral oxygen for $T=5000$~K calculated using approximate formulae following Allende~Prieto~et~al.~(\\cite{allende03}), and the ratios of these rate coefficients to our rate coefficients in Table~\\ref{tab:ratecoeffs} ($\\langle \\sigma v \\rangle_\\mathrm{approx} / \\langle \\sigma v \\rangle_\\mathrm{this\\,work}$).  Note, $a(b) = a \\times 10^b$.} \\label{tab:allende} \\begin{tabular}{lrrrrrr} \\hline \\hline Initial & \\multicolumn{6}{c}{Final state} \\\\ state  & 2p$^4$ $^1D  $&2p$^4$ $^1S  $&    3s $^5S^o$&    3s $^3S^o$&    3p $^5P  $&     3p $^3P  $             \\\\ \\hline & \\multicolumn{6}{c}{\\underline{$\\langle \\sigma v \\rangle$ from approximate formulae}}\\\\ 2p$^4$ $^3P  $&   $7.04(-12)$&   $4.04(-14)$&   $4.06(-19)$&   $4.91(-19)$&   $1.00(-20)$&    $5.62(-21)$             \\\\ 2p$^4$ $^1D  $&              &   $7.01(-12)$&   $7.04(-17)$&   $2.94(-17)$&   $1.74(-18)$&    $9.75(-19)$             \\\\ 2p$^4$ $^1S  $&              &              &   $6.13(-14)$&   $2.56(-14)$&   $1.51(-15)$&    $8.50(-16)$             \\\\ 3s $^5S^o$&              &              &              &   $5.11(-10)$&   $2.18(-08)$&    $1.69(-11)$             \\\\ 3s $^3S^o$&              &              &              &              &   $1.20(-10)$&    $3.79(-08)$             \\\\ 3p $^5P  $&              &              &              &              &              &    $2.29(-10)$             \\\\ &&&&&&\\\\ & \\multicolumn{6}{c}{\\underline{$\\langle \\sigma v \\rangle_\\mathrm{approx} / \\langle \\sigma v \\rangle_\\mathrm{this\\,work}$}}\\\\ 2p$^4$ $^3P  $&   $5.10(-01)$&   $1.28(+00)$&   $3.10(-01)$&   $6.48(-01)$&   $6.48(-01)$&    $4.43(-01)$             \\\\ 2p$^4$ $^1D  $&              &   $5.53(-01)$&             -&   $4.79(+01)$&             -&    $9.77(+00)$             \\\\ 2p$^4$ $^1S  $&              &              &             -&             -&             -&    $5.98(+01)$             \\\\ 3s $^5S^o$&              &              &              &   $1.31(-02)$&   $2.14(+00)$&    $2.02(-02)$             \\\\ 3s $^3S^o$&              &              &              &              &   $1.09(-02)$&    $2.94(+00)$             \\\\ 3p $^5P  $&              &              &              &              &              &    $3.92(-03)$             \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}  {\\ch Zatsarinny \\& Tayal~(\\cite{zatsarinny03}) have published effective collision strengths based on calculations using the $B$-spline $R$-matrix approach with non-orthoganal orbitals.  They provide data for transitions between the $^3P$, $^1D$ and $^1S$ states of the ground configuration and from these states to 23 other excited states.  We compare our results with theirs at 5000~K in Table~\\ref{tab:zt}.  The data are in good agreement, agreeing within a factor of 3 at worst, often much better, the mean ratio $\\langle \\sigma v \\rangle_\\mathrm{ZT} / \\langle \\sigma v \\rangle_\\mathrm{this\\,work}$ being 1.23 with a standard deviation of 0.67.}  Estimating the errors in the computed rate coefficients is difficult.  The generally good agreement of our cross section results with experiment, particularly at low energy, would indicate that our rate coefficients are of quite acceptable accuracy.   Errors due to neglect of coupling to the continuum and effects of pseudo-resonances at energies above the ionisation threshold are expected to be small as these energies hardly contribute to the rate coefficient for the temperatures of interest.  {\\ch Perhaps the best estimate of the error is obtained from the scatter among different $R$-matrix calculations where available. From the comparison with the rates of Zatsarinny \\& Tayal~(\\cite{zatsarinny03}) and the rates of Berrington \\& Burke~(\\cite{berrington81}) we would estimate that the rate coefficients have errors of the order of 70\\% at the temperatures of interest.}   \\begin{table*} \\begin{center} \\caption{{\\ch Rate coefficients $\\langle \\sigma v \\rangle$ in units of cm$^3$~s$^{-1}$ for electron-impact excitation of neutral oxygen for $T=5000$~K from Zatsarinny \\& Tayal~(\\cite{zatsarinny03}) for transitions from the ground configuration, and the ratios of these rate coefficients to our rate coefficients in Table~\\ref{tab:ratecoeffs} ($\\langle \\sigma v \\rangle_\\mathrm{ZT} / \\langle \\sigma v \\rangle_\\mathrm{this\\,work}$).  Note, $a(b) = a \\times 10^b$.}} \\label{tab:zt} \\begin{tabular}{lrrrrrr} \\hline \\hline Initial & \\multicolumn{6}{c}{Final state} \\\\ state  & 2p$^4$ $^1D  $&2p$^4$ $^1S  $&    3s $^5S^o$&    3s $^3S^o$&    3p $^5P  $&     3p $^3P  $             \\\\ \\hline & \\multicolumn{6}{c}{\\underline{$\\langle \\sigma v \\rangle$ from Zatsarinny \\& Tayal}}\\\\ 2p$^4$ $^3P  $&   $2.05(-11)$&   $1.16(-14)$&   $1.94(-18)$&   $1.15(-18)$&   $3.42(-20)$&    $2.86(-20)$             \\\\ 2p$^4$ $^1D  $&              &   $7.79(-12)$&         $ 0.$&   $3.45(-19)$&         $ 0.$&    $7.29(-20)$             \\\\ 2p$^4$ $^1S  $&              &              &         $ 0.$&         $ 0.$&         $ 0.$&    $1.63(-17)$             \\\\ &&&&&&\\\\ & \\multicolumn{6}{c}{\\underline{$\\langle \\sigma v \\rangle_\\mathrm{ZT} / \\langle \\sigma v \\rangle_\\mathrm{this\\,work}$}}\\\\ 2p$^4$ $^3P  $&   $1.49(+00)$&   $3.67(-01)$&   $1.48(+00)$&   $1.52(+00)$&   $2.21(+00)$&    $2.25(+00)$             \\\\ 2p$^4$ $^1D  $&              &   $6.15(-01)$&             -&   $5.62(-01)$&             -&    $7.31(-01)$             \\\\ 2p$^4$ $^1S  $&              &              &             -&             -&             -&    $1.14(+00)$             \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}  It is difficult to predict the impact that these new data will have on the results of non-LTE calculations in stellar atmospheres.  The impact on abundance analysis of solar-type stars, including metal-poor stars, is presently being investigated by Fabbian~et~al.~(in preparation).  It would also be of interest to investigate the impact on predicted solar centre-to-limb behaviour of the oxygen triplet lines (see e.g.\\ Allende Prieto~et~al.~\\cite{allende04}).  {\\ch We have presented estimates of cross sections and rate coefficients for all transitions between the seven lowest $LS$ states of neutral oxygen.  It should be noted that the calculations produce data for transitions between all target states listed in Table~\\ref{tab:states}.  We restricted the data presented here to the lowest seven states for two reasons.  First, the quality of the results for transitions involving higher states is certainly lower.  In particular, the effect of pseudo-resonances on the rate coefficients would be expected to be larger, since the thresholds are nearer the ionisation energy.  Secondly, the transition rates among the low-lying states are expected to be most important for the non-LTE problem of interest, in particular the modelling of the 777~nm triplet 3s $^5S^o \\rightarrow$ 3p $^5P$.  However, the computed rates for the transitions involving the higher states should give at least order of magnitude estimates and may be obtained on request from the author.}  More calculations and experimental data for electron-impact excitation of oxygen would be valuable, particularly from excited states where few or no other comparisons are available.  {\\ch In particular calculations for transitions between excited states using the $B$-spline $R$-matrix method with non-orthogonal orbitals, with more careful consideration of short range correlation and long range polarisation effects would be important.}  Calculations for excitation by hydrogen collisions are also needed.  Recent work by Krems~et~al.~(\\cite{krems06}) calculated rates for a single transition, but again {\\ch as for electrons}, for stellar atmospheres applications, data are {\\ch needed} for more low-lying states."
    },
    "0609/astro-ph0609367_arXiv.txt": {
        "abstract": "{We present our analysis of the extensive monitoring of SS433 by the RXTE observatory collected over the period 1996-2005. The difference between energy spectra taken at different precessional and orbital phases shows the presence of strong photoabsorption ($N_{\\rm H}>10^{23}$\\,cm$^{-2}$) near the optical star, probably due to its powerful, dense wind. Therefore the size of the secondary deduced from analysis of X-ray orbital eclipses might be significantly larger than its Roche lobe size, which must be taken into account when evaluating the mass ratio from analysis of X-ray eclipses. Assuming that a precessing accretion disk is geometrically thick, we recover the temperature profile in the X-ray emitting jet that best fits the observed precessional variations in the X-ray emission temperature. The hottest visible part of the X-ray jet is located at a distance of $l_0/a\\sim0.06-0.09$, or $\\sim2-3\\times10^{11}$cm from the central compact object, and has a temperature of about $T_{\\rm max}\\sim30$ keV. We discovered appreciable orbital X-ray eclipses at the ``crossover'' precessional phases (jets are in the plane of the sky, disk is edge-on), which under model assumptions put a lower limit on the size of the optical component $R/a\\ga0.5$ and an upper limit on a mass ratio of binary companions $q=M_{\\rm x}/M_{\\rm opt}\\la0.3-0.35$, if the X-ray opaque size of the star is not larger than $1.2R_{\\rm Roche, secondary}$. ",
        "introduction": "The source SS433 is the only Galactic X-ray binary with X-ray emission from optically thin thermal plasma originating in hot jets outflowing with a sub-relativistic velocity of 0.26 $c$. The binary system is thought to consist of a compact object (probably a black hole) accreting matter at a super-Eddington rate from a high-mass star filling its Roche lobe (\\citealt{margon84}, \\citealt{cherep02}, see \\citealt{fabrika04} for a recent review). Apparently, the innermost accretion flow and the supercritical accretion disk wind completely screen the region of the main energy release, and most of the energy from SS433 is observed in the optical and UV (Cherepashchuk et al., 1982; Dolan et al., 1997).  The system demonstrates a complex variability including periodicities - the precessional ($\\sim$162 days), orbital ($\\sim$13 days), and nutational (nodding,$\\sim 6$ days) periods - which can be used to tightly constrain the binary system parameters. In particular, the kinematical model of the system leads to a binary inclination angle  of $i\\sim 78.05^\\circ$, and the jet precession angle is $\\sim 20.92^\\circ$ \\cite{fabrika04}.   Mildly relativistic ($v\\sim0.26$c) jets launched in the vicinity of the compact object consist of protons and electrons (``heavy'' jet) with a high temperature (1-30 keV). Plasma moving along the jet gradually cools down so that, at distances $\\sim 10^{13}-10^{14}$ cm from the central object, a thermal instability develops and clumps start forming. At distances $10^{14}-10^{15}$ cm, the temperature of matter drops to $\\sim2\\times10^{4}$K so that optical line emission appears. At larger distances the jets shine in the radio diapason.  The main properties of the X-ray emission observed from SS433 can be summarized as follows:  \\begin{itemize}  \\item The X-ray emission mechanism is thermal bremsstrahlung radiation \\citep[e.g.][]{marshall79,w96,matsuoka86,brinkmann91}. Based on measurements of the doppler shifts of the emission lines from the thermal plasma, it is believed that the total (or majority of) the X-ray emission originates in the jets \\citep{fabrika04}.   \\item The temperature of optically thin plasma decreases along the jet. The hottest regions near the jet base have temperatures up to $\\sim$20-30 keV \\citep[e.g.][]{kotani96,marshall02,cherep03,namiki03}.  \\item A detailed study of X-ray emission lines in SS433 shows that the full opening angle of the jet should be very small, not larger than $\\sim 1-2^\\circ$. Such an opening angle roughly corresponds to free expansion of the jet material in the direction perpendicular to the jet motion \\citep{marshall02,namiki03}.  \\item There are X-ray eclipses caused by the donor star. The temperature of the observed radiation drops significantly during the eclipse, and the depth of the eclipse increases at high energies \\citep{stewart87,kawai89,brinkmann91}.  \\item Precession of the geometrically thick accretion disk is accompanied by strong variations in both temperature and the flux of X-ray emission \\citep{yuan95}. Amplitudes of precessional and orbital variations in X-ray emission strongly increase with energy \\citep[e.g.][]{cherep05}. The precession variations are likely to be due to partial eclipse of the X-ray jets by the thick precessing accretion disk.  \\end{itemize}   Over the life time of the RXTE observatory many observations of SS433 were performed at different precessional and orbital phases (see e.g.\\citealt{gies02a,nandi05}). The relatively complete sampling of the source at different orbital and precessional phases provides us with an opportunity to perform a tomographic study of the X-ray jet in SS433. Examining of orbital eclipses (caused by the optical star) and precessional variability (caused by the geometrically thick accretion disk) allows us to separate contributions of different parts of the jet to the total X-ray emission. ",
        "conclusions": "X-ray emission of SS433 demonstrates both chaotic (aperiodic) and periodic variations. In this paper, we used extensive observational data obtained by RXTE to studing systematic variations in the X-ray emission of SS433 caused by precessional and orbital motions in the system. \\begin{itemize} \\item By comparing X-ray spectra of SS433 near and in the eclipse, we obtained strong signatures of a significant ($N_H>10^{23}$ cm$^{-2}$) photoabsorption of X-rays near the companion star. We argue that this might be caused by the presence of a dense stellar wind from the companion star. The mass loss rate in the wind needed to explain the observed photoabsortion is $\\dot{M}\\sim 10^{-6}$ M$_{\\odot}$/yr.  \\item From detection of strong photoabsorption near the secondary star, we conclude that the X-ray opaque size of the secondary might be significantly larger than the Roche lobe size of the star. Therefore one should be cautious when estimating the Roche lobe size of the secondary from X-ray orbital eclipses.    \\item By both assuming the shape of the geometrically-thick accretion disk restricted by the Roche lobe size of the compact star and using the RXTE observations of SS433 at different precessional phases, we recovered the temperature profile of plasma along the jet. The obtained maximum temperature visible in the jet is $T\\sim30$ keV at a distance of $l/a=0.06-0.09$ from the compact object. The radius of the jet at this distance is $r_0/a=0.01-0.016$. Here the jet was assumed to have a constant opening angle.  \\item We reliably detected orbital X-ray eclipses in SS433 during the ``crossover'' precessional phases when the X-ray jets and the axis of the accretion disk lie exactly on the plane of the sky. The observed depth of the X-ray eclipse combined with our assumptions about the accretion disk thickness implies that the size of the star is larger than $R/a\\ga0.5$, yielding an upper bound on the mass ratio of the components in SS433 $q<0.3-0.35$, assuming that the radius of the eclipsing region (star plus inner wind) cannot be much larger than $1.2R_{\\rm Roche, secondary}$.   \\end{itemize}"
    },
    "0609/astro-ph0609826_arXiv.txt": {
        "abstract": "% Ultra Luminous Infra-Red Galaxies (ULIRGs) -- gas-rich mergers harboring the most extreme star-forming conditions encountered in the local Universe -- are thought to re-enact the galaxy formation processes we are only barely able to glimpse in the distant Universe. Combining new single-dish molecular line observations of $^{12}$CO, $^{13}$CO, HCO$^+$, HCN, and CS towards the two ULIRGs Arp\\,220 and NGC\\,6240 with existing data in the literature, we have compiled the most extensive molecular line data-sets to date of such galaxies. The data allow us to put strong constraints on the properties of the dense star forming gas in these two systems, and compare the relative usefulness of CS, HCN and HCO$^+$ as tracers of dense gas. In addition, we have build molecular line templates based on our observations, and demonstrate that Herschel/HI-FI will be able to detect the high-$J$ transitions of most of the above molecules in a large sample of ULIRGs out to $z\\ls 0.5$, assuming Arp\\,220 and NGC\\,6240 are representative of the ULIRG population at these redshifts. ",
        "introduction": "% The coming years will witness the emergence of a new generation of ground-based and space-born mm/sub-mm telescopes with sensitivities and frequency coverage that will allow us to study the interstellar medium (ISM) in high-$z$ proto-galaxies in unprecedented detail. Unfortunately, our expectations of how such objects will look in mm/sub-mm line emission are ill-informed due to the dearth of similar data locally. Virtually all studies of local (U)LIRGs to date have been of low-$J$ CO transitions or of single HCN/HCO$^+$ lines, and as a result we know very little about the properties of their dense, star-forming gas phase. While pioneering multi-line studies of low-intensity starbursts taught us the value of having many independent dense gas tracers (Nguyen-Q-Rieu et al.\\ 1992; Jackson et al.\\ 1995; Paglione, Jackson \\& Ishizuki 1997), the relatively benign star formation rates of such galaxies suggest they may not be optimal guides for the state of the ISM in (U)LIRGs and high-$z$ starbursts. In the local Universe, Arp\\,220 and NGC\\,6240 are ideally suited for studies of the dense gas in extreme starburst galaxies as they have the most well-sampled FIR/sub-mm spectral energy distributions (SEDs) of their category (Dopita et al.\\ 2005) and a large number of rotational lines of CO -- and some of HCN, CS and HCO$^+$ -- have already been observed. With this in mind, we used the 15-m James Clerk Maxwell Telescope and the IRAM 30-m Telescope, to observe HCN, CS, and HCO$^+$ molecular line emission towards these two sources (Greve et al.\\ 2006). ",
        "conclusions": ""
    },
    "0609/astro-ph0609017_arXiv.txt": {
        "abstract": "An open question remains whether Ultraluminous X-ray Sources (ULXs) really contain intermediate-mass black holes (IMBHs). We carefully investigated the XMM-Newton EPIC spectra of the four ULXs that were claimed to be strong candidates of IMBHs by several authors. We first tried fitting by the standard spectral model of disk blackbody (DBB) + power-law (PL), finding good fits to all of the data , in agreement with others. We, however, found that the PL component dominates  the DBB component at $\\sim$ 0.3 to 10 keV.  Thus, the black hole parameters derived solely from the minor DBB component are questionable. Next, we tried to fit the same data by the ``$p$-free disk model'' without the PL component, assuming the effective temperature profile of $T_{\\rm eff} \\propto r^{-p}$ where $r$ is the disk radius. Interestingly, in spite of one less free model parameters, we obtained similarly good fits with much higher innermost disk temperatures, $1.8 < kT_{\\rm in} < 3.2$ keV. More importantly, we obtained $p \\sim 0.5$, just the value predicted by the slim (super-critical) disk theory, rather than $p = 0.75$  that is expected from the standard disk model. The estimated black hole masses from the $p$-free disk model are much smaller; $M\\ltsim 40 M_\\odot$. Furthermore, we applied a more sophisticated slim disk model by Kawaguchi (2003, ApJ, 593,69), and obtained good fits with roughly consistent black hole masses. We thus conclude that the central engines of these ULXs are super-critical accretion flows to stellar-mass black holes. ",
        "introduction": "X-ray data analysis has become one of the most widely studied subjects in these days, % and has made great contributions in the progress of astrophysics. This is particularly true in the field of astrophysical black holes. Thanks to rapid improvement of the sensitivity of X-ray detectors, new types of X-ray sources have been discovered. Some of them seem to be black holes which were not known before.  Recent X-ray observations on many nearby spiral galaxies showed that the X-ray radiation consists of emission from several discrete X-ray sources, such as hot gaseous media, active galactic nucleus (AGN) and the point-like off-center X-ray sources whose X-ray luminosities significantly exceed the Eddington luminosity of a neutron star (\\cite{15}). The bright point-like off-center X-ray sources which cannot be identified with young supernova remnants are known as ultraluminous X-ray sources (ULXs; e.g., \\cite{25}; \\cite{34}). The ULXs, despite the fact of being very luminous ($L_{\\rm x}\\sim 10^{39 - 41}{\\rm erg} {\\rm s}^{-1}$), cannot be explained as collections of many sources with each luminosity less than the Eddington limit, since many of the ULXs exhibit significant time variabilities (e.g., \\cite{25}). Therefore, one reasonable assumption is that the ULXs are single compact objects powered by accretion flow. Then, how can we understand the extremely large luminosities exceeding the Eddington luminosity for a mass of $\\sim $10 $M_\\odot$?  A possibility that ULXs are only bright toward our line of sights (beaming model) was proposed  (e.g. \\cite{22}, \\cite{23}, \\cite{35}), but is not generally accepted now. In fact, Wang (2002) has shown the existence of a bright nebula surrounding the ULX M81 X-9,  and some ULXs also exhibit extended optical nebulae (\\cite{51}). There is an evidence that some ULXs are associated with a spreading wave of star formation,  which supports the idea that many ULXs, in particular those located in star forming galaxies, are high mass X-ray binaries (\\cite{21}). These facts indicate that the ULX surrounding nebulae are  powered by the central high mass X-ray sources,  and that strong X-ray beaming toward our direction is  unlikely.   Then, the question is, how massive the central compact objects  are. At present, there are basically two distinct lines of thoughts: sub-Eddington accretion onto an intermediate (several hundreds $M_\\odot$) mass black hole (IMBH), and super-Eddington (super-critical) accretion onto a stellar-mass black hole. The success in fitting several ULXs spectra with a multicolor disk blackbody (DBB) and a power-law (PL) model  supports the idea of IMBHs for the ULXs (\\cite{29}, \\cite{30}; \\cite{9}, \\cite{36}), since the cool inner disk temperatures and  the large innermost radii, obtained through model fitting, suggests a black hole mass within the IMBH range. The other line of thoughts, supporting a super-Eddington accretion flow, is the slim disk model with stellar mass black holes (\\cite{48}; \\cite{20}, \\cite{14}; \\cite{33}). The present article concerns with the super-Eddington thick accretion disks introduced in the early 1980s by the Warsaw group and their collaborators (e.g. \\cite{1}; \\cite{2}) and with the slim accretion disks, introduced in the late 1980s by the Warsaw and the Kyoto groups (\\cite{3}, \\cite{4}). In these papers, the very possibility of the super-Eddington beaming in the thick accretion disk funnels, crucial for arguments presented here, was explicitly recognized, discussed and stressed.   The IMBH notion faces a serious problem that the formation of IMBH itself remains unsettled (e.g., \\cite{52}). In contrast, the alternative possibility of the super-critical accretion has been poorly investigated so far due mainly to our limited knowledge about properties of the super-critical flow. The situation has been remarkably improved in the past few years, however, since the basic tools for investigating high-luminosity accreting systems are now available. In the present paper, we investigate the XMM-Newton EPIC spectra of the four ULXs which had been claimed as strong IMBH candidates, using theoretical super-critical models to see whether the super-critical accretion takes place in these ULXs or not. Our final goal is to determine the mass of the central black holes and settle down the issue regarding the origin of ULXs.   We applied accretion disk spectral models to the observed ULX energy spectra, constrain the innermost disk radius, and estimate the black hole mass. A very similar method as the one used in the present paper was successfully applied to Galactic black hole binaries (e.g., \\cite{12}; \\cite{13}; \\cite{17}). Several authors have recently done elaborate mass and angular momentum estimates by spectral fitting for microquasars, i.e. by a very similar method as the one used in the present paper (\\cite{38}; \\cite{27}; \\cite{10}). In the present study, however, we focus on the mass estimation and do not attempt to estimate the spin of the black hole.  The mass of the ULXs can also be accurately estimated by a different method than that described in our article. For this other independent estimate, double peak QPOs in the 3:2 resonance should be detected in ULXs (\\cite{5}). This is because these QPOs scale inversely with the mass of the source, which was first shown for microquasars by McClintock and Remillard (2003), and more recently for low-mass Seyfert galaxies by Lachowicz et al. (2006). There is also an indication that the same scaling is true for a (possible) detection of double peak 3:2 QPOs in SgrA* (\\cite{41}), while for the ULXs, in general, it has not been detected yet. According to Mucciarelli et al. (2006), the only ULX where a QPO has been discovered at present is M82 X-1. Scaling the frequency inversely to the BH mass, the observed QPO frequency range (from a few tens to a few hundreds mHz) would yield a black hole mass anywhere in the interval from a few tens to a few thousands solar mass (\\cite{40}; \\cite{32}).  The Plan of this paper is as follows: We first describe the X-ray data which we used in the present study and our methods of fitting in Section 2. In Section 3, we then give the results of fitting, together with their implications on the black hole mass and mass accretion rates. Section 4 is devoted to discussion, and the final Section 5 concludes the paper. ",
        "conclusions": "We have investigated the XMM-Newton EPIC spectra of the four ULXs which were claimed to be strong candidates of IMBHs by several authors. We found that these spectra are successfully fitted with slim disk models, and the derived masses are $\\sim $11 to 60 $M_\\odot$ depending on different assumptions.  We have not seen the evidence of the ``intermediate mass'' black holes having several hundreds of $M_\\odot$. When the mass accretion rate is close to or exceeding the critical rates, the slim disk is theoretically predicted. We suggest that ULXs are stellar mass black holes, at most several tens of $M_\\odot$ under super-critical accretion rates shining at super-Eddington luminosities.    \\bigskip \\bigskip We gratefully thank Marek Abramowicz for valuable comments and suggestions. This work was supported in part by the Grants-in-Aid of the Ministry of Education, Science, Culture, and Sport  (14079205, 16340057 S.M.), by the Grant-in-Aid for the 21st Century COE ``Center for Diversity and Universality in Physics'' from the Ministry of Education, Culture, Sports, Science and Technology (MEXT) of Japan. K.V. gratefully  thank Ken-ya Watarai for the discussions at some beginnings of this work and also MEXT scholarship. T.K. thanks the financial supports from JSPS Postdoctoral Fellowship (01879)."
    },
    "0609/astro-ph0609221_arXiv.txt": {
        "abstract": "{The X--ray burster \\src, located in the Galactic globular cluster \\glob, is an ultracompact binary (orbital period $\\sim$21 minutes), likely harbouring a degenerate companion.} {The source has been observed at soft $\\gamma$-rays several times with the \\inte\\ satellite, during the monitoring of the Galactic plane, with an unprecedented exposure time. We analysed all available \\inte\\ observations, with the main aim of studying the long-term behaviour of this Galactic bulge X--ray burster.} {The spectral results are based on the systematic analysis of all \\inte\\ observations covering the source position performed between March 2003 and November 2005.} {The source X--ray emission is hard and is observed, for the first time, up to 100~keV. A broad-band  spectrum obtained combining the  \\inte\\ spectrum together with a quasi-simultaneous \\xmm\\ observation performed in September 2003 is well modeled with a disk-blackbody emission (with an inner disk temperature of $\\sim$0.8~keV) together with a power-law (with a photon index, $\\Gamma$, of 2.1). The 2--100~keV luminosity is 1.5$\\times$10$^{36}$~erg~s$^{-1}$ (assuming a distance of 6.8~kpc). } {\\inte\\ observations  reveal for the first time that this X--ray burster displays a very hard X--ray spectrum, with a cut-off at energies higher than 100~keV, and that the source spends most of the time in this low luminosity and hard state. Indeed, a previous \\sax\\ observation in April 1997 observed high energy emission from \\src\\ only up to 50~keV. } ",
        "introduction": "\\src\\ (Swank et al. 1976) is an X--ray burster located in the galactic globular cluster \\glob. The likely optical counterpart displays an UV modulation (Anderson et al.~\\cite{a:93}) with a period of 20.6 minutes (interpreted as the orbital period), implying a degenerate companion of 0.04$M_\\odot$ (Homer et al.~\\cite{h:96}).   The first broad-band spectrum from this source was obtained with \\sax\\ in the energy range 0.3--50~keV during a pointing performed in 1997 (Sidoli et al. \\cite{s:01}). The spectrum was well described with a disk-blackbody and Comptonized continuum with $N_{\\rm H}$ = $3.9 \\times 10^{21}$~\\hcm, an inner disk temperature, $kT_{\\rm in}$, of $\\sim$0.6 keV, an inner projected radius of $\\sim$5~km (for an assumed \\glob\\ distance of 6.8~kpc, Harris~\\cite{ha:96}), a temperature, $kT_{0}$, of the input ``seed'' photons, of 0.8$^{+0.1} _{-0.2}$~keV, an electron temperature, $kT_{\\rm e}$, of 70~keV, and an optical depth, $\\tau$, of 1.7. The estimated 0.1--100~keV luminosity was 1.9$\\times10^{36}$~erg~s$^{-1}$ (at 6.8~kpc).  Other X--ray observations are reported for \\src\\ limited to the energy range below 10~keV ($ASCA$, Juett et al.~\\cite{j:01}; \\xmm, Sidoli et al.~\\cite{s:05}; $Chandra$, Juett \\& Chakrabarty~\\cite{j:05}).  We report here the first detection of \\src\\ above 50~keV with the \\inte\\ satellite, obtained with an unprecedented exposure time collected during the monitoring of the Galactic plane. ",
        "conclusions": "Previously, the source broad-band spectrum was observed only with \\sax\\ in April 1997. At that time, the PDS non-imaging instrument, covering the range from 20 to 200 keV, detected the source only up to 50~keV (Sidoli et al. 2001).  We report here for the first time the discovery of hard (50--100~keV) X--ray emission from the X--ray burster \\src\\ and a long-term study of its X--ray spectral behaviour. It is now possible to directly measure the high energy luminosities of this source: they are $\\sim$1.3$\\times$10$^{36}$~erg~s$^{-1}$ and 7$\\times$10$^{35}$~erg~s$^{-1}$ respectively in the 1--20~keV and 20--200~keV energy ranges (at 6.8~kpc). This implies that \\src\\ falls into the so-called ``burster box'' (Barret et al., 2000).   In order to allow a proper comparison with the \\inte\\ spectrum we re-analysed the \\sax\\ observation (April 1997; refer to Sidoli et al. 2001 for the details of the data reduction) adopting the same model used here, a disk-blackbody emission together with a power-law. This resulted in the following parameters ($\\chi^2$/dof=170.4/158): an absorbing column density of (0.46$\\pm{0.03}$)$\\times$10$^{22}$~cm$^{-2}$, an inner disk blackbody temperature, kT$_{\\rm in}$, of 0.66$\\pm{0.03}$~keV, and a powerlaw photon index of 1.96$\\pm{0.06}$. Thus, there is a good agreement with the \\inte\\ spectroscopy.  The new observations with \\inte\\ allow us to put much more stringent limits on the presence of a high energy cut-off. In Fig.~\\ref{fig:cutoff} we compare the confidence contour levels for the high energy cut-off and the power-law photon index, obtained with the EPIC--IBIS joint spectrum (solid contours) and with \\sax\\ (dashed contours). During the \\sax\\ observation a lower limit to the high energy cut-off could be placed at E$_{c}$$>$60~keV (90\\% level), while now with \\inte\\ we can interestingly shift it towards much higher energies, with  E$_{c}$$>$100~keV. This allows us to include \\src\\ among the few tens of low-luminosity low-mass X--ray binaries  (LMXRBs) for which the spectrum has been observed up to 100~keV (Barret et al. 2000, Di Salvo \\& Stella 2002, Bazzano et al. 2006). More interestingly, \\src\\ is now among the hardest type I X--ray bursters in our Galaxy.  The \\inte\\ spectrum displays properties which fits well into the classification of low luminosity ($\\sim$0.01 L$_{\\rm Edd}$) weakly magnetized neutron stars with ``hard spectra'' (Barret et al. 2001): a broad-band spectrum extending up to 100~keV, well described by a soft disk-blackbody emission together with a hard Comptonized component, with a optical depth of the Comptonizing corona of $\\sim$2, seed photons temperature below 1~keV, a high electron temperature, above 40~keV. A few LMXRBs appear to spend most of the time in this ``hard state'' (Di Salvo \\& Stella 2002), like 4U~0614+091 (Piraino et al. 1999), which also displays spectral parameters similar to \\src, with a photon index in the range 2.3--2.4. Also for \\src\\ the four different \\inte\\ spectra suggest that the source spend most of its lifetime in this spectral state (on  timescales of months or years, if compared with \\sax).  The very hard emission from \\src\\ is clearly evident also when compared (see Fig.~\\ref{fig:comp}) to a typical spectrum of a Z-source, GX~349+2, reported in an \\inte\\ study of the spectral behaviour of the persistent LMXRBs hosting a neutron star (Paizis et al. 2006).    \\begin{figure}[ht!] \\centering \\includegraphics[angle=-90,width=8.0cm]{5759fig4.ps} \\caption{Comparison of the confidence contour levels for the high energy cutoff (in units of keV): solid contours have been derived analysing EPIC--IBIS joint spectrum, while the dashed contours mark the \\sax\\ results. } \\label{fig:cutoff} \\end{figure}    \\begin{figure}[ht!] \\centering \\includegraphics[angle=-90,width=10.0cm]{5759fig5.ps} \\caption{Comparison of the \\inte\\ spectrum of the source \\src\\ (data-set 2, the same dispayed in Fig.~\\ref{fig:broad}), with a spectrum of a typical Z-source, GX~349+2 (data from Paizis et al. 2006). } \\label{fig:comp} \\end{figure}"
    },
    "0609/astro-ph0609292_arXiv.txt": {
        "abstract": "{Rotational evolution in the pre-main sequence is described with new sets of pre-MS evolutionary tracks including rotation, non-gray boundary conditions (BCs) and either low (LCE) or high convection efficiency (HCE). } {Using observational data and our theoretical predictions, we aim at constraining (1) the differences obtained for the rotational evolution of stars within the ONC by means of these different sets of new models; (2) the initial angular momentum of low mass stars, by means of their templates in the ONC. } {We discuss the reliability of current stellar models for the pre-MS. While the 2D radiation hydrodynamic simulations predict HCE in pre-MS, semi-empirical calibrations either seem to require that convection is less efficient in pre-MS than in the following MS phase (lithium depletion) or are still contradictory (binary masses). We derive stellar masses and ages for the ONC by using both LCE and HCE. } {The resulting mass distribution for the bulk of the ONC population is in the range 0.2$-$0.4{\\msun} for our new non-gray models and,  as in previous analyse, in the range 0.1$-$0.3{\\msun} for models having gray BCs. In agreement with Herbst et al.\\ (2002) we find that a large percentage ($\\sim$70\\%) of low-mass stars (M\\simlt0.5{\\msun} for LCE; M\\simlt0.35{\\msun} for HCE) in the ONC appears to be fast rotators (P$<$4days). Three possibilities are open: 1) $\\sim$70\\% of the ONC low mass stars lose their disk at early evolutionary phases; 2)their ``locking period'' is shorter; 3) the period evolution is linked to a different morphology of the magnetic fields of the two groups of stars. We also estimate the range of initial angular momentum consistent with the observed periods. } {The comparisons made indicate that a second parameter is needed to describe convection in the pre-MS, possibly related to the structural effect of a dynamo magnetic field.}  \\keywords { Stars: evolution -- Stars: interiors -- Stars: rotation -- Stars: Hertzsprung Russell and C-M diagrams }  \\authorrunning {Landin et al.} \\titlerunning {Non-gray rotating stellar models and the evolution in the ONC} ",
        "introduction": "We compute new sets of pre-MS models with the ATON code for stellar evolution, the version presented in Mendes et al. (\\cite{mendes}), which includes stellar rotation according to the description by Endal \\& Sofia (\\cite{endal}), updated for the present work to employ non-gray boundary conditions by Allard \\& Hauschildt (\\cite{allard1}) and  Allard et al.\\ (\\cite{allard00}).  Our main goal is to improve our understanding of the appropriate physical constraints to be used for a general description of the evolution of stellar structure and its angular momentum with time. We are particularly interested on the choice of the stellar initial angular momentum and its variation with time, and, also, on the importance of convection efficiency during the pre-main sequence. To do this, we check our choices with respect to sets of relevant observations.  Since the pioneering work by Henyey et al. (\\cite{henyey}) and Hayashi (\\cite{hayashi}) it is commonly accepted that pre-MS stars derive their luminosity by gravitational contraction, with the only exception of the short deuterium burning phase. Derivation of masses of young stellar associations has then generally been made by standard hydrostatic stellar models including deuterium burning, an approach that we will use in this paper. This procedure rests on the assumption that neither the residual accretion after the protostellar phase nor the uncertainty in the zero point of ages affect the results in a strong way.  Nevertheless, the theoretical description of moderately low and low mass objects is affected by the first principle uncertainties in the description of some physical inputs, in particular opacities, convection, equation of state (D'Antona \\cite{franca5}) and treatment of boundary conditions (BCs, Chabrier \\& Baraffe \\cite{chabrier}). Since the low mass stars in the pre-MS are fully convective and over-adiabatic, any change in the convective model substantially alters the location of the track in the theoretical HR plane. The use of a less efficient treatment of convection leads to larger temperature gradients, so that, for a given luminosity, the structure readjusts on a more expanded configuration, with a consequent shift of the track to lower effective temperatures ($T_{\\rm eff}$) (e.g. D'Antona \\& Mazzitelli \\cite{franca1}, \\cite{franca3}; D'Antona \\& Montalb\\'an \\cite{franca2}). The path followed by the theoretical pre-MS tracks on the HR diagram is also dependent on the boundary conditions used to fit the numerical integration of the structural equations of the interior with the atmosphere. The use of a non-gray atmospheric treatment shifts the tracks to cooler $T_{\\rm eff}$s within an extended interval of masses and ages (Montalb\\'an et al.\\ \\cite{josefina}). The effect due to the non-grayness of the atmosphere is, in many cases, overwhelmed by the uncertainties related to the treatment of convection, that has a similar, but even stronger effect on the tracks, with only a few exceptions that are relevant for this work, as will be discussed in \\S \\ref{ngconv}.  Young stellar clusters provide a unique opportunity to test stellar pre% -MS models. Many studies in the past have been focused on the Orion Nebula Cluster (ONC), because it contains thousands of pre-MS objects. Hillenbrand (\\cite{hillen97}) measured the V and I magnitudes (and colors) of $\\sim$900 stars and located them in the theoretical HR diagram by using bolometric corrections and taking into account the interstellar extinction for the determination of colors and bolometric magnitudes. In spite of the non-% negligible uncertainties weighting on the derivation of these stellar parameters (Hillenbrand et al. \\cite{hillen97}), such a work has been widely used, in connection with theoretical pre-MS tracks, to infer important information concerning the cluster itself. The mass and age distributions and the slope of the mass function can all be estimated by inferring the appropriate mass and age for each observed star. It is clear that the results obtained with this approach will depend to a certain extent on the set of tracks used to perform the analysis, and on the physical inputs adopted to calculate the evolution. This is confirmed by the fact that, specifically for the ONC, studies that used different sets of evolutionary tracks reached significantly different conclusions, particularly with respect to the mass and age distribution, and, more important, to the age spread and, consequently, to the evolutionary history of the star formation process within the cluster (Palla \\& Stahler \\cite{palla3}).  In the past few years, new detailed observational studies of the ONC have been undertaken, focused on the rotational properties of the stars. Stassun et al.\\ (\\cite{stassun}) and Herbst et al.\\ (\\cite{herbst02}) measured the rotational periods of $\\sim$400 stars belonging to the ONC. All of them are in the Hillenbrand (\\cite{hillen97}) sample, so they can be located on the HR diagram. More recently, Stassun et al.\\ (\\cite{stassun04}) and Flaccomio et al.\\ (\\cite{flacco03a}, \\cite{flacco03b}) reanalyzed all the archival {\\it Chandra} observations of the ONC studying in great detail the X-ray properties of the observed objects, in an attempt to elucidate the origin of X-ray emission in pre-MS stars. All this available information renders the ONC an excellent laboratory to test stellar evolution theories of the pre-MS phase.  None of the previous analyses of the ONC rotational database has been done using non-gray models, and the effect of using different convection efficiencies has not been extensively tested. We therefore decided to use this database to test and calibrate our new sets of non-gray tracks for rotating stellar models, with masses in the range 0.085$\\leq$$M/M_{\\odot}$$\\leq$3.8. As a byproduct, we will have insight on how much the rotational properties of this population depend on the choice of the evolutionary tracks. At present, there is no definitive observational constraint that can be used to choose among the tracks obtained by using different convection inputs. In this work we derive masses and ages with different sets, in order to appreciate how much the results we are presently interested in, namely those on the angular momentum evolution, depend on the choice. In particular, we will test the role of (a) boundary conditions (gray and non-gray), (b) different convection efficiencies in the framework of the MLT, and (c) rotating and non-rotating models.  As indicated by the non-negligible differences among the set of tracks adopted, we show that the detailed period distribution as a function of mass and age is dependent on the physical inputs. However, the qualitative information on the rotational distribution of stars of different mass in the ONC remains similar, and we confirm that the distribution of periods is bimodal only for masses larger than a ``transition\" mass depending on the convection model.  In \\S \\ref{sect2} we summarize the uncertainties in the theoretical models, and also try to clarify some issues left ambiguous in the recent literature on pre-MS evolution. In \\S \\ref{models} we provide a brief description of the latest rotational, non-gray version of the ATON stellar evolution code, and describe the choice of the main physical and chemical inputs used for the present work in \\S \\ref{input}. The role played by the convective modeling on the location of pre-MS tracks on the HR diagram, the role played by non-grayness, and the lithium pre-MS depletion are discussed in \\S \\ref{sect5}. In \\S \\ref{sect6} we present the photometric and rotational data related to the ONC stellar population, and the derivation of masses and ages from our new sets of tracks. The results on the rotation period distribution and its possible interpretation(s) are discussed in \\S \\ref{poponc}. Conclusions are given in \\S \\ref{sect8}. ",
        "conclusions": "\\label{sect8} \\begin{figure*}[htb] \\centering{ \\includegraphics[width=5.9cm]{landinf26.eps} \\includegraphics[width=5.9cm]{landinf27.eps} \\includegraphics[width=5.9cm]{landinf28.eps} } \\caption{The fractional X-ray luminosity as a function of mass for the same stars and models shown in Fig.~\\ref{lxlbol}. \\label{lxlbrat}} \\end{figure*} We use the observed stellar population of the ONC to test our pre-MS evolutionary tracks, to better understand the main physical properties characterizing the evolution of young stars. By comparing the location of the tracks in the HR diagram with the position of the observed objects, we assign to any single star a mass and an age, for three different convection efficient models.  We find the well known result that the treatment of convection is generally the most relevant physical input in determining the T$_{\\rm eff}$ of pre-MS tracks in the HR diagram. The boundary conditions adopted also play a non-negligible role in determining the path followed by the evolutionary sequences on the HR plane. Gray models are systematically hotter than their non-gray counterparts. This effect, for the range of masses and ages at which most of the ONC population is found, has a similar quantitative effect as a change of the convective model. The use of non-gray models is recommended to describe these early evolutionary phases.  On the observational side, we find that the bulk of the observed stars in the ONC have masses in the range 0.2M$_{\\odot}$$\\leq$M$\\leq$0.4M$_{\\odot}$, for all non-gray models. The age distributions are more affected by the choice of the MLT parameter $\\alpha$. Ages are 1$-$2 Myr from the $\\alpha1.0$ set, 0.6$-$2.5Myr for $\\alpha2.0$ set and 0.4$-$1.6Myr for $\\alpha$2.2 models. This study confirms the presence of a dichotomy in the rotational properties between the objects with M$<$M$_{\\rm tr}$, whose period distribution peaks at short values, and stars with M$>$M$_{\\rm tr}$, that present a secondary peak at $P$$\\sim$8d. The transition value of mass between the two populations is at M$\\sim$0.5M$_{\\odot}$ for LCE and at M$\\sim$0.35M$_{\\odot}$ for HCE models. If disk-locking is responsible for the secondary peak observed in the overall period distribution, these results can be interpreted by assuming either that the masses M$<$M$_{\\rm tr}$ lose their disk earlier, or that their locking period is shorter. The X-ray emission shows no correlation with period, supporting the suggestion that these stars are indeed in the super-saturated regime of the rotation-X ray luminosity relationship. The correlation of the X-ray flux with mass appears to be the consequence of the increased average luminosity of more massive objects.  For the low-mass stars that presumably evolved without a disk, we find that our results are consistent with an evolution-conserving angular momentum. The comparison between the model period evolution and the observed values suggests that initial angular momenta at least 3 times larger than those found by means of the Kawaler (\\cite{kawaler}) law are needed for stars with mass in the range 0.2$\\,M_{\\odot}$$\\leq$$M$$\\leq$0.4$\\,M_{\\odot}$. This analysis was not possible for the higher masses sample, due to the lack of a statistically meaningful sample of rapidly rotating stars.   The idea that the double population of the ONC can be explained on the basis of a different morphology of the stellar magnetic fields seems to be ruled out by the fact that almost all the observed sources are fully convective, according to our interpretation. However, a mechanism (e.g. the magnetic field itself) inhibiting convection might favour an earlier appearance of the radiative core, at least in some of the stars. In our analysis we found other indications that convection in the pre-MS may be affected by other parameter(s): although 2D hydrodynamic simulations predict HCE in the pre-MS, we find two results in favour of LCE: (1) the lithium depletion in HCE models is too large to be consistent with the pre-MS depletion shown in young open clusters; (2) the age distribution derived from HCE models for two groups of smaller and higher masses is very different. It may be that the lower convection efficiency needed in the pre-MS is due to the structural effect of the dynamo-induced magnetic field, as suggested by Ventura et al. (\\cite{ventura}) and D'Antona et al. (\\cite{dantonaetal2000}). This possibility leads us not to dismiss the idea of an earlier appearance of a radiative core in the M$>$M$_{\\rm tr}$ group.  The evolutionary tracks (from 0.085 to 1.6M$_{\\odot}$) and isochrones (from 2$\\times$10$^5$ to 1$\\times$10$^7$Myr) are available from the following web site:  www.mporzio.astro.it/$\\sim$tsa"
    },
    "0609/astro-ph0609771_arXiv.txt": {
        "abstract": "Self-similar shock solutions in spherically symmetric polytropic gas flows are constructed and analyzed in contexts of proto-star formation processes. Among other possible solutions, we model a similarity shock across the sonic critical curve with an inner free-fall core collapse and a simultaneous outer expansion of the extended envelope; the separation or stagnation surface between these two flow zones travels outwards in a self-similar manner at a variable speed. One readily obtains accretion shock solutions. Semi-complete self-similar solutions across the sonic critical curve either once or twice without shocks can also be constructed. Features of star formation clouds of our polytropic model include the mass density scaling in the outer flow zone $\\rho\\propto r^{-2/(2-\\gamma)}$, the temperature scalings of the inner flow zone $T\\propto r^{-3(\\gamma-1)/2}$ and of the outer flow zone $T\\propto r^{-2(\\gamma-1)/(2-\\gamma)}$, and the variable central mass accretion rate $\\dot{M}=k^{3/2}t^{(3-3\\gamma )}m_0/G$ where $\\gamma$ is the polytropic index, $k$ is a constant, $m_0$ is the core mass, and $G$ is the gravitational constant. Spectral line profiles characteristic of the `envelope expansion with core collapse' (EECC) shock solutions are expected. Random magnetic field permeated in a partially ionized cloud can be incorporated into this theoretical polytropic model framework. We discuss briefly our results in context of the oft-observed starless B335 cloud system as an example. ",
        "introduction": "In astrophysical and cosmological contexts, the dynamic evolution of a spherical gas flow has been studied both numerically and analytically since late 1960s \\cite{larson69,penston69}. Sufficiently far away from initial and boundary conditions, such flows may gradually evolve into a self-similar phase (Sedov 1959). Shu (1977) constructed the `expansion wave collapse solution' (EWCS) in a self-gravitating isothermal gas cloud and advocated the so-called inside out collapse scenario for forming low-mass stars \\cite{shu77,shu87}. Such a similarity isothermal cloud collapse process involves a diverging mass density $\\rho\\propto r^{-3/2}$ and a diverging radial infall speed $v\\propto -r^{-1/2}$ near the collapsing core yet with a $\\rho \\propto r^{-2}$ far away in a static envelope characterized by the outer part of a singular isothermal sphere (SIS). The boundary of the infall region expands at a constant sound speed. For a more general polytropic gas, similarity solutions have also been derived and analyzed from various perspectives \\cite{cheng78,goldreich80,yahil83,suto88,mclaughlin97,fatuzzo04,lou06}. With effects of unknown energetic processes relegated to the polytropic index $\\gamma$, these polytropic solutions, among others, become more flexible and versatile in modelling core collapsing processes in a molecular cloud. Shocks can be constructed in an isothermal flow \\cite{tsai95,shu02,shenlou04,bianlou05,Yu06}. These similarity shock solutions contain important realistic features to model large-scale flow dynamics around star-forming regions. While there are two- or three-dimensional numerical simulations to model sub-structures in star-forming regions \\cite{truelove98}, we explore in this paper the full extent of possible one-dimensional features under spherical symmetry that are extendable to two- or three-dimensional flow modelling in principle.  Recently, Lou \\& Shen (2004) obtained a novel class of semi-complete isothermal flows, including those referred to as `envelope expansions with core collapses' (EECC) that may or may not involve the sonic critical line. Such solutions can model simultaneous inner collapse and outer flow expansion with or without shocks and are applied to star-formation regions such as B335 cloud system \\cite{zhou93,shenlou04}. In particular, Shen \\& Lou (2004) emphasized that a cloud envelope, with a central collapsing core to form a star or a star cluster, may either expand or contract depending on initial environmental conditions. Fatuzzo et al. (2004) considered a polytropic gas collapse under self-gravity in reference to observationally inferred far-away inflows. They adopted different polytropic indices for equations of state for an initial static cloud and a cloud in dynamic evolution in a more general manner. Central mass accretion rates are also examined as compared with observations (Fatuzzo et al. 2004). In this paper, we study polytropic gas flow solutions across the sonic critical curve either with or without shocks, which was not considered by Fatuzzo et al. (2004). Moreover, semi-complete EECC solutions crossing the sonic critical curve twice in a polytropic gas are also constructed as an extension of the isothermal result of Lou \\& Shen (2004). Both shock and EECC features are expected in clouds with collapsing cores surrounding proto-star forming regions. This theoretical development, together with shock modelling (Bian \\& Lou 2005; Yu \\& Lou 2005; Yu et al. 2006) and adaptations, may also bear physical implications for the dynamic evolution phase with a timescale of $\\sim 10^3$yrs linking AGBs and PNe (e.g., Balick \\& Frank 2002 and extensive references therein)  Observations of star-forming regions have been accumulating over years. Various optical, infrared, sub-millimeter and more recent X-ray observations reveal structures on different scales of hundreds of star-forming regions, a few of which are well observed \\cite{zhou93,shirley02,swift05} and become valuable sources for checking and testing theoretical model development. By simulating characteristic spectral line profiles given an underlying model of dynamic or magnetohydrodynamic (MHD) flows, one can infer information of a star-forming region by estimating the density and flow profiles, velocity dispersion, temperature scaling and central mass accretion rate and so forth.  With these considerations in mind, we present in \\S 2 our polytropic similarity flow model, including the polytropic EECC and shock solutions. Several aspects of probing star-forming processes are discussed in \\S 3. Further discussion of our model application is presented in \\S 4. Several details of mathematical derivations are summarized in Appendices A, B and C for the convenience of reference. ",
        "conclusions": "For star-forming cloud systems, the outer mass density profile, the inner and outer temperature profiles, and the core mass accretion rate are crucial to test or constrain a polytropic cloud model. Given a self-similar dynamic profile in a polytropic gas, spectral line profiles can be simulated. These spectral line profiles are very important to show different evolution epochs in forming proto-stars. In addition to forming proto-stars, the variable core mass accretion rate $\\dot M$ by equation (\\ref{accretion rate1}) also bears physical implications for the dynamic evolution phase with a timescale of $\\sim 10^3$yrs linking AGBs and PNe \\cite{balickfrank02}. Further observations and analysis are needed especially for post-PNe of different ages. We also expect adaptations of our polytropic shock model to a certain epoch of supernova explosions (Lou \\& Wang 2006), with radiation and neutrino pressures also taken into account.  Shock phenomena are common in forming proto-stars, formation and evolution of PNe, and supernova explosions; they are important for the gas heating, the production of high-energy photons, the synthesis of heavy elements, and the acceleration of cosmic ray particles. In these contexts as well as others, our polytropic shock model has wide applicability. Properties of polytropic shocks should be explored further to include radiation processes, magnetic field effects (Yu et al. 2006; Wang \\& Lou 2006) and atomic/nuclear processes associated with shocks.  Radiative transfer processes will be incorporated into our model to make the underlying dynamic processes `visible'. Spectral line profiles of certain atoms and molecules can reveal profiles of mass density, flow velocity and gas temperature. The EECC profiles and shocks are the most important features. Certain spectral line features derived from various underlying dynamic solutions should be compared with observations. Recent and further high-resolution spectral IR and microwave telescopes will surely stimulate the development of dynamic study of star-formation, formation/evolution of PNe, SNe and other relevant astrophysical processes."
    },
    "0609/astro-ph0609401_arXiv.txt": {
        "abstract": "We present a search for CO($1 \\rightarrow 0$) emission in three Local Group dwarf irregular galaxies: IC5152, the Phoenix dwarf, and UGCA438, using the ATNF Mopra radio telescope. Our scans largely cover the optical extent of the galaxies and the stripped HI cloud West of the Phoenix dwarf. Apart from a tentative but non-significant emission peak at one position in the Phoenix dwarf, no significant emission was detected in the CO spectra of these galaxies. For a velocity width of 6\\,\\kms, we derive 4$\\sigma$ upper limits of 0.03\\,K\\,\\kms, 0.04\\,K\\,\\kms\\ and 0.06\\,K\\,\\kms\\ for IC5152, the Phoenix dwarf and UGCA438, respectively. This is an improvement of over a factor of 10 compared with previous observations of IC5152; the other two galaxies had not yet been observed at millimeter wavelengths. Assuming a Galactic CO-to-\\Htwo\\ conversion factor, we derive upper limits on the molecular gas mass of $6.2\\times 10^4$\\,\\Msol, $3.7\\times 10^3$\\,\\Msol\\ and $1.4\\times 10^5$\\,\\Msol\\ for IC5152, the Phoenix dwarf and UGCA438, respectively. We investigate two possible causes for the lack of CO emission in these galaxies. On the one hand, there may be a genuine lack of molecular gas in these systems, in spite of the presence of large amounts of neutral gas. However, in the case of IC5152 which is actively forming stars, molecular gas is at least expected to be present in the star forming regions. On the other hand, there may be a large increase in the CO-to-\\Htwo\\ conversion factor in very low-metallicity dwarfs ($-2 \\le [$Fe/H$] \\le -1$), making CO a poor tracer of the molecular gas content in dwarf galaxies. ",
        "introduction": "Dwarf galaxies form the most abundant galaxy population in the nearby universe.  They represent a wide class of low-mass, low-metallicity stellar systems, with some being gas-rich and presently actively forming stars (dwarf irregulars or dIrrs), while others are gas-poor and apparently formed their stars long ago (dwarf spheroidals or dSph). In between the above extremes exists a subclass of gas-rich dwarfs, called transition-type galaxies, that are not currently for\\-ming stars. These objects occupy an evolutionary state intermediate between dIrrs and dSphs \\citep{mateo1998,grebel2001}.  Most studies of the interstellar medium (ISM) of dwarf galaxies up to now have focused on their atomic gas (HI) content \\citep{young2003,conselice03,buyle05,bouchard05,tarchi05}. Nevertheless if, in accordance with massive galaxies (e.g. the Milky Way), the star formation predominantly takes place in molecular clouds, that are observed to form within cold HI regions (with velocity dispersions of the order of 3\\,\\kms), one would expect a correlation between the star formation rate and the CO emission rather than between the star formation rate and the amount of cold HI gas. However, so far no such relation has been observed for dIrrs, neither for CO or HI. Generally individual HI clouds are found close to star formation sites, with usually an offset by a few 100\\,pc. Also, some very dense HI clouds, which surely exceed the density threshold above which star formation is expected to be possible, show no signs of ongoing or recent star formation \\citep{mateo1998,begum06}. These results show that the presence of a cold HI component may be a necessary ingredient for star formation, but not a sufficient one: it seems that sometimes a trigger is needed to initiate star formation in dense gas clouds \\citep{taylor1994,hernquist1995,taylor1997,pisano1999}. Moreover, the LG (Local Group) dwarf galaxies show a very wide variety of gas masses and star formation rates. Some are quite gas-rich but have little ongoing star formation while others are gas-poor but show evidence for recent star formation \\citep{mateo1998}. In those LG dIrrs that have been detected in CO, the emission usually comes from individual clouds, about 50~pc wide \\citep[see][and references therein]{israel1997,mateo1998,derijcke06}. The locations of star formation sites, traced by HII regions, correlate very well with regions of CO emission. This shows that the emission of the molecular ISM is a much better tracer of star formation than that of the cold atomic ISM. In order to investigate this interplay between the properties of the atomic and molecular ISM and star formation, we observed a sample of nearby dIrrs:~IC5152, the Phoenix dwarf and UGCA438. These galaxies have published HI masses but, with the exception of IC5152, have not yet been surveyed for the presence of a molecular ISM.  CO is often used to trace the molecular gas in galaxies because it is the most abundant (after H$_2$) and brightest molecule. It is related to the FIR luminosity of the galaxies that is emitted by the warm dust heated by embedded high-mass OB stars, indicating that it indirectly correlates with the current star formation rate \\citep[e.g.][]{gay2004,murgia2005}. It has also been found that dwarf galaxies below a metallicity of $12 + \\log({\\rm O/H}) \\sim 7.9$ are not detected in CO emission \\citep{taylor1998}. This suggests that the CO to H$_2$ conversion factor ($X_{\\rm CO}$) depends on metallicity and that the CO-H$_2$ relation becomes sharply non-linear around this metallicity \\citep{taylor1998,taylor2001}. However, In studying these low-metallicity dIrrs we might be able to shed new light on this cut-off in CO detections.  In the next section we present more details on the three galaxies in our sample. Section~\\ref{sec_obs} describes the new CO observations, and the results are presented and discussed in sections~\\ref{sec_results} and \\ref{sec_disc} with a comparison with recent HI observations. Finally we summarise our conclusions in section~\\ref{sec_concl}. ",
        "conclusions": "\\label{sec_concl}  We have obtained very deep CO emission maps of three Local Group dwarf irregular galaxies: IC5152, the Phoenix dwarf, and UGCA438. We did not detect CO emission in these galaxies, with 4$\\sigma$ upper limits of $0.03-0.06$~K~km~s$^{-1}$ over a 6~km/s velocity width. Our results are summarized in Table \\ref{tab_results}. As is obvious from the H{\\sc i} maps of these galaxies, our maps cover both their stellar and H{\\sc i} distributions.  There could be two possible explanations for the lack of CO emission in these galaxies. On the one hand, there could be a genuine lack of molecular gas in these systems, in spite of the presence of large amounts of neutral gas. On the other hand, the CO-to-\\Htwo\\ conversion factor might become very large in these very low-metallicity dwarfs, making CO a poor tracer of the molecular gas content. If this is indeed the case, this has implications for galaxy formation scenarios in the early universe, since molecules are an important coolant of the ISM. In the absence of CO, star formation might not proceed as we observe it today in our Galaxy. We find no correlation between the composition of the interstellar medium and the star-formation rate. All galaxies have H{\\sc i} associated with them and Phoenix is the only one with a tentative CO detection. However, only IC5152 is actively forming stars. In the case of the Phoenix dwarf, we are most likely witnessing the process of the transition of a dIrr to a dSph galaxy, where the neutral gas has been stripped from the galaxy. There might be a hint that some molecular gas is still present in the galaxy, but this needs further confirmation."
    },
    "0609/astro-ph0609637_arXiv.txt": {
        "abstract": "This paper presents a novel method for determining the probability that a supernova candidate belongs to a known supernova type (such as Ia, Ibc, IIL, \\emph{etc.}), using its photometric information alone. It is validated with Monte Carlo, and both space- and ground- based data. We examine the application of the method to well-sampled as well as poorly sampled supernova light curves and investigate to what extent the best currently available supernova models can be used for typing supernova candidates. Central to the method is the assumption that a supernova candidate belongs to a group of objects that can be modeled; we therefore discuss possible ways of removing anomalous or less well understood events from the sample.  This method is particularly advantageous for analyses where the purity of the supernova sample is of the essence, or for those where it is important to know the number of the supernova candidates of a certain type (\\emph{e.g.}, in supernova rate studies). ",
        "introduction": "Type Ia supernovae, empirically established to be standardized candles, are a staple of experimental cosmology.  A number of future cosmological probes (\\emph{e.g.}, DES, Pan-STARRS, LSST, JDEM) are planning massive sky surveys that will collect very large samples of supernovae, which will have to be classified into several known types. Supernova candidates are most easily classified by their spectra, but for very large ground-based surveys it may be  impractical to obtain a spectral confirmation of the supernova type for each candidate.  Therefore, it is imperative to develop reliable methods of supernova classification based on photometric information alone.  Photometric typing has been described in~\\cite{pozn02},~\\cite{riess2004},~\\cite{john} and ~\\cite{sull1}, among others. Most of the existing methods rely on color-color or color-magnitude diagrams for supernova classification.   In this paper, we propose a novel approach to the photometric typing of supernova candidates.  This method is based on a probability derived using a Bayesian approach, and is well suited for extracting the maximum amount of information out of limited data.  Unlike methods relying on color-color or color-magnitude diagrams, which require a comparison of the candidate's color(s) with pre-existing tables or plots -- a comparison that requires a good understanding of the errors and assumptions that went into the making of the literature data -- our approach calls for a calculation of a single number and automatically takes into account all of the information that is available for a given candidate while incorporating the currently best-known models for supernova behavior.   A Bayesian method in the context of supernova light curve fitting has been used in~\\cite{barris}; however, it was applied specifically to type Ia supernovae to deduce redshift-independent distance moduli. In this paper, on the other hand, we describe a probabilistic approach to typing photometrically surveyed supernovae. Such an approach allows the possibility of the marginalization (integration) of the unknown, or nuisance, parameters. In contrast to a traditional $\\chi^2$ calculation, this technique is simply a calculation of a probability and does not involve fitting or minimization. However, it does require that the candidate sample be well understood; or, in other words, that each candidate be one of a number of hypothesized objects whose behavior can be well modeled.   The purpose of this paper is twofold.  Apart from introducing a methodology that can be easily applied ``as is'' or extended as needed, we would like to test the extent to which applying the best currently available supernova models helps in typing supernova candidates.  There is no doubt that within the next several years new and improved supernova models will be constructed; when they are, they can be easily worked into the method.  The paper is structured as follows.  In Section~\\ref{sec:like}, we derive the probability that a given candidate is a Type Ia supernova.  In Section~\\ref{sec:valid}, we discuss how well the method works when applied to poorly sampled, Hubble Space Telescope GOODS data, and to well-sampled, ground-based SNLS data. We suggest further improvements to the method in Section~\\ref{sec:improve}, discuss its possible application to ``anomalous'' objects in Section~\\ref{sec:anomalies}; and to fitting for supernova parameters, in Section~\\ref{sec:fitting}. Conclusions are given in Section~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl} We introduced a novel method to determine the probability that a supernova candidate is indeed a certain type of supernova sing photometric information alone. The probability is derived using a Bayesian approach. We have tested the method on both poorly sampled HST GOODS space-based data and well-populated SNLS ground-based light curves, with good results even for data where very few epochs available. While we considered primarily the application of this method to identifying type Ia supernova candidates, it can of course be used to identify any other known supernova types (see Fig.~\\ref{fig:mc}).  The results of these studies show the method to be promising.   We have assumed that the candidate sample consists entirely of supernovae.  The method naturally incorporates a number of possible hypotheses for the supernova types, allowing one to introduce prior knowledge of the probability distributions for various supernova parameters, as well as to marginalize them.  Because it is crucial that the sample consist of candidates that have reliable models (\\emph{i.e.}, well understood supernova types such as Branch-normal Ia, Ibc, \\emph{etc.}), we also proposed a possible approach for eliminating anomalous or less well understood candidates from the data.  As more models for different supernova types become available, more candidates will pass this preliminary selection to be further considered for typing.   We have used the best currently available supernova models for all of the supernova types considered. While aware of their limitations, we note that using these models already yields good discriminating power; as better supernova models become available from the ongoing and upcoming ground- and space- based supernova surveys, classifying supernovae with this approach will become proportionally more reliable.    The method described in this paper could be useful for a number of studies. An obvious application is selecting a high-purity sample of type Ia's for further analysis -- \\emph{e.g.}, creating a Hubble diagram and extracting cosmology.  In order to do so, one would re-parametrize the likelihood in terms of the distance modulus $\\mu$ and maximize $l(\\mu,z,{\\rm Ia}|{\\rm candidate})$ to obtain the best estimates of the distance modulus and the redshift of the candidate, as well as the errors on these parameters, as described in Section~\\ref{sec:fitting}. A pure sample of type $T$ supernovae could also be used for calculating the rates of this class of objects, provided one performs careful Monte Carlo studies of the method's rejection factor for type $T$ supernovae. Another potentially useful application would be adapting the method for early-time light curve parameters in such a way as to make it a trigger for type Ia supernovae in large supernova surveys. Other possible extensions of the method would include introducing prior information on supernova rates for the different supernova types, as well as building a Bayes factor to measure the relative likelihood that a given candidate is more or less likely to be a Ia supernova than another type."
    },
    "0609/astro-ph0609346_arXiv.txt": {
        "abstract": "We report the discovery of  hard X-ray delays in the  X-ray emission of the Seyfert 1 galaxy Mrk~110, based on a long XMM-Newton observation. Cross correlation between the X-ray light curves of different energy bands reveals an energy dependent delay ranging from a few minutes to an hour. We find that the energy spectrum can be modeled by Comptonization of disk blackbody photons. The energy dependent delay can be modeled as due to the effect of Comptonization in a hot plasma confined within 10 Schwarzschild radius of the black hole. We discuss our results in the context of inverse Comptonization of the soft photons by highly energetic plasma. ",
        "introduction": "The dominant X-ray radiation mechanism in accreting black holes is commonly thought to be inverse Compton scattering of low energy photons by a cloud of hot ($T\\sim 10^{8} - 10^{9}~K$) electron plasma near the black hole \\citep{st80}. The physical geometry of the coronal plasma responsible for scattering the photons is not well constrained by the current observations. The ubiquity of a disk blackbody component accompanied by a power-law tail in the overall spectra of galactic black hole candidates (GBHCs) motivated several workers to develop the so called two phase accretion disk-corona models \\citep[see eg.,][]{hm91,hm93, pkr97,sz94,ste95,bel98}. In these models, the blackbody radiation from the cold disk enter the hot corona and is Comptonized into X-rays. A part of the hard X-rays from the corona, being reprocessed in the disk, produce the reflection hump. The geometry of the corona controls this feedback mechanism which in turn determines the spectral slope of the escaping radiation. The Kompaneets y-parameter (and hence the temperature) is determined by the energy balance between heating and cooling mechanism inside the plasma. One important question in all such models is the method by which the gravitational energy is converted to the energy of electrons. The ideas explored in the literature include magnetic flares \\citep{hmg94,ste95,pou96,bel99a,bel99b}, advection dominated disk very close to the black hole \\citep{ny94}, Bondi type free fall beyond a shock region \\citep{ct95}, but, as yet, there is no consensus on the exact mechanism. A general prediction of the above models is that the hard X-ray variations should lag behind those in softer bands, as hard photons undergo more number of scatterings in the plasma before escape. Measurement of such time-lags has the potential of constraining the size of the region which in turn will help us to understand the mechanism by which the electron cloud is energized.   Frequency dependent time lags have been observed in Galactic black hole candidates \\citep{miy88,mkk91} as well as in a few active galactic nuclei \\citep[AGNs;][]{pnk01,zh02}. Interpreting these lags as due to Comptonization requires the size of the emitting region to be very large, typically several thousand Schwarzschild radii (compatible to the lowest frequency at which the lag is observed), leading to the problem of heating the electron cloud at large distances from the black hole. Hence, sometimes these lags are interpreted as due to the energy dependent asymmetries in random shots \\citep{mk89}; that is interpreting the lags as due to the production of the variability itself.  One of the problems of detecting lags at higher frequencies and relating them to the Comptonization at the innermost regions of the accretion disk around black holes could be observational. If the Comptonization process occurs at 10 -- 20 Schwarzschild radii and if the Comptonization process gives a lag which is a factor of few larger than the light travel time in these regions, one expects a lag of a few tens of milliseconds in a black hole of mass 10 M$_\\odot$ and  about a day in a AGN of mass 10$^8$ M$_\\odot$. Detection of such delays is observationally a difficult task. Bright nearby AGNs with a black hole mass of 10$^7$ M$_\\odot$, where one expects a delay of about an hour, are the ideal sources to look for delays due to the process of Comptonization. With this motivation, we have searched for delays in one of the bright low mass AGN Mrk~110, based on a long observation using the XMM-Newton observatory. In this {\\it Letter} we present the cross-correlation analysis in different energy bands of this source and show that hard X-rays are delayed by a significant amount of time (from few minutes to an hour). The 0.3 - 12 keV X-ray spectrum can be represented by a  Comptonization model which can explain the measured hard X-ray delay.  \\begin{figure} \\centering \\includegraphics[scale=0.35,angle=-90]{f1.eps} \\caption{EPIC-PN light curves (1000 s bin) of Mrk~110 in different energy range. Start of variability is indicated by downward arrows (see text)}. \\label{lc} \\end{figure} ",
        "conclusions": "We find significant energy dependent delay between the hard and soft X-ray emission in the Seyfert 1 galaxy Mrk 110. Similar delays were also found in the bright Seyfert 1 galaxy MCG-6-30-15 \\citep{pon04}. \\citet{gal04} found alternating leads and lags in the NLS1 IRAS 13224-3809. Though other interpretations like geometric effects and energy dependent shape of shots can also explain the observed energy dependent delays in Mrk 110, we try to interpret the results as due to Comptonization, particularly because of the fact that the spectral analysis favours a two component Comptonization model. We consider a static Compton cloud with optical depth $\\tau_T$, and electron temperature $\\Theta=kT_e/m_e c^2$. A soft seed photon of energy $E_0$ injected into the cloud increases its energy by a factor of $A=1+4\\Theta+16\\Theta^2$, on an average after each scattering, so that after n scattering its energy is $E_n=A^n E_0$. The photon mean free path is $\\lambda \\approx R/max(1,\\tau_T)$ (where R is the size of the X-ray emitting region). The time difference between successive scatterings is then $t_c = (R/c)/max(1,\\tau_T)$, where $\\tau_T= N \\sigma_T R$ is the Thompson optical depth, and $\\sigma_T$ is the cross section of Thompson scattering, N is the electron number density of the scattering medium. The time needed to reach the energy $E_n$ is $t_n=nt_c$ \\citep{st80,pay80}. We have calculated time lags of different high energy bands with respect to 0.2 -- 0.3 keV band. Hence in our case $E_0=0.25$ keV. We fit the above equation ($t_n=nt_c$) to the result (hard lag as a function of energy) of our analysis and find that the data is in good agreement with the equation (Fig.~\\ref{delay}). Using the values of the best fit parameters of the above result and the parameters of the Comptonization model fitted to spectrum we get $R\\sim2\\times10^{11}~m~\\sim 10~R_S$. This value of R is physically realistic as most of the energy is dissipated within 10 $R_S$. \\citet{bl98} pointed out that any scenario in which the observed hard time lags are purely due to static Comptonization requires that the radial extent of the hot corona exceeds $\\sim 10^{4}~R_S$ of a solar mass black hole ($R>10^{8}~m$). This is incompatible with current models of accretion flows onto galactic black holes \\citep[see the review by][]{lia98} even from simple energy arguments (see section~1). But the results reported here pertains to lags at very short times scales and hence is consistent with the static Comptonization model. The harder power-law emission extending to 12 keV (and presumably beyond) can also be produced by Comptonization, either in another purely thermal plasma or non-thermal electrons in a plasma with a hybrid thermal/non-thermal distribution.  The origin of the hot plasma above the accretion disk is quite debatable. If the radiation pressure inside the accretion disk very close to the black hole is very high then the accretion disk can be puffed up and the hot thermal plasma can be formed just above the accretion disk. Hybrid thermal/non-thermal plasmas have often been successfully used to model the observed data \\citep{gier99,pc98}. A possible origin of the non-thermal component is in process of magnetohydrodynamic turbulence occurring in the corona \\citep{lm97}. Stochastic acceleration of particles by stochastic gyroresonent acceleration in accreting plasma can also accelerate particles to higher energies \\citep{dml96,lkl96}.  We conclude that inverse Compton scattering of soft photons by highly energetic electron distribution provides a satisfactory explanation of the hard X-ray time lag observed in the XMM-Newton observation. The energy spectrum in the 0.3 -- 12 keV band can be modelled either by two component Comptonization or by a hybrid thermal/non-thermal Comptonization. It is not possible to distinguish between them without good quality high energy data.  This work is based on observations obtained with the XMM-Newton, an ESA science mission with instruments and contributions directly funded by ESA member states and the USA (NASA). Authors are grateful to H. Netzer, G. Dewangan and S. Mandal for useful suggestions. We thank the referee for useful comments."
    },
    "0609/astro-ph0609493_arXiv.txt": {
        "abstract": "We define three requirements for accurate simulations that attempt to model circumstellar disks and the formation of collapsed objects (e.g. planets) within them. First, we define a resolution requirement based on the wavelength for neutral stability of self gravitating waves in the disk, where a Jeans analysis does not apply. For particle based or grid based simulations, this criterion takes the form, respectively, of a minimum number of particles per critical (`Toomre') mass or maximum value of a `Toomre number', $T= \\delta x/\\lambda_T$, where the wavelength, $\\lambda_T$, is the wavelength for neutral stability for waves in disks. The requirements are analogues of the conditions for cloud collapse simulations as discussed in \\citet{BB97} and \\citet{Truelove97}, where the required minimum resolution was shown to be twice the number of neighbors per Jeans mass or 4-5 times the local Jeans wavelength, $\\lambda_J$, for particle or grid simulations, respectively.  We apply our criterion to particle simulations of disk evolution and find that in order to prevent numerically induced fragmentation of the disk, the Toomre mass must be resolved by a minimum of six times the average number of neighbor particles used. We investigate the origin of the apparent discrepancy between the number of particles required by the cloud and disk fragmentation criteria and find that it is due largely to ambiguities in the definition of the Jeans mass, as used by different authors. We reconcile the various definitions, and when an identical definition of the Jeans mass is used, the condition that $J\\la 1/4$ in the Truelove condition is equivalent to requiring about 10-12 times the average number of neighbor particles per Jeans mass in an SPH simulation, reducing the difference between simulations of disks and clouds to about two. While the numbers of particles per critical mass are similar for both the Jeans and Toomre formalisms, the Toomre requirement is more restrictive than the Jeans requirement when the local value of the Toomre stability parameter $Q$ falls below about one half.  Second, we require that particle based simulations with self gravity use a variable gravitational softening, in order to avoid inducing fragmentation by an inappropriate choice of softening length. We show that using a fixed gravitational softening length for all particles can lead either to artificial suppression or enhancement of structure (including fragmentation) in a given disk, or both in different locations of the same disk, depending on the value chosen for the softening length. Unphysical behavior can occur whether or not the system is properly resolved by the new Toomre criterion.  Third, we require that three dimensional SPH simulations resolve the vertical structure with at least $\\sim4$ particle smoothing lengths per scale height at the disk midplane, a value which implies a substantially larger number per vertical column because the disk itself extends over many scale heights. We suggest that a similar criterion applies to grid based simulations. We demonstrate that failure to meet this criterion leads to underestimates in the midplane density of up to 30--50\\% at resolutions common in the literature. As a direct consequence, gas pressures will be dramatically underestimated and simulations of self gravitating systems may artificially and erroneously inflate the likelihood of fragmentation. We outline an additional condition on the vertical resolution in simulations that include radiative transfer in order to ensure a correct description of the cooling, specifically that the temperature structure near the disk photosphere must be well resolved. As an example, we demonstrate that for an isentropic vertical structure, the criterion translates to resolution comparable to $H/20$ near the disk photosphere, to avoid serious errors in transfer rates of thermal energy in and out of the disk.  Finally, we discuss results in the literature that purport to form collapsed objects and conclude that many are likely to have violated one or more of our criteria, and have therefore made incorrect conclusions regarding the likelihood for fragmentation and planet formation. ",
        "introduction": "The formation of collapsed objects, both in the context of the collapse of molecular cloud cores and in the later context of the clumping of material in a circumstellar disk is very difficult to model numerically because of the huge range of spatial scales involved. Accurate simulation of the collapse usually requires that all of these scales be well resolved, if the result is not to be contaminated by numerically induced fragmentation.  For three dimensional (3D) simulations, \\citet{Truelove97} defined a minimum resolution condition for the numerical validity of a simulation that models the collapse of a molecular cloud core using a grid based hydrodynamic code. Contemporary work by \\citet[][hereafter BB97]{BB97} has examined necessary resolution conditions for the collapse of a similar system in the context of Smoothed Particle Hydrodynamics (SPH) simulations, and also discussed the influences that choices of the form of gravitational softening may have on the results. Both of these works define minimum resolution criteria in the context of a Jeans collapse of a gas cloud resolved in 3D, but while \\citet{Truelove97} observe that fragmentation is enhanced by a failure of the criterion, BB97 only observe enhanced fragmentation if the gravitational softening used in their particle based simulations is smaller than the hydrodynamic smoothing. Otherwise, they observe that fragmentation can be unphysically delayed in a system where it is known to occur. Later work of \\citet{HGW06} explores the problem of Jeans collapse in isolation. They find that SPH simulations of systems with initial conditions like that of the original linearized analysis do not fragment artificially, even when under resolved.  A number of recent works in the field of planet and brown dwarf formation \\citep{DynI,DynII,Pick98,Pick00a,Pick03,Boss98,Boss00, Boss02,Boss04,Mayer02,Mayer04,RABB,ArmHan99,Luf04} have discussed the formation of planets and brown dwarfs via gravitational fragmentation in disks. In disks systems like those discussed, the Jeans formalism used to develop the \\citealt{Truelove97} and BB97 criteria is not valid both because the disk scale height is usually small compared to the Jeans wavelength and because of the existence of shear, which plays a role as important as self gravity for the growing structures. As yet however, no analogous resolution criterion exists for disks, which may be modeled in either fully in two dimensions (2D), effectively integrating over the vertical coordinate, or in 3D, for which no vertical integration is assumed, but for which the Jeans wavelength based criterion still may not apply.  In addition to requirements that simulations resolve the wavelengths of the relevant instabilities sufficiently, simulations must also satisfy a number of other criteria if they are to be believed. Approximations made outside the realm of the physical model, perhaps used to model the behavior of unresolved or poorly resolved phenomena, must not drive the results of the simulations themselves. In the context of simulations of disks, particularly those using particle based hydrodynamic methods such as SPH, an important consideration will be the implementation of gravitational softening. At a more fundamental level, the algorithm used to solve the equations used to model physical system must do so accurately, without becoming unstable or generating large errors through some other sort of deficiency. A quick glance through the literature \\citep[see e.g. textbooks of][]{hockeast88,fletcher97,leveque02} will demonstrate to even a casual observer that the study of numerical methods in relation to their stability is one in which extensive studies on many topics have been performed. Verification of methods on physically relevant test problems however, is comparatively more widespread but may often suffer from insufficient detail or generality. While many studies discuss the fidelity of the numerical solution on some simple or contrived test problem, often no test problems sufficiently similar to the physical system under study can even be devised.  In section \\ref{sec:numerical-factors}, we first extend the previous 3D work of \\citet{Truelove97} and BB97 to self gravitating thin disk systems, then outline alternatives for gravitational softening in particle simulations, and the possible consequences each choice may have on results. In section \\ref{sec:testing}, we define a test problem with which we determine appropriate values for resolution for particle simulations. We continue in section \\ref{sec:softening} with a discussion of the gravitational softening assumed in particle simulations, and the specific numerical issues encountered in a study of disks. Next, in section \\ref{sec:disks-3d}, we discuss the application of the criterion to simulations done in 3D, and define a test problem to validate the accuracy of numerical codes attempting to model disks in 3D. Using this test problem, we demonstrate that failure to resolve the vertical structure of disks will lead to large errors in the numerical solution for the evolution of the entire physical system. Finally, in section \\ref{sec:remarks}, we summarize our results and discuss them in the context of the models presented in the literature. ",
        "conclusions": "\\label{sec:remarks}  After exploring the conditions required for simulations to produce numerically valid results, we turn now to summarizing those conditions, to comparing the work done here with previous work in the literature, and to commenting on additional requirements beyond numerics that are required for simulations to accurately reproduce the evolution of real systems.  \\subsection{Summary of our main results}\\label{sec:summary}  Our first result is to extend the numerical criterion for the validity of simulations involving collapse in clouds and based on the analytic Jeans collapse formalism (BB97, \\citet{Truelove97}), to the case of disk systems where the geometry of the mass distribution is both flattened and rotating so that the Jeans treatment is inapplicable. Whereas a minimum spatial resolution per Jeans wavelength or minimum number of particles per Jeans mass is required in the cloud collapse case for grid or particle based simulations, respectively, in simulations of disks we require a minimum spatial resolution per `Toomre wavelength' or minimum number of particles per `Toomre mass', for grid based or particle base simulations, respectively. The two forms of the new criterion are given in equations \\ref{eq:Toom-cond} and \\ref{eq:max-surfdens}.  We also note that simulations failing to resolve one or the other condition yield quite different outcomes. In the case of the Jeans condition, BB97 conclude that insufficient resolution will delay or suppress numerically induced fragmentation with softening and smoothing equal, and only enhance it when smoothing is much larger than softening. Our results demonstrate that a simulation that fails to resolve the Toomre condition is characterized by enhanced fragmentation compared to one that is well resolved, even when softening and smoothing are equal. This fact allows an empirical determination of the minimum resolution required in a simulation in order to ensure fragmentation that develops is not of numerical origin. For 2D SPH simulations, we have determined the minimum number of particles per Toomre mass as six times the average number of neighbors ($\\sim20$) used for the realization of the hydrodynamic quantities themselves.  We have not determined a value of \\nreso\\  for 3D models, but its seems unlikely to be dramatically different from 2D, since the same physical system is being modeled with the same numerical technique. To the extent that differences may be present, we expect that the criterion will be more conservative (more particles required) based on two arguments. First, unstable waves in the Toomre analysis are intrinsically two dimensional, so resolving their third dimension is meaningless: small particle separations or small grid spacing in that coordinate do not contribute to increasing the resolution of the waves themselves. Using the same empirical fragmentation/non-fragmentation limiting condition used in section \\ref{sec:test-crit} to establish a required particle count, will therefore overestimate the required resolution because some neighbors will be counted only for purposes of establishing a value for \\nreso, but will not actually serve to increase the resolution insofar as the Toomre condition is concerned. Secondly, although the Toomre criterion implicitly assumes validity of the hydrodynamic equations, in practice the numerical method may fail to do so (section \\ref{sec:disks-3d}), with the consequence that pressure forces are underestimated. Although a vertical resolution requirement is nominally an entirely different condition from resolving instability wavelengths, we see no practical manner in which they might be separated from each other in the empirical tests we describe. Therefore, calibration of the requirements for the Toomre condition in 3D will implicitly incorporate both requirements from the wave analysis and from the vertical resolution condition, yielding a resolution requirement defined by contributions from both sources.  A second important result of our work is to reconcile a variety of mutually inconsistent definitions of the Jean collapse stability criteria made by different authors, due to the use of different definitions of the Jeans mass itself. We show that using the same definition of the Jeans mass, a particle based simulation will require approximately 600 particles to resolve one Jeans mass, while a grid based simulation will require about 64 zones (i.e. $1/J^3$ where $J$ is defined in equation \\ref{eq:Jeans-cond}, and $J\\sim 1/4$). The minimum resolution of 600 particles per Jeans mass corresponds to approximately 12 times the average number of neighbor particles ($\\sim50$) used by BB97 in their SPH code. This value is a factor of two higher than the factor of six times the average number of neighbors as derived from the 2D disk simulations explored in section \\ref{sec:2dtest-evo}.  Third, we discuss the importance of the implementation of interparticle forces in particle based simulations of self gravitating disks. We show that a naive translation of the 3D method for gravitational softening in SPH into 2D leads to an unphysical, finite gravitational force at zero separations between particles. In consequence, significant particle `pairing' occurs as simulations proceed, resulting in an effectively time dependent resolution. We consider two alternatives to remedy the behavior. First, we artificially modify the derivative of the SPH kernel according to the TC92 prescription, generating a finite, repulsive pressure force at zero separation which compensates for the attractive gravitational force. Second, we artificially modify the gravitational softening itself, to produce instead a zero force at zero separation, as is physically appropriate.  Both alternatives are effective in reducing the likelihood of disk simulations to produce fragmentation and of particles in those simulations to become paired. Although the TC92 variant appears to be slightly more effective (see e.g. section \\ref{sec:enhancement}), we view it as marginally the less preferable of the alternatives because it too produces a finite force at zero separation, albeit of opposite sign so that a sum of gravitational and pressure forces yields a small net force. Finite forces at zero separation, whether or not they are combined with others to produce a near zero sum, must be considered a less physical and therefore less desirable alternative to a solution which retains the collisionless nature of the model more directly.  Fourth, we show that force imbalances on the spatial scale of the smoothing and softening develop when the two length scales are not equal. While acting only at very small spatial scales, these imbalances can induce large scale changes in the outcome of simulations where they exist, including artificially induced fragmentation in otherwise stable systems. To avoid such induced outcomes, we conclude that the gravitational softening length and the hydrodynamic smoothing length of each particle must be very nearly identical. As noted above, our conclusion for disk simulations is similar to that made by BB97 in their study of collapsing clouds, but stronger in the sense that it holds even when the systems are, and are expected to remain, properly resolved according to the Toomre criterion.  In neither the disk or cloud case however, does our conclusion come without price. Specifically, in order to better resolve the flow in both high and lower density regions of the same simulation, most modern SPH codes set the smoothing length dynamically as a function of the local flow. Setting the smoothing and softening lengths equal means that the softening length also must vary in time. This is important because allowing temporally varying softening is equivalent to explicitly allowing a violation of conservation of energy in the simulation. While we believe the magnitude of violation to be small, its effect must be considered when evaluating self gravitating particle based simulations.  Finally, we discuss the importance of adequately resolving the vertical structure of 3D disk simulations. We find that insufficient resolution of the vertical structure leads to substantial deficits in the realized midplane densities compared to the predictions of analytical models. We find that resolution comparable to at least $\\sim4$ smoothing lengths per disk scale height in the disk midplane are required to accurately fit the coefficients in analytic formulae for the vertical structure of non self gravitating isothermal disks in particle simulations. Restated in terms of the full thickness of the disk, the criterion is equivalent to the statement that at least $\\sim17-20$ particles are required per vertical column. While we have not studied the effects of resolution in 3D grid based simulations, we expect similar criterion to hold for them as well.  The consequences of failures to resolve the vertical structure in self gravitating systems will be that pressure forces will be incorrectly determined. Any balance between them and self gravitation will therefore be falsely biased in favor of gravitational fragmentation. Presumably, the bias will distort the evolution of waves in the disk (whether self gravitating or not) even when fragmentation does not occur, because the density and pressure deviations imply a loss of fidelity of the solution to the equations of hydrodynamics provided by the numerical method.  We discuss the fact that while satisfying the criterion above is necessary for simulations to accurately evolve circumstellar disks, it may not be sufficient for them to do so alone, depending on the goals of the researcher. Simulations for which the criterion is satisfied everywhere may still under resolve high altitude structure of the disk, a fact which will be of great importance for models in which radiative transport is included. Such models must also resolve the temperature structure at the disk photosphere in order not to be contaminated by unacceptably large numerical errors. We provide an example of the requirements when an isentropic vertical structure is present, for which we derive the computationally extremely demanding condition that the photosphere must be resolved at a spatial scale smaller than $\\sim H/20$ to avoid serious errors. We point out however that the exact requirement will be highly sensitive to the details of the model and its effects on the temperature, density and opacity structures at high altitudes.  \\subsection{Comparisons to particle based simulations in the literature}\\label{sec:part-comps}  Our test problem clearly shows that the fragmentation processes seen in \\citet{DynI} were due to insufficient resolution, and we believe that many other works in the literature both before and since may be similarly affected. Of the work discussing particle based simulations, the simulations done by \\citet{Mayer02,Mayer04} are among the highest resolution simulations so far published, evolving either $2\\times 10^5$ or $10^6$ particles (in various runs) in three dimensions. It may therefore be particularly useful to contrast our test problems with the conditions employed by them, their results and conclusions.  They perform a series of SPH simulations with several disk masses between \\mrat=0.075 and \\mrat=0.125, each with minimum Toomre $Q$ values between 0.8 and 1.9. They show that the more massive disks fragment on a time scale of a few hundred years (about 10 orbits in the region where the clumps form) when the minimum $Q$ values in the disk are below \\qmin$\\approx1.4$. They also find that clump formation is more vigorous with higher resolution simulations of the same initial condition that with lower. An isothermal gas equation of state was used initially for most of the models, and after a critical over density was exceeded, was switched over a small transition time to an adiabatic evolution that included heating due to compression, viscosity and shocks. Other models used an adiabatic equation of state for all times, and resulted in fragmentation that was much weaker or non-existent.  We first attempt to obtain some indications of the degree to which the flow is well resolved by the condition in equation \\ref{eq:max-surfdens} at different times during the evolution, using their discussion of the physical features of the flow and resolution. For example, prior to clump formation the most unstable wavelength is quoted as $\\sim2$~AU, including a correction factor for 3D effects. The flow is clearly well resolved by the simulation at this early stage, but the value of the Toomre wavelength\\footnote{Note that their definition of the `Toomre wavelength' is that of equation \\ref{eq:disk-critwavel}, but that their later definition of the most unstable wavelength matches that of equation \\ref{eq:Toom-wavel}.} at this time is only representative of the mildly non-linear regime of the background flow. It is not indicative of the later, more strongly non-linear regime near the forming clump, for which both the locally applied Jeans based or Toomre based resolution criteria are most useful.  Well after a clump has formed, when the clump has become a well defined, condensed and independent entity rather than a part of the disk, we expect that the Toomre based criterion will become invalid. At this time, \\citealt{Mayer04} state that the resolution employed in their simulations was sufficient to resolve the Jeans mass (for which the underlying analysis is valid if the clump is much smaller than a disk scale height) with a few hundred to a few thousand particles, even though the densities there reached as much as $\\sim10^6$ higher than the background. An important detail of the implementation of \\citealt{Mayer04} is that the number of neighbors used for realizing the hydrodynamic quantities is fixed at 32, a value common throughout the cosmological SPH literature, but significantly smaller than the value of $\\sim50$ used (but not specifically recommended) by BB97. Correcting for this systematic difference would not appear to alter the statement that the Jeans mass is properly resolved however, because while the number of particles per Jeans mass would be lower, it would still remain above the $\\sim100$ quoted by BB97.  We can make relatively few direct inferences about intermediate times when the clump has begun to form but is still unbound or only weakly bound. At such times the local value of $Q$ may fall below one half, and we expect a much stronger resolution constraint from the Toomre criterion than from the Jeans criterion (see equation \\ref{eq:TJ-relative}), especially if the approximation that $\\Omega\\approx\\kappa$ is not made. Therefore, while the clump may be well resolved by the Jeans criterion, it may not be similarly well resolved by the Toomre criterion. Based solely on the number of particles and the similarity of the initial conditions to our own test problem, we conclude that the high resolution (1 million particle) simulations of \\citealt{Mayer04} appear to be well resolved by our Toomre based resolution criterion. We can make a similar, but much more tentative conclusion about the lower resolution ($2\\times10^5$ particle) versions in the same work. A more detailed examination of the simulations at both resolutions in the context of the Toomre resolution criterion would be desirable, particularly in the context of the sensitive dependence of the resolution requirement to the initial minimum value of $Q$ and on the local value of $Q$ in the fragmenting region.  \\citealt{Mayer04} advocate a fixed gravitational softening whose length scale is half of the initial hydrodynamic smoothing length of the particles. With this ratio of softening to smoothing, BB97 show that the gravitational force can exceed the pressure force by a factor of as much as seven on size scales of $h/2$, and by a factor of 2--3 on scales of $h$. The actual length scale used is determined (Mayer 2004 private communication) as an average of the smoothing lengths in the `flat' temperature region of the disk, corresponding approximately to the outer half of its radial extent, meaning that the outer half of their disks will have small scale force imbalances that originate in the numerical treatment of softening and smoothing.  The results of our simulations show the consequences of such imbalances in disk simulations. In regions where the hydrodynamic smoothing length exceeded the gravitational softening length, fragmentation was strongly enhanced over that of the variable softening case, even when sufficient numerical resolution was used according to the Toomre criterion. Since the force imbalance is not due to any physical origin, but only to a breakdown in the simulation method itself, we conclude that clumping observed in the simulations may not be an accurate depiction of an underlying physical system, and should be reexamined to ensure that they are not subject to this numerical flaw.  Of still greater concern is the resolution of their simulations in the context of the discussion in section \\ref{sec:3dsph-nosg-signif}. Both the morphology of the system modeled by \\citealt{Mayer04} and the physics itself are different than those employed by our test models, however it seems unlikely that these modifications will alter the required resolution by the large factors needed to conclude that their vertical structure is well resolved, especially for their most unstable (lowest $Q$) models. The extent to which any violation of the criterion actually affects fragmentation is unclear however. It is likely that a definitive statement of the effects will require a complete study of the problem at extremely high resolution, making the question quite costly to address at present.  Several other recent papers have discussed fragmentation in disks. Similar concerns regarding the vertical resolution will apply to many of them. For example, in \\citet{Boss04}, as well as his earlier, similar work, the vertical coordinate (the spherical coordinate $\\theta$ in his case) is resolved with a total of 23 zones. Although they are asymmetrically arranged to concentrate resolution where they are needed most, only some $\\sim$10-12 lie within ten degrees ($H/R\\approx0.17$) of the midplane, where most disk matter will be located. Although our vertical resolution requirement has been tested only with SPH simulations, there is no reason to believe that an analogous requirement will be required for grid based simulations as well, since they must solve the same set of equations. Naively applying the $\\sim4$ smoothing lengths per scale height as a $\\sim4$ grid cells per scale height criterion directly to the Boss simulations, we observe that the vertical structure may be very near the required limits. The exact number of cells will undoubtedly be different than a straightforward carryover however, so the extent to which Boss's or any other any grid based simulations will be affected by the errors arising from deficient vertical resolution is unclear.  \\citet{RABB} showed that for sufficiently fast cooling rates, clumping could occur in disks resolved with $2.5\\times10^5$ particles. Due to their cooling prescription, disks in their work dropped to globally averaged $Q$ values near unity, well below those in our simulations, for which we determined erroneously low midplane densities developed. In later work discussing similar systems (with overall morphologies identical to those in our study, but with additional physics), \\citet{LR04} claim that their vertical structure is resolved with $\\sim5$ smoothing lengths per scale height. It is difficult to reconcile this result with our own figure \\ref{fig:hiterat-rhodist}, especially given the much lower $Q$ values (see their figure 3) realized in their study and in the earlier work of \\citealt{RABB}. It appears likely however, that low vertical resolution has likely influenced their results through the artificially low pressure forces induced by density deficits.  Apart from the issue of vertical resolution, we note that their overall resolution (in 3D) is similar to that in our highest resolution 2D test problem. The authors have not discussed resolution issues in the context of either Jeans or Toomre stability, so we can draw few inferences about the viability of their results under either the criteria of equation \\ref{eq:max-voldens} or \\ref{eq:max-surfdens}. In the context of the latter condition however, the $Q$ values throughout their simulations rapidly approach values near unity. This is important because equation \\ref{eq:max-surfdens-Q} shows that our resolution requirement is quite sensitive to the value of $Q$, so that much higher resolution will be required for later evolution even if the initial condition is known to be well resolved by the Toomre condition.  Other simulations of \\citet{Luf04} perform a large parameter study of disks already containing one planet, in order to explore the possibility that its present may trigger the formation of further objects. For their parameter study, they use simulations with $10^5$ particles, but with varying minimum $Q$ or initial planet mass. The initial conditions of these disks were in most respects quite similar to those of Mayer \\etal, however the resolution employed was far lower in order to facilitate the large parameter study. Given the much lower resolution, it appears likely that these authors have under resolved the vertical structure of the disk, resulting in incorrectly low densities and pressures and the possibility of artificially enhanced fragmentation. While an attempt to duplicate their study is beyond the scope of this work, we have run selected simulations sampling their parameter space and have observed that they resolve (or fail to resolve) the Toomre condition at similar points in the parameter space for which they conclude that clumping was induced or not induced by the passage of the seed planet. Due to the probable violation of both criteria, we therefore believe that their conclusions have been contaminated by insufficient resolution and must be reexamined.  \\subsection{Beyond numerics: what is required for the viability of the disk instability model for Jovian planet formation to be either verified or falsified}\\label{sec:beyond}  In this work, it has not been our purpose to discuss all of the requirements for simulations modeling gravitational fragmentation in disks to produce correct and physically relevant results. Such a description of a real physical system requires not only numerically valid simulations, but also relevant initial conditions and a correct and complete physical model. Numerically valid simulations will be interesting only to the extent that they model real systems with real physics.  Of particular interest for the physical understanding of disk evolution is the fact that our highest resolution SPH test model did not produce {\\it any} clumps although it was evolved for much longer in time than its lower resolution cousins. This is an important physical result because the test problem employed a locally isothermal equation of state, which is usually (and erroneously-see the discussion in \\citet{DynII,Pick03}) equated with the strong cooling limit where fragmentation is nearly inevitable. The fact that clump formation did not occur in this limit, in a properly resolved simulation, may place the disk instability model for planet and brown dwarf formation on considerably less solid ground.  At the same time, the Boss simulation, implementing conditions that are much less biased toward clumping, {\\it did} produce clumps while remaining numerically valid according to the Toomre condition, though some doubt may remain regarding the resolution of the vertical structure. His simulations include a description of radiative cooling processes in 3D, and he concludes that the clumping is due both to the efficient radiative cooling and to the fact that convection in the disk's vertical direction is efficient. Convection is able to transport thermal energy out of the optically thick disk midplane to its photosphere surface at higher altitudes, so that even optically thick regions like forming clumps cool efficiently. On the other hand, previous 2D work of \\citet{DynII}, which assumed efficient vertical convection as a consequence of the radiative cooling prescription employed, did not produce clumps. Moreover, the total radiated energy emitted by the disks modeled in those simulations seriously underestimated the values observed in real systems because they were too cold. In a more complete physical model of the disk energy budget of heating and cooling processes, the disks would have to be warmer in order to reproduce the observations, making clumping even less likely.  Other fully 3D radiative transfer \\citep{MDPC} models confirm neither the Boss conclusion about the efficiency of vertical convection nor the clumping that results from it. The authors continue to investigate the origin of the discrepancies, but no certainty yet exists. Possibilities include both differences in the numerical treatments, for example of the boundaries, the resolution of both the Poisson solver and the hydrodynamic codes themselves, and of the physical models, for example the treatment of external radiation field, what opacities are used in the radiative transport and the equation of state for the gas. The important point to take from these discrepant results is that not only do questions of numerical validity remain to be addressed, but also the physical models. Although progress has certainly been made, to date an insufficient understanding of what physical processes are important for fragmentation has been developed. Determination of the true viability of the disk instability model for Jovian planet formation therefore still lies in the future."
    },
    "0609/astro-ph0609170_arXiv.txt": {
        "abstract": "We present optical and near-IR observations of the host galaxy of GRB\\,020127, for which we measure $R\\!-\\!K_s=6.2$ mag.  This is only the second GRB host to date, which is classified as an ERO.  The spectral energy distribution is typical of a dusty starburst galaxy, with a redshift, $z\\approx 1.9$, a luminosity, $L\\approx 5$ L$^*$, and an inferred stellar mass of $M_*\\sim 10^{11}-10^{12}$ M$_\\odot$, two orders of magnitude more massive than typical GRB hosts.  A comparison to the $z\\sim 2$ mass-metallicity ($M$-$Z$) relation indicates that the host metallicity is about $0.5$ Z$_\\odot$.  This result shows that at least some GRBs occur in massive, metal-enriched galaxies, and that the proposed low metallicity bias of GRB progenitors is not as severe as previously claimed.  Instead we propose that the blue colors and sub-$L^*$ luminosities of most GRB hosts reflect their young starburst populations.  This, along with the locally increased fraction of starbursts at lower stellar mass, may in fact be the underlying reason for the claimed low metallicity bias of $z\\lesssim 0.2$ GRB hosts. Consequently, GRBs and their host galaxies may serve as a reliable tracer of cosmic star formation, particularly at $z\\gtrsim 1$ where the $M$-$Z$ relation is systematically lower, and the star formation rate is dominated by sub-$L^*$ galaxies. ",
        "introduction": "\\label{sec:intro}  Observations of the host galaxies of long-duration $\\gamma$-ray bursts (GRBs) have led to the general consensus that they are faint and blue. This has been interpreted as evidence for intense star formation activity, as well as low metallicity and stellar mass \\citep{fjm+03,fdm+03,chg04,fls+06,sgb+06}.  In particular, a comparison of $z\\lesssim 0.2$ GRB hosts to Sloan Digital Sky Survey galaxies suggests that GRBs occur preferentially in low metallicity galaxies with $Z\\sim 0.1-0.5$ Z$_\\odot$ \\citep{sgb+06}.  These authors also propose an upper metallicity cutoff of about 0.15 Z$_\\odot$ for typical GRBs (i.e., with $E\\sim 10^{51}$ erg).  At higher redshifts, $z\\sim 0.5-1$, a comparison to the host galaxies of GOODS supernovae indicates that GRBs tend to occur in lower luminosity galaxies (by about 1 mag) and therefore presumably at lower metallicity \\citep{fls+06}.  Metallicities measured from afterglow absorption spectra (at $z>2$) span a wider range, $Z\\sim 0.05-0.5$ Z$_\\odot$ (e.g., \\citealt{bpc+06,pro06}).  Theoretical studies have also argued for ``low'' metallicity of GRB progenitors \\citep{mw99}, although we stress that the definition is not well-quantified, and black hole remnants can apparently be formed even above solar metallicity \\citep{hfw+03}.  These observations raise two crucial and related questions: (i) {\\it Is there in fact a low metallicity bias for GRB progenitors, in both an absolute sense and relative to the average metallicity at any redshift?}  (ii) {\\it Can GRBs and their host galaxies be used as a representative tracer of star formation at high redshift ($z\\gtrsim 1$)?}  The answer to these questions is of great importance given the unique ability of GRBs to probe the interstellar medium and star forming environments of galaxies over a wide redshift range, extending to $z>6$ \\citep{bcc+06}.  It is also crucial to understand the effect of metallicity in light of the redshift evolution of the average metallicity and mass function of galaxies (e.g., \\citealt{esp+06}).  Equally important, the metallicity range for GRB progenitors impacts our understanding of the progenitor population and GRB formation scenarios \\citep{mw99,hfw+03}.  The prevalence of GRBs in blue galaxies may also reflect an observational bias against dusty galaxies due to obscuration of the optical afterglow.  Two lines of evidence suggest that this may not be a significant problem.  First, the host galaxies of the ``optically dark'' GRBs (those with clear evidence for dust obscuration in the optical band) are typically not redder than the hosts of optically-bright GRBs \\citep{bck+03,fdm+03}.  A notable exception is GRB\\,030115, with a dust-reddened afterglow and a host galaxy with $R\\!-\\!K\\approx 5.4$ mag \\citep{lfr+06}.  Second, radio, submillimeter, and IR observations have not led to a preferential detection of the dark GRB hosts \\citep{bbt+03,bck+03,lcf+06}, despite an expectation that these galaxies should have obscured star formation.  Conversely, the host galaxies with long-wavelength detections (and hence obscured star formation), have blue colors typical of the GRB host population as a whole \\citep{bck+03}.  Here we present optical and NIR observations of GRB\\,020127, which reveal that the host galaxy is an extremely red object (ERO) with $R\\!-\\!K_s=6.2$ mag.  The GRB position is obtained from X-ray and radio observations of the afterglow.  The host SED is best fit by an obscured starburst galaxy template at $z\\approx 1.9$, with a stellar mass\\footnotemark\\footnotetext{We use Vega magnitudes (unless otherwise noted), and the standard cosmology with $H_0=71$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m=0.27$, and $\\Omega_\\Lambda= 0.73$.} of $\\sim 10^{11.5}$ M$_\\odot$ and an inferred metallicity \\citep{esp+06} of about 0.5 Z$_\\odot$, suggesting that GRB progenitors can in fact exist at near-solar metallicity, and may occur in the most massive galaxies at $z>1$. ",
        "conclusions": "\\label{sec:disc}  The host galaxy of \\grb, identified through X-ray and radio afterglow observations, is the reddest, and most luminous and massive GRB host discovered to date.  Unfortunately, the early optical limits (\\S\\ref{sec:obs}) are too shallow to ascertain whether the afterglow was significantly extinguished by dust; from the X-ray flux at 4.14 d we expect $R\\sim 21$ mag at the time of the early optical observations, at the level of the available limits.  For GRB\\,030115, the only other burst with an ERO host galaxy, the afterglow itself was shown to be significantly dust extinguished, with $R\\!-\\!K\\approx 6$ mag \\citep{lfr+06}.  It is therefore possible that dusty galaxies are under-represented in the current GRB host sample, due to obscuration of the optical afterglow, but we stress that the hosts of several dark bursts are not significantly redder than the median of the sample \\citep{bck+03}.  The discovery of a GRB in a massive, and hence metal-enriched, starburst galaxy indicates that the proposed low metallicity bias of GRB progenitors ($\\lesssim 0.1$ Z$_\\odot$; \\citealt{fls+06,sgb+06}) is not likely to be as severe as previously claimed, particularly at $z\\gtrsim 1$.  This appears to be true not only in terms of an absolute metallicity cutoff, but also relative to the mass-metallicity relation of galaxies at high redshift.  In particular, while the hosts of GRBs at $z\\lesssim 0.2$ appear to have metallicities at the low end of the distribution for local galaxies \\citep{sgb+06}, at least some GRBs at $z\\sim 2$ occur in the most metal-enriched galaxies {\\it at that redshift}.  This conclusion is also supported by the metallicities derived from afterglow absorption spectra (at $z>2$), which range up to solar values \\citep{bpc+06,pro06}, as well as by the detection of submillimeter emission from some GRB hosts \\citep{bck+03}.  Since GRB progenitors appear to occur at least up to $\\sim 0.5$ Z$_\\odot$, and given that the $M$-$Z$ relations at $z\\sim 1$ and $\\sim 2$ are systematically lower by about 0.3 and 0.6 dex, respectively, compared to the local relation \\citep{thk+04,sgl+05,esp+06}, we conclude that even if a slight low metallicity bias does exist, its effect will diminish beyond $z\\sim 1$.  Instead, the blue colors of GRB host galaxies may reflect their young stellar populations.  This is likely related to the fact that GRB progenitors are massive stars, which explode within a few million years of formation, thereby leading to the selection of young starburst galaxies.  In fact, for $z\\sim 1-2$, the average $R-K$ color of a 0.1 Gyr population is about 2.5 mag, in good agreement with the observed colors of GRB hosts, while for 0.3 and 1 Gyr populations it is about 1 and 2.5 magnitudes redder, respectively (Figure~\\ref{fig:r-k}).  This explanation might also account for the transition at low redshift ($z\\lesssim 0.2$) from a representative to a predominantly low luminosity host population, since locally young starburst activity occurs primarily in low mass galaxies.  \\citet{bcw+04} show that locally the fraction of galaxies with recent starburst activity increases strongly with decreased mass.  Similarly, \\citet{bpw+05} show that while at $z\\sim 0.7$ nearly half of all galaxies with $M\\gtrsim 10^{10}$ M$_\\odot$ undergo intense star formation activity, locally, this number is less than $1\\%$.  In this scenario, the driving parameter is stellar population age, which may be misidentified as a low metallicity bias.  Clearly, a study of the stellar population ages of the low redshift GRB hosts is required to assess whether age, and not metallicity, in fact plays the underlying role.  Thus, we conclude that at least at $z\\gtrsim 1$, GRBs and their host galaxies are likely to trace star formation in a relatively unbiased way, with possibly a preference for younger starburst populations.  At $z\\sim 0$ there seems to be a bias in favor of low luminosity/mass galaxies \\citep{sgb+06}, but this may be driven by stellar population age and not metallicity.  The luminosity function of GRB hosts (see also \\citealt{jbf+05}) therefore indicates that the bulk of the star formation at $z\\gtrsim 1$ takes place in galaxies fainter than L$^*$, in good agreement with the steep faint-end slope of the luminosity function of high redshift galaxies."
    },
    "0609/astro-ph0609200_arXiv.txt": {
        "abstract": "We derive the second and most stringent limit to date of the X-ray/radio flux ratio (\\fxfr) for the radio bursts associated with the recently identified source class, the Rotating Radio Transients (RRATs).  We analyze 20.1 hr of \\rxte/PCA observations of \\rrat\\ -- a period during which 350\\ppm23 RRAT radio bursts occurred, based on the previously observed average radio burst rate. No X-ray bursts were detected, implying an upper-limit on the X-ray flux for RRAT-bursts of $\\fxlim<$1.5\\tee{-8} \\cgsflux\\ (2-10\\,\\keV), or a luminosity $<2.3\\tee{37} (d/3.6\\kpc)^2$\\cgslum.  The time-average burst flux is $<$2\\tee{-13} \\cgsflux\\ (0.5-8 keV) -- a factor of 10 below that of the previously identified persistent X-ray counterpart. Thus, X-ray bursts from the RRAT are energetically unimportant compared with the persistent X-ray emission.  From the previously observed burst radio flux, we derive an upper-limit $\\fxfr\\leq4.2\\tee{-12} \\cgsflux\\perval{\\rm mJy}{-1}$ for the radio bursts from this RRAT, the most stringent to date, due to the high radio flux of bursts from this source.  The \\fxfr\\ ratio is a factor $\\sim$80 larger than that of the millisecond pulsar PSR~B1821-24; thus emission processes of X-ray/radio efficiency comparable to MSP pulses cannot be ruled out. However, if the RRAT burst emission mechanism is identical to the \\msec\\ bursts of magnetars, then the \\msec\\ bursts of magnetars should be easily detected with radio instrumentation; yet none have been reported to date. ",
        "introduction": "Recently, a new observational class of neutron stars has been proposed (\\citealt{mclaughlin06}, M06 hereafter), Comprised of eleven sources, the class is defined by the phenomenon of fast (2-30 ms) radio bursts, which repeat on timescales of 4\\,min - 3\\,hr.  Periodicities were claimed to be found in the range of $P=$0.4-7\\,s for ten of the eleven sources, through period folding; this suggested association with rotating neutron stars.  In three of the eleven sources, period derivatives were claimed, with one suggesting a magnetic field strength of 5\\tee{13} G -- a possible magnetar.  Because the $P-\\dot{P}$ analysis remains unpublished and undescribed, we regard association of the radio-bursting population with neutron stars to be an unsupported hypothesis; nonetheless, we adopt the hypothesis and nomenclature of referring to the sources as Rotating RAdio Transients (RRATs), and consider their association with neutron stars to be the leading hypothesis.  One of these -- \\rrat -- has a $7\\arcsec\\ \\times32\\arcsec$ error ellipse (M06)\\footnote{It is unclear why the positional uncertainty from this source is so much smaller than, for example, the $15\\arcmin\\times7\\arcmin$ error region for RRAT J1911+00, obtained in the same study with the same instrumentation.}, which permitted discovery of an X-ray counterpart, CXOU J181934.1$-$145804 \\citep{reynolds06}, with absorption-corrected flux of 2\\tee{-12} \\cgsflux\\ (0.5-8 keV), or 3.1\\tee{33} \\cgslum\\ at the $d=3.6$\\,\\kpc\\ dispersion-measure-indicated distance.  Two-hundred twenty nine radio bursts at $\\nu=$1.4 GHz, with average half-duration $w_{50}$=3\\,\\msec\\ were observed in 13 hours of radio integration, the brightest with a peak flux of $S_\\nu=3.6\\,\\Jy$.  If we take $w_{50}$ to be the light crossing time of the emission region, this implies a source brightness temperature (e.g., \\citealt{guedel02}):  \\begin{equation} T_b = 3\\times10^{23} \\left(\\frac{S_\\nu}{3.6\\, {\\rm Jy}}\\right) \\left( \\frac{ 1.4\\, {\\rm GHz}}{\\nu}\\right)^2 \\left( \\frac{3\\,{\\rm ms}}{w_{\\rm 50}} \\right)^2 \\left( \\frac{d}{3.6\\, {\\rm kpc}} \\right)^2 \\, {\\rm K } \\end{equation}  \\noindent  This is significantly above the brightness temperature limit for inverse Compton scattering ($T_B\\sim\\ee{12} K$) which likely requires a coherent emission mechanism for this and other RRATs, possibly accompanied by X-ray emission.  In this paper, we analyze \\rxte/PCA data, and search for X-ray bursts which may accompany the observed radio bursts, finding none.  In \\S~\\ref{sec:obs}, we describe the observations and analysis; in \\S~\\ref{sec:results}, we present the results of the X-ray burst search; and in \\S~\\ref{sec:con}, we discuss the implications of these results and conclude. ",
        "conclusions": "\\label{sec:con}  \\rrat\\ is the most prolific radio-burster of all the reported RRATs -- 229 radio bursts detected in 13\\,hr of radio observations.  On average, we would expect 350\\ppm23 such radio bursts to have occurred during these observations, assuming the radio burst rates do not change with time.  None of these had a corresponding X-ray flux $\\geq$1.5\\tee{-8} \\cgsflux (2-10 keV).  We expect 2, on average, to have been as radio bright as the brightest previously observed (3.6 Jy), which implies that $\\geq$1 such bursts were emitted (95\\% confidence).  This, in combination with our derived burst flux limit, results in an X-ray to radio flux ratio limit for bursts from \\rrat\\ of approximately:  $$ \\fxfr \\leq 4.2\\tee{-12}\\, \\cgsflux\\,{\\rm mJy}^{-1} $$  \\noindent This limit is a factor of 14 lower than that derived for RRAT~J1911+00 \\cite{hoffman06}, due to the factor of 14 greater radio flux observed from the brightest radio burst from \\rrat\\ (M06).  It may be instructive to compare this ratio with that from other transient coherent emission processes.  The limit corresponds to \\fxfr$\\leq\\ee{14.7}$ Hz, comparable to the average ratio  observed from stellar coronal activity \\citep{guedel93}; however the timescales for coronal flares (for example) are typically $\\sim10^3\\,\\sec$, much longer than the \\msec\\ flares of RRATs.  Compared to MSPs, this upper-limit to \\fxfr\\ is a factor 84 above the observed ratio for (for example) the millisecond pulsar PSR~B1821-24 ($\\fxfr = 5\\tee{-14} \\cgsflux \\perval{\\rm mJy}{-1}$; \\citealt{hoffman06}).   Thus, this limit does not exclude an emission mechanism for RRAT bursts with comparable X-ray to radio efficiency to MSPs emission mechanism.  The claim of measured $P-\\dot{P}$ values (M06) which imply magnetic field strengths approaching those of magnetars, makes comparison between the RRAT-bursts and those from magnetars a natural one, with the expectation that the emission mechanisms may be similar.  The absence of X-ray/radio burst flux limits for the \\msec\\ X-ray bursts of magnetars \\citep{gogus99,gogus01,gavriil02,kaspi03,gavriil04,woods05} precludes direct comparison with the emission mechanism powering those sources. However, a radio observation of such would be valuable to compare with the derived \\fxfr\\ limits for RRATs.  Taking 1E~2259+586 as an example \\citep{gavriil04}, with median duration of $\\sim$100 \\msec, and peak fluences of \\ee{-9} \\erg\\perval{\\rm\\,cm}{-2}, magnetar bursts would produce coincident radio bursts of $\\geq2500\\mJy$, if the magnetar bursts were due to a similar emission mechanism as RRATs. This radio flux is comparable to the brightest RRAT radio bursts yet observed (M06), and are therefore readily detectable, yet no such bursts have been reported.  However, the burst rate of the \\msec\\ X-ray bursts of magnetars is not constant in time; whether this is true for the radio bursts of RRATs has not be described in the reported observations (M06).  On the other hand, a non-detection of radio bursts from magnetars to flux limits below this value would require a different \\fxfr\\ ratio for the \\msec\\ bursts from magnetars, and may therefore imply a distinct emission mechanism.  A burst flux limit $F_{\\rm lim}<$1.5\\tee{-8} \\cgsflux (2-10 keV), averaged over 3\\,\\msec, combined with an average burst rate of 228/13 hr (M06), results in time-average burst flux of $<$2.2\\tee{-13} \\cgsflux\\ (2-10 keV), or, extrapolating to the energy band for the X-ray counterpart \\citep{reynolds06} $<$2.1\\tee{-13} \\cgsflux (0.5-8 keV).  The time-average burst X-ray flux is a factor of 10 less than the persistent X-ray counterpart in this pass-band \\cite{reynolds06}.   Thus, X-ray bursts related to the radio bursts of this RRAT are energetically unimportant compared to the persistent X-ray flux.  In addition, the time-average burst X-ray efficiency $\\eta=L_{X}(0.5-8\\,\\keV)/\\dot{E}$, compared with the reported dipolar emission $\\dot{E}=24.94(5)\\tee{31} \\cgslum$ (M06)  is $\\eta\\leq7\\tee{-3} (d/3.6\\,kpc)^2$.  It is unlikely that great improvement will be made in the X-ray flux limit presented here using existing instrumentation.  An improved flux limit cannot be obtained with {\\em Chandra}: 2 counts in the HRC-S detected within 2\\,\\msec\\ corresponds to 3\\tee{-8} \\cgsflux (extrapolated to 2-10 keV).  On-source \\rxte/PCA observations could improve on the present flux sensitivity by a factor of $\\times2$."
    },
    "0609/astro-ph0609036_arXiv.txt": {
        "abstract": "We investigate the conditions of star formation in the Large Magellanic Cloud. We have conducted a survey for water maser emission arising from massive young stellar objects in the 30 Doradus region (N\\,157) and several other H\\,{\\sc ii} regions in the Large Magellanic Cloud (N\\,105A, N\\,113 and N\\,160A). We have identified a new maser source in 30\\,Dor at the systemic velocity of the LMC. We have obtained 3$-$4\\,$\\mu$m spectra, with the ESO Very Large Telescope, of two candidate young stellar objects. N\\,105A\\,IRS1 shows H recombination line emission and its Spectral Energy Distribution (SED) and mid-infrared colours are consistent with a massive young star ionising the molecular cloud. N\\,157B\\,IRS1 is identified as an embedded young object, based on its SED and a tentative detection of water ice. The data on these four H\\,{\\sc ii} regions are combined with mid-infrared archival images from the Spitzer Space Telescope to study the location and nature of the embedded massive young stellar objects and signatures of stellar feedback. Our analysis of 30\\,Dor, N\\,113 and N\\,160A confirms the picture that the feedback from the massive O and B-type stars, which creates the H\\,{\\sc ii} regions, also triggers further star formation on the interfaces of the ionised gas and the surrounding molecular cloud. Although in the dense cloud N\\,105A star formation seems to occur without evidence of massive star feedback, the general conditions in the LMC seem favourable for sequential star formation as a result of feedback. In an appendix we present water maser observations of the galactic red giants R\\,Doradus and W\\,Hydrae. ",
        "introduction": "One of the great unknowns in our understanding of star formation concerns the role played by galactic environmental parameters. For instance, molecular cloud processes, e.g. cooling and magnetic field diffusion, depend on the presence of metals, thus it is unlikely star formation is not affected by metallicity. The Magellanic Clouds have metallicity and density of the interstellar media (ISM) that are lower than in the Milky Way. Thus, they provide a star formation template that is more representative of the formation of stars at higher redshifts and offer a step closer towards exploring the formation of the first generation of stars at zero metallicity.  Young pre-main-sequence populations have recently been identified in H\\,{\\sc ii} regions in the Magellanic Clouds \\citep{nota06,gouliermis06}. However, relatively little work has been done in investigating the earlier, more embedded stages of star formation in the Magellanic Clouds. At the distance of the Large Magellanic Cloud (LMC, $\\sim$\\,50\\,kpc), one is limited to study only the most massive stellar embryos that will form massive O and B stars. Infrared (IR) observations are an excellent way to investigate the properties of such objects, both because they are heavily embedded and thus invisible at shorter wavelengths and also because they are surrounded by dense envelopes of gas, dust and ices that are revealed at these wavelengths \\citep{vandishoeck04}. Recently, a few massive young stellar objects (YSOs) in the LMC have been investigated \\citep{vanloon05b,jones05}.  Molecular masers (in particular water, hydroxyl and methanol) are very bright emission lines that seem to be closely associated with the earliest stages of massive star formation \\citep{debuizer05} when the YSOs are difficult to detect even at IR wavelengths. Maser emission is found in the vicinity of embedded YSOs and thus they are powerful beacons of current star formation in a molecular cloud. Several water maser sources have been discovered in the Magellanic Clouds (see below). In this contribution we describe a new survey of water masing sources in the LMC at 22\\,GHz with the Parkes Telescope.  In recent years, the 30 Doradus Nebula (N157) has become the paradigm of large scale triggered star formation. As the largest and nearest extragalactic H\\,{\\sc ii} region in the Local Group it offers a rare insight into the spatial and temporal properties of starbursts. \\citet{walborn99} identified the so-called star-formation fronts, the interfaces between the energetic outflows from the central compact cluster R\\,136 (first generation star formation episode) and the surrounding molecular clouds, in which there is evidence of ongoing triggered star formation. Water maser emission has also been detected at various locations in 30\\,Dor \\citep{whiteoak83,vanloon01,lazendic02}. The wealth of bright protostars identified in the region \\citep{rubio98,brandner01} hints at the possibility that more masing activity hitherto undetected might be present. We describe our mapping survey of the inner region of 30\\,Dor, in which we have identified 4 distinct maser sources, one of them new.  Three other H\\,{\\sc ii} regions in the LMC, known to show water maser emission \\citep[N\\,113, N\\,105A and N\\,160A;][]{whiteoak83,whiteoak86} are also observed at 22\\,GHz, as well as most of the previous non-detections for water maser emission in the LMC and SMC \\citep{scalise82,whiteoak83}. Since our observations, \\citet{lazendic02} also reported 22\\,GHz water maser emission towards N\\,159.  The H\\,{\\sc ii} region N\\,113 \\citep{henize56} shows a complex structure of H$\\alpha$ bubbles with a rich molecular gas and dust morphology. Several young clusters are associated with N\\,113 \\citep*{bica92} and current star formation activity is occuring within continuum sources in the central area of the nebula \\citep{brooks97,wong06}. The H\\,{\\sc ii} region N\\,105 \\citep{henize56} is a complex of evacuated bubbles and dense molecular material with several young clusters associated with it \\citep{ambrocio98}; current star formation as indicated by maser activity seems to concentrate in the denser central part of the region, N\\,105A. N\\,160A is the brightest component in the N\\,160 H\\,{\\sc ii} complex \\citep{henize56}. Besides maser emission, both an embedded protostar \\citep{henning98} and a dense molecular core \\citep{bolatto00} point at ongoing star formation. \\citet{malayeri02} and \\citet{nakajima05} describe the gas morphology and stellar content of N\\,160.  In order to investigate the relationship between the masers and the gas kinematics and star formation activity, we have obtained narrow-band H$\\alpha$ images at the AAT of N\\,113, and L$^\\prime$-band images and 3$-$4\\,$\\mu$m spectroscopy at the ESO-VLT in N\\,157B and N\\,105A. These new observations, combined with archival IRAC/Spitzer images are used to relate the embedded population to the local gas and dust morphology. ",
        "conclusions": "We have conducted a survey for water maser emission in 30\\,Doradus and a sample of H\\,{\\sc ii} regions in the Large Magellanic Cloud, to investigate the conditions under which current star formation occurs. The locations of the maser sources are compared with infrared images from the ESO/VLT (at 3.8\\,$\\mu$m) and the Spitzer Space Telescope, and an H$\\alpha$ image in the case of N\\,113. We also present 3$-$4\\,$\\mu$m spectroscopy to investigate the nature of two protostar candidates, in the N\\,157B and N\\,105A regions, showing evidence for the onset of the stellar feedback process in one of them.  Our water maser survey of 30\\,Dor uncovered one new source, 0539$-$691C, with a velocity consistent with the systemic velocity of the LMC. All detected water masers are located in the densest part of the nebula at the interface between neutral and ionised gas. N\\,105A\\,IRS1 shows strong H recombination line emission and its SED and IR colours are consistent with an embedded young massive star ionising its immediate surroundings. It also shows evidence for outflows. A nearby diffuse IR source, N\\,105A\\,blob, shows both H recombination and PAH emission. We identify N\\,157B\\,IRS1 as an embedded protostar based on the analysis of its SED and a tentative detection of the 3.1\\,$\\mu$m water ice feature.  In the well-developed H\\,{\\sc ii} regions 30\\,Dor, N\\,113 and N\\,160A, no water masers have been detected deep within the molecular cloud complexes. They are always found at the interfaces between molecular cloud and H\\,{\\sc ii} region. This provides strong evidence that feedback from massive stars triggers subsequent star formation. Although in the dense cloud N\\,105A star formation seems to occur without evidence of massive star feedback, the general conditions in the LMC seem favourable for sequential star formation as a result of feedback.  The wind speed and mass-loss rate of O and B stars decrease at lower metallicity (Mokiem \\& de Koter, in preparation), thus massive star feedback could be weak in metal-poor environments but we see no evidence for this in the LMC (metallicity $\\sim$40 per cent solar). This is important as the duration and intensity of star formation epochs in galaxies may depend on the efficiency of the local massive star feedback to trigger further star formation."
    },
    "0609/cond-mat0609598_arXiv.txt": {
        "abstract": "The new mathematical framework based on the free energy of pure classical fluids presented in [R. D. Rohrmann, Physica A {\\bf 347}, 221 (2005)] is extended to multi-component systems to determine thermodynamic and structural properties of chemically complex fluids. Presently, the theory focuses on $D$-dimensional mixtures in the low-density limit (packing factor $\\eta < 0.01$). The formalism combines the free-energy minimization technique with space partitions that assign an available volume $v$ to each particle. $v$ is related to the closeness of the nearest neighbor and provides an useful tool to evaluate the perturbations experimented by particles in a fluid. The theory shows a close relationship between statistical geometry and statistical mechanics. New, unconventional thermodynamic variables and mathematical identities are derived as a result of the space division. Thermodynamic potentials $\\mu_{il}$, conjugate variable of the populations $N_{il}$ of particles class $i$ with the nearest neighbors of class $l$ are defined and their relationships with the usual chemical potentials $\\mu_i$ are established. Systems of hard spheres are treated as illustrative examples and their thermodynamics functions are derived analytically. The low-density expressions obtained agree nicely with those of scaled-particle theory and Percus-Yevick approximation. Several pair distribution functions are introduced and evaluated. Analytical expressions are also presented for hard spheres with attractive forces due to K\\^ac-tails and square-well potentials. Finally, we derive general chemical equilibrium conditions. ",
        "introduction": "\\label{s.intr}  The detailed knowledge of atomic and molecular po\\-pu\\-la\\-tions is essential to analyze and evaluate accurately thermodynamic properties and chemical processes in fluids as well as their monochromatic opacities. In particular, they are required in studies of stellar interiors, atmospheres and circumstellar envelopes. Gas densities in stellar atmospheres and circumstellar structures are low, they are hardly higher than 0.001 g\\,cm$^{-3}$. Nevertheless, in these cases it is necessary to deal with a high variety of particles and electronic configurations and therefore, with a series of physical processes to calculate the radiation emitted by the star accurately \\cite{kurucz,hubeny}.\\par Among the most difficult challenges in determining gas models is to obtain self-consistently atomic densities and opacity absorption coefficients by taking into account the non-ideal effects. The source of non-ideal effects is the presence of Coulombic interactions between charged particles as well as due to the short-ranged interactions between neutral-neutral and neutral-charged particles. Due to its relevance and complexity, the elaboration of gas models from mechanical statistical approaches is still an active area of research  today \\cite{fontaine,magni,robnik,mihalas,rogers, saumon,potekhin,lisal,juranek,wini}.\\par The present study is motivated by the difficulties of predicting quantitatively the optical properties of stellar atmospheres. In typical gas models the opacity data are obtained through a complimentary calculation once the equations of state and mean abundances of bound states are obtained. This procedure can produce, however, internal inconsistencies in a model, mainly when the abundances of atoms are calculated using internal energies that are different from those assumed to obtain the radiative transitions \\cite{paper1,rohrmann}. This can lead to serious disagreements between predictions and observations or experiments (e.g., the unphysical Lyman opacity \\cite{bergeron}). Therefore, it is highly desirable to develop a theoretical formalism that treats atoms subjected to different perturbations by considering them in separate groups and evaluate accordingly the populations and opacities in a consistent way. Such a gas model has then to consider explicitly not only all atomic excited states but also do they need to discriminate different perturbation states with the use of an appropriate particle-state variable. Our work is then oriented to develop a model with these characteristics.\\par With the aim of establishing a detailed and accurate equation of state for fluids, in a previous work we have developed a gas statistical formalism which combines free energy minimization methods with space partitions \\cite{paper1}. The novel feature of this formalism, as compared to other more standard and well-established statistical-mechanical theories of the fluid/liquid state, is the attempt of keeping trace, in a thermodynamically self-consistent way, of a more refined structural information on the local environment of a given particle through an extra parameter, the ``available volume'', which is determined by the distance of the reference particle to its closest neighbor. For noninteracting particles, the available volume ``$v$'' is equivalent to the spherical volume whose radius is the distance between the respective centers of the particle and its nearest neighbor. This volume is considerably smaller for particles with repulsive interactions. The theory \\cite{paper1} was formulated for one-component dilute fluids, where the Helmholtz free energy is written in terms of the occupation number distribution $N_v$ of the $v$-variable and where the particle interactions are introduced using pair potentials. The result thus obtained represents a unified treatment of thermodynamics and the structure of fluids. The formulation was applied to hydrogen atoms in an electrically neutral medium of charged particles encompassing ions and electrons. The atomic populations were derived by establishing groups of atoms having different plasma perturbations measured according to the size of their available space volumes $v$. Although the results obtained are only illustrative of its possible performances, this theory represents the first attempt of evaluating in detail the populations by combining consistently the plasma interaction effects.\\par Approaches based on space partitions were attempted in the past with the so-called ``cell\" theories to obtain equations of state for fluids. However, the character of these space divisions is quite different as compared to that employed in \\cite{paper1} and in the present paper. The cell models \\cite{hcb,hill,kirk} were originally developed by Eyring and Hirschfelder \\cite{eyring} and by Lennard-Jones and Devonshire \\cite{lennard}. They were inspired by the lattice-like configurations acquired by some systems at high densities. In such theories the fluid volume is divided into an imaginary lattice of cells. Each particle is then confined into a cell, which implies that a space partition is assumed beforehand, instead of being deduced from the model proper. Besides, by ignoring the interchange of particles between cells, an error is introduced in the system-entropy calculation, which is usually corrected using an arbitrary ``communal entropy'' factor.\\par The perspective of the theory presented in \\cite{paper1} is new, as it determines a space division by assigning the $v$-volume to the particles. This is achieved by minimizing the system free energy under the condition that the sum of all $v$-volumes be equivalent to the entire volume of the fluid. In this procedure, no particle confinement is established, because the assignment of an individual volume to each particle does not introduce any ownership connection between space regions and particles, i.e., $v$ does not carry any information on the particle coordinates. Therefore, the formulation presented in \\cite{paper1} is free from difficulties that are inherent to the cell theories, i.e. entropy corrections, and it can provide new results, such as the combined thermodynamical and structural information of a fluid, which cannot be derived from the cell models.\\par In the present paper we generalize the formalism to $D$-dimensional multi-component fluids that was initially developed in \\cite{paper1} only for pure or single-component substances. It aims at giving particularly simple analytic forms to provide a unified derivation of thermodynamic and geometric properties that characterize them. The formulation of the theory is general, but attention is focused on its application to hard sphere models, which represent the repulsive interactions among atoms and geometrically simple molecules \\cite{barker}. We also include the analysis of hard spheres with weak long-range attraction (K\\^ac type) and hard spheres with square-well potential. The approximations discussed here should be suitable for more realistic short-ranged potentials, like those governing many atomic and colloidal fluids. Our introduction of space partitions in statistical thermodynamics marks a considerable progress in the description of gases with respect to previous theories, since it gives the possibility of obtaining thermodynamic and structural information on fluids simultaneously and self-consistently. The present extension to multicomponent fluids enables us to treat arbitrarily complex chemical mixtures by including systematically the inter-particle interaction potentials. The theory can have many applications, not only in Astrophysics, but also in other fields such as the analysis of colloidal and protein solutions in biochemistry.\\par In the following sections, we first introduce a description of states of a gas mixture based on space partitions. General expressions for the Helmholtz free energy are presented in Sect. \\ref{s.Helm} and the thermodynamic equilibrium states are derived in Sect. \\ref{s.equ_sta}. We analyze there some non-conventional thermodynamic variables that are generated by the theory and derive general expressions to calculate chemical potentials. Then, applications of the formalism to a few multi-component systems based on hard particles with and without attractive forces are given in Sect. \\ref{s.hard}. The results obtained are compared with those from Percus-Yevick and scale particle theories. In Sect. \\ref{s.chem} we extend the theory to include chemical reactions in the gas. The conclusions are presented in Sect. \\ref{s.concl}.\\par ",
        "conclusions": "\\label{s.concl}  We have developed a mathematical framework in terms of a Helmholtz free-energy model which aims at describing equilibrium properties of $D$-dimensional multi-component fluids. This theory is an extension of a former model devoted to one-component fluids and combines free-energy minimization methods with space partitions based on the available volume $v$ to a particle and the occupation number distributions of this variable. Thus we obtain an unified formalism that treats thermodynamics and structure of complex fluids simultaneously.\\par The theory is presently developed for the case of dilute fluids, where inter-particle interactions are evaluated from pair potentials. Two closure relations (\\ref{e.nilninl}) and (\\ref{e.vid}) are required to obtain a closed formalism. The former is associated to the particle number conservation (\\ref{e.Nil}) and provides the probability that a particle had a specific species as the closest neighbor. The second closure relation (\\ref{e.vid}) specifies the available-volume variable and, therefore, is linked to the filling space condition (\\ref{e.Vcons}). Consequently, unconventional thermodynamic variables were derived by introducing space partitions in the thermo-statistical analysis of fluids. They are Lagrange's multipliers related to the specific particle number and volume conservations, which describe thus chemical potentials and the pressure, respectively. In addition, we have obtained general expressions for the excess chemical potentials and chemical equilibrium conditions of non-ideal fluids, as well as several pair distribution functions.\\par The theory provides a straightforward way to calculate thermodynamic quantities and structure properties of gas mixtures. As illustrative examples, we have considered several hard sphere systems for which simple analytical expressions were derived. The equations of state obtained for pure hard-sphere systems are in good agreement with the scaled-particle and Percus-Yevick solutions valid for low-densities. We reckon that this agreement confirms a posteriori the validity of the several approaches adopted. Besides, well-known solutions for van der Waals mixtures were recovered easily by using generalized K\\^ac potentials. We also have derived closed analytical forms for thermodynamic quantities of square-well mixture fluids. These evaluations constitute a reference for future studies of short ranged attractive systems. The influence of repulsive and attractive forces on the space partition among particles was analyzed. The results thus obtained express a close relationship between statistical geometry and statistical mechanics.\\par The formalism introduced in the present work provide a useful tool to investigate the structure and the effective interactions in chemical complex fluids. It is particularly well suited to study some astrophysical gases where detailed descriptions of species and interactions are required. This occurs, in particular, in atmospheres and circumstellar envelopes, where the density can be hardly higher than some $\\eta=0.01$, so that the approximations for low densities given here can be used safely. We consider that present theory has other many applications in applied problems of current interest, such as studies of equilibrium properties of colloidal and protein solutions.  Among the most important results obtained in this work, we can quote :\\par \\begin{itemize} \\item A functional (or series of functionals for mixing) contains the entire thermodynamical information of the system and leads to a classification of particles according to their distance to the closest neighbor through the assignment of a given volume available to each particle. So, it is thus possible to distinguish the atoms according to the different degrees of perturbations carried by their neighbors. To our knowledge, there is no other model able to make this distinction up to now. \\item The current formalism is able to answer to the following inquiries :\\par \\begin{itemize} \\item  What is the probability of finding a particle of class $j$ at a distance $r$ from another particle of species $i$ whose closest neighbor is of class $k$, knowing that $r$, $i$, $j$ and $k$ are arbitrary ?\\par \\item Idem as before, but for a neighbor $k$ situated at a given (arbitrary) distance from the particle of class $i$ ?\\par \\end{itemize} \\item  The chemical equilibrium  of a fluid and the classification of particles according to the distance to the respective closest neighbor, are done simultaneously and self-consistently with the laws of thermodynamics. \\item As far as we know, it is the first time that a theory is able to provide thermodynamical variables associated with the distribution of the space between the particles of different chemical species directly. The variables are related to the conventional thermodynamical variables explicitly, such as pressure, chemical potential, etc.\\par \\end{itemize} The description level attained by our formalism overrides whatever known pre-existing model and constitutes an unparalleled tool to represent chemical mixing at low densities. It also provides conceptual elements on the connection between statistical mechanics and spatial statistics, which has never been presented up to now.\\par"
    },
    "0609/astro-ph0609750_arXiv.txt": {
        "abstract": "We report here the results of 2  observations performed simultaneously  with \\integral, \\rxte, the Ryle and Nan\\c{c}ay radio telescopes. These observations  belong to the so-called $\\nu$ and $\\lambda$ classes of variability during which  a high level of correlated X-ray and  radio variability is observed. We study the connection between the accretion processes seen in the X-rays, and the ejections seen in radio. By observing an ejection during class $\\lambda$, we generalise the fact that the discrete ejections in GRS 1915+105 occur after sequences of soft X-ray dips/spikes. We then identify the most likely trigger of the ejection through a spectral approach to our \\integral\\ data.  We show that each ejection is very probably the result of the ejection of a Comptonising medium responsible for the hard X-ray emission seen above 15 keV with \\integral. ",
        "introduction": "GRS~1915+105 is the most active microquasar of our Galaxy. An extensive review on this source can be found in Fender \\& Belloni (2004). To summarize, \\grs\\ hosts a black hole (BH) of  14.0 $\\pm$ 4.4 M$_{\\odot}$ (Harlaftis \\& Greiner 2004), it is one of the brightest X-ray sources in the sky and it is a source of superluminal ejection (Mirabel \\& Rodriguez 1994), with true velocity of the jets $\\geq 0.9c$. The source is also known to show a compact jet during its periods of low steady levels of emission (e.g. Fuchs et al.\\ 2003). Multi-wavelength coverages from radio to X-ray  have shown a clear but complex association between the soft X-rays and radio/IR behaviours. Of particular relevance is the existence  of radio QPO in the range 20--40 min associated with the X-ray variations on the same time scale (e.g. Mirabel et al.\\ 98). These so called ``30-minute cycles'' were interpreted as being due to small ejections of material from the system, and were found to correlate with the disc instability, as observed in the X-ray band.\\\\ \\indent Extensive  observations at X-ray energies with  {\\it RXTE} allowed Belloni et al.\\ (2000) to  classify all the observations into 12 separate classes (labelled with greek letters), which could be interpreted as transitions between three basic states (A-B-C): a hard state and two softer states.  These spectral changes are, in most of the classes, interpreted as reflecting the rapid disappearance of the inner portions of an accretion disc, followed by a slower refilling of the emptied region (Belloni et al.\\ 1997).  \\\\ \\indent The link between the accretion and ejection processes is, however, far from being understood and different kind of models are proposed to explain all observational facts including the X-ray low (0.1-10 Hz) frequency QPOs. It should be added that until the launch of \\integral, all studies had been made below 20 keV with the PCA onboard \\rxte, bringing, thus, only few constraints to the behaviour of  the hard X-ray emitter (hereafter called corona). Since the launch of \\integral\\ in late 2002 we have monitored  GRS 1915+105 with long exposure ($\\sim100$ ks) pointings. All the observations have been conducted simultaneously with other instruments, in particular \\rxte\\ and the Ryle Telescope, and in some cases with others (Spitzer, Nan\\c{c}ay, GMRT, Suzaku,...), with the aim of understanding the physics of the accretion-ejection phenomena, including,  for the first time, the behaviour of the source seen above 20 keV up to few hundred keV. We report here the results obtained during observations showing sequences of  X-ray hard dips/soft spikes (hereafter cycles), followed by radio flares. ",
        "conclusions": "Altough the 2 observations show clear differences, visible in the shape of the light curves, the evolution of the hardness ratio (Fig. \\ref{fig:oct04} \\& \\ref{fig:Nov04}), that of the spectral parameters, the similarity of the cycles in their sequences,  indicate that the physics connecting accretion and ejection is the same. In class $\\nu$ the disc seems to approach the compact object and gets brigther, throughout the A-B-C-D sequence, it does the same in class $\\lambda$. In both cases the evolution seems to be marked by a real change in the physical properties of the Comptonizing plasma (corona) which result in  a change of the spectral state. More importantly, when studying the evolution of the relative fluxes of the 2 components in the various intervals, while in all cases the 3-50 keV disc flux increases from the dip to the main spike, the Compton flux decreases by a factor of $\\sim2.4$ in class $\\nu$, and higher than 2.7 in class $\\lambda$. This behaviour seems to suggest that the dip following the precursor spike in the 2 classes, is due to the disappearance of the so-called ``corona'', and not simply to a pivoting of the spectrum. \\\\ \\indent Given the disappearance of one emitting medium on the one hand, and the appearance of ejected material on the other hand, it is tempting to consider that the ejection occurs at the precursor spike, and is therefore the result of the ejection of the corona. In class $\\beta$ it has been shown that the X-ray spike half way through the dip is the trigger of the ejection (Mirabel et al.\\ 1998). In addition, Chaty (1998)  has further proposed that during class $\\beta$, the spike corresponds to the disappearance of the corona, suggesting that the coronal material is ejected in this particular class too. Rodriguez et al.\\ (2002) reached the same conclusions comparing the behaviour of the source at low and hard X-ray energies during a class $\\alpha$ observation.  The case of GRS 1915+105, although spectacular, is not unique. Rodriguez et al.\\ (2003) have come to the conclusion that in the microquasar XTE J1550$-$564 during its 2000 outburst the corona is ejected at the peak of the outburst and further detected in radio. The main difference with \\grs\\ are the constants of time over which the ejections occur. While in XTE J1550$-$564, one ejection only is seen after the peak of the outburst at soft X-ray energies, in GRS 1915+105 those kind of events can occur several times within few hours (Fig. \\ref{fig:oct04}). The general behaviour, however, seems to be the same. The ejection is coincident with a disappearance of one of the emitting medium, leading to the conclusion that it is this media, the corona, that is blown away and detected later at radio wavelengths.\\\\ \\indent In addition to confirming this proposition by including a precise analysis of the hard X-rays, our observations allow us to generalise the connection of the accretion-ejection phenomena in GRS 1915+105 by observing an ejection during a class $\\lambda$ observation for the first time. We also provide evidence that in at least two cases showing dip-spike sequences, the return to a high level of soft X-ray, or the major spike, has a complex structure, and can be subdivided into a precursor spike followed by a rapid dip and a major flare. Our analysis suggests that the ejection is triggered at  the precursor spike, similarly to class $\\beta$ in which the ejection is triggered at the soft spike half way through the soft X-ray dip (Mirabel et al.\\ 1998). Preliminary analysis of an \\integral\\ observation during which ejections are observed show that the time delay between the soft X-ray spike and the peak of the radio flare at 15 GHz is of the same order of magnitude in class $\\beta$ (about 0.3 hours Rodriguez et al.\\ in prep.), vs. $\\sim$ 0.27 and 0.31 hours in classes $\\nu$ and $\\lambda$ respectively. The relative constancy of this delay further strenghtens the association of the spike and the radio flare in all classes.\\\\ \\indent Although the mechanisms giving rise to the ejections is far from being understood, our analysis answers the important question regarding  the moment of the ejection which, in turns,  brings interesting constraints on the existing models. The sequences of X-ray dip/spike could be due to a ``magnetic flood'' (Tagger et al.\\ 2004) during which a magnetic instability develops during the spectrally hard dips, advects magnetic flux in the inner regions of the disc. The X-ray spike and associated ejection would be due to a reconnection event allowing the system to get rid of the accumulated magnetic flux, blowing away the corona, and leading to the evolution to a soft X-ray state. Many points remain unsolved, though. In particular, our analysis poses the question on the relation between the approach of the disc and the ejection of the corona. Is one consequence of the other? \\indent Our multiwavelength monitoring will continue in the future with the main goal to obtain more sequences of dip-spike radio flare sequences which will therefore allow us to have better statistics on the occurence of radio flares and their connection to the X-rays, and hard X-rays behaviour, and hopefully better constrain  the models."
    },
    "0609/astro-ph0609616_arXiv.txt": {
        "abstract": "The acceleration of the expansion of the universe, and the need for Dark Energy, were inferred from the observations of Type Ia supernovae (SNe Ia)\\citep{1998AJ....116.1009R,1999ApJ...517..565P}. There is consensus that SNe Ia are thermonuclear explosions that destroy carbon-oxygen white dwarf stars that accrete matter from a companion star\\citep{1960ApJ...132..565H}, although the nature of this companion remains uncertain.  SNe Ia are thought to be reliable distance indicators because they have a standard amount of fuel and a uniform trigger --- they are predicted to explode when the mass of the white dwarf nears the Chandrasekhar mass\\citep{1931ApJ....74...81C} --- 1.4 solar masses.  Here we show that the high redshift supernova SNLS-03D3bb has an exceptionally high luminosity and low kinetic energy that both imply a {\\em super}-Chandrasekhar mass progenitor. Super-Chandrasekhar mass SNe Ia should preferentially occur in a young stellar population, so this may provide an explanation for the observed trend that overluminous SNe Ia only occur in young environments\\citep{1996AJ....112.2391H,2006sullivan}.  Since this supernova does not obey the relations that allow them to be calibrated as standard candles, and since no counterparts have been found at low redshift, future cosmology studies will have to consider contamination from such events. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/nucl-th0609052_arXiv.txt": {
        "abstract": "We construct a class of nuclear equations of state based on a schematic potential model, that originates from the work of Prakash et. al. \\cite{Prakash-88}, which reproduce the results of most microscopic calculations. The equations of state are used as input for solving the Tolman-Oppenheimer-Volkov  equations for corresponding neutron stars. The potential part contribution of the symmetry energy to the total energy is parameterized in a generalized form both for low and high values of the baryon density. Special attention is devoted to the construction of the symmetry energy in order to reproduce the results of most microscopic calculations of dense nuclear matter. The obtained nuclear equations of state are applied for the systematic study of the global properties of a neutron star (masses, radii and composition). The calculated masses and radii of the neutron stars are plotted as a function of the potential part parameters of the symmetry energy. A linear relation between these parameters, the radius and the maximum mass of the neutron star is obtained. In addition, a linear relation between the radius and the derivative of the symmetry energy near the saturation density is found. We also address on the problem of the existence of correlation between the pressure near the saturation density and the radius. ",
        "introduction": "Neutron stars (NS) are some of the densest manifestations of massive objects in the universe which provide very rich information for testing theories of dense matter physics and also provide a connection among nuclear physics, particle physics, statistical physics and astrophysics \\cite{Lattimer-04,Shapiro-83,Glendenning-96,Weber,Prakash-94,Heiselberg-00, Steiner-05,Prakash-97,Lattimer-00,Sahakian74}. The global aspects of neutron stars, such as the masses, radii and composition are determined by solving the so-called Tolman-Oppenheimer-Volkov (TOV) equations \\cite{Tolman-39,Volkov-39}. However there are large variations in predicted radii and maximum masses because of the uncertainties in the nuclear equation of state (EOS) near and mainly above the saturation density $n_s$ \\cite{Steiner-05,Dieperink-03,Dieperink-05,Stone-02,Stone-03,Stone-07, Shetty-05,Shetty-06,Klahn-06,Danielewicz,Danielewicz-06,Danielewicz-04,Wiringa-88,Li-05,Zuo-02,Douchin-01,Baldo-97,Lee-98, Liu-04,Krastev-06,Li-06,Li-06-2,Li-04,Vretenar-03}. The total energy of neutron rich matter (the case of a neutron star) can be written as a sum of two parts. The first one is the contribution of the symmetric nuclear matter (which is well known) and the second is the symmetry energy (SE) which still is uncertain although several constraints exist from ground state masses (binding energies) and giant dipole resonances of laboratory nuclei. A consequence of this uncertainty is that different models predict up to a factor of 6, variations in the pressure of neutron star matter near $n_s$, even though the pressure of symmetric matter is better known, being nearly zero at the same density. This pressure variation accounts for the nearly 50\\% variation in the predictions of neutron star radii \\cite{Lattimer-04,Lattimer-01}.  In general, the value of the SE at nuclear saturation density and mainly the density dependence of the SE are both difficult to be determined in the laboratory. The motivation of the present work is to propose a new parameterization for the potential part of the symmetry energy $E_{sym}(n)$ in order to be able to reproduce the results of a variety of microscopic models both in low and high values of the baryon density. Especially the trend of the symmetry energy just above the equilibrium density $n_s$ is a critical factor in determining the neutron star radius.  In order to calculate the global properties of neutron stars (mass, radius ets.) the hydrostatic equilibrium equations of Tolman, Oppenheimer and Volkov have to be solved once the equation of state is specified. However, the composition of a neutron star still remains uncertain and the construction of the EOS, which is based on the ingredients of the NS's and the kind of interactions which characterize them, is subjected to several assumptions. In any case the calculated EOS has to satisfy the following requirements \\cite{Baldo-97}: i) It must display the correct saturation point for symmetric nuclear matter (SNM); ii) it must give a SE compatible with nuclear phenomenology especially at high densities; iii) for SNM the incompressibility at saturation must be compatible with the values extracted from phenomenology; iv) both for neutron matter and SNM the speed of sound must not exceed the speed of light (causality condition), at least up to the relevant densities.  In the present work we consider that the neutron star core is composed only by an uncharged mixture of neurons, protons and electrons in equilibrium with respect to the weak interaction ($\\beta$-stable matter). However in general, in the range of densities $n \\geq n_s$, the hadronic phase of superdense matter with a rich spectrum of particles (hyperons, baryonic resonances, $\\pi^-$ and $K^-$ mesons, and a small portion of leptons) is realized. The model which is used for the construction of the EOS is a generalization of a schematic potential model based on a previous work of Prakash et al. \\cite{Prakash-88}. The model reproduces the results of most microscopic calculations of dense matter \\cite{Wiringa-88}. It is worthwhile to notice that there many ways to determine the equation of state through the many-body approach of interacting hadrons. Some of the most recent ones are based on variational methods \\cite{Akmal97,Akmal98,Morales02} and some are based on microscopic calculations \\cite{Zhou04}. In order to face the problem that stems from the uncertain behavior of the SE at high densities we perform a suitable parameterization both in low and high densities. In the previous work of Prakash et al \\cite{Prakash-88,Prakash-97} the parameterization of the potential term of the SE is achieved by the introduction of three different choices of the potential contribution to the SE. To advance, in the present work, we suggest a more generalized parameterization of the potential term of the SE which is more flexible and efficient, reproduces the predictions of most microscopic calculations of dense matter \\cite{Wiringa-88} and confirms the results of various empirical data.  The most striking feature of the proposed parameterization is the different form of the parameterization function $F(u)$ for densities below saturation point and for densities above this point. This is not surprising since this was already entailed in microscopic calculations. Although  the behavior of the SE for densities below the saturation point still remains unknown, significant progress has been made only most recently in constraining the SE at subnormal densities and around the normal density from the isospin diffusion data in heavy-ion collisions \\cite{Wen05, Bao05}. This has led to a significantly more refined constraint on neutron-skin thickness of heavy nuclei \\cite{Steiner05, Chen05} and the mass-radius correlation of neutron stars \\cite{Li-06,Li-06-2,Li-04}. For densities above the saturation point the trend of the SE is model dependent and exhibits completely different behavior.  The above characteristic of the SE is well reflected in our proposed models. In view of the previous comment, the proposed parameterization of the potential term of the SE has the advantage to be able to reproduce microscopic calculations in cases where the SE, at low densities, increases along with the density and then begins to fall although the density continues to increase. This is a well known characteristic of a class of Skyrme interactions \\cite{Stone-02,Stone-03,Stone-07} and of Gogny Hartree-Fock calculations \\cite{Li-06,Li-06-2,Li-04}. Special effort has been devoted to find analytical relations between the radius $R$ and the pressure $P$ which correspond to a special density $n$ for a fixed value of the mass $M$ of the neutron star. So an accurate determination of a neutron star radius will permit evaluation of the pressure of neutron star matter. All the above will provide a direct determination of the density dependence of the nuclear SE at these densities \\cite{Lattimer-01,Lattimer-03}.  Finally, we also address  on the problem of neutron star cooling \\cite{Pethick-92,Lattimer-91}. It is well known that the direct Urca process can occur in neutron stars if the proton concentration exceeds some critical value in the range of 11-15 \\%. The proton concentration can be  determined by the trend of the SE especially just above the equilibrium density. So, the detailed knowledge of the SE behavior is crucial for the existence of the direct Urca process.  The plan of the paper is as follows. In section $2$ the proposed model and the relatives formulas are discussed and analyzed. Results are reported and discussed in section $3$, while the summary of the work is given in section $4$. ",
        "conclusions": "First we apply our model in a simple case where the potential part of the SE is parameterized as $F(u)=u^c$ and the total SE contribution can be written as follows \\begin{equation} E_{sym}(u)=13u^{2/3}+17u^c .\\label{Esym-c-simpler} \\end{equation}  The potential parameter $c$ varies between $0.4\\leq c \\leq 1.5$ which gives reliable values of the SE. The total pressure of the cold beta-stable nucleonic matter is given by \\begin{equation} P(n,x)=n^2\\left[\\frac{E'(n,\\frac{1}{2})}{A}+E_{sym}'(n)(1-2x)^2\\right]+P_{e^-}(n) , \\label{Pres-tot} \\end{equation} where $E_{sym}'(n)$ takes the form \\begin{equation} E_{sym}'(n)\\equiv \\frac{d E_{sym}(n)}{d n} =\\frac{1}{n_s} \\left[\\frac{26}{3}u^{-1/3}+17cu^{c-1}\\right] . \\label{Esym-Fuc} \\end{equation}  We are interested for the total pressure at the saturation density $n_s$. Considering that the electron pressure $P_{e^-}(n)$ is \\cite{Lattimer-01} \\begin{equation} P_{e^-}(n)\\cong nx(1-2x)E_{sym}(n) \\label{Pres-electron} \\end{equation} then the total pressure at $n_s$ is given by the expression \\cite{Lattimer-01} \\begin{equation} P_{s}(n_s,x_{s})=n_{s}(1-2x_{s})\\left[n_{s}E_{sym}'(n_s)(1-2x_{s})+E_{sym}(n_{s})x_{s}\\right] , \\label{Pres-sat} \\end{equation} where $E_{sym}'(n_s)=\\left[\\frac{d E_{sym}(n)}{d n}\\right]_{n=n_s}$ and the equilibrium proton fraction at $n_s$ is given \\begin{equation} x_s \\simeq (3 \\pi^2 n_s)^{-1}(4E_{sym}(n_s)/\\hbar c)^3 \\simeq 0.04 . \\label{pr-fr-ns} \\end{equation}  For small values of $x_s$ we find that \\begin{equation} P_s(n_s,x_s)\\simeq n_s^2 E_{sym}'(n_s) . \\label{Pres-sat-2} \\end{equation}  From the former expression it is obvious that the pressure is mostly sensitive to the density dependence of the SE at the saturation point $n_s$. Using our model from equation (\\ref{Esym-Fuc}) we get \\begin{equation} E_{sym}'(n_s)\\simeq  \\frac{1}{n_s}\\left(\\frac{26}{3}+17c\\right) . \\label{Esym-sat} \\end{equation}  From equations (\\ref{Pres-sat-2}) and (\\ref{Esym-sat}) it is concluded that the relation between the pressure $P_s$ and the potential parameter $c$ is \\begin{equation} P_{s}(n_s,x_{s})\\simeq n_s \\left(\\frac{26}{3}+17c\\right) . \\label{Pre-c1} \\end{equation}  In order to calculate the global properties of the neutron star, radius and mass we solved numerically the TOV equations (\\ref{TOV-1}) with the given equations of state constructed with the present model. For very low densities ($n<0.08$ fm$^{-3}$) we used the equation of state taken from Feynman, Metropolis and Teller \\cite{Feynman-49} and also from Baym, Bethe and Sutherland \\cite{Baym-71}.  In order to illustrate the density dependence trend of the SE proposed in our model we display in figure 1a $E_{sym}(n)$ as a function of the density $n$ for various values of the potential parameter $c$. It is obvious that the parameter $c$ affects decisively the trend of the SE, especially at high values of the density. So, it is very interesting to study how the values of the parameter $c$, and consequently the potential contribution, affect the gross properties of the NS. In the same figure, results of Ref. \\cite{Akmal98} (the case A18+$\\delta$u+UIX$^{*}$, see TABLE VI and VII of Ref. \\cite{Akmal98}) are included. It is found that the use of the phenomenological equation (\\ref{Esym-3}) with proper value ($c=0.9$) of function (\\ref{fu-1}) reproduces the results of the above microscopic calculations.  Figure 2a demonstrates the linear dependence between the radius $R_{max}$ and the parameter $c$. From our analysis it is concluded that there is a direct relation between $R_{max}$ and the parameterization of the potential part of the SE. The star symbol corresponds to the case  A18+$\\delta$u+UIX$^{*}$ with $c_1=0.9$. The corresponding relation was derived with the least-squares fit method and has the form \\begin{equation} R_{max}=9.10195+2.08304c .\\label{Rmax-c-fit} \\end{equation}  In addition, in figure 2b we indicate the behavior of the neutron star radius $R_{1.4}$, which corresponds to a neutron star mass $M=1.4 M_{\\odot}$, versus the parameter $c$. The star symbol corresponds to the case  A18+$\\delta$u+UIX$^{*}$ with $c_1=0.9$. It is obvious that there is also a linear relation between $R_{1.4}$ and $c$ which is \\begin{equation} R_{1.4}=10.52114+3.56746c . \\label{R14-c-fit} \\end{equation}  Figure 2c displays the correlation between the maximum mass of the neutron star $M_{max}$ and the parameter $c$. The star symbol corresponds to the case  A18+$\\delta$u+UIX$^{*}$ with $c_1=0.9$. We found an almost linear relation between $M_{max}$ and $c$ which has the form \\begin{equation} M_{max}=1.87442+0.12344c . \\label{Mmax-c-fit} \\end{equation}  By combining equations (\\ref{Esym-sat}) and (\\ref{R14-c-fit}) we found a linear relation between $R_{1.4}$ and $E'_{sym}(n_s)$ which has the form \\begin{equation} R_{1.4}=8.702+0.0336E'_{sym}(n_s) , \\label{R14-Esym-fit-ns} \\end{equation} and vice-versa \\begin{equation} E'_{sym}(n_s)=-259.185+29.783R_{1.4} . \\label{Esym-R14-fit-ns} \\end{equation}  In order to illustrate further the relation between the radius $R_{1.4 }$ and the trend of the SE we plot in figure 3a the radius $R_{1.4 }$ versus the derivative of the symmetry energy $E'_{sym}(3n_s/2)$ at the baryon density $n=3n_s/2$. A linear relation is found, which has the form \\begin{equation} R_{1.4}=10.1798+0.0242E'_{sym}(3n_s/2) . \\label{R14-Esym-fit} \\end{equation}  It is concluded that there is a direct relation between the radius $R_{1.4 }$ and the trend of the SE, close to the saturation point $n_s$. In addition in figure 3b we plot $E'_{sym}(3n_s/2)$ versus $R_{1.4 }$ with the linear correlation \\begin{equation} E'_{sym}(3n_s/2)=-420.40895+41.3066R_{1.4} . \\label{Esym-R14-fit} \\end{equation}  To ensure the relativistic causality in the present model we display figure 4a, where the ratio of the speed of sound to speed of light $c_s/c_l$ is plotted versus the baryon density for various values of the parameter $c$. Evidently, in all cases the speed of sound does not exceed that of light even at high values of the baryon density. In addition in figure 4b we plot the proton fraction $x_p$ calculated from expression (\\ref{prf-1}) as a function of the baryon density $n$. It is obvious that only in the cases where $c>0.5$ the proton fraction, after a specific density, exceed the critical value $x^{Urca}\\simeq0.11$ which ensures the beginning of the Urca process.  We also tried to find the correlation between the pressure $P$ (and consequently the radius R) and the SE for other values of the density $n$. In order to clarify the problem of the expected relation between the radius  and the pressure we present a more simplified model of a non-relativistic equation with a polytrope type of EOS. Thus the EOS has the form \\cite{Shapiro-83,Lattimer-01} \\begin{equation} P=K \\rho^{\\gamma} , \\qquad \\gamma=1+\\frac{1}{\\lambda} . \\label{EOS-Polyt} \\end{equation} and the  radius of the star is given by \\begin{equation} R=\\left[\\frac{(\\lambda+1)K}{4 \\pi G} \\right]^{1/2}\\rho_c^{(1-\\lambda)/2\\lambda} \\xi_1 , \\label{Rad-L-E} \\end{equation} where $\\rho_c$ is the central density and $\\xi_1$ is the solution of the equation $\\theta(\\xi_1)=0$, where the function $\\theta(\\xi)$ is the solution of the differential Lane-Emden equation \\begin{equation} \\frac{1}{\\xi^2}\\frac{d}{d \\xi}\\xi^2\\frac{d \\theta}{d \\xi}=-\\theta^{\\lambda} . \\label{L-E} \\end{equation}  Now it is obvious from equations (\\ref{EOS-Polyt}) and (\\ref{Rad-L-E}) that in case $\\lambda=1$ (or $\\gamma=2$) where $\\xi=\\pi$ we have \\begin{equation} \\frac{R}{P^{1/2}}=\\left[\\frac{2 G}{\\pi}\\right]^{1/2}\\frac{1}{\\rho} \\ . \\label{R-P-LE} \\end{equation}  Thus, from equation (\\ref{R-P-LE}) we concluded that in the case of a polytrope with $\\gamma=2$ there is a universal relation of the ratio $R/P^{1/2}$ calculated for a specific value of the density $\\rho$. However if general relativity effects are included in the above analysis the exponent $1/2$ of the pressure is found to be smaller \\cite{Lattimer-01}.  Following the above statement we plot in figure 5a the quantity $R_{1.4}P^{-a}$ as a function of the radius $R_{1.4}$. One can see that there is a correlation between the radii $R_{1.4}$ and the pressure evaluated at densities $1n_s$, $3n_s/2$ and $2n_s$. The values of the parameters $a$ and $C(n)$ have been defined by least-squares fit of the expression $R_{1.4}=C(n)P^{a}$.  In addition, in figure 5b we plot the quantity $R_{1.4}R^{-1/4}$ as a function of the $R_{1.4}$ to compare it with the previous work of Lattimer et. al \\cite{Lattimer-01}. It is worthwhile to notice that the quantity $R_{1.4}R^{-1/4}$ is a mild increasing function of the radius $R_{1.4}$. This effect is more evident for densities far from the saturation( $n=3n_s/2,2n_s$). So, from our study it is concluded that there is a slight dependence of the quantity $R_{1.4}R^{-1/4}$ from the potential parameter $c$ and consequently from the trend of the SE.  We proceed now in the more complicated case where the function $F(u)$ is given by expression (\\ref{FU-2}). In that case the derivative of the $E_{sym}(n)$ is given by \\begin{eqnarray} E_{sym}'(n)&=&\\left\\{ \\begin{array}{ll} \\frac{1}{n_s}\\left[\\frac{26}{3}u^{-1/3}+17c_1u^{c_1-1}\\right] &    \\mbox{$u \\leq 1$} \\\\ \\\\ \\frac{1}{n_s}\\left[ \\frac{26}{3}u^{-1/3}+17\\left(c_1+1-c_2+ e^{1-u}u^{c_2}(\\frac{c_2}{u}-1)\\right) \\right] & \\mbox{$u \\geq 1$} \\ . \\end{array} \\right. \\label{DESYM-2} \\end{eqnarray}  The potential parameters $c_1$ and $c_2$ varied between $0.5 \\leq c_1 \\leq 1.2 $ and $0 \\leq c_2 \\leq 2$ in order to get a reliable density dependent SE.  In figure 1b we display $E_{sym}(n)$ as a function of the density $n$ for various values of the potential parameters $c_1$ and $c_2$. In general the case is as follows, for fixed values of the parameter $c_2$, the SE is an increasing function of $c_1$. In addition, for fixed values of the parameter $c_1$ the increase of the parameter $c_2$ leads to a decrease of the SE. It is seen that, within the present model, the stiff or soft behavior of $E_{sym}(n)$ found in various microscopic calculations, is reproduced. As a comparison, similar to figure 1a, results of Ref. \\cite{Akmal98} (the case A18+$\\delta$u+UIX$^{*}$) are included. It is found that the use of the phenomenological equation (\\ref{Esym-3}) with proper values ($c_1=0.77$ and $c_2=1.09$) of function (\\ref{FU-2}) reproduces the results of the above microscopic calculations.  Figure 6a illustrates the behavior of the radius $R_{1.4}$ as a function of the second potential parameter $c_2$ for various values of the first potential parameter $c_1$. The calculated points for various values of $c_1$ can be reproduced by a second order polynomial.  \\begin{eqnarray} R_{1.4}&=&12.58956+0.05378c_2-0.46172c_2^2, \\qquad c_1=0.5 \\nonumber \\\\ R_{1.4}&=&13.04786-0.02752c_2-0.32340c_2^2, \\qquad c_1=0.7 \\nonumber \\\\ R_{1.4}&=&13.85705-0.08087c_2-0.21393c_2^2, \\qquad c_1=1.0 \\nonumber \\\\ R_{1.4}&=&14.61946-0.19441c_2-0.14536c_2^2, \\qquad c_1=1.2 \\ . \\label{R-c1-c2-a} \\end{eqnarray}  In all examined cases, the radius $R_{1.4}$ is a decreasing function of the potential parameter $c_2$. This is a direct consequence of the softening of the equation of state due to increase of the parameter $c_2$.  In addition, in figure 6b the behavior of the radius $R_{1.4}$ as a function of the first potential parameter $c_1$ is reproduced for various values of the second potential parameter $c_2$. The least-squares fit values are given for the following linear equations \\begin{eqnarray} R_{1.4}&=&11.09603+2.85172c_1, \\qquad c_2=0.0 \\nonumber \\\\ R_{1.4}&=&11.01810+2.85517c_1, \\qquad c_2=0.5 \\nonumber \\\\ R_{1.4}&=&10.42034+3.06724c_1, \\qquad c_2=1.2 \\ . \\label{R-c1-c2} \\end{eqnarray}  Unlike the previous case, $R_{1.4}$ is an increasing function of the potential parameter $c_1$. The increase of the parameter $c_1$ leads to the stiffness of the SE as indicated in figure 6b. It is worthwhile to note that the slopes of the best fit lines are almost the same and there is just a shift of the lines depending on the values of the parameter $c_2$.  Also, from figure 6a and 6b we conclude that the radius $R_{1.4}$ depends mainly on the parameter $c_1$ which determines the derivative of the $E_{sym}(n)$ and also the pressure $P_{sat}$ at the saturation density $n_s$. However, there is a small dependence on the parameter $c_2$ which is connected with the trend of $E_{sym}(n)$ at higher values  of the density $n_s$.  Figures 6c and 6d demonstrate the dependence of the radius $R_{max}$ and the mass $M_{max}$ respectively on the parameter $c_1$ (for fixed values of the parameter $c_2$). The star symbol corresponds the the case A18+$\\delta$u+UIX$^{*}$ ($c_1=0.77$ and $c_2=1.09$). In both cases a linear relations holds between $R_{max}$, $M_{max}$ and the parameter $c_1$. The lines correspond to the least-squares fit values.  It is of interest to compare the NS properties ($R_{1.4}$, $R_{max}$ and $M_{max}$) which originated from the use of equations  (\\ref{fu-1}) and (\\ref{FU-2}). As an example we use the parameterization of equations  (\\ref{fu-1}) and (\\ref{FU-2}) which reproduce very well the $E_{sym}$ trend of the case A18+$\\delta$u+UIX$^{*}$. As a result it is found that $R_{1.4}(c=0.9)=13.42$ and $R_{1.4}(c_1=0.77, c_2=1.09)=12.76$ (difference ~5 \\%) , $R_{max}(c=0.9)=11$ and $R_{max}(c_1=0.77, c_2=1.09)=10.59$ (difference ~3.7 \\%), $M_{max}(c=0.9)=1.978$ and $M_{max}(c_1=0.77, c_2=1.09)=1.974$ (difference ~0.2 \\%). It is obvious that in the case  A18+$\\delta$u+UIX$^{*}$ we receive almost identical results for $M_{max}$ while there is a small difference for $R_{1.4}$ and $R_{max}$. The differentiation of the values of $R_{1.4}$ is a consequence of the linear relation which hold between the radius $R_{1.4}$ and the derivative of the $E_{sym}$ (and consequently according to (\\ref{Pres-sat-2}) to the pressure), close to the saturation point (see figures $3$ and $7$). More specifically, we receive for the two cases, $P_s(c=0.9)=3.83467$ MeV fm$^{-3}$ and $P_s(c_1=0.77,c_2=1.09)=3.48107$ MeV fm$^{-3}$. In general, the small differentiation on the radii  is not surprising since the trend of equations  (\\ref{fu-1}) and (\\ref{FU-2}), due to suitable parametrization, are similar. However, it is worth to point out that equation  (\\ref{FU-2}) is a generalization of equation (\\ref{fu-1}), in the meaning that while equation (\\ref{fu-1}) describes well the case where $E_{sym}$ is a increasing function of the density, equation (\\ref{FU-2}) is sufficiently flexible  to describe in addition the case where $E_{sym}$ at low densities increases along with the density and then begins to decreases  although the density continues to increases.  To illustrate further  this point, we studied  the correlations between the derivative of the symmetry energy $E'_{sym}$ and the radius $R_{1.4}$ close to the saturation point $n=3n_s/2$. In figure 7a we plot the radius $R_{1.4}$ versus the derivative of the symmetry energy $E'_{sym}(3n_s/2)$ for fixed values of the potential parameter $c_2$. One can see that there is a linear relation between $R_{1.4}$ and $E'_{sym}(3n_s/2)$ just like in figure 3a. The effect of the parameter $c_2$ is to induce a parallel shift of the best fit lines. In figure 7b we indicate the inverse relation, that means $E'_{sym}(3n_s/2)$ versus $R_{1.4}$. The least-squares fit values for both cases and for various values of the parameter $c_2$ are given for the following equations \\begin{eqnarray} R_{1.4}&=&8.70394+0.02684 E_{sym}'(3n_s/2), \\qquad c_2=0 \\nonumber \\\\ E_{sym}'(3n_s/2)&=&-318.5026+36.82989 R_{1.4} \\ , \\label{R-Esym-0} \\end{eqnarray} \\begin{eqnarray} R_{1.4}&=&9.73291+0.02687 E_{sym}'(3n_s/2), \\qquad c_2=0.5 \\nonumber \\\\ E_{sym}'(3n_s/2)&=&-359.07613+36.98138 R_{1.4} \\ , \\label{R-Esym-05} \\end{eqnarray} \\begin{eqnarray} R_{1.4}&=&10.27305+0.02887 E_{sym}'(3n_s/2), \\qquad c_2=1.2 \\nonumber \\\\ E_{sym}'(3n_s/2)&=&-354.58213+34.54212 R_{1.4}  \\ . \\label{R-Esym-12} \\end{eqnarray}  In figure 8a, likewise with the figure 4a we display the ratio $c_s/c_l$ as a function of the density for various cases. It is obvious that the relativistic causality is satisfied once again. In figure 8b we display the proton fraction $x_p$ as a function of the density for various cases. A more systematic study of the $x_p$ leads to the conclusion that the potential parameter $c_1$ plays the most critical role for the occurence of the Urca process. Specifically a higher value of the $c_1$ leads to the beggining Urca process in smaller values of the baryon density.  Figure 9 illustrates the behavior of the quantity $R_{1.4}P^{-a}$ as a function of the radii $R_{1.4}$ for pressure determined at $n=n_s$, $3n_s/2$, $2n_s$, and also for $c_2=0$ (figure 9a), $c_2=0.5$ (figure 9b) and $c_2=1.2$ (figure 9c).  In figure 10 we plot the quantity $R_{1.4}P^{-1/4}$ as a function of $R_{1.4}$ for the pressure determined at $n=n_s$, $3n_s/2$, $2n_s$ and for $c_2=0$ (figure 10a), $c_2=0.5$ (figure 10b) and $c_2=1.2$ (figure 10c). It is obvious once again that the quantity $R_{1.4}P^{-1/4}$ is almost constant only when the pressure is calculated at the saturation point $n_s$. When the pressure is calculated at densities $n=3n_s/2$ and $n=2n_s$ the quantity $R_{1.4}P^{-1/4}$ is an increasing function of the radius $R_{1.4}$. Thus, as in the case of the simple parameterization of the SE, it is concluded that there is a dependence of the quantity $R_{1.4}P^{-1/4}$ from the first potential parameter $c_1$ as well as from the second potential parameter $c_2$ and consequently from the trend of the SE both for low and high values of the baryon density."
    },
    "0609/astro-ph0609434_arXiv.txt": {
        "abstract": "% We provide a library of some 7000 SEDs for starbursts and ultra luminous galaxies ({\\tt http://www.eso.org/\\~\\/rsiebenm/sb\\_models}). Its purpose is to quickly obtain estimates of the basic parameters, such as luminosity, size and dust or gas mass and to predict the flux at yet unobserved wavelengths.  The procedure is simple and consists of finding an element in the library that matches the observations.  The objects may be in the local universe or at high $z$.  We calculate the radiative transfer in spherical symmetry for a stellar cluster permeated by an interstellar medium with standard (Milky Way) dust properties.  The cluster contains two stellar populations: old bulge stars and OB stars.  Because the latter are young, a certain fraction of them will be embedded in compact clouds which constitute hot spots that determine the MIR fluxes.  We present SEDs for a broad range of luminosities, sizes and obscurations.  We argue that the assumption of spherical symmetry and the neglect of clumpiness of the medium are not severe shortcomings for computing the dust emission.  The validity of the approach is demonstrated by matching the SED of the best studied galaxies, including M82 and Arp220, by library elements; one example is shown for a galaxy at high redshift ($z \\sim 3$). Generally, one finds an element which fits the observed SED very well, and the parameters defining the element are in full accord with what is known about the galaxy from detailed studies. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609602_arXiv.txt": {
        "abstract": "We present numerical models of the circumstellar disks of Be stars, and we describe the resulting synthetic H$\\alpha$ emission lines and maps of the wavelength-integrated emission flux projected onto the sky.  We demonstrate that there are monotonic relationships between the emission line equivalent width and the ratio of the angular half-width at half maximum of the projected disk major axis to the radius of the star. These relationships depend mainly upon the temperatures of the disk and star, the inclination of the disk normal to the line of sight, and the adopted outer boundary for the disk radius. We show that the predicted H$\\alpha$ disk radii are consistent with those observed directly through long baseline interferometry of nearby Be stars (especially once allowance is made for disk truncation in binaries and for dilution of the observed H$\\alpha$ equivalent width by continuum disk flux in the $V$-band). ",
        "introduction": "%  Be stars are rapidly rotating B-type stars that lose mass into an equatorial, circumstellar disk \\citep{por03}.  The disk flux is observed in emission lines (especially the hydrogen H$\\alpha$ line), an infrared excess, and the polarization of starlight scattered in the disk.  It is now possible to investigate the actual extent and geometry of the circumstellar disks through optical long baseline interferometry of nearby Be stars \\citep{ste95,qui97,che05,tyc05,tyc06}. Most of these investigations have focused on the angular resolution of the disks in a narrow spectral band centered on the H$\\alpha$ emission line.  The long baseline interferometric observations are limited to the nearby sample of Be stars at present, so we need other indirect methods to help estimate the sizes of Be star disks in more distant targets.  Here we explore how the disk radius may be estimated from measurements of the disk H$\\alpha$ emission flux. We begin with an idealized case to illustrate some of the important parameters that define the relationship between equivalent width and radius.  Suppose that the disk is isothermal and geometrically thin and that the line emission is dominated by an inner optically thick region that appears as an ellipse projected onto the plane of the sky.  We also make the simplifying assumption that the specific intensity of the disk emission $I_\\lambda^d$ is isotropic. Then the wavelength integrated H$\\alpha$ flux will depend on the product of the surface flux, the wavelength range where H$\\alpha$ is optically thick, and the projected solid angle, \\begin{equation} F({\\rm H}\\alpha) = \\pi I_\\lambda^d <\\triangle\\lambda > \\cos i \\left({R_d \\over D}\\right)^2 \\end{equation} where $<\\triangle\\lambda>$ is the wavelength interval over which the line is optically thick (averaged over the visible disk), $i$ is the disk normal inclination angle, $R_d$ is the boundary radius of the optically thick disk, and $D$ is the distance to the star. We measure this flux through the emission equivalent width $W_\\lambda$, which is given in wavelength units relative to the local continuum flux, \\begin{equation} F({\\rm H}\\alpha) = |W_\\lambda| (1+\\epsilon) \\pi I_\\lambda^s \\left({R_s \\over D}\\right)^2 \\end{equation} where $\\epsilon$ is the ratio of disk continuum flux to stellar flux in the H$\\alpha$ region, $I_\\lambda^s$ is the stellar specific intensity near H$\\alpha$ (assumed isotropic), and $R_s$ is the radius of the star.  We can then equate these two expressions to find the predicted relationship between the ratio of disk to stellar radius and H$\\alpha$ emission equivalent width, \\begin{equation} {R_d \\over R_s} = \\sqrt{ {I_\\lambda^s \\over I_\\lambda^d} {{W_\\lambda (1+\\epsilon)}\\over {<\\triangle\\lambda> \\cos i}}}. \\end{equation}  This special case is characterized by a disk radius that varies as the square root of the emission equivalent width, and this expectation is supported by the interferometric observations of \\citet{tyc05} who find that the emitted H$\\alpha$ flux scales with the square of the disk radius.   However, the coefficient of the relationship also depends on the ratio of the stellar to disk specific intensities that is dependent upon the stellar and disk temperatures.  Furthermore, there is an inclination dependence in both the $\\cos i$ projection factor and the $<\\triangle\\lambda>$ term (through the range in velocities and Doppler shifts viewed through the disk).  The $<\\triangle\\lambda>$ term also depends upon the assumed disk density since the optical depth is based on the product of neutral hydrogen density and the emission line profile.  Finally, there must also be some emission contribution from the optically thin outer disk that is ignored in expression above.  Thus, in order to predict accurately the relationship between equivalent width and disk radius, we need a more complete model that accounts for the specific dependences on disk temperature, inclination, and density.  We have adopted the Be disk model approach of \\citet{hum00} that is based upon models developed by \\citet{hor86} and \\citet{hor95} for accretion disks in cataclysmic variables. The disk is assumed to be axisymmetric and centered over the equator of the underlying star, and the gas density varies as \\begin{equation} \\rho (R,Z) = \\rho_0 R^{-n} \\exp \\left[-{1\\over2}\\left({Z\\over{H(R)}}\\right)^2\\right] \\end{equation} where $R$ and $Z$ are the radial and vertical cylindrical coordinates (in units of stellar radii), $\\rho_0$ is the base density at the stellar equator, $n$ is a radial density exponent, and $H(R)$ is the disk vertical scale height. The neutral hydrogen population within the disk is found by equating the photoionization and recombination rates \\citep{gie06}. The disk gas is assumed to be isothermal and related to the stellar effective temperature $T_{\\rm eff}$ by $T_d = 0.6 T_{\\rm eff}$ \\citep{car06}.  The numerical model represents the disk by a large grid of azimuthal and radial surface elements, and the equation of transfer is solved along a ray through the center of each element according to \\begin{equation} I_\\lambda = S_\\lambda^L (1 - e^{-\\tau_\\lambda}) + I_\\lambda^s e^{-\\tau_\\lambda} \\end{equation} where $I_\\lambda$ is the derived specific intensity, $S_\\lambda^L$ is the source function for the disk gas (taken as the Planck function for the disk temperature $T_d$), $I_\\lambda^s$ is the specific intensity for a uniform disk star (taken as the product of the Planck function for $T_{\\rm eff}$ and an H$\\alpha$ photospheric absorption line derived from the grid of \\citealt{mar05}), and $\\tau_\\lambda$ is the integrated optical depth along the ray (for assumed Keplerian rotation; see eq.\\ [4] in \\citealt{hor95} and eq.\\ [7] in \\citealt{hum00}). The first term applies to all the disk area elements that are unocculted by the star (see eq.\\ [12] in \\citealt{hum00}) while the second term applies to all elements that correspond to the projected photospheric disk of the star. The absorption line adopted in $I_\\lambda^s$ is Doppler shifted according to solid body rotation for the photospheric position in a star that is rotating at $90\\%$ of the critical value.  The model code calculates a synthetic line profile over the range from $-2000$ to $+2000$ km~s$^{-1}$ at 10 km~s$^{-1}$ intervals by summing the product of projected area and specific intensity over the disk grid.   We then integrate the line profile relative to the continuum to form a predicted equivalent width (taken as positive for net absorption or negative for net emission). The model also forms a synthetic image of the star plus disk in the plane of the sky by summing the intensity over a 2.8~nm band centered on H$\\alpha$.  We collapsed this image along the projected major axis to get the summed spatial intensity, and then we determined the radius where the summed intensity drops to half its maximum value (excluding the spatial range corresponding to photospheric intensity).   We adopt this half-maximum summed intensity radius as the effective disk radius in the following discussion.  We need to specify the stellar mass and radius in the code in order to determine the disk Keplerian velocities, and we adopted these as a function of stellar effective temperature (or spectral subtype) using the eclipsing binary results for B-stars from \\citet{har88}.  We determined the relationship between model equivalent width and disk radius by running a sequence of models defined by the spectral subtype of the star, the disk inclination, the disk radial density exponent $n$, and the selected value for the outer boundary of the disk.   The final variable for any sequence is the base density $\\rho_0$, and we computed models over a range in $\\rho_0$ that corresponds to H$\\alpha$ profiles with no visible emission up to strong emission cases with $|W_\\lambda| > 50$~\\AA .  We show the predicted relationship between the disk radius and H$\\alpha$ equivalent width in Figure~1 for the case of a B2~Ve star ($T_{\\rm eff} = 23100$~K) assuming $n=3.0$ and an outer boundary at $100 R_s$. The lower solid line shows the pole-on case ($i=0^\\circ$).  At higher inclination angles ($i=50^\\circ$ and $i=80^\\circ$ for the middle and top lines, respectively) the projected disk area is smaller ($\\propto \\cos i$), and consequently the emission equivalent width is smaller for a given disk radius.  We do not show the $i=90^\\circ$ case because the shear broadening approximation for the emission profile width breaks down for an edge-on orientation \\citep{hor86}. The symbols along each sequence indicate positions of fiducial base density values.  Note that all three curves show the expected square root dependence predicted for the simple case (eq.\\ [3] above).  \\placefigure{fig1}     %  The three dashed lines in Figure~1 show the corresponding inclination cases for a B8~Ve star ($T_{\\rm eff} = 11600$~K) again made assuming $n=3.0$ and an outer boundary at $100 R_s$. These curves are all shifted to the left relative to the B2~Ve sequences because the photospheric absorption component is larger in cooler B-stars.  The model profiles of the B8~Ve sequence have less H$\\alpha$ emission at a given disk radius than their B2~Ve counterparts because at this lower temperature the disk-to-star intensity ratio, the vertical scale height, and the thermal emission line broadening are all smaller.  The dotted line in Figure~1 shows how the H$\\alpha$ equivalent width is reduced when the outer boundary for the disk grid is moved inwards from $100 R_s$ to $25 R_s$ for the case of a B2~Ve star with a disk inclination of $80^\\circ$ and a disk density exponent of $n=3.0$. The outer disk region contains low brightness, optically thin gas, and although the emission contribution from any particular outer radius area element is small, the projected area of such elements increases with radius so that the selection of the outer boundary condition is important. We found that increasing the outer boundary from $100 R_s$ to $200 R_s$ for this parameter set increased the equivalent widths by only $10\\%$, so we adopted the $100 R_s$ boundary as our nominal choice. The selection of the boundary radius will mainly be important for Be stars in binary systems where the outer disk will be truncated by the gravitational influence of the companion.  Finally, the dot-dashed line in Figure~1 shows another sequence for the same B2~Ve, $i=80^\\circ$, and $100 R_s$ boundary case, but calculated this time with a larger disk density exponent of $n=3.5$ (i.e., a steeper drop off in density). The overall shape of this curve is almost the same as the one for $n=3.0$ (although the fiducial base density points have very different locations), and this suggests that the particular choice of density law is not important for the derived equivalent width -- radius relationships. Investigations of the infrared flux excess in Be stars suggest that the exponent falls in the range $n = 2 - 4$ \\citep{por03}. We caution that the results become more sensitive to the outer disk boundary for $n<3$ where the emission from the optically thin, outer regions assumes more significance.  We have constructed such model sequences for subtypes B0, B2, B4, B6, and B8~Ve, for inclinations $i=0^\\circ$, $50^\\circ$, and $80^\\circ$, and for outer boundary radii of $17 R_s$, $25 R_s$, $50 R_s$, and $100 R_s$. These numerical results, an interpolation program, and our code are available to interested readers at our Web site\\footnote{http://www.chara.gsu.edu/$^\\sim$gies/Idlpro/BeDisk.tar}. ",
        "conclusions": ""
    },
    "0609/astro-ph0609328_arXiv.txt": {
        "abstract": "Observations of the high-mass X-ray binary \\src\\ with the Swift Burst Alert Telescope (BAT) do not show modulation at the previously reported period of 9.6 days found from observations made with the Rossi X-ray Timing Explorer (RXTE) All-Sky Monitor (ASM).  Instead, the strongest peak in the power spectrum of the BAT light curve occurs at a period of 19.25 $\\pm$ 0.08 days, twice the period found with the RXTE ASM.  The maximum of the folded BAT light curve is also delayed compared to the maximum of the folded ASM light curve.  The most recent ASM data folded on twice the 9.6 day period show similar morphology to the folded BAT light curve. This suggests that the apparent period doubling is a recent secular change rather than an energy-dependent effect.  The 9.6 day period is thus not a permanent strong feature of the light curve. We suggest that the orbital period of \\src\\ may be twice the previously proposed value. ",
        "introduction": "The Uhuru X-ray source \\src\\ (Giacconi et al. 1972) was identified with the optical counterpart BD +53\\degrees2790 by Steiner et al. (1984). Optical photometry and the presence of a double peaked \\halpha\\ line in the spectrum led Steiner et al. (1984) to conclude that BD +53\\degrees2790 is a Be star. However, it has been argued that the optical properties, in particular the behavior of the H$\\alpha$ emission line,  show that this star is actually a very peculiar active O type star (Negueruela \\& Reig 2001, Rib\\'o et al. 2006, Blay et al. 2006).  Saraswat \\& Apparao (1992) reported the detection of X-ray pulsations from \\src\\ with a period of 392s from EXOSAT observations. However, Corbet \\& Peele (2001) argued that the period search technique used by Saraswat \\& Apparao (1992) was inappropriate in searching for coherent variability. Corbet \\& Peele (2001) reanalyzed the same EXOSAT data set and did not find any evidence for pulsations. In addition, Corbet \\& Peele (2001) and Negueruela \\& Reig (2001) analyzed Rossi X-ray Timing Explorer (RXTE) Proportional Counter Array (PCA) observations of \\src\\ made in 1997 and again failed to find any pulsations. The X-ray spectrum can be well fitted with a ``standard'' high-mass X-ray binary spectrum (White, Swank, \\& Holt 1983) of a power law modified by a high-energy cut off (e.g. Corbet \\& Peele 2001; Masetti et al. 2004). The lack of an iron line in the spectrum (Negueruela \\& Reig 2001; Masetti et al. 2004) makes it similar to high-mass X-ray binaries containing Be stars rather than supergiant donors which often have significant emission near 6.4 keV (e.g. Nagase 1989). Despite the lack of pulsations, evidence for the presence of a neutron star in \\src\\ has been claimed to come from X-ray spectroscopy. Masetti et al. (2004) analyzed a  BeppoSAX observation of \\src\\ made in 1998 and found tentative evidence for the presence of a cyclotron resonance absorption feature. Torrej\\'on et al. (2004) analyzed the same data set together with RXTE observations made in 1997 and 2001 and also found possible evidence for the presence of a cyclotron absorption feature around 30 keV. Blay et al. (2005) also report the presence of an absorption feature near 32 keV in spectra obtained with INTEGRAL.  From an analysis of approximately 5.5 years of observations of \\src\\ with the  RXTE All-Sky Monitor (ASM) Corbet \\& Peele (2001) found modulation at a period of 9.568 $\\pm$ 0.004 days. If this modulation is due to orbital variability then this would be one of the shortest orbital periods known for a Be type system (Raguzova \\& Popov 2005). Corbet \\& Peele (2001) looked for differences between ``odd'' and ``even'' outbursts that would indicate that the underlying period was actually twice this value but did not find any evidence for this. Rib\\'o et al. (2006) analyzed additional RXTE ASM data and derived a refined period of 9.5591 $\\pm$ 0.0007 days. Overall, the X-ray light curve of \\src\\ differs from the majority of Be/neutron star systems in that the source does not display either periodic Type I outbursts that occur near periastron passage separated by quiescent intervals or less frequent but larger Type II events which are not modulated on the orbital period (e.g. Charles \\& Coe 2003). Instead, although \\src\\ is very variable, particularly on shorter timescales, it is essentially persistent on long timescales.  In summary, although initial results suggested that \\src\\ was a member of the common Be/neutron star binary class of objects, further observations have called this into question. Optical spectroscopy has cast doubt on the Be star classification, and the lack of pulsations means that the nature of the presumed accreting compact object in the system is not yet securely determined. Because of the lack of pulsations from \\src\\ Corbet \\& Peele (2001) considered alternative models for \\src\\ in which the accreting object might be a white dwarf, black hole, or the magnetosphere of a neutron star rather than the surface of a neutron star.  In this paper we present an analysis of observations of \\src\\ made with the Burst Alert Telescope (BAT) on board the Swift satellite. These data do not show modulation on the 9.6 day period but, unexpectedly, show a maximum in the power spectrum at twice this period. We update the ASM results and compare the folded BAT and ASM light curves and show that recent ASM data show similar behavior to that found with the BAT. We find that the phase of the BAT maxima in the light curve does not exactly coincide with either the ``odd'' or ``even'' outbursts on the 9.6 day period. We also present light curves of \\src\\ obtained with the RXTE Proportional Counter Array (PCA) and High Energy Timing Experiment (HEXTE) in 2001 which cover orbital phases around the maximum in the folded BAT light curve. We make a comparison of \\src\\ with other systems with somewhat similar properties. While there is some similarity with GRO J2058+42, the overall properties of \\src\\ so far appear to be unique. ",
        "conclusions": "The ASM and BAT light curves show that strong modulation at a period of 9.6 days is not a persistent feature of the light curve of \\src. The BAT light curve of \\src\\ shows a modulation at twice the period previously found from RXTE ASM observations. A detailed examination of the ASM light curve also demonstrates a change in the nature of the modulation. The orbital period of \\src\\ might thus be twice the originally proposed value of 9.6 days. Although the mechanism causing the variable modulation in \\src\\ is not yet clear, a model where there is an equatorial concentration of the wind from the primary star and the equatorial and orbital planes are offset may be able to explain at least some of the observed properties. It would be very valuable if a radial velocity orbit could be obtained from optical observations as this should conclusively give the value of the orbital period, show the relative orientation of the binary components at the times of X-ray maxima, and give the orbital eccentricity. If the period doubling is caused by a change in the circumstellar environment in the system, then this may also show effects in the optical spectrum such as the strength or shape of the H$\\alpha$ emission line. Continued high-energy monitoring of \\src\\ with the BAT and/or ASM will show how long the change to the period doubled state continues which may assist in understanding its cause.  The nature of \\src\\ continues to be puzzling and the determination that the 9.6 day modulation is not a persistent feature of the light curve adds to the peculiarities of the source. The exceptional properties of \\src\\ may be related to the unusual nature of the mass donor in the system. \\src\\ may thus provide a new parameter regime where models of wind accretion may be tested."
    },
    "0609/astro-ph0609091_arXiv.txt": {
        "abstract": "We use kinematic measurements of a large sample of galaxies from the Team Keck Redshift Survey in the GOODS-N field to measure evolution in the optical and near-IR Tully-Fisher relations to $z = 1.2$.  We construct Tully-Fisher relations with integrated line-of-sight velocity widths of $\\sim 1000$ galaxies in $B$ and $\\sim 670$ in $J$; these relations have large scatter, and we derive a maximum-likelihood least squares method for fitting in the presence of scatter.  The $B$-band Tully-Fisher relations, from $z=0.4$ to $z=1.2$, show evolution of $\\sim 1.0-1.5$ mag internal to our sample without requiring calibration to a local TF relation.  There is evolution in both Tully-Fisher intercept and slope, suggesting differential luminosity evolution.  In $J$-band, there is evolution in slope but little evolution in overall luminosity.  The slope measurements imply that bright, massive blue galaxies fade {\\it more strongly} than fainter blue galaxies from $z \\sim 1.2$ to now.  This conclusion runs counter to some previous measurements and to our naive expectations, but we present a simple set of star formation histories to show that it arises naturally if massive galaxies have shorter timescales of star formation, forming most of their stars before $z\\sim 1$, while less massive galaxies form stars at more slowly declining rates.  This model predicts that the higher global star formation rate at $z\\sim 1$ is mostly due to higher SFR in massive galaxies.  The amount of fading in $B$ constrains star formation timescale more strongly than redshift of formation.  Tully-Fisher and color-magnitude relations can provide global constraints on the luminosity evolution and star formation history of blue galaxies. ",
        "introduction": "Surveys of large samples of galaxies at significant lookback times are a powerful instrument for measuring the evolutionary history of galaxies.  First steps in these probes include measuring number distributions of galaxies, such as the galaxy luminosity function and color-magnitude distribution.  A second step is the measurement of scaling relations of galaxy properties.  Locally, relations between galaxy luminosity and characteristic internal kinematic velocity -- the Tully-Fisher (TF) relation for disks, the $D_n-\\sigma$ relation for spheroidals and its descendant, the Fundamental Plane -- are among the tightest correlations and strongest tools for characterizing galaxies.  These relations can be used to probe evolution in galaxy properties, and simultaneously, a successful scenario for galaxy formation and evolution will have to reproduce these relations, their scatter, and their evolution with cosmic time.  However, measuring galaxy internal kinematics at intermediate to high redshift is difficult since the sources are faint and high spectral resolution is needed.  In Paper I (Weiner \\etal\\ 2006) we presented measurements of galaxy kinematics from emission lines, using Keck/DEIMOS spectra from the Team Keck Redshift Survey (TKRS) in the GOODS-N (Great Observatories Origins Deep Survey) field (Wirth \\etal\\ 2004; Giavalisco \\etal\\ 2004). The sample with kinematics from Paper I contains $\\sim 1000$ emission line velocity dispersions from integrated emission, with median $<z> = 0.637$, and reaches beyond $z>1$. We use integrated linewidths for fitting Tully-Fisher relations in restframe $B$ and $J$ bands.  For a subsample of 380 galaxies, we fit the 2-d spectra to model spatially resolved rotation and dispersion.  Paper I compared the properties of the spatially resolved rotation and dispersion to validate the integrated linewidths.  A number of previous studies have measured luminosity-velocity relations from $0.1<z<1.0$.  Forbes \\etal\\ (1996), Rix \\etal\\ (1997) and Mall{\\' e}n-Ornelas \\etal\\ (1999) used velocity dispersions from integrated emission for $\\sim 20$ galaxies each at median redshifts 0.48, 0.25 and 0.55. The pioneering studies of Vogt \\etal\\ (1996, 1997) modeled rotation curves for 17 galaxies of disky morphology with median $<z> = 0.47$ by combining Keck/LRIS slitlet spectra with structural information from HST photometry.  Subsequent studies of rotation curves with similar modeling procedures include Simard \\& Pritchet (1998); Vogt (2000); Ziegler \\etal\\ (2002) and Bohm \\etal\\ (2004); Milvang-Jensen \\etal\\ (2003), Bamford \\etal\\ (2005, 2006), and Nakamura \\etal\\ (2006); Metevier \\etal\\ (2006); and Conselice \\etal\\ (2005) who used $K$-band imaging of galaxies in the HST Groth Strip. These generally contained 20--100 galaxies with median redshifts $\\sim 0.4 - 0.5$, and include field and cluster samples.  These studies have produced a range of results for evolution in the zeropoint of the Tully-Fisher relation; we discuss comparisons of TF results in Section \\ref{sec-othertf}.  TF measurements from resolved rotation curves have mostly indicated a relatively mild amount of magnitude evolution.  Previous samples are drawn from a variety of populations and selection criteria.  In general, samples for measuring rotation curves have been selected to be inclined disky objects with measurable and reasonably orderly rotation.  By contrast, the sample reported here is essentially selected only on magnitude and emission line strength.  The lack of selection may be important, given evidence that peculiar and disturbed galaxies are more common at high redshift (e.g. Abraham \\etal\\ 1996). The TKRS sample is also large and spans redshifts from 0.1 to 1.6, which allows a measurement of Tully-Fisher relation evolution internal to the sample rather than by comparison to a fiducial local Tully-Fisher relation.  Our primary tool for measuring kinematics is one-dimensional line-of-sight velocity dispersion (linewidth), and secondarily the spatially resolved rotation profiles. Paper I discussed the properties of these velocity measures, and showed that integrated linewidth as a measure of the characteristic velocity of a galaxyis fairly robust against observational effects and the details of galaxy kinematics, and that we can measure velocities even for galaxies that are kinematically anomalous or not in orderly rotation, albeit with scatter due to the geometries of velocity fields and inclinations.  We adopt an LCDM cosmology with $h=0.7$, $\\Omega_M = 0.3$, and $\\Omega_\\Lambda = 0.7$.  Magnitudes quoted in this paper are in the AB system unless explicitly indicated as Vega. The sample and observations are discussed in Paper I; Section \\ref{sec-sample} briefly introduces the data. Section \\ref{sec-lwtf} presents Tully-Fisher relations for integrated linewidths and for rotation curves, and measures their evolution. Section \\ref{sec-othertf} compares to other TF measurements, Section \\ref{sec-evolcause} discusses the causes of evolution in TF relation, and Section \\ref{sec-models} uses luminosity evolution models to interpret the TF evolution. In the Appendix we derive a maximum likelihood method for fitting relations, such as our Tully-Fisher relation, that have a substantial scatter. ",
        "conclusions": "\\label{sec-conclude}  We measured line-of-sight integrated kinematic linewidths \\sigoned\\ and spatially resolved line-of-sight rotation \\vrot\\ and disperson \\sigtwod\\ from the galaxy spectra of the Team Keck survey of the GOODS-N field. In Paper I we showed that linewidths are a measure of internal kinematics that is relatively robust against observational effects, although there is significant scatter between any individual galaxy linewidth and true circular velocity or dynamical mass. We use 968 linewidths of galaxies with $M_B<-18$ and 677 linewidths of galaxies with $M_J<-19$ to measure evolution of the restframe $B$ and $J$-band Tully-Fisher relations from $z\\sim 0.4$ to 1.2. This is the largest sample of high-redshift galaxies with kinematics to date, and samples the blue galaxy population without morphological or other pre-selection. It allows both an internal comparison, without having to fix to a local TF relation, and the measurement of intercept and slope evolution.  The intrinsic scatter in the Tully-Fisher relation for linewidths is large, 0.18 dex in \\logsigoned, which requires special care in fitting methods; we derive a maximum likelihood fitting method in the Appendix.  The intrinsic scatter is partly due to the lack of inclination and extinction corrections; kinematic pecularity is likely to both contribute to the scatter and to reduce its non-gaussianity by reducing the number of orderly low-inclination objects with low observed velocity.  Correlations of TF residual with other parameters are weak and do not suggest obvious second-parameter sources of scatter or evolution.  In the $B$-band linewidth TF relation, there is very significant intercept evolution and $3\\sigma$ evidence for slope evolution. The sense of the slope evolution is such that the high-redshift inverse TF relation, velocity as a function of magnitude, $V(M)$, is shallower than at low redshift. The $J$-band linewidth TF relation shows relatively little evolution in intercept, but a similar evolution in slope to that found in $B$-band. The $B$-band TF relation for rotation velocities \\vrot\\ also supports intercept evolution, but the sample is too small and noisy to measure slope evolution.  The sense of slope evolution is most easily interpreted as differential luminosity evolution: the most luminous galaxies at $z>1$ fade by more than the less luminous galaxies.  There is more fading in $B$ than in $J$ band.  In the mean, a galaxy with $M_B=-21.5$ at $z=1.2$ fades by $\\sim 1.5$ $B$ magnitudes to $z=0.4$, but only by $\\sim 0.5$ mag in $J$. The $B$-band evolution is larger than some previous TF measurements; some differences are that our sample has many more galaxies at high redshift, that it is not morphologically selected to favor normal disks, and that it does not reject galaxies with non-rotating kinematics.  Because the evolution is differential, the amount of fading seen is also dependent on the magnitudes of the different TF samples.  The fading of $1-1.5$ mag in $B$ at fixed velocity that we see is similar to the fading of $1.3 \\pm 0.3$ mag in $L_B^*$ in the blue galaxy LF measured by Willmer \\etal\\ (2006), and the fading of luminosity at fixed radius seen by Melbourne \\etal\\ (2006).  A simple, but not unique, model for these results is that most of the blue galaxy population evolves unaltered by merging or density evolution since $z \\sim 1$, with stellar population fading the primary driver of luminosity evolution.  We present a simplistic model to explain differential luminosity evolution: we construct galaxy histories with star formation rate declining exponentially with time (\\taumods, ${\\rm SFR} \\propto e^{-t/\\tau}$), and make the SFR timescale $\\tau$ a function of mass along the Tully-Fisher relation.  In this scenario, massive galaxies have short $\\tau$ and form the bulk of their stars early, before $z=1$, while low mass galaxies have long $\\tau$ and build up stellar mass slowly.  This arrangment of \\taumods\\ has long been indicated by the color-magnitude relation for blue galaxies (e.g. Tully \\etal 1982). The model correctly predicts that the massive galaxies fade substantially from $z \\sim 1$ to now, while low mass galaxies fade little since they are still building stellar mass.  The decrease of the global SFR since $z\\sim 1$ is thus dominated by the decrease in SFR in massive galaxies, while the activity of low-mass galaxies changes relatively little in the mean.  Thus the Tully-Fisher slope evolution, as differential luminosity evolution, is a natural outcome of an appealing model for star formation histories of blue galaxies.  Measurements of Tully-Fisher evolution help us to constrain models of galaxy evolution by relating galaxies at one epoch to another, tracking the evolutionary descent of galaxies.  In this paper we have taken the simplest approach of assuming pure luminosity evolution, i.e. that galaxies' characteristic velocities evolve much less than their luminosities, and that the blue galaxy population is essentially constant in number. Future work on this subject can be more sophisticated. Theoretical ideas about the evolution of velocity with halo mass can improve this approach.  The color-magnitude relation, infrared magnitudes, stellar mass estimates, and the luminosity and velocity functions provide a wealth of data which models must confront.  The scatter about these relations should be related to episodic variations about the mean star formation histories.  Together these may be used to understand the mass assembly history of galaxies."
    },
    "0609/astro-ph0609058_arXiv.txt": {
        "abstract": "We present a study of the optical (BRI) and near-infrared (JHK) luminosity fuctions (LFs) of the GEMS sample of 60 nearby groups of galaxies between $0.01\\!<\\! z \\!<\\! 0.04$, with our optical CCD photometry and near-IR photometry from the 2MASS survey.  The LFs in all filters show a depletion of galaxies of intermediate luminosity, two magnitudes fainter than $L_\\ast$, within $0.3\\,R_{500}$ from the centres of X-ray faint groups. This feature is not as pronounced in X-ray bright gropus, and vanishes when LFs are found out to $R_{500}$, even in the X-ray dim groups.  We argue that this feature arises due to the enhanced merging of intermediate-mass galaxies in the dynamically sluggish environment of low velocity-dispersion groups, indicating that merging is important in galaxy evolution even at $z\\!\\sim\\! 0$. ",
        "introduction": "\\label{sec:intro} Observational studies of the environmental dependence of galaxy evolution mostly concentrate on galaxies in rich clusters, even though $<5$\\% of galaxies are found in such environments. Most galaxies are found in groups, and arguably the group environment plays an important role in their evolution. In groups, galaxies could be transformed by their interaction with the intragroup medium, or with other galaxies by means of a variety of processes (e.g. stripping, tidal interaction) or through  merging with other galaxies. Here we seek to study the optical and near-IR properties of galaxies in groups to find whether they support the relative importance of any of these processes.  \\begin{figure} \\centering \\includegraphics[height=5.3cm,angle=-90]{tm4.ps} \\caption{Cumulative $B$-band Luminosity Function of 25 GEMS galaxy groups, divided into two categories: X-ray bright groups ($L_X>10^{41.7}$ erg~s$^{-1}$, plotted as triangles and  X-ray dim groups ($L_X<10^{41.7}$ erg~s$^{-1}$), plotted as circles. Only the LF of X-ray bright groups can be fit with a single Schechter function. } \\label{fig:cumulf}       % \\end{figure} ",
        "conclusions": "We argue that the missing intermediate-luminosity galaxies in the optical and near-infrared luminosity functions of X-ray dim groups indicate that, in the dynamically sluggish environment of such groups (which have low velocity disperison), dynamical friction would facilitate more rapid merging, thus depleting intermediate-luminosity galaxies to form a few giant central galaxies.  We also show that this effect is seen only in the interior regions of the groups ($R<0.3\\,R_{500}$), and vanishes as one approaches $R_{500}$. rather than a bright-end enhancement caused by excess star formation. In \\cite{miles06a}, we show that this feature cannot arise due to enhanced star formation in the brightest galaxies, or due to a varying morphological mix of galaxies in various groups.  It is often suggested (e.g. \\cite{con06}) that mergers \\index{mergers} are not an important ingredient of galaxy evolution in the recent Universe ($z\\!<\\!1$). Here we have shown that in nearby poor groups, merging is still an important process.  This picture of galaxy evolution leads to a definite prediction.  If X-ray dim groups are indeed systems undergoing rapid dynamical evolution, the stellar populations in their galaxies would be significantly younger than those in X-ray bright groups in case of dissipative merging.  This can be observationally verified.   \\subparagraph{Acknowledgements} We thank Trevor Ponman, Duncan Forbes, Paul Goudfrooij and other members of the GEMS team for their contributions to this work, and the organisers for a very stimulating meeting."
    },
    "0609/astro-ph0609744_arXiv.txt": {
        "abstract": "{} {To describe the evaporation status of the extrasolar planets, we propose to consider an energy diagram in which the potential energy of the planets is plotted versus the energy received by the upper atmosphere.} {Here we present a basic method to estimate these quantities. For the potential energy, we include the modification of the gravity field by the tidal forces from the parent stars.} {This description allows a quick estimate of both the escape rate of the atmospheric gas and the lifetime of a planet against the evaporation process. In the energy diagram, we find an evaporation-forbidden region in which a gaseous planet would evaporate in less than 5 billion years. With their observed characteristics, all extrasolar planets are found outside this evaporation-forbidden region. The escape rates are estimated to be in the range $10^{5}$\\,g\\,s$^{-1}$ to $10^{12}$\\,g\\,s$^{-1}$, with few cases above $10^{11}$\\,g\\,s$^{-1}$. The estimated escape rate for HD\\,209458\\,b is found to be consistent with the lower limit of $10^{10}$\\,g\\,s$^{-1}$ obtained from interpretation of the H\\,{\\sc i} Lyman-$\\alpha$ observations.  Finally, this diagram suggests possibilities for the nature of the recently discovered Neptune-mass planets. We find that GJ\\,436\\,b, 55\\,Cnc\\,e and HD\\,69830\\,b cannot be low mass gaseous planets. With density necessarily above 0.5\\,g\\,cm$^{-3}$ to survive evaporation, these planets must contain a large fraction of solid/liquid material. Concerning GJ\\,876\\,d, we find that it must have a density larger than $\\sim$3\\,g\\,cm$^{-3}$ to survive the strong EUV energy flux from its nearby parent star. GJ\\,876\\,d must contain a large fraction of massive elements. } {} ",
        "introduction": "\\label{Introduction}  Among the nearby two hundreds extrasolar planets identified today, there is a particular class of planets very close to their parent stars with semi major axis $\\la$0.1\\,AU. Among them, there are massive hot-Jupiters (also named ``Pegasides'', Guillot et al.\\ 1996) and small mass planets (see for instance, Santos et al.\\ 2004; Rivera et al.\\ 2005). These planets constitute about 20\\% and 5\\% respectively, of the planets identified today. Since the discovery of the first of them, 51\\,Peg\\,b (Mayor \\& Queloz 1995), the issue of their evaporation status has been raised (Burrows \\& Lunine 1995; Guillot et al.\\ 1996). Since these planets exist and orbit stars which are not particularly young, it is clear that the evaporation rate of the observed massive hot-Jupiters should be modest enough to not dramatically impact their evolution. At least they can survive billion years at the place where they are observed. It is extremely unlikely that we are witnessing a transient phase in the evolution of extrasolar planets (Baraffe et al.\\ 2004). In brief, the discoveries of a large number of massive hot-Jupiters\\footnote{usually defined as massive planets orbiting in less than 10~days} and very-hot-Jupiters\\footnote{usually defined as planets orbiting in less than 3~days ({\\it e.g.},  Gaudi et al.\\ 2005)} led to the conclusion that the evaporation {of massive planets} has to be modest.  In that frame the discovery that the transiting extrasolar planet HD\\,209458\\,b is indeed losing mass was mostly unexpected (Vidal-Madjar et al.\\ 2003). From the observation of 15$\\pm$4\\% absorption in Lyman-$\\alpha$ during the transit, a lower limit of the H\\,{\\sc i} escape rate was estimated to be of the order of $\\sim10^{10}$\\,g\\,s$^{-1}$ (Vidal-Madjar et al.\\ 2003; Vidal-Madjar \\& Lecavelier des Etangs 2004). This mass loss rate corresponds to a loss of $\\sim$0.1\\% of the total mass of the planet over 10 billion years (10 Gyr). Models of the time variation of this escape rate lead to slightly larger values for the total fraction of mass loss over the whole stellar life (see e.g., Lecavelier des Etangs et al.\\ 2004; Baraffe et al.\\ 2004).  Different detailed models have been developed to understand the observed evaporation process and to estimate the escape rate (e.g., Yelle 2004; Tian et al. 2005,  Garc\\'\\i a Mu\\~noz 2006), to be compared with the observational constrains (Vidal-Madjar et al.\\ 2003, 2004). Finally, it has been suggested that evaporation can lead to significant modification of the planets nature in the case of planets with a mass only a fraction of the mass of the observed hot-Jupiter and orbiting close to their parent stars (Lecavelier des Etangs et al.\\ 2004; Baraffe et al.\\ 2005).  However, models were applied to specific cases; it is not yet possible to conclude on the order of magnitude of the possible evaporation of planets which are being discovered without the effort to model in detail each particular case. The purpose of the present paper is not to review previous works and extended modeling efforts. Here we propose an alternative approach to obtain general characteristics from basic observed parameters. This approach allow statistical conclusion without the need to model in detail every individual planet for which detailed physical characteristics are not available. We propose to outline the main characteristics of a given planet that allow concluding on the general trend of its evaporation status, escape rate and lifetime related to the erosion through evaporation.  Here we present the idea of an energy diagram to address these issues (Sect.~\\ref{Potential energy of planets versus stellar energy flux}), and the way to simply estimate the energetic characteristics of a given planet (Sect.~\\ref{The extreme ultraviolet illumination from the central stars} and~\\ref{Potential energy}). Then an evaluation of the lifetime and escape rates is presented in Sect.~\\ref{Lifetime and escape rate}. The resulting diagram is shown in Sect.~\\ref{The energy diagram}. A diagram corresponding to the escape rate is given in Sect.~\\ref{Diagram of the potential energy of the atmosphere and the corresponding escape rate}. Uncertainties and evolutionary tracks are presented in Sect.~\\ref{Uncertainties} and~\\ref{Evolution tracks}. Finally we conclude on the possible nature of the recently discovered Neptune mass planets (Sect.~\\ref{Neptune mass planets in the diagram}). ",
        "conclusions": "To describe the evaporation status of the extrasolar planets, we proposed to consider an energy diagram in which the potential energy of the planets is plotted versus the energy received by the upper atmosphere. We presented a basic method to estimate these quantities. For the energy flux, we proposed a first approach using the type of the parent stars or, if available, their rotation period. For the potential energy, we included the modification of the gravity field by the tidal forces from the parent stars. This description allows a quick estimate of both the escape rate of the atmospheric gas and the lifetime of a planet against the evaporation process. We found that the region of lifetime below 5\\,Gyr is devoid of planets, which reinforces the belief that this diagram is an appropriate description of the evaporation. In the evaporation-forbidden region, the evaporation can modify the nature and the characteristics of the planets such that they move to region of the diagram corresponding to longer lifetime and consequently become long-lived.  The escape rates estimated from the diagram are found to be in the range $10^{5}$\\,g\\,s$^{-1}$ to $10^{12}$\\,g\\,s$^{-1}$, with few cases above $10^{11}$\\,g\\,s$^{-1}$. Although this is a crude estimate, the estimated escape rate for HD\\,209458\\,b is found to be consistent with the lower limit of $10^{10}$\\,g\\,s$^{-1}$ obtained from interpretation of the H\\,{\\sc i} Lyman-$\\alpha$ observations.  Finally, this diagram allows the description of the possible nature of the recently discovered Neptune-mass planets. We found that GJ\\,436\\,b, 55\\,Cnc\\,e and HD\\,69830\\,b cannot be low mass gaseous planets. These planets must contain a large fraction of solid/liquid material. Concerning GJ\\,876\\,d, we argued that it must be dense enough to survive the strong EUV energy flux from its nearby parent star. GJ\\,876\\,d must contain a large fraction of massive elements."
    },
    "0609/astro-ph0609813_arXiv.txt": {
        "abstract": "{}{To model the abundance gradients in the disk  of the Milky Way for several chemical elements (O, Mg, Si, S, Ca, Sc, Ti, Co, V, Fe, Ni, Zn, Cu, Mn, Cr, Ba, La and Eu), and compare our results with the most recent and  homogeneous observational data.} {We adopt a chemical evolution model able to well reproduce the main properties of the solar vicinity. The model assumes that the disk formed inside-out with a timescale for the formation of the thin disk of 7 Gyr in the solar vicinity, whereas the halo formed on a timescale of 0.8 Gyr. We also adopt new empirical stellar yields derived to best fit the abundances and the abundance ratios of the solar vicinity.} {We compute, for the first time, the abundance gradients for all the above mentioned elements in the galactocentric distance range 4 - 22 kpc. The comparison with the observed data on Cepheids in the galactocentric distance range 5-17 kpc gives a very good agreement for many of the studied elements. In addition, we fit very well the data for the evolution of Lanthanum in the solar vicinity for which we present results here for the first time. We explore, also for the first time, the behaviour of the abundance gradients at large galactocentric distances by comparing our results with data relative to distant open clusters and red giants and select the best chemical evolution model model on the basis of that.} {We find a very good fit to the observed abundance gradients, as traced by Cepheids, for most of the elements, thus confirming the validity of the inside-out scenario for the formation of the Milky Way disk as well as the adopted nucleosynthesis prescriptions. } ",
        "introduction": "Understanding the formation and the evolution of the Milky Way is fundamental to improve the knowledge of the formation of spiral galaxies in general. Many are the observational constraints in the Milky Way. The most important ones are represented by the evolution of the abundances of the chemical  elements. Recently, many chemical evolution models have been developed to explain the chemical composition of the solar vicinity and the [el/Fe] vs. [Fe/H] patterns (e.g. Henry et al. 2000; Liang et al. 2001; Chiappini et al. 2003a, 2003b; Akerman et al. 2004, Fran\\c cois et al. 2004). Other important constraints, which are connected  with the evolution of the Galaxy disk, are the abundance gradients of the elements along the disk of the Milky Way. In general term, abundance gradients are a feature commonly observed in many galaxies with their metallicities decreasing outward from the galactic centers. The study of the gradients provides strong constraints to the mechanism of galaxy formation; in fact, the star formation and the accretion history as function of the galactocentric distance in the galactic disk influence strongly the formation and the development of the abundance gradients. (see Matteucci\\& Fran\\c cois 1989, Boisser \\& Prantzos 1999, Chiappini et al. 2001). Many models have been already computed to explain the behaviours of abundances and abundance ratios as functions of galactocentric radius (e.g  Hou et al. 2000; Chang et al. 1999; Chiappini et al. 2003b; Alib\\'es et al. 2001) but they restrict their predictions only to a small number of chemical elements and do not consider very heavy elements.  We base our work on the chemical evolution model described in Chiappini et al. (2001), which is able to well reproduce the gradients in the Milky Way for N, O, S and Fe.  In this work we calculate the behaviour of the largest  number of heavy elements (O, Mg, Si, S, Ca, Sc, Ti, Co, V, Fe, Ni, Zn, Cu, Mn, Cr, Ba, La and Eu) ever considered in this kind of models. In this way we are also able to test the the recent  nucleosynthesis prescriptions described in Fran\\c cois et al. (2004) for the $\\alpha$ and iron peak elements and in Cescutti et al. (2005) for the Eu and Ba, whereas the prescriptions for Lanthanum are newly calculated in this paper following the same approach adopted for Barium in the paper by Cescutti et al. (2005).  Chemical evolution models adopting the above nucleosynthesis prescriptions have already been shown to reproduce the evolution of the abundances in the solar neighborhood. Here we extend our predictions to the whole disk to check if these models can also reproduce the abundance gradients. We compare our model predictions  with new  observational data collected by Andrievsky et al.(2002abc,2004) and Luck et al.(2003) (hereafter 4AL). They measured the abundances of all the selected elements (except Ba) in a  sample of 130 galactic Cepheids found in the galactocentric distance range from 5 to 17 kpc. In addition to the data by 4AL we also compare our theoretical predictions with abundance measurements in giants and open cluster located at even larger galactocentric distances.  The paper is organized as follows: in Section 2 we present the observational data, in Section 3 the chemical evolution model is presented and in Section 4 the adopted nucleosynthesis prescriptions  are described. In Section 5  we present the results  and in Section 6 some conclusions are drawn. ",
        "conclusions": "The aim of this work was to compare new observational data on the radial gradients for  17 chemical elements with the predictions of our chemical evolution model for the Milky Way. This model has been already tested on the properties of the solar vicinity and contains a set of yields which best fit the abundances and abundance ratios in the solar vicinity, as shown in Fran\\c cois et al. (2004).  The bulk of observational data comes from the abundances derived in a large number of Cepheids observed by 4AL. For the first time, it is possible to verify the predictions for many heavy elements with a statistical validity.  The comparison between model predictions and observational data showed that our model well reproduces the gradients of almost all the elements that we analyzed. Since abundance gradients can impose strong constraints  both on the mechanism of galaxy formation, in particular of the galactic disk, and the nucleosynthesis prescriptions, we can conclude that:  \\begin{itemize}  \\item The model for the Milky Way disk formation, assuming an inside-out building-up of the disk, as suggested originally by Matteucci \\& Fran\\c cois (1989), can be considered successful; in fact, for almost all the considered elements, we find a good fit to the observational data ranging from 5 to 17 kpc. In particular, the model assuming a constant total surface mass density for the halo best fits the data of Cepheids. In fact, at large galactocentric distances the halo mass distribution influences the abundance gradients (see Chiappini et al. 2001).  \\item In our chemical evolution model we adopt a threshold in the gas density for star formation in the disk of $7M_{\\odot}pc^{-2}$, whereas for the halo phase we have several options with and without threshold. The threshold in the halo is $4M_{\\odot}pc^{-2}$. We also assume a constant surface mass density for the halo or variable with galactocentric distance. This is important for the gradients at very large galactocentric distances, where the enrichment from the halo predominates over the enrichment occurring in the thin disk thus influencing the abundances at such large distances.  We conclude that to reproduce the flat gradients suggested by the abundance measurements at large galactocentric distances, we need to assume a constant density distribution and a threshold in the star formation during the halo phase. However, there are still many uncertainties in the data at very large galactocentric distances and only more data will allow us to draw firm conclusions on this important point.   \\item The chosen nucleosynthesis prescriptions (empirical yields by Fran\\c cois et al 2004) successfully reproduce the abundances gradients  of each specific element, besides reproducing the [el/Fe] vs [Fe/H] relations in the solar neighborhood, as shown already by Fran\\c cois et al. (2004).    \\item We also presented new results concerning the [La/Fe] vs. [Fe/H] relation in the solar neighbourhood. The data are the new ones by Fran\\c cois et al. (2006). We conclude that La has the same origin as Ba: the bulk of La originates from low mass stars in the range 1-3$M_{\\odot}$ as an s-process elements, but a fraction of  La originates, as an r-process element, from stars in the mass range 12-30$M_{\\odot}$.   \\end {itemize}"
    },
    "0609/astro-ph0609508_arXiv.txt": {
        "abstract": "We have developed a new, quasi-Lagrangian approach for numerical modeling of magnetohydrodynamics in low to moderate $\\beta$ plasmas such as the solar corona. We introduce the concept of a {}``fluxon'', a discretized field line. Fluxon models represent the magnetic field as a skeleton of such discrete field lines, and interpolate field values from the geometry of the skeleton where needed, reversing the usual direction of the field line transform. The fluxon skeleton forms the grid for a collection of 1-D Eulerian models of plasma along individual flux tubes. Fluxon models have no numerical resistivity, because they preserve topology explicitly. Our prototype code, \\emph{FLUX}, is currently able to find 3-D nonlinear force-free field solutions with a specified field topology, and work is ongoing to validate and extend the code to full magnetohydrodynamics. FLUX has significant scaling advantages over conventional models: for {}``magnetic carpet'' models, with photospheric line-tied boundary conditions, FLUX simulations scale in complexity like a conventional 2-D grid although the full 3-D field is represented. The code is free software and is available online. In this current paper we introduce fluxons and our prototype code, and describe the course of future work with the code.\\\\ ~\\\\ ~ ",
        "introduction": "MHD modeling is key to understanding solar eruptive events and their effect on the heliospheric enviroment. Solar flares are known to be driven by magnetic reconnection (e.g. \\citealt{Sturrock1984}) and coronal mass ejections (CMES) are generally tied to magnetic instability of one kind or another (e.g. \\citealt{Sturrock1989,Amari1999,Antiochos1999,Chen2000,Fan2003,Roussev2003}). Understanding and predicting such events requires numeric modeling of the plasma in the solar corona as it evolves under the line-tied boundary conditions imposed by the solar photosphere. Current MHD modeling software is not able to reproduce the conditions of the solar corona under controlled conditions, because numerical effects in the simulation dominate over their counterparts in the real corona for many physical situations. This is evident in the difficulty of maintaining strong current sheets or other stressed flux systems, such as filaments, for long periods of time; and in the difficulty of reproducing the rapidly varying heating rate in coronal loops and bright points.  Conventional magnetohydrodynamic models typically use Eulerian (fixed) grids, introducing nonphysical dissipative effects such as viscosity and resistivity. These dissipative effects dominate the behavior of quiescent structures with current sheets and may even destabilize simulated CME-bearing systems \\citet{Lin2002}, making modeling of solar evolution difficult at best. Further, numerical diffusion of both momentum ('numerical viscosity') and magnetic field ('numerical resistivity') are dependent on grid speed, so that Eulerian grid size is typically chosen to oversample the physical structure. The extra resolution is used to minimize the unwanted diffusion and preserve the plasma system for as much simulated time as possible. Oversampling by a factor of 10 in each dimension yields a factor of $10^{4}$ slowdown in the overall simulation, requiring huge facilities to simulate even simple systems in 3-D. Adaptive-mesh refinement improves performance significantly (e.g. \\citealt{Welsch2004,Lynch2003}) but still requires oversampling.  Lagrangian grids eliminate numerical resistivity but add additional problems. For example, as the grid distorts with the plasma motion, the fidelity of discrete differential operators degrades. Fully Lagrangian treatments of the corona shear rapidly because fluid motion is decoupled in the cross-field direction.  To maximize the advantages of both the Eulerian and Lagrangian approaches to MHD modeling, we have developed a prototype fluxon modeling code, \\emph{FLUX} (the {}``Field Line Universal relaXer''), that is a hybrid between the two. In the low-$\\beta$ regime, all forces are negligible when compared to the Lorenz force components; FLUX is essentially a force-free field solver that can support an independent plasma density parameter at each location in the simulation. FLUX demonstrates the fluxon method, and work is ongoing to add plasma pressure and related physical phenomena to the simulation framework. In the following sections, we briefly describe the basis of the fluxon numerical approach (\\S\\ref{sec:Fluxon-theory}), describe our code and its performance (\\S\\ref{sec:FLUX-implementation}), and discuss the direction of future work both on code development and on applications (\\S\\ref{sec:Future-work}). ",
        "conclusions": ""
    },
    "0609/astro-ph0609022_arXiv.txt": {
        "abstract": "We present new results of a spectroscopic survey of circumstellar \\HI in the direction of evolved stars made with the Nan\\c cay Radiotelescope. The \\HI line at 21 cm has been detected in the circumstellar shells of a variety of evolved stars\\,: AGB stars, oxygen-rich and carbon-rich, Semi-Regular and Miras, and Planetary Nebulae. The emissions are generally spatially resolved, i.e. larger than 4$'$, indicating shell sizes of the order of 1 pc which opens the possibility to trace the history of mass loss over the past $\\sim 10^4-10^5$ years. The line-profiles are sometimes composite. The individual components have generally a quasi-Gaussian shape; in particular they seldom show the double-horn profile that would be expected from the spatially resolved optically thin emission of a uniformly expanding shell. This probably implies that the expansion velocity decreases outwards in the external shells (0.1--1 pc) of these evolved stars.  The \\HI line-profiles do not necessarily match those of the CO rotational lines. Furthermore, the centroid velocities do not always agree with those measured in the CO lines and/or the stellar radial velocities. The \\HI emissions may also be shifted in position with respect to the central stars. Without excluding the possibility of asymmetric mass ejection, we suggest that these two effects could also be related to a non-isotropic interaction with the local interstellar medium.  \\HI was detected in emission towards several sources ($\\rho$ Per, $\\alpha$ Her, $\\delta^2$ Lyr, U CMi) that otherwise have not been detected in any radio lines. Conversely it was not detected in the two oxygen-rich stars with substantial mass-loss rate, NML Tau and WX Psc, possibly because these sources are young with hydrogen in molecular form, and/or because the temperature of the circumstellar \\HI gas is very low ($<$ 5 K). ",
        "introduction": "Hydrogen is the most abundant element in the atmospheres of Asymptotic Giant Branch (AGB) stars and in their outflows. It may be found in the form of atomic hydrogen, H\\,{\\sc {i}}, or molecular hydrogen, H$_2$, besides relatively minor (in abundance) species. The fraction of hydrogen in each of these two forms is an important factor controlling the atmospheric structures of red giants and also the winds which develop from their atmospheres (e.g. Lamers \\& Cassinelli 1999).  It is therefore of interest to study the circumstellar environments of AGB stars in the low excitation hydrogen lines. Unfortunately molecular hydrogen is presently difficult to detect because its lowest rotational lines lie in a region of the spectrum which cannot be accessed from the ground, and because a high spectral resolution is needed.  On the other hand atomic hydrogen has a line at 21 cm in an easily accessible, and well protected, region of the radio spectrum. However the observations are difficult because the intensity is weak and above all because there is a contamination by interstellar hydrogen along the same lines of sight. For a long time these difficulties have limited the \\HI observations of AGB sources to $o$ Cet (Bowers \\& Knapp 1988, BK1988). Nevertheless we recently readdressed this question with the upgraded Nan\\c{c}ay Radiotelescope and could detect this line in emission towards several AGB sources. We have already presented our results on four of them: RS Cnc (G\\'erard \\& Le~Bertre 2003, Paper I), EP Aqr and Y CVn (Le~Bertre \\& G\\'erard 2004, Paper II) and X Her (Gardan et al. 2006, Paper III). We detected these four AGB sources with a good signal-to-noise ratio and showed that the line profiles are composite. The individual components have Gaussian-like profiles, although in some cases a rectangular profile could be fitted as well. The \\HI emissions are spread over a velocity range confined to the total CO velocity extent. However, we noted that sometimes the central velocities in CO and \\HI do not coincide exactly. Furthermore, in Paper I, we argued that atomic hydrogen is present already in the stellar atmosphere of the central star (RS Cnc) which might even be devoid of molecular hydrogen. In Paper II we could resolve angularly two circumstellar shells, and show that \\HI may be present at large distances from the central stars, 0.5 pc or more for EP Aqr. The \\HI data display a large scale asymmetry suggesting that asymmetric structures may show up early in the AGB phase (EP Aqr \\& X Her), well before the planetary nebula phase. In the case of the source associated to Y CVn, the \\HI data even seem to probe a large region extending out to where the stellar outflow is interacting with the Interstellar Medium (ISM).  It thus appears that the \\HI emission line profiles, when observed with a good signal-to-noise ratio and a fair removal of the ISM contamination, may reveal important clues on the physics of stellar outflows and their kinematics. Also, as atomic hydrogen has a long lifetime in the ISM, it can be used to trace circumstellar envelopes out to large radial distances. In the present paper, we give results on 22 new sources for which we have now good quality spectra, or good upper limits. The main emphasis of our programme is put on AGB sources and we have tried to select a variety of representative specimens in terms of variability types, chemical composition, mass loss rates, etc. In addition, we have also included two planetary nebulae in order to sample more completely the post-main-sequence evolution of stars of low and intermediate masses. Two of our sources have already been observed in \\HI with the Very Large Array (VLA), $o$ Ceti (BK1988) and NGC\\,7293 (Rodr{\\'\\i}guez et al. 2002, R2002). This allows a useful comparison between single-dish and interferometric data. ",
        "conclusions": "We have undertaken a search for the \\HI line at 21 cm in the direction of AGB stars and related sources. We have detected circumstellar \\HI in a variety of sources sampling the post main-sequence evolution of low and intermediate mass stars. The emissions are extended indicating shell sizes from 0.2 pc to more than 2 pc. The \\HI line can therefore be used to trace the mass loss history for long periods, up to several 10$^5$ years. However, we may have missed atomic hydrogen at low temperature ($<$~10~K) and hydrogen locked into H$_2$.  For some sources ($\\rho$ Per, $\\alpha$ Her, $\\delta^2$ Lyr, U CMi) our \\HI detection is the first radio detection.  The \\HI line-profiles indicate that the outflows are slowing down, probably due to the interaction with surrounding matter. Furthermore, we observe offsets in velocity and spatial asymmetries that indicate a distortion of the outflows. These effects were already noted for EP Aqr (Paper II) and for X Her (Paper III), and seem frequent. They could be connected to the motion of the sources with respect to the local ISM, or to its inhomogeneity, as well as to the mass ejection process itself."
    },
    "0609/astro-ph0609214_arXiv.txt": {
        "abstract": "We have developed a position-sensitive scintillation camera with a large area absorber for use as an advanced Compton gamma-ray camera. At first we tested GSO(Ce) crystals. We compared light output from the GSO(Ce) crystals under various conditions: the method of surface polishing, the concentration of Ce, and co-doping Zr. As a result, we chose the GSO(Ce) crystals doped with only 0.5 mol\\% Ce, and its surface polished by chemical etching as the scintillator of our camera. We also made a 16$\\times$16 cm$^2$ scintillation camera which consisted of 9 position-sensitive PMTs (PSPMTs Hamamatsu flat-panel H8500 ), the each of which had 8$\\times$8 anodes with a pitch of 6 mm and coupled to  8$\\times$8 arrays of pixelated 6$\\times$6$\\times$13 mm$^3$ GSO(Ce) scintillators. For the readout system of the 576 anodes of the PMTs, we used chained resistors to reduce the number of readout channels down to 48 to reduce power consumption. The camera has a position resolution of less than 6mm and a typical energy resolution of 10.5\\% (FWHM) at 662 keV at each pixel in a large area of 16$\\times$16 cm$^2$.  Furthermore we constructed a 16$\\times$16 array of 3$\\times$3$\\times$13 mm$^3$ pixelated GSO(Ce) scintillators, and glued it to a PMT H8500. This camera had the position resolution of less than 3mm, over an area of 5$\\times$5 cm$^2$, except for some of the edge pixels; the energy resolution was typically 13\\% (FWHM) at 662 keV. ",
        "introduction": "We have developed an advanced Compton camera for gamma-ray astronomy in the range of 100 keV to 20 MeV \\cite{takada}. It needs a scintillation camera as a detector for Compton-scattered gamma rays with a good energy and position resolution and a large area because the resolution and the efficiency for the scattered gamma rays contribute to the angular resolution and the efficiency of the advanced Compton camera. In addition, radiation hardness and a high counting-rate performance of the scintillation camera are required. For these reasons, we chose a GSO(Ce) (Gd$_2$SiO$_5$:Ce) crystal as a scintillator, and the PSPMT H8500 (Flat Panel PMT produced by Hamamatsu\\cite{hamamatsu}) as a position-sensitive photon device.   GSO(Ce) has advantages in astronomical use, such as having a higher-Z, faster decay time than NaI(Tl), a higher light output than BGO, greater radiation hardness and less radioactivation than most other known scintillators. Furthermore, GSO(Ce) can be easily cut and polished, since it is nonhydroscopic.  The PSPMT H8500 was recently developed as a promising device for nuclear physics and medicine, for example, PET and SPECT \\cite{pani} \\cite{Gim} \\cite{Herber} \\cite{houiken}. It has 8$\\times$8 anodes with a 6 mm pitch and 12-stage metal channel dynodes. The advantage of this PMT is  % that it has a much smaller dead space and a larger effective area than that of the previous multi-anode PMTs. The effective area of this PMT is 49$\\times$49 mm$^2$, which is 89\\% of the package size.   In this paper we report on the results of measurements with pixelated GSO(Ce) scintillators and the performance (energy resolution and position resolution) of the scintillation camera developed by us. ",
        "conclusions": ""
    },
    "0609/astro-ph0609164_arXiv.txt": {
        "abstract": "Irregular structure in planetary rings is often attributed to the intrinsic instabilities of a homogeneous state undergoing Keplerian shear. Previously these have been analysed with simple hydrodynamic models. We instead employ a kinetic theory, in which we solve the linearised moment equations derived in Shu and Stewart 1985 for a dilute ring. This facilitates an examination of velocity anisotropy and non-Newtonian stress, and their effects on the viscous and viscous/gravitational instabilities thought to occur in Saturn's rings. Because we adopt a dilute gas model, the applicability of our results to the actual dense rings of Saturn are significantly curtailled. Nevertheless this study is a necessary preliminary before an attack on the difficult problem of dense ring dynamics.  We find the Shu and Stewart formalism admits analytic stability criteria for the viscous overstability, viscous instability, and thermal instability. These criteria are compared with those of a hydrodynamic model incorporating the effective viscosity and cooling function computed from the kinetic steady state. We find the two agree in the `hydrodynamic limit' (i.e. many collisions per orbit) but disagree when collisions are less frequent, when we expect the viscous stress to be increasingly non-Newtonian and the velocity distribution increasingly anisotropic. In particular, hydrodynamics predicts viscous overstability for a larger portion of parameter space. We also numerically solve the linearised equations of the more accurate Goldreich and Tremaine 1978 kinetic model and discover its linear stability to be qualitatively the same as that of Shu and Stewart's. Thus the simple collision operator adopted in the latter would appear to be an adequate approximation for dilute rings, at least in the linear regime. ",
        "introduction": "This paper examines the collisional dynamics of differentially rotating, particulate disks, and addresses specifically the intrinsic mechanisms which might induce the generation of irregular, fine-scale structure, such as that exhibited by Saturn's B-ring.  \\emph{Voyager}, and \\emph{Cassini} more recently, revealed that these essentially axisymmetric variations appear on a very wide range of length-scale-- from the limit of resolution ($\\sim$ 100m) to several hundred kilometres (Horn and Cuzzi 1995, Porco \\emph{et al.} 2005 ). Though some proportion of B-ring structure may be caused by non-dynamical effects such as variations in albedo and phase function- see Cuzzi and Estrada 1996, for example) and from gravitational interactions with neighbouring satellites (Thiessenhusen \\emph{et al.} 1995), it is believed that the majority corresponds to changes in optical thickness resulting from a collective dynamics (Tremaine 2003).  This work studies \\emph{dilute} particulate systems, and as such may not be directly applicable to the saturnian rings. It is nevertheless an informative first stop before we embark on the challenging stability analysis of a dense granular gas. Other than presenting a general framework in which such a calculation can be tackled, the dilute gas formalism permits us to isolate analytically the interesting effects of anisotropy and non-Newtonian viscous stress upon the local axisymmetric instabilities thought to be manifest in planetary rings.  Since \\emph{Voyager} first reported the radial stratification of Saturn's rings, theoreticians have advocated a number of possible causes. These have included: ballistic transport (Durisen 1995),  electromagnetic effects (Goertz and Morfill 1988), and interleaved shearing and shear-free zones (Tremaine 2003). Others have looked to the local instabilities of viscous fluid disks, interpreting the disordered state observed as the saturated endpoint of their nonlinear evolution. The `viscous instability' was the first proposed (Lin and Bodenheimer 1981, Ward 1981, and Lukkari 1981). Essentially a monotonic `clumping', it is associated with an outward angular momentum flux which decreases with surface density: $d(\\nu \\sigma)/d\\sigma <0$ (where $\\nu$ is kinematic viscosity and $\\sigma$ surface mass density). Though a dilute ring's viscosity depends on optical depth in a manner which promises the existence of such an instability (Goldreich and Tremaine 1978, Shu and Stewart 1985), Saturn's rings are most likely `dense', and theoretical and numerical N-body studies have revealed that such rings do not manifest the appropriate viscosity for this instability to develop (Araki and Tremaine 1986, and Wisdom and Tremaine 1988).  The `viscous overstability' was first examined in the context of accretion disks by Kato (1978), and, as the name suggests, originates in overcompensation by the system's restoring forces: the stress oscillation which accompanies the epicyclic response in an acoustic-inertial wave will force the system back to equilibrium so strongly that it will `overshoot'. The mechanism relies on: a) the synchronisation of the viscous stress's oscillations with those of density, and b) the viscous stress increasing sufficiently in the compressed phase. In hydrodynamics only the latter consideration is relevant, which furnishes the criterion for overstability: $\\beta\\equiv (d\\ln\\nu/d\\ln \\sigma)>\\beta^*$, where $\\beta^*$ is a number dependent on the thermal properties of the ring (Schmit and Tscharnuter 1995).  The viscous overstability has been a favoured explanation for smaller scale B-ring structure in recent years, a status stemming, primarily, from the viscosity profiles computed by Araki and Tremaine's 1986 dense gas model and Wisdom and Tremaine's 1988 particle simulations. Both appear to satisfy the above criterion. Consequently the linear behaviour of the instability has been thoroughly examined, though only within a hydrodynamic framework (Schmit and Tscharnuter 1995, Spahn \\emph{et. al.} 2000, and Schmidt \\emph{et al.} 2001). In addition, Schmidt and Salo (2003) have constructed a weakly nonlinear theory, and the overstability's long-term, nonlinear behaviour has been numerically studied by Schmit and Tscharnuter (1999). An isothermal model was adopted in both cases. \\vskip0.3cm  However, fluid mechanics is only one of several theoretical approaches thathave been deployed to capture the behaviour of planetary rings: N-body simulations, generalisations of stellar dynamics, and kinetic theory. This plurality is surely indicative of how the field of planetary ring dynamics falls uncomfortably between the more familiar frameworks of classical physics. Our analysis will predominantly focus upon the kinetic approach, and solve the set of moment equations derived from the kinetic theories proposed in Shu and Stewart 1985 (hereafter referred to as `SS85' ) and in Goldreich and Tremaine 1978 (referred to as `GT78'). These formalisms have hitherto been used mainly to establish equilibrium solutions. This paper goes the next step by exploring their linear theory.  Such an analysis is more involved than the analogous hydrodynamic calculation; however, the latter's adoption of the Navier-Stokes stress model introduces two assumptions which may be inappropriate in the ring context and whose consequences are instructive to investigate. Firstly, the Navier-Stokes model presumes the distribution of the particles' velocity dispersion to be nearly isotropic. In the regime of many collisions per orbit this is an acceptable supposition, as collisions scatter particles randomly on the average. However if the collision rate, $\\omega_c$, is of the same order as the orbital frequency, $\\Omega$, (as it is presumed to be in Saturn's rings) this need not be true. Secondly, hydrodynamics assumes an `instantaneous' (local in time) relationship between stress and strain. This may not hold when $\\omega_c\\sim\\Omega$. Generally the viscous stress possesses a relaxation time of order $1/\\omega_c$, which in this regime will be comparable to the dynamic time scale. Thus the immediate history of the stress cannot be ignored. Its inclusion should have the most impact on the stability of oscillating modes, especially the overstability, it depending on the synchronisation of the stress and density oscillations. A kinetic model can address both issues, accounting for anisotropy within an appropriate collision term and providing a straightforward way, by the taking of moments, to generate dynamical equations for the viscous stress. Another advantage is that a kinetic model lets us explicitly include the microphysics of particle-particle interactions and thence to potentially model a larger set of the physical mechanisms at play (such as collisions, irregular surfaces, spin , size distribution, etc). It also narrows the scope of our simplifying assumptions to the particulars of collisions between spheres of ice, which have been observed in the laboratory (see Bridges \\emph{et al.} 1984, Hatzes \\emph{et al.} 1988, amongst others) . \\vskip0.3cm The paper is organised as follows. In Section 2 we present a brief derivation of the set of closed, vertically averaged moment equations in a local model, the shearing sheet, for a general kinetic equation. We then specialise to the SS85 kinetic model and briefly exhibit its equilibrium characteristics. Section 3 contains the linear theory of a general kinetic equation and proves results pertaining to self-gravity and the thermal modes when in the limit of length-scales much longer than the disk height. We then derive stability criteria for the viscous instability and overstability specific to the SS85 model. Comparisons with an analogous calculation for a hydrodynamic model and the GT78 model are performed in Sections 4 and 5. In the former we isolate the effects of non-Newtonian stress on the linear theory, and in the latter observe how the SS85 collision term approximates satisfactorily that of the more accurate triaxial Gaussian model. We draw our conclusions in Section 6. ",
        "conclusions": "In summary, we have calculated the linear stability of a patch of a dilute planetary ring in a shearing sheet employing a second-order kinetic theory. We have defined `diluteness' as the insignificance of non-local effects, whose importance we can roughly quantify by $a\\Omega/c$ for lower optical depths. But throughout Saturn's rings the velocity dispersion is low enough and particle size large enough, for this quantity to be non-negligible, thus we recognise that the applicability of a dilute model is restricted. Nevertheless this work reveals the interesting effects anisotropy and non-Newtonian stress have on the viscous instability and overstability. It also provides a preliminary framework in which a linear analysis of a dense gas may be undertaken.  We adopted a system of vertically integrated moment equations derived from Shu and Stewart's inelastic generalisation of the BGK equation in their 1985 paper. The ensuing system manifests both the viscous instability and overstability in the long-wavelength limit, and their extensions induced by self-gravity on intervals of intermediate wavelength. It predicts viscous or viscous/gravitational instability in dilute rings where particles collide about 3.5 times or more per orbit (optical thicknesses greater than 0.2, or less depending on the importance of self-gravity on the collision frequency). Viscous overstability does not feature in dilute regions.  These results were compared with a hydrodynamic model incorporating the kinetic theory's effective viscosity (evaluated at equilibrium) and cooling function. We find that the criterion for viscous instability is precisely the same as that for the kinetic model, but the two criteria diverge in the intermediate range of collision frequency for the viscous overstability. According to the kinetics, overstability is suppressed in this regime, a result in contradiction with the fluid model which predicts the instability to be widespread. We conclude that hydrodynamics presents a misleading impression of the viscous overstability's prevalence on small and intermediate ranges of $\\omega_c/\\Omega_0$, for a dilute ring. This arises from the increasingly non-Newtonian nature of the stress in this regime, an effect the kinetic account captures but which the hydrodynamical cannot. We stress that this incongruence will be most marked for a dilute ring; as the collisional stress tensor is local in time, we assume the denser the ring the better the hydrodynamic approximation.  We also undertook a linear stability analysis of the Goldreich and Tremaine 1978 kinetic model, employing a Boltzmann-type collision term with a triaxial Gaussian, which should factor in anisotropic effects more thoroughly. The criteria for viscous instability and overstability mirror qualitatively that of the SS85 model, and thus suggest that the linear dynamics of a disk are insensitive to the the precise details of the collision term. In particular, the criterion of viscous instability matches closely that of the SS85 model. The mechanism of this instability would thus seem relatively robust, not only resistant to the effects of non-locality in time but also of anisotropy. The discrepancy between the marginal curves of viscous overstability computed by the SS85 and GT78 models is a little more pronounced. We conclude the viscous overstability is relatively fragile for small and intermediate $\\omega_c/\\Omega_0$--- very susceptible to non-locality in time and slightly altered by anisotropy  On intermediate and short scales the SS85 model also predicted the emergence of a quasi-thermal instability which the hydrodynamical analysis revealed to be associated with a mode orienting the thermal and density perturbations in such a way to keep the equation of motion undisturbed. This mechanism thus mimicks that operating in Field's thermal instability (Field 1965). We find that the instability should occur on optical thicknesses less than 2. However hydrodynamics shows that this instability is severely sensitive to heat conduction; $\\kappa$ need not be large to extinguish it entirely. Thus the omission of third-order moments in the kinetic model bestows on it a prominence that it does not deserve. When self-gravity is included intermediate scale instability can occur for much larger $\\kappa$ which in this case we refer to as a `viscous/gravitational instability'. Self-gravity, in addition to its important dynamic role on intermediate scales, will enhance the collision frequency, and render the dependence of $\\omega_c$ on $N$ as steeply as $N^{3/2}$. In general this will make the system more susceptible to both the viscous instability and overstability on long lengthscales.  Two possibly fruitful avenues which proceed from this work are: linear stability calculations of dense and/or spinning particle disks, and nonlinear studies of the evolution of the viscous instability and overstability using the second-order moment equations.  In principle the techniques employed here can be set upon kinetic models of dense and spinning rings. Detailed formulations have been developed in Shukhman 1984, Araki and Tremaine 1986, Araki 1988 and Araki 1991, but, unmodified, their algebraical truculence surely prohibits all but equilibrium calculations. The success of the SS85 model encourages us that a simplification of the dense disk collision terms could provide a model that captures the essential physics, and also expedites linear and nonlinear analysis. The results obtained with such a model would be directly pertinent to N-body simulations and Saturn's B-ring. It could also ascertain the validity of the fluid description in the dense regime.  It would be instructive also to determine the impact of the nonlinear evolution of the second-order moments on the development of the instabilities discussed, as compared to the Newtonian fluid model, which so far has only been explored in the isothermal case (Schmit and Tscharnuter 1999). They may elicit interesting nonlinear behaviour from the viscous overstability. But an understanding of the small-scale instabilities, even if unphysical, is essential for the sucessful numerical implementation of these equations.  \\textbf"
    },
    "0609/astro-ph0609597_arXiv.txt": {
        "abstract": "In this work, we confront ten cosmological models with observational $H(z)$ data. The possible interaction between dark energy and dust matter is allowed in some of these models. Also, we consider the possibility of (effective) equation-of-state parameter (EoS) crossing $-1$. We find that the best models have an oscillating feature for both $H(z)$ and EoS, with the EoS crossing $-1$ around redshift $z\\sim 1.5$.\\\\  \\noindent PACS numbers: 95.36.+x, 98.80.Es, 98.80.-k ",
        "introduction": "\\label{sec1}  A lot of cosmological observations, such as SNe Ia~\\cite{r1,r2}, WMAP~\\cite{r3}, SDSS~\\cite{r4}, Chandra X-ray Observatory~\\cite{r5} etc., find that our universe is experiencing an accelerated expansion. These results also suggest that our universe is spatially flat, and consists of about $70\\%$ dark energy with negative pressure, $30\\%$ dust matter (cold dark matters plus baryons), and negligible radiation. Dark energy study has been one of the most active fields in modern cosmology~\\cite{r6}.  In the past years, many cosmological models are constructed to interpret the present accelerated expansion. One of the important tasks is to confront them with observational data. The most frequent method to constrain the model parameters is fitting them to the luminosity distance \\be{eq1} d_L(z)=(1+z)\\int_0^z \\frac{d\\tilde{z}}{H(\\tilde{z})}\\,, \\ee which is an integral of Hubble parameter $H\\equiv\\dot{a}/a$, where $a=(1+z)^{-1}$ is the scale factor ($z$ is the redshift); a dot denotes the derivative with respect to cosmic time $t$. In this work, we use the observational $H(z)$ data directly, rather than the luminosity distance $d_L(z)$. The observational $H(z)$ data are based on differential ages of the oldest galaxies~\\cite{r7}. In~\\cite{r8}, Jimenez {\\it et al.} obtained an independent estimate for the Hubble constant by the method developed in~\\cite{r7}, and used it to constraint the equation-of-state parameter (EoS) of dark energy. The Hubble parameter depends on the differential age as a function of redshift $z$ in the form \\be{eq2} H(z)=-\\frac{1}{1+z}\\frac{dz}{dt}\\,. \\ee Therefore, a determination of $dz/dt$ directly measures $H(z)$~\\cite{r9}. By using the differential ages of passively evolving galaxies determined from the Gemini Deep Deep Survey (GDDS)~\\cite{r10} and archival data~\\cite{r11}, Simon {\\it et al.} determined $H(z)$ in the range $0\\,\\lsim\\, z\\,\\lsim\\, 1.8$~\\cite{r9}. The observational $H(z)$ data from~\\cite{r9} are given in Table~\\ref{tab1} and shown in Figs.~\\ref{fig2}--\\ref{fig5}.  \\begin{table}[htbp] \\begin{center} \\begin{tabular}{c|lllllllll}\\hline\\hline $z$ &\\ 0.09 & 0.17 & 0.27 & 0.40 & 0.88 & 1.30 & 1.43 & 1.53 & 1.75\\\\ \\hline $H(z)\\ ({\\rm km~s^{-1}\\,Mpc^{-1})}$ &\\ 69 & 83 & 70 & 87 & 117 & 168 & 177 & 140 & 202\\\\ \\hline $1 \\sigma$ uncertainty &\\ $\\pm 12$ & $\\pm 8.3$ & $\\pm 14$ & $\\pm 17.4$ & $\\pm 23.4$ & $\\pm 13.4$ & $\\pm 14.2$ & $\\pm 14$ &  $\\pm 40.4$\\\\ \\hline\\hline \\end{tabular} \\end{center} \\caption{\\label{tab1} The observational $H(z)$ data~\\cite{r8,r9} (see~\\cite{r15} also).} \\end{table}  These observational $H(z)$ data have been used to constrain the dark energy potential and its redshift dependence by Simon {\\it et al.} in~\\cite{r9}. Yi and Zhang used them to constrain the parameters of holographic dark energy model in~\\cite{r12}. Some relevant works also include~\\cite{r13,r14}. Recently, in~\\cite{r15}, Samushia and Ratra have used these observational $H(z)$ data to constrain the parameters of $\\Lambda$CDM, XCDM and $\\phi$CDM models. In this work, we will revisit these observational $H(z)$ data and compare them with some cosmological models.  By looking carefully on the observational $H(z)$ data given in Table~\\ref{tab1} and shown in Figs.~\\ref{fig2}--\\ref{fig5}, we notice that two data points near $z\\sim 1.5$ and $0.3$ are very special. They deviate from the main trend and dip sharply, especially the one near $z\\sim 1.5$; the $H(z)$ decreases and then increases around them. This hints that the effective EoS crossed $-1$ there. This possibility has not been discussed in previous works~(e.g.~\\cite{r12,r15}). In the present work, we will seriously explore this possibility.   \\begin{center} \\begin{figure}[htbp] \\centering \\includegraphics[width=0.45\\textwidth]{int.eps} \\caption{\\label{fig1} The quantity $L(z)\\equiv H^2(z)/H_0^2-\\Omega_{m0}(1+z)^3$ versus redshift $z$, for the fiducial parameters $H_0=72~{\\rm km~s^{-1}\\,Mpc^{-1}}$ and $\\Omega_{m0}=0.3$.} \\end{figure} \\end{center}   On the other hand, we also consider the possible interaction between dark energy and dust matter. This is inspired by the data point near $z\\sim 1.5$, which dips so sharply and stays clearly outside of the best-fit of the $\\Lambda$CDM, XCDM and $\\phi$CDM models studied in~\\cite{r15}. In Fig.~\\ref{fig1}, we show the quantity $L(z)\\equiv H^2(z)/H_0^2-\\Omega_{m0}(1+z)^3$ versus redshift $z$, which is associated with the fractional energy density of dark energy, for the fiducial parameters $H_0=72~{\\rm km~s^{-1}\\,Mpc^{-1}}$ and $\\Omega_{m0}=0.3$, where the subscript ``0'' indicates the present value of the corresponding quantity. It is easy to see that the fractional energy density of dark energy of the point near $z\\sim 1.5$ is negative (beyond $1\\sigma$ significance). To avoid this, one can decrease the corresponding $\\Omega_{m0}$ or make the matter decrease with the expansion of our universe slower than $a^{-3}$. Inspired by this, it is natural to consider the possibility of exchanging energy between dark energy and dust matter through interaction.  We assume that dark energy and dust matter exchange energy through interaction term $C$, namely \\bea &&\\dot{\\rho}_{de}+3H\\left(\\rho_{de}+p_{de}\\right)=-C,\\label{eq3}\\\\ &&\\dot{\\rho}_m+3H\\rho_m=C,\\label{eq4} \\eea which preserves the total energy conservation equation $\\dot{\\rho}_{tot}+3H\\left(\\rho_{tot}+p_{tot}\\right)=0$. The interaction forms extensively considered in the literature (see~\\cite{r16,r17,r18,r19,r20,r21,r22,r23,r24,r25,r26,r39} for instance) are $$C\\propto H\\rho_m,\\ H\\rho_{tot},\\ H\\rho_{de}, \\ \\kappa\\rho_m\\dot{\\phi},\\ \\ldots$$ In this work, we consider the simplest case for convenience, i.e. \\be{eq5} C=3\\alpha H\\rho_m, \\ee where $\\alpha$ is a dimensionless constant. Combining Eqs.~(\\ref{eq4}) and~(\\ref{eq5}), it is easy to get \\be{eq6} \\rho_m=\\rho_{m0}\\,a^{-3(1-\\alpha)}=\\rho_{m0}(1+z)^{3(1-\\alpha)}. \\ee  In the following sections, we will compare the observational $H(z)$ data with some cosmological models with/without interaction between dark energy and dust matter. We consider a spatially flat FRW universe throughout. We adopt the prior $H_0=72~{\\rm km~s^{-1}\\,Mpc^{-1}}$, which is exactly the median value of  the result from the Hubble Space Telescope (HST) key project~\\cite{r27}, and is also well consistent with the one from WMAP 3-year result~\\cite{r3}. Since there are only 9 observational $H(z)$ data points and their errors are fairly large, they cannot severely constrain model parameters alone. We perform a $\\chi^2$ analysis and compare the cosmological models to find out the one which catches the main features of the observational $H(z)$ data. We determine the best-fit values for the model parameters by minimizing \\be{eq7} \\chi^2(parameters)=\\sum\\limits_{i=1}^9\\frac{\\left[H_{mod}(parameters;z_i) -H_{obs}(z_i)\\right]^2}{\\sigma^2(z_i)}\\,, \\ee where $H_{mod}$ is the predicted value for the Hubble parameter in the assumed model, $H_{obs}$ is the observed value, $\\sigma$ is the corresponding $1\\sigma$ uncertainty, and the summation is over the 9 observational $H(z)$ data points at redshift $z_i$. ",
        "conclusions": "\\label{sec5}  In Table~\\ref{tab2}, we summarize all ten models considered in this work. It is obvious that they divide into two groups. The first six models fail to catch the main features of the observational $H(z)$ data, and hence have large $\\chi_{min}^2$ and $\\chi_{min}^2/dof$. The last four models with oscillating $H(z)$ ansatz, on the contrary, are well compatible with the observational $H(z)$ data. They can reproduce the sharp dip around $z\\sim 1.5$, and their effective EoS crossed $-1$ also. Therefore, it is not surprising that their corresponding $\\chi_{min}^2$ and $\\chi_{min}^2/dof$ are significantly smaller than the ones of the first six models.  \\begin{table}[htbp] \\begin{center} \\begin{tabular}{c|c|c|c|c}\\hline\\hline Model & $\\hspace{4mm}\\chi_{min}^2\\hspace{4mm}$ &$\\ \\chi_{min}^2/dof\\ $ & $P\\left(\\chi^2>\\chi_{min}^2\\right)$ & $\\begin{array}{c} {\\rm Ranked\\ by}\\\\ P\\left(\\chi^2>\\chi_{min}^2\\right) \\end{array}$\\\\ \\hline $\\Lambda$CDM & 9.04 & 1.13 & 0.34 & 5 \\\\ Int$\\Lambda$CDM & 8.89 & 1.27 & 0.26 & 6 \\\\ XCDM & 9.02 & 1.29 & 0.25 & 7 \\\\ IntXCDM & 8.48 & 1.41 & 0.21 & 9 \\\\ VecDE & 9.05 & 1.29 & 0.25 & 8 \\\\ IntVecDE & 9.04 & 1.51 & 0.17 & 10 \\\\ \\hline OA1 & 4.27 & 0.85 & 0.51 & 3 \\\\ IntOA1 & 3.54 & 0.89 & 0.47 & 4 \\\\ OA2 & 2.81 & 0.56 & 0.73 & 1 \\\\ IntOA2 & 2.80 & 0.70 & 0.59 & 2 \\\\ \\hline\\hline \\end{tabular} \\end{center} \\caption{\\label{tab2} Summarizing all ten models considered in this work.} \\end{table}  Some remarks are in order. The first one is on the failure of the vector-like dark energy model to reproduce the main features of the observational $H(z)$ data, since the EoS of vector-like dark energy and then the effective EoS can cross $-1$ in principle. According to~\\cite{r16}, for the case of interaction term $C=3\\alpha H\\rho_m$, the vector-like dark energy behaves like a pure cosmological constant in the late time. So, it is not surprising that the vector-like dark energy model considered in this work is very close to the $\\Lambda$CDM model. However, when the interaction term $C$ takes other forms, the late time behavior of vector-like dark energy can be considerablely different, as shown explicitly in~\\cite{r16}. Therefore, it is interesting to compare the observational $H(z)$ data with the interacting vector-like dark energy model with other interaction forms and different potentials. We leave this to future works.  Secondly, we note that the best-fit value $\\Omega_{m0}=1.002$ of IntOA1 model is inconsistent with the results from clusters of galaxies~\\cite{r38} and 3-year WMAP~\\cite{r3} etc. In this sense, the IntOA1 model may be ruled out, although it fits the observational $H(z)$ data fairly well.  The third remark is concerned with EoS crossing $-1$. There are many pieces of other observational evidence for this in the literature~\\cite{r30,r31,r32,r33}. Also, a lot of theoretical models whose EoS can cross $-1$ have been built (see for examples~\\cite{r16,r17,r24,r32,r34,r40} and references therein). In this work, we present independent evidence for EoS crossed $-1$, from the observational $H(z)$ data.  Fourthly, we see that the last four models fit the data much better than the first six models. We note that the last four models have an oscillating feature, for both predicted $H(z)$ and EoS, as shown in Figs.~\\ref{fig4} and~\\ref{fig5}. In fact, this is also the main ideas of~\\cite{r35,r36,r37,r41}. However, the last four models studied here and the ones considered in~\\cite{r35,r36,r37,r41} are all {\\em parameterized} models. We consider that it is important to build some physically motivated models in which the oscillating EoS arises naturally.  Fifthly, we admit that although the first six models have considerably larger $\\chi^2_{min}/dof$, they are not unacceptable. For instance, the $\\Lambda$CDM model has $\\chi^2_{min}=9.04$ for eight degrees of freedom with $34\\%$ probability, which is not unacceptably low. Also, before the new and improved $H(z)$ data are available, it is too early to talk about {\\em strong observational evidence} for non-monotonic behavior of $H(z)$, since one can wonder that the dip around $z\\sim 1.5$ could be due to unknown measurement errors and so on. Therefore, to firmly rule out non-oscillating models~(e.g. the first six models considered here), the more and better observational $H(z)$ data are required.  Finally, the observational $H(z)$ data provide an independent approach to constrain the cosmological models. However, by now, the observational $H(z)$ data only have 9 data points and their error bars are fairly large. Hence, they cannot severely constrain the cosmological models so far. A good news from~\\cite{r15} is that a large amount of $H(z)$ data is expected to become available in the next few years. These include data from the AGN and Galaxy Survey (AGES) and the Atacama Cosmology Telescope (ACT), and by 2009 an order of magnitude increase in $H(z)$ data is anticipated."
    },
    "0609/astro-ph0609074_arXiv.txt": {
        "abstract": "We consider the corrugation instability of the self--similar flow with an accelerating shock in the highly relativistic regime. We derive the correct dispersion relation for the proper modes in the self--similar regime, and conclude that this solution is unstable. ",
        "introduction": "{}  It has long been known that the propagation of a shock through the outer layers of a star may, for sufficiently steep density distributions, lead to shock acceleration. Self--similar solutions (Gandel'man and Frank-Kamenetskii 1956, Sakurai 1960, Raizer 1964, Grover and Hardy 1966, Hayes 1968) notoriously provide excellent approximations to the generic solutions, for a reason most clearly discussed by Raizer (1964): behind the shock, a sonic point is formed, which prevents causal connection between the region immediately behind the shock with the more distant regions, which are sensitive to boundary conditions. Disconnection from initial conditions limits the number of parameters on which the solution may depend, while still allowing the existence of a self--similar solution. Because of the independence from initial conditions, it is expected that, shortly after their formation, all solutions with accelerating shocks will quickly approach the self--similar solution.  Because of the possible relevance of these solutions to the evolution of Gamma Ray Bursts, there has been recently a strong surge of interest in their properties, especially in the relativistic regime, which, at least asymptotically, should be most appropriate to GRBs. A highly relativistic, self--similar solution has been presented by Perna and Vietri (2002) in planar geometry, for which accelerating, self--similar solutions appear for exponentially decreasing density stratifications outside the star. In spherical geometry, it is instead possible to find a richer spectrum of self--similar solutions, for basically any type of power--law density stratification (Best and Sari 2000, Pan and Sari 2005, Sari 2005). These solutions are often referred to as Type II solutions (Barenblatt 1996), meaning that the behaviour of the full solution cannot be determined from dimensional analysis alone, like in the well--known Sedov solution, but is fixed instead by the conditions of regularity at some particularly difficult point, like the sonic point.  The stability of relativistic hydrodynamic solutions with shocks is rarely considered in the literature. Gruzinov (2001) considered the linear stability of the spherically symmetric, relativistic explosions of Blandford and McKee (1976), but nothing is currently known about the stability of accelerating solutions, in the relativistic regime. In the Newtonian regime, the self--similar solutions are known to be unstable, from both a linear analysis (Chevalier 1990), and a numerical, non--linear one (Luo and Chevalier 1994). The relativistic solutions differ from this analysis in two distinct respects: on the one hand, they include relativistic effects, on the other one, since the shock is supposed to be highly relativistic, they are obliged to assume the validity of the ultra--relativistic equation of state, $p = e/3$. Both of these circumstances lead to a greater susceptibility to corrugational instabilities. In fact, softening the equation of state leads to corrugational instability even for shocks propagating into a {\\it constant density} environment: Ryu and Vishniac (1987) found a range of instabilities provided the adiabatic index is $\\gamma \\lesssim 1.2$. Also, assuming an adiabatic index close to $\\gamma =5/3$ leads to both a narrower instability range in the parameter $k$, and in smaller growth rates than determined by Chevalier (1990) for $\\gamma = 4/3$ in the Newtonian, accelerating self--similar flow (Vietri, unpublished).  On the other hand, relativistic effects tend to make pressure waves slow when compared to bulk motion, and thus to decouple nearby layer of fluid. Thus, if a purely kinematic difference exists in the motion of adjacent layers, it is more likely to grow because of the lack of restoring effects.  {}  The plan of this paper is as follows. In the next section, we shall briefly summarize the properties of the zero--th order solution of Perna and Vietri (2002). In the following section we shall perform a linear stability analysis, showing that accelerating, self--similar shocks are indeed unstable. In the last section, we shall discuss our results. ",
        "conclusions": "To summarize, we stress here that the hydrodynamical quantities grow as follows. Let us call $u_x$ the $x$-component of four-velocity. For $\\delta\\!\\Gamma \\neq 0$: \\begin{eqnarray}{} \\delta u_x\\approx\\delta\\!\\gamma &\\propto& \\Gamma^{1} \\\\ \\delta e &\\propto& \\Gamma^{2+\\alpha} \\\\ \\delta n &\\propto& \\Gamma^{2+\\alpha} \\\\ \\delta\\!X &\\propto& \\Gamma^{-2}\\;, \\end{eqnarray} while for the other, stronger mode, $\\delta\\!X_0 \\neq 0$, we find: \\begin{eqnarray}{} \\delta u_x\\approx\\delta\\!\\gamma &\\propto& \\Gamma^{3} \\\\ \\delta e &\\propto& \\Gamma^{4+\\alpha} \\\\ \\delta n &\\propto& \\Gamma^{4+\\alpha} \\\\ \\delta\\!X &\\propto& \\Gamma^0\\;. \\end{eqnarray} Obviously, the most surprising result is that, in this last mode, the concentrations of energy and baryon number $\\delta e$ and $\\delta n$ grow very fast, as $\\propto \\Gamma^{4+\\alpha}$, or, stated otherwise, \\begin{equation} \\frac{\\delta e}{e} \\propto \\frac{\\delta n}{n} \\propto \\Gamma^2\\;. \\end{equation} Obviously, these concentrations will quickly become nonlinear, and their subsequent fate can only be ascertained through a numerical analysis. The same conclusion is reached when one considers the quick growth of $\\delta u_x \\propto \\Gamma^{3}$.  The linear growth rates are so high that we can expect basically all small perturbations to become nonlinear during the acceleration of the ejecta of a Hypernova, when $\\Gamma$ changes from $\\Gamma_{\\textrm{i}}-1 \\ll 1$ to $\\Gamma_{\\textrm{f}} \\approx 100$.  We should remark that the independence of these results from $k$ is a peculiarity of the relativistic solution which does not exist in the Newtonian counterpart of the problem (Chevalier 1990). This independence is derived in an explicit way in the Appendix.    Our conclusions differ from those of Wang, Loeb and Waxman (2003), who were unable to find self--similar perturbations to a similar problem in spherical geometry, reporting only a marginal growth for perturbations in the linear regime for an accelerating shock in spherical symmetry. Though not exactly identical, the growth rates should however coincide in the limit of large wavenumbers, where curvature effects become negligible. In support of our work we can make two different points. On the one hand, there is the the coincidence of the growth rates when these are computed in two, very different ways: the simple, physical one (presented above) and the more detailed, mathematical one to be presented in the Appendix.  On the other hand, there is the similarity between the stability analysis in the Newtonian (Chevalier 1990) and relativistic regime (this paper): like us, Chevalier found two independent modes, for sufficiently small wavenumbers (the limit in which we can compare directly the two sets of computations), with very simple indices, $s=0,1$ (but please notice that they are defined in a different way from ours). Of these, the mode $s=0$ is of course only marginally unstable, while the other one, corresponding to a time growth $\\propto |t|^{-1}$ for $t\\rightarrow 0$, is identified by Chevalier (1990) as the physically relevant one. This is similar to our result of two independent modes, with distinct indices $\\delta\\!e/e \\propto \\Gamma^s$, $s=0,2$, the one with $s=2$ resulting in the most severe (and thus physically relevant) instability.  In a future paper, we will investigate the nonlinear development of the corrugational instability discussed here.  Thanks are due to the referee, Dr. R. Sari, for helpful comments."
    },
    "0609/astro-ph0609242_arXiv.txt": {
        "abstract": "We describe a model that accounts for the complex morphology of spiral density waves raised in Saturn's rings by the co-orbital satellites, Janus and Epimetheus.  Our model may be corroborated by future Cassini observations of these time-variable wave patterns. ",
        "introduction": "Owing to typical astronomical timescales, a galaxy's spiral arms are often considered as a fixed pattern.  So, too, for the numerous, tightly wound spiral waves detected in Saturn's rings.  In fact, both systems are dynamically active, with waves traveling away from resonant sites.  This is manifest only in Saturn's case, where a pair of moons -- Janus and Epimetheus -- occupy nearly identical orbits that are interchanged every 4 years, causing the resonance locations in the rings to skip back and forth by tens of km.  Since spiral density waves are initiated in Saturn's rings at locations where ring particle orbits are in a Lindblad resonance with a perturbing moon, the starting points of waves jump as well, allowing wave trains to interfere in complex ways \\citep{PorcoSSR04}.  High-resolution images of the rings were obtained by the Cassini spacecraft's Imaging Science Subsystem (ISS) on 2004~July~1 and on 2005~May~20/21.  The calibration and image processing of these data, resulting in a series of brightness scans with orbital radius, along with a catalog of important satellite gravitational resonances falling within the rings, are presented by \\citet[hereafter \\Pone]{PorcoSOIScans06}.  Further analysis of these data, employing techniques derived from the wavelet transform, is presented by \\citet[hereafter \\Ptwo]{soirings}.  Examination of the Cassini images to date shows that density waves raised by the co-orbitals are both unusual and variable in their morphology.  In this paper, we describe a model that accounts for much of the observed structure.  We proceed to use this model to predict the future morphology of selected waves at the times and locations of planned Cassini observations, which we expect will test our predictions.  \\begin{figure}[!t] \\begin{center} \\includegraphics[width=8cm,keepaspectratio=true]{jeorbits8.eps} \\caption{The orbits of Janus and Epimetheus (with epicycles removed) in a frame rotating with their mass-averaged mean motion.  The radial separation of the librations is exaggerated by a factor of 500.  Numbered points mark the moons' positions on 1)~2004~July~1, 2)~2005~May~21, and 3)~2006~September~9.  Arrow lengths indicate motion accomplished in 100 days.  \\label{JanEpiOrb}} \\end{center} \\end{figure} ",
        "conclusions": "Saturn's rings are a nearby and accessible natural laboratory \\citep{BC06}, in which we can observe phenomena of broader interest, such as how a disk responds to variable forcing.  Our model is a first step towards a new understanding of the observed complex and time-variable waveforms.  We will revisit this topic in more detail, after the accuracy of our predictions (at least for data that will be obtained this year) becomes apparent."
    },
    "0609/astro-ph0609418_arXiv.txt": {
        "abstract": "We present a fast and efficient hybrid algorithm for selecting exoplanetary candidates from wide-field transit surveys. Our method is based on the widely-used {\\sc SysRem} and Box Least-Squares (BLS) algorithms. Patterns of systematic error that are common to all stars on the frame are mapped and eliminated using the {\\sc SysRem} algorithm. The remaining systematic errors caused by spatially localised flat-fielding and other errors are quantified using a boxcar-smoothing method. We show that the dimensions of the search-parameter space can be reduced greatly by carrying out an initial BLS search on a coarse grid of reduced dimensions, followed by Newton-Raphson refinement of the transit parameters in the vicinity of the most significant solutions. We illustrate the method's operation by applying it to data from one field of the SuperWASP survey, comprising 2300 observations of 7840 stars brighter than $V=13.0$. We identify 11 likely transit candidates. We reject stars that exhibit significant ellipsoidal variations indicative of a stellar-mass companion. We use colours and proper motions from the 2MASS and USNO-B1.0 surveys to estimate the stellar parameters and the companion radius. We find that two stars showing unambiguous transit signals  pass all these tests, and so qualify for detailed high-resolution spectroscopic follow-up. ",
        "introduction": "Among the 194 extra-solar planets discovered in the last decade, the ``hot Jupiters\" currently present the greatest challenges to understanding and the greatest observational rewards. Efforts to explain their origin have transformed theories of planetary-system formation. Do these planets form via gravitational instability in cold discs, or must they form by core accretion beyond the ice boundary, where  ice mantles on dust grains allow rapid agglomeration of a massive core of heavy elements? How rapidly  do they migrate inwards through a massive protoplanetary disc, and what mechanism halts the migration?  The subset of these planets that transit their parent stars are of key importance in addressing these questions, because they are the only planets whose radii and masses can be determined directly. The first such discovery \\citep{charbonneau2000} confirmed the gas-giant nature of the hot Jupiters; indeed the inflated radius of HD 209458b continues to challenge our understanding of exoplanetary interior structure. Subsequent discoveries have presented further surprises. The apparent correlation between planet mass and minimum survivable orbital separation among the known transiting planets offers important clues to the ``stopping mechanism\" for inward orbital migration of newly-formed giant planets in protoplanetary discs \\citep{mazeh2005}. The high core mass of the recently-discovered Saturn-mass planet orbiting HD 149026 (Sato et al. 2005)\\nocite{sato2005} presents difficulties for models of planet formation via gravitational instability. {\\em Spitzer} observations of the secondary eclipses of TReS-1b, HD 209458b and HD189733b have provided the first estimates of the temperatures in these planets' cloud decks (\\citealt{charbonneau2005}; \\citealt{deming2005}; \\citealt{deming2006}).  To date, however, only ten such planets have been discovered. Three of the ten have been found via targeted radial-velocity searches, the other six through large-scale photometric surveys for transit events. The primary goal of the SuperWASP project (Pollacco et al. 2006)\\nocite{pollacco2006} is to discover bright, new transiting exoplanets in sufficiently large numbers that we can place studies of their mass-radius relation, and the suspected relationship between minimum orbital distance and planet mass, on a secure statistical footing. SuperWASP's  ``shallow-but-wide\" approach to transit hunting is designed to find planets that are not only sufficiently bright ($10 < V < 13.0$) for high-precision radial-velocity follow-up studies to be feasible on telescopes of modest aperture, but also for detailed follow-up studies such as transmission spectroscopy during transits, and Spitzer secondary-eclipse observations. Only five of the ten presently-known transiting planets orbit stars brighter than $V=14$ and so have such strong follow-up potential.  Here we present a methodology used in the search for exoplanetary transit candidates in data from the first year of SuperWASP operation. We employ the {\\sc SysRem} algorithm of \\citet{tamuz2005} to identify and remove patterns of correlated systematic error from the stellar light-curves. We present a refined version of the  Box Least-Squares (BLS) algorithm of \\citet{kovacs2002}, which permits a fast grid search and efficient refinement of the most promising solutions without binning the data. Using simulated transits injected into real SuperWASP data we develop a filtering strategy to optimise and quantify the recovery rate and false-alarm probability as functions of stellar magnitude and transit depth. We develop simple plausibility tests for transit candidates, using the transit duration and depth to estimate the mass of the parent star and the radius of the planet. We mine publicly-available catalogues to obtain the $V-K$ colours, proper motions and other properties of the host stars. Combined with the transit durations and depths, these give improved physical parameters of each system and help us to  identify the most promising candidates for spectroscopic follow-up observations. ",
        "conclusions": "In this paper we have described the methodology that we have adopted for searching for transits in the large body of data produced by the SuperWASP camera array on La Palma during its first few months of operation, and applied it to the 7840 stars brighter than $V=13.0$ in a survey field centred at RA  01h 43m, Dec +31$^\\circ$ 26'.  We find that an initial search using the BLS method on a coarse grid of transit epochs and orbital frequencies is sufficient for us to eliminate more than 95 percent of the stars searched. This allows us to perform a finer grid search on only the remaining 198 stars in the field under consideration. We have adapted the analytic Newton-Raphson method of \\citet{protopapas2005} to refine the orbital solutions around periodogram peaks found with the BLS method. This method quickly yields the depth and duration of the transits, and the frequency and phase of the photometric orbit, while keeping the dimensionality of the search grid (and hence the processing time required) as low as possible.  We use the ellipsoidal-variation methodology of \\citet{drake2003} and \\citet{sirko2003} in our light-curve modelling to eliminate probable stellar binaries. We use the publicly-available 2MASS and USNO-B1.0 catalogues to obtain colours and proper motions for candidates, and to estimate the stellar and planetary radii using methods similar to those of \\citet{seager2003} and \\citet{tingley2005}. These methods confirm that two of our most significant transit detections in this field, and one more marginal one, have transit properties fully consistent with those expected for planets with radii comparable to or somewhat smaller than Jupiter.  The survey field chosen to illustrate the method is only one of more than 100 such regions surveyed during the course of 2004. Candidates from the other fields will be presented and discussed in subsequent papers. Together with the three candidates presented here, all likely planetary-transit candidates will be subjected to high-resolution spectroscopic  follow-up in the latter half of 2006, using the methodology employed successfully by \\citet{bouchy2005} for OGLE-III follow-up."
    },
    "0609/astro-ph0609132_arXiv.txt": {
        "abstract": "\\ion{Cl}{1} is the atomic species most directly coupled to molecular hydrogen due to its chemistry. Its weakest lines are thereby probably the best tracer of optically thick H$_2$ components in diffuse clouds. We report on the empirical determination of the oscillator strengths for four \\ion{Cl}{1} absorption lines predicted to be weak and often detected toward moderately reddened sight lines observed with the {\\it Far Ultraviolet Spectroscopic Explorer} ({\\it FUSE}). We compared our oscillator strength estimates with the oscillator strength calculations listed in Morton (2003). We find that our empirical oscillator strength values for the \\ion{Cl}{1} 1004, 1079, 1090 and 1094 \\AA\\ lines differ from the theoretical predictions by factors of $\\sim$3.1, 1.2, 2.4 and 0.42, respectively. We briefly discuss the value of \\ion{Cl}{1} as tracer of molecular gas for our star sample. ",
        "introduction": "Due to its unique chemistry, chlorine plays an important role in characterizing the neutral gas components (Jura 1974; Dalgarno et al. 1974) in the interstellar medium (ISM). In diffuse clouds, the detection of chlorine in its various ionization states allows us to disentangle the diffuse molecular gas components from the neutral atomic gas components (Jura \\& York 1978) along sight lines exhibiting complex structures that are often unresolved with the current instrumentation (e.g., Rachford et al. 2002). Most importantly, \\ion{Cl}{1} offers an alternate and more reliable way of tracing cold H$_2$ gas components in diffuse clouds since, unlike CO, the chlorine abundance remains unaffected by molecular processes such as chemical fractionation and photodissociation (Sonnentrucker et al. 2002).  \\ion{Cl}{2} exhibits transitions solely below 1080 \\AA, all with laboratory-determined $f$-values (Schectman et al. 1993, 2005). \\ion{Cl}{1}, on the other hand, is predicted to have over 50 transitions between $\\lambda$920 and $\\lambda$1390 \\AA\\  (Morton 2003 and references therein). To date, only three lines, $\\lambda$$\\lambda$1088, 1097 and 1347 \\AA, possess experimentally-determined $f$-values. Because of its strength the \\ion{Cl}{1} 1347 \\AA\\ line is almost always optically thick and its use requires careful treatment of saturation effects (Jura \\& York 1978; Jenkins et al. 1986; Harris \\& Bromage 1984). The $\\lambda$1088 \\AA\\ line often suffers from severe blending with an adjacent CO band (Federman 1986) while the $\\lambda$1097 \\AA\\ line, due to its intrinsic weakness, requires significant reddening and high S/N ratio to be detectable along a given sight line. Consequently, accurate determination of the \\ion{Cl}{1} column density and the chlorine cosmic abundance remained difficult with prior space-based missions such as {\\it Copernicus} and {\\it IUE} (Keenan \\& Dufton 1990).  The {\\it FUSE} wavelength range and especially its greater sensitivity -- compared to previous missions-- allowed us to revisit the role that chlorine can play in increasing our understanding of the structure in the diffuse atomic and diffuse molecular ISM. Since a reliable atomic database is necessary for such an endeavor, we proposed to empirically determine the oscillator strength ($f$-value) for 4 \\ion{Cl}{1} lines particularly well suited for our study. The \\ion{Cl}{1} 1004, 1079, 1090, 1094 \\AA\\ lines are predicted to be weak, are often detected toward low-to-moderately reddened stars and do not significantly suffer from blends with other species. We present the results of our $f$-value investigation for these 4 \\ion{Cl}{1} lines toward the moderately reddened stars, HD91597, HD220057 and HD108 and three additional stars from the {\\it FUSE} archive. Section 2 summarizes our observations and data reduction technique. Section 3 describes our analysis, compares our empirical $f$-value estimates to the theoretical calculations listed in Morton (2003), and discusses the value of \\ion{Cl}{1} as a tracer of H$_2$. ",
        "conclusions": ""
    },
    "0609/astro-ph0609304_arXiv.txt": {
        "abstract": "We present the first high spatial resolution X-ray study of the massive star forming region NGC 6357, obtained in a 38 ks {\\it Chandra}/ACIS observation. Inside the brightest constituent of this large HII region complex is the massive open cluster Pismis 24. It contains two of the brightest and bluest stars known, yet remains poorly studied; only a handful of optically bright stellar members have been identified. We investigate the cluster extent and Initial Mass Function and detect $\\sim$800 X-ray sources with a limiting sensitivity of $\\sim 10^{30}$ ergs s$^{-1}$; this provides the first reliable probe of the rich intermediate-mass and low-mass population of this massive cluster, increasing the number of known members from optical study by a factor of $\\sim 50$. The high luminosity end ($\\log L_h$[2-8 keV]$\\ge 30.3$ ergs s$^{-1}$) of the observed X-ray luminosity function in NGC 6357 is clearly consistent with a power law relation as seen in the Orion Nebula Cluster and Cepheus B, yielding the first estimate of NGC 6357's total cluster population, a few times the known Orion population.  We investigate the structure of the cluster, finding small-scale substructures superposed on a spherical cluster with 6 pc extent, and discuss its relationship to the nebular morphology. The long-standing $L_X - 10^{-7}~L_{bol}$ correlation for O stars is confirmed.  Twenty-four candidate O stars and one possible new obscured massive YSO or Wolf-Rayet star are presented. Many cluster members are estimated to be intermediate-mass stars from available infrared photometry (assuming an age of $\\sim 1$ Myr), but only a few exhibit $K$-band excess.  We report the first detection of X-ray emission from an Evaporating Gaseous Globule at the tip of a molecular pillar; this source is likely a B0-B2 protostar. ",
        "introduction": "\\object{NGC 6357} (= W 22 = RCW 131 = S 11) is a large \\hii\\ region complex in the southern sky excited by the OB association Pismis 24, at a distance of $\\sim 2$ kpc \\citep{Neckel78,Felli90,Massi97,Bohigas04}. This young stellar cluster and its surrounding gaseous environment are relatively poorly studied; only a handful of stellar members have been identified and characterized, and the molecular cloud has not been well-mapped.  In optical and near-infrared images, the exciting cluster is not prominent against the Galactic field population except for a few bright members (Figure \\ref{fig:dss2mass}).  Yet, two of these stars are remarkably luminous and massive with spectral types O3 If (M$\\sim 200-300$ M$_\\odot$) and O3 III \\citep[M$\\sim 100$ M$_\\odot$;][]{Massey01, Walborn02}.  Such massive stars are usually found in only the richest Galactic clusters and the giant starburst regions such as Carina, 30 Doradus, and NGC 3603.  This raises the question is Pismis 24 much richer than it appears, or is it a poor cluster with a top-heavy Initial Mass Function (IMF)? This would be highly unusual, as top-heavy IMFs are seen only in extreme starburst environments such as the super-star clusters of M 82 \\citep{Elmegreen05}.  We investigate this question by measuring, for the first time, the low-mass pre-main sequence (pre-MS) population of the Pismis 24 cluster and its environs.  This is achieved with a sensitive X-ray observation with the {\\it Chandra} X-ray Observatory.  {\\it Chandra} observations can penetrate heavy absorption up to $A_V \\simeq 500$ \\citep{Grosso05} and resolve point sources with sub-arcsecond resolution.  In most cases, the X-ray detected stars appear on existing near-infrared (NIR) images; the advantage is that the strong X-ray emission of pre-MS stars \\citep[elevated $10^1-10^4$ above MS stars][]{Preibisch05a} traces magnetic activity rather than photospheric or circumstellar disk emission, therefore effectively discriminating cluster members from older Galactic field stars.  We estimate that the cluster has $\\simeq 10,000$ stars, $\\simeq 5$ times richer than the known Orion Nebula Cluster. The presence of two O3 stars is compatible with a cluster of this population.  The known massive stars and the large number of low mass X-ray stars suggest that Pismis 24 is a very rich cluster with a standard IMF.  Our study also gives considerable new information about the cluster and its X-ray properties.  We provide a census of $\\sim$750 cluster members with sub-arcsecond positions, 613 of which have NIR counterparts.  Although not complete, the survey includes pre-MS stars with masses extending down to $\\sim 0.3 M_{\\odot}$; most do not have circumstellar disks revealed by $K$-band excess.  We study the spatial distribution of the cluster members and identify a possible subcluster. We examine the ringlike morphology of the NGC 6357 nebula and investigate the possible feedback from the massive stars to their environs. We measure X-ray properties of a dozen known O stars, and list over 20 new candidate O stars.  A faint X-ray source is found associated with a protostar in an Evaporating Gaseous Globule (EGG) for the first time.  A few pre-MS stars are seen with very powerful X-ray flares with peak $L_x \\sim 10^{32}$ ergs s$^{-1}$.  Section \\ref{overview.sec} reviews past study of the Pismis 24 cluster and NGC 6357 region.  The {\\it Chandra} observation and its analysis are described in \\S \\ref{obs.sec}, and a NIR observation using the Simultaneous 3 color InfraRed Imager for Unbiased Surveys (SIRIUS) camera on the Infrared Survey Facility (IRSF) telescope is summarized in \\S \\ref{sirius.sec}. The association between X-ray sources and optical and near-infrared (ONIR) stars, as well as the infrared colors of X-ray counterparts are presented in \\S \\ref{counterparts.sec}.  The cluster population inferred from the X-ray emitting stars, their X-ray luminosity function (XLF), the spatial distribution of the cluster members, and the morphology of the nebula are discussed in \\S \\ref{XLF.sec}.  X-ray properties of O stars and intermediate- to low-mass pre-MS stars are presented in \\S \\ref{ximf.sec}. The possibility of diffuse X-ray emission is examined in \\S \\ref{diffuse.sec}. A brief summary of our findings is given in \\S \\ref{summary.sec}. This study is part of a series of X-ray investigations of high-mass star formation regions that includes the Rosette Nebula (Townsley et al.\\ 2003; Wang et al., in preparation), M 17 (Townsley et al.\\ 2003; Broos et al., in preparation), the Orion Nebula (Feigelson et al.\\ 2003; Getman et al.\\ 2005 and associated articles), 30 Doradus (Townsley et al.\\ 2006a, 2006b), and Cepheus B/OB3b (Getman et al.\\ 2006). ",
        "conclusions": "}  This 38 ks ACIS observation of the NGC 6357 field has provided the first high spatial resolution X-ray image of this massive star-forming complex. We summarize the main results of our study as follows:  1. We detect 779 X-ray sources with a limiting X-ray sensitivity of $\\sim 10^{30}$ ergs s$^{-1}$. There are 445 ACIS sources matched to 2MASS point sources with $\\le 2^{\\prime\\prime}$ separation. Further visual inspection of an archival HST image and SIRIUS $JHK$ images yield a total of 616 ONIR counterparts. We estimate the extragalactic contaminants via a careful simulation of detection efficiency of the background AGN population, and conclude that no more than 4\\% of our detections are AGNs. Similarly we adopt the Besan\\c{c}on stellar population sythesis model of the Galactic disk and suggest 11 stars as foreground objects based on their X-ray softness and NIR color. Excluding $\\sim 20$ X-ray sources without ONIR identifications that are possible AGNs, the rest of the $\\sim 140$ X-ray sources without ONIR counterparts are likely new Pismis 24 members that are clustered around the massive stars and/or deeply embedded in the cloud. The X-ray detected population provides the first deep probe of the rich population of this massive cluster, increasing the number of known members from optical study by a factor of $\\sim 50$.  2. On the NIR color-color diagram (Figure~\\ref{fig:ccd}), it is clear that ACIS detects more Class III sources than Class II/I sources. In the color-magnitude plot (Figure~\\ref{fig:cmd}), the locations of most of the ACIS sources are consistent with being high mass stars and pre-MS intermediate mass stars ($M \\ge 2M_{\\odot}$), subjected to a visual extinction of $5<A_V<15$ at 2.5 kpc. The new high mass stars, if spectroscopically confirmed, double the known high mass population of the cluster. Dozens of intermediate mass stars are newly recognized members of the cluster. The ratio between Class III and Class II objects in the intermediate- to high-mass range is $\\sim 15:1$. This is an important new sample for stellar and disk studies of intermediate mass pre-MS stars. The brightness of several K-excess sources makes them extremely favorable for ground based follow-up. We find a very luminous X-ray emitting infrared source subjected to significant absorption, which may be a new obscured massive YSO or a Wolf-Rayet star.  3. We determine the spatial distribution of over 700 new members to sub-arcsecond precision for the first time. The cluster is roughly spherically symmetric, with the highest stellar density around the massive O3 stars. A density enhancement centered at the tip of the molecular filament deserves further investigation (Figure~\\ref{fig:stellar_density}). A radial profile of the stellar surface density shows that the physical size of the cluster spans over 2 pc in radius and extends beyond 6 pc (Figure~\\ref{fig:offaxis_hist}).  4. XLFs are derived for lightly obscured and heavily obscured stellar populations in NGC 6357 (Figure~\\ref{fig:XLF}). The high X-ray luminosity end ($\\log L_X \\ge 30.3$ ergs s$^{-1}$) of the NGC 6357 XLF is clearly consistent with a power law relation as seen in COUP and Cep B, allowing a scaling comparison with the COUP XLF. The first estimate of the total cluster population of Pismis 24 is a few times the known Orion population. The presence of two O3 stars is consistent with the standard IMF.  5. We detect all 10 known OB stars with spectral type earlier than B1 in the Pismis 24 cluster, including the WC7+O7 binary WR 93. Their $L_X/L_{bol}$ ratios are consistent with the canonical $10^{-7}$ value with less than an order of magnitude scatter. The O3 stars Pis 24-1 (O3.5If) and Pis 24-17 (O3.5IIIf) exhibit X-ray luminosities $L_X\\sim 10^{33}$ ergs s$^{-1}$. The X-ray spectrum of Pis 24-1 shows a soft emission component ($kT\\sim 0.5$ keV) and a harder component ($kT\\sim 1.7$ keV), while Pis 24-17 exhibits an unusally strong \\sixiii\\ line with a soft $0.7$ keV thermal plasma. The WC7+O7 binary WR93 is the brightest X-ray source in the field ($L_X\\simeq 2 \\times 10^{33}$). Its hard spectrum suggests that the X-ray emission is generated from colliding winds. The lightcurves of these massive stars are constant through the observation.  We assembled a list of candidate O stars from sources with high X-ray and infrared luminosity to facilitate future spectroscopy follow-up. Four candidate HAeBe stars are also identified.  6. We report the detection of X-ray emission from an EGG at the tip of a molecular pillar, which was previously found to be an IR source and the peak of the radio continuum. It is in the early evolutionary of stages of star formation. The spectral type estimated from its SED is B0--B2. The non-detection of EGGs in the Eagle Nebula may be the result of lower stellar mass and earlier evolutionary phases. We assemble a sample of 16 deeply embedded objects.  7. Several flaring sources exhibit luminosities above $\\log L_{t,c} \\ge 32.0$ ergs s$^{-1}$, which is exceptionally high for pre-MS stars. One powerful flare seen in ACIS \\#672 (CXOU J172457.87-341203.9) is comparable to the strongest X-ray flares known in pre-MS stars ($L_{t,c}\\sim 32.5$ ergs s$^{-1}$). Its lightcurve rises 15 times above the quiescent level during the flare, with a slow rise phase that is distinct from the fast rise flares seen pre-MS stars.  8. Soft unresolved X-ray emission in the NGC 6357 region is present. However, the luminosity derived from spectral fitting is consistent with the estimated level of integrated emission from the unresolved pre-MS stars.  We sincerely thank the anonymous referee for his/her careful review and many constructive suggestions to improve the clarity of our paper. We are grateful to Steve Strom for an enlightening discussion on disk fractions among massive stars.  We thank Takahiro Nagayama, Takahiro Kawadu, and the IRSF/SIRIUS team for taking the SIRIUS images of NGC 6357.  Support for this work was provided to Gordon Garmire, the ACIS Principal Investigator, by the National Aeronautics and Space Administration (NASA) through NASA Contract NAS8-38252 and Chandra Contract SV4-74018 issued by the Chandra X-ray Observatory Center, which is operated by the Smithsonian Astrophysical Observatory for and on behalf of NASA under contract NAS8-03060. 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 NASA and the National Science Foundation.  This research has made use of the SIMBAD database and the VizieR catalogue access tool, operated at CDS, Strasbourg, France. This research made use of data products from the Midcourse Space Experiment.  Processing of the data was funded by the Ballistic Missile Defense Organization with additional support from NASA Office of Space Science.  This research has also made use of the NASA/ IPAC Infrared Science Archive, which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with NASA.  {\\it Facility}: \\facility{CXO (ACIS)}"
    },
    "0609/astro-ph0609020.txt": {
        "abstract": "We present a new method for studying the star formation history of late-type, cluster galaxies undergoing gas starvation or a ram-pressure stripping event by combining bidimensional multifrequency observations with multi-zones models of galactic chemical and spectrophotometric evolution. This method is applied to the Virgo cluster anemic galaxy NGC~4569. We extract radial profiles from recently obtained UV GALEX images at 1530 and 2310 \\AA, from visible and near-IR narrow (H$\\alpha$) and broad band images at different wavelengths ($u$, B, $g$, V, $r, i, z$, J, H, K), from Spitzer IRAC and MIPS images and from atomic and molecular gas maps. The model in the absence of interaction (characterized by its rotation velocity and spin parameter) is constrained by the unperturbed H band light profile and by the H$\\alpha$ rotation curve. We can reconstruct the observed total-gas radial-density profile and the light surface-brightness profiles at all wavelengths in a ram-pressure stripping scenario by making simple assumptions about the gas removal process and the orbit of NGC~4569 inside the cluster. The observed profiles cannot be reproduced by simply stopping gas infall, thus mimicing starvation. Gas removal is required, which is more efficient in the outer disk, inducing a radial quenching in the star formation activity, as observed and reproduced by the model. This observational result, consistent with theoretical predictions that a galaxy-cluster IGM interaction is able to modify structural disk parameters without gravitational perturbations, is discussed in the framework of the origin of lenticulars in clusters. ",
        "introduction": "\\setcounter{footnote}{0} Perturbations induced by the harsh cluster environment makes the evolution of cluster galaxies different than that of their counterparts in the field. Besides the well known morphology segregation effect (Dressler 1980; Whitmore et al. 1993), there is also evidence clearly indicating that the cluster late-type galaxy population is systematically different from the field: cluster spirals are generally depleted in their total HI content and are less active in forming stars than their isolated counterparts. This is most recently and extensively reviewed in Boselli \\& Gavazzi (2006). The nature of the perturbing phenomenon has not yet been unambiguously identified. Several physical processes have been proposed to explain gas removal in clusters. Some of them are related to the dynamical interactions of cluster galaxies with the hot ($T_{IGM}$ $\\sim$ 10$^7$ K) and dense ($\\rho_{IGM}$ $\\sim$ 10$^{-3}$ atom cm$^{-3}$) intergalactic medium (IGM) (ram-pressure, Gunn \\& Gott 1972; viscous stripping, Nulsen 1982; thermal evaporation, Cowie \\& Songaila 1977). Others are due to the gravitational interactions with nearby companions (Merritt 1983) or with the wider potential of the cluster (Byrd \\& Valtonen 1990). %The lack of a sufficient gas reservoir needed to %feed star formation is then at the origin of the decrease of the %activity observed in cluster galaxies. %\\\\ %All these processes, however, {\\textbf are active only nearby the core of the %clusters, where the IGM gas and galaxy densities, the gas temperature and the %cluster potential are at their maximum. %NOT TRUE FOR LOW VELOCITY INTERACTIONS! %YOU MUST SPECIFY IS LOW OR HIGH VELOCITY INTERACTIONS OR BOTH.} %Recent, complete spectroscopic surveys of the nearby universe such as the SDSS or the %2dF, able to trace the radial variation of the star formation activity %up to large cluster-centric distances, have shown that the activity of %late-type galaxies begins to decrease at $\\sim$ 1-2 virial radii from %the cluster center (Gomez et al. 2003; Lewis et al. 2002, Tanaka et %al. 2004, Nichol 2004), scale-lengths comparable to those observed in %the HI (Gavazzi et al. 2005,2006a). Since these distances (1-2 virial %radii) are significantly larger (a factor of 5-10) than those where %galaxy-IGM or galaxy-galaxy interactions {\\textbf WHAT KIND OF TIDAL INTERACTION} %are expected to be at place ($\\sim$ inside a %core radius), other processes more active at the periphery of clusters %such as galaxy harassment (Moore et al. 1996) {\\textbf I'M NOT SURE YOU ARE RIGHT SAYING %THAT HARASSMENT IS MORE ACTIVE AT THE CLUSTER PERIFERY}, starvation (Balogh et %al. 2000) or pre-processing (Fujita 2004) have been proposed to %explain the radial decrease of the star formation activity with the %cluster-centric distance. %\\\\ The detailed analysis of several well known resolved galaxies in the Virgo and in the Coma clusters, generally favor a galaxy-IGM interaction scenario (Vollmer et al. 1999, 2000, 2001a, 2004b; Kenney et al. 2004; Yoshida et al. 2004) in these instance. Using complete, multifrequency datasets combined with model predictions several studies have unambigously shown that ram-pressure stripping occurs even at the periphery of clusters (e.g. CGCG 97-73 and 97-79 in A1367, Gavazzi et al. 2001), and that galaxies recently ($\\leq$ 500 Myr) stripped by ram-pressure populate nearby clusters well outside the cluster core (Boselli \\& Gavazzi 2006). \\\\ %This apparent discrepancy {\\textbf WHICH DISCREPANCY, HAI APPENA DETTO CHE GAV\\&BOS HANNO %DIMOSTRATO CHE LA RAM PRESSURE E' ATTIVA ANCHE FUORI DAL CLUSTER CORE} %poses major limits to the understanding of the evolution of galaxies in %different environments, i.e. Ram-pressure stripping has also been invoked to explain the origin of the lenticular galaxy population inhabiting rich clusters. One early suggestion was that, as a consequence of gas removal through ram-pressure stripping, cluster galaxies first become ``anemic'' and then passively evolved into lenticulars (van den Bergh 1976). Statistical considerations, as well as structural, kinematical and spectroscopic properties of lenticulars, however, do not appear to confirm this simple scenario. As discussed in Dressler (1980, 2004) the morphological type variation in clusters is too weak a function of the local galaxy density. Furthermore, the disks of lenticulars have, on average, surface brightnesses and bulge-to-disk ratios that are significantly higher than early-type spirals (Dressler 1980; Christlein \\& Zabuldoff 2004; Boselli \\& Gavazzi 2006) making the formation of lenticulars through gas sweeping in spirals very unlikely.  The larger scatter and a small zero-point offset in the Tully-Fisher relation observed in the Virgo and Coma cluster lenticulars also indicate that S0s can hardly be formed by simple gas removal from healthy spirals (Hinz et al. 2003).  Furthermore cluster lenticulars generally have spectroscopic signatures of recent, post-starburst activity (Poggianti et al. 2004) difficult to trigger in a simple ram-pressure stripping scenario.\\\\  The origin of lenticulars is still an open problem and limit our understanding of the evolution of galaxies in different environments. This riddle can be solved only by first realistically quantifying the physical effects of ram-pressure stripping by combining model predictions with multifrequency observations of representative samples. %This promising exercice has been successfully %accomplished by Vollmer and collaborators %in a recent series of papers joining kinematical observations with updated dynamical models %(Vollmer et al. 1999, 2000, 2004a,b,c, 2005, Vollmer 2003). %These works have clearly shown the potentiality in constraining not only the nature of the %acting phenomenon, but also the time scale since the interaction started. %As stated in Boselli \\& Gavazzi (2006), times scales are in fact strong constraints %for discriminating the different processes, being short in galaxy-IGM %interactions ($\\sim$ 5 10$^8$ yr) and relatively long in those processes %acting at the periphery of clusters such as harassment, starvation or pre-processing %(some 10 $^9$ yr). %\\\\ %Despite the on-going physical processes, it is clear that %the fuel supply needed to feed star formation is more efficiently %removed where the host-galaxy potential well is weakest, i.e., in the %outer disk. Given the strong relation between the gas surface density %and the star formation activity in spiral disks, commonly known as the %Schmidt law (Kennicutt 1998; Boissier et al$.$ 2003), it is expected %that star formation will be quenched in the outer (lower density) %portions of the galaxy producing truncated disks. While interferometric observations of galaxies %in Virgo have clearly shown that HI disks are less extended in those %objects located close to the cluster center (Cayatte et %al$.$ 1990), the observational evidence for a truncation of the star %forming disks has been proven by H$\\alpha$ imaging (Koopmann \\& Kenney 2004a,b; Boselli \\& %Gavazzi 2006). The degree of truncation should depend both on the %ages of the interaction and of the galaxy disk stellar populations, %and could thus be used as a clock to date the perturbing event, giving a new, strong constraint %for discriminating among those physical processes acting on galaxies. %Although a truncation of the disk profile has been %predicted (Larson et al$.$ 1980), we still do not know what the %passive evolution of a stellar disk is once its gas is removed. In %particular, it is unclear whether the progressive radial suppression %of star formation is able to reproduce the structural properties of %lenticulars, generally characterized as having higher surface %brightness of their stellar disks and higher bulge-to-disk ratios %than spirals (Dressler 1980). \\\\ We have been collecting multi-frequency data for a large sample of late-type galaxies in nearby clusters and in the field in order to undertake the comparative statistical analysis of any systematic differences between cluster and field objects. The most important results obtained from the analysis done so far are reported in Boselli \\& Gavazzi (2006). Combined with multi-zone models for the chemical and spectrophotometic evolution of galaxies (Boissier \\& Prantzos 2000), this unique database is helping us to understand the evolution of cluster spirals. \\\\ The aim of the present paper is to give a complete description of the multi-zone spectrophotometric models of galaxy evolution used to study the effects of the environment on cluster galaxies. %These are %essentially those of Boissier \\& Prantzos (2000) modified to take care %of a time-dependent gas stripping event as predicted by the dynamical %models of Vollmer et al. (2001b) (ram pressure scenario) or %by simply suppressing gas infall (starvation scenario).  As a first application we present a detailed study of the radial profiles of the Virgo cluster galaxy NGC~4569 (M90).  NGC~4569, the prototype anemic galaxy as defined by van den Bergh (1976), is extremely deficient in HI, having only about one tenth of the atomic gas of a comparable field galaxy of similar type and dimensions. This galaxy has a truncated H$\\alpha$ and HI radial-density profile (at a radius of $\\sim$ 5 kpc; see Fig.~\\ref{images} and \\ref{FIGmodelREF}) as firstly noticed by Koopmann \\& Kenney (2004a,b) and Cayatte et al$.$ (1994), bearing witness to a recent interaction with the cluster environment. NGC~4569 is located close ($\\sim$ 2 degrees, $\\sim$ 240 kpc for a virial radius of 1.7 Mpc) to the cluster center. Being one of the largest galaxies ($\\sim$ 10 arcmin, $\\sim$ 50 kpc) in the Virgo cluster, NGC~4569 is the ideal candidate for our study since it is  spatially resolved at all the wavelengths considered here. NGC 4569 has also been selected because it has been previously  studied using dynamical models by Vollmer et al. (2004a). A direct comparison of the results obtained using totally independent techniques (in particular, dating the interaction)  is thus possible. We just remind that multifrequency and kinematical observation of the inner part of NGC 4569 done by Jogee et al. (2005) seem to indicate that the nucleus and the inner stellar bar of the galaxy might have been triggered by a recent tidal interaction. There are however no observational evidences indicating that the HI disk has been truncated by any gravitational interaction with nearby companions. %Our %assumption of a ram-pressure stripping event at the origin of the %radial truncation of the gas and stellar disk of NGC 4569 is not in %contradiction with the Jogee et al. result for two reasons: as %discussed in Boselli \\& Gavazzi (2006) tidal interactions might induce %some gas infall in the nuclei of cluster galaxies, but can hardly %remove the outer disk gas {\\textbf E' VERO CHE LA TIDAL NON TRONCA MA POTREBBE %MODIFICARE LA NORMALIZZAZIONE DEL PROFILO PORTANDO GAS AL CENTRO. IO LASCEREI SOLO %LA SECONDA MOTIVAZIONE. INOLTRE NON \u00c8 ANCORA CHIARO SE NGC4569 HA AVUTO TIDAL, INFATTI %COME DICE VOLLMER, POTREBBE ESSERE ANCHE SOLO UNO STARBURST DOVUTO A SECULAR EVOLUTION ED ALLA %PRODUZIONE DELLA BARRA AL CENTRO}. Since our models are adapted for disk galaxies, we will however limit our study to the disk component excluding the central nucleus.\\\\ The combination of the multi-frequency 2-D data with our spectrophotometric models allow us to study, for the first time, the radial evolution of the different stellar populations in this prototype, gas-stripped cluster galaxy with the aim of understanding whether its structural properties can evolve, even in principle,  into those of a typical cluster lenticular (S0) galaxy.  A more complete analysis of a statistically significant sample of cluster galaxies will be presented in a future communication. ",
        "conclusions": "\\begin{figure} \\centering \\epsscale{1.0} \\includegraphics[width=8cm,angle=0]{f10.eps} \\small{\\caption{ Variation of the effective surface brightness (mean surface brightness within $R_e$, the radius containing half of the total light) and radius due to differential variation of the star formation history of NGC~4569 in a) a starvation and b) ram-pressure stripping scenario. Open triangles are for the unperturbed model, the other symbols for different ages of the interaction. } \\label{figmuchange}} \\end{figure} The present work gives the first quantitative estimate of the structural evolution of stellar disks in cluster galaxies due to gas removal caused by a dynamical interaction of the galaxy with the IGM, delivering a strong message concerning the passive stellar evolution of stripped disks. %Although the model only qualitatively reproduces the observed %multi-wavelength radial profiles (the mismatch being attributed to %resolution effects and large uncertainties in the extinction %correction) it delivers a strong message concerning the passive %stellar evolution of stripped disks. %NGC 4569 seems to be a representative %example of HI-deficient, anemic cluster galaxies, with %, as it seems to be the %case (Boselli \\& Gavazzi 2006), %the results obtained from this analysis can be extended to a larger %galaxy population. First of all it is clear that the truncation of the total gas disk profile is soon reflected on the young population stellar disk, confirming the predictions of Larson et al$.$ (1980). As a consequence gas-stripped galaxies have color gradients opposite to that of normal, isolated spirals, which generally have bluer colors in their outer disks. NGC~4569 is bluest towards the center (see {Fig. \\ref{figcolorsbprof} and \\ref{RGB}}). The trend is especially true for colors tracing the relatively young populations ($< \\sim$ 10$^8$ yr); colors tracing populations older than the interaction event show the usual gradient (i.e., redder towards the center, as the B-$i$ color index). The inversion of the color gradient, here observed for the first time in a cluster galaxy, is well reproduced by our model.\\\\ The second major conclusion of the present analysis is that the perturbation which induced the truncation of the stellar disk of NGC 4569 is relatively recent event ($\\sim$100~Myr, in any case $\\leq$ 500 Myr). Ram-pressure is favored with respect to starvation since the latter is simply not able to reproduce the observed truncation of the gas and star forming (H$\\alpha$, 8 $\\mu$m, FUV) disk. \\\\ Since NGC 4569 seems to be typical of the HI-deficient galaxy population inhabiting nearby clusters, characterized by truncated HI and star forming disks (Cayatte et al. 1994; Koopmann et al. 2006), with largely unperturbed older stellar populations\\footnote{We remind that there is a strong relationship between the ratio of the optical to H$\\alpha$ radii and the HI deficiency parameter in cluster galaxies (see Fig. 11 of Boselli \\& Gavazzi 2006).}, we generalize this statement by saying that gas stripping in cluster galaxies is due to relatively recent perturbations. This result is thus a major constraint on the evolution of cluster galaxies since it rejects long time-scale phenomena such as harassment and starvation.\\\\ % %The consequence of these findings in the interpretation of the %evolution of cluster spiral galaxies is significant. The detailed %analysis presented in Boselli \\& Gavazzi (2006) has shown that the %degree of truncation of the star forming disk in cluster spirals is %proportional to their HI-deficiency parameter. At the same time the %stellar disks observed in the R band are not truncated, as in NGC %4569. This observational evidence allow us to conclude that NGC 4569 %is indeed representative of the HI-deficient galaxy population in %nearby clusters and to generalize that the perturbation responsible %for their gas depletion, whatever its nature is, is a relatively %recent event ($\\leq$ 500 Myr). %If gas stripping in %earby clusters is primarily due to the interaction of galaxies with %the cluster IGM (consistently with the recent work by Boselli \\& Gavazzi 2006 and %with the specific case of NGC 4569 as indicated by this work or by Vollmer et al. 2004a) we can conclude %that these galaxy-IGM interactions took place less than half a Gyr ago. %those phenomenon invoked to explain %the decrease of the star formation activity (and HI gas deficiency) %observed at large cluster-centric distances (harassment, starvation and %pre-processing) just because their timescales are too long by a factor %of $\\sim$ 10.\\\\ The ram-pressure stripping scenario has recently been criticized as unable to reproduce the radial decrease with the cluster-centric distance of the star formation activity of galaxies in the nearby Universe.  Recent, complete spectroscopic surveys of the nearby universe such as the SDSS or the 2dF have shown that the activity of late-type galaxies begins to decrease at 1-2 virial radii from the cluster center (Gomez et al. 2003; Lewis et al. 2002, Tanaka et al. 2004, Nichol 2004), scale-lengths comparable to those observed in the HI (Gavazzi et al. 2005, 2006a) and/or H$\\alpha$ (Gavazzi et al. 2002a, 2006b) in nearby clusters. Since these distances (1-2 virial radii) are significantly larger (a factor of 5-10) than those where galaxy-IGM interactions are expected to be more efficient ($\\sim$ inside a core radius), other processes already active at the periphery of clusters [such as galaxy harassment (Moore et al. 1996), starvation (Balogh et al. 2000) or pre-processing (Fujita 2004; Cortese et al. 2006b)] have been proposed to explain the radial decrease of the star formation activity with the cluster-centric distance.\\\\ As extensively discussed in Boselli \\& Gavazzi (2006), however, the presence of galaxies with clear signs of ongoing ram-pressure stripping at the periphery of nearby clusters and the large velocity dispersion, combined with the radial orbits (Dressler 1986; Solanes et al. 2001) of the late-type galaxy population can explain the observed decrease of the star forming activity at large cluster-centric radii in a ram-pressure stripping scenario.  NGC 4569, for instance, might have travelled $\\sim$0.6 Mpc (1/3 of the virial radius) since it was stripped, a distance significantly larger than the core radius of the cluster (130 kpc). We recall that in Virgo the drop of the star formation in bright galaxies is observed at $\\sim$0.7 virial radii (Gavazzi et al. 2006b), while the increase of the HI-deficiency parameter occurs at about 1 virial radius (Gavazzi et al. 2005). The higher velocity dispersion and IGM gas density of rich, evolved clusters, as well as a clumpy IGM distribution, might thus be at the origin of ram-pressure stripped galaxies up to $\\sim$ one virial radius. \\\\     %However if the galaxy-IGM interaction is able to perturb %galaxies only in the core of the cluster, how can we reproduce the %observed cluster-centric radial variation of the star formation %activity and of the HI gas content observed by Gavazzi et %al. (2005,2006a)? %If these late-type galaxies entering the cluster are %moving on radial orbits, as seems to be the case (Dressler 1986; %Solanes et al. 2001), then at an average speed of $\\sim$ 1150 km %s$^{-1}$ along the line of sight (the velocity dispersion of the %late-type component in the Virgo cluster, Boselli \\& Gavazzi 2006) %these galaxies might have traveled $\\sim$ 0.6 Mpc, and thus be present %up to $\\sim$ one third of the virial radius of the cluster. This %scale-length is just slightly smaller than that of the drop of the star formation %activity in bright galaxies ($\\sim$ 0.7 virial radii) and of the %increase of the HI-deficiency parameter ($\\sim$ 1 virial radius; %Boselli \\& Gavazzi 2006) {\\textbf NON E' VERO!!! 1/3 E' MENO DELLA META' DI 0.7 ED E' UN %TERZO DI 1! QUINDI 1/3 NON E' SLIGHTLY SMALLER THAN 1!\\\\}. %{\\textbf DA QUI A FINE PARAGRAFO NON MI E' CHIARO COSA VUOI DIRE} %The consistency with the results obtained by %the SDSS and the 2dF surveys, where thresholds on the star formation %activity are observed at slightly larger scalengths (1-2 virial radii) %is guaranteed by the fact that i) ram-pressure stripping might be %efficient not only in the cluster center but up to $\\sim$ 0.75 virial %radii (Boselli \\& Gavazzi 2006) and ii) the velocity dispersion %and the the density of the intracluster medium in %more relaxed and evolved clusters might be larger than in Virgo %making ram-pressure a more efficient process %(Boselli \\& Gavazzi 2006).\\\\    One of the most intriguing (and still open) questions regarding the effects of the environment on the evolution of galaxies is that of the origin of lenticulars, and their overabundance in the centers of rich clusters. Are lenticulars an independent population of galaxies formed in the primordial high-density environments, or were they spiral disks whose star formation activity had been quenched once their gas reservoir was removed by the unfavorable cluster environment? %Answering this question unequivocally is a challenge. The present work has shown for the first time how a galaxy-cluster IGM interaction is able to remove most of the gas reservoir, inducing important structural modifications in the stellar disk properties. We have in fact shown that, because of the differential radial stellar evolution of spiral disks, we can expect that cluster spirals have (at least at short wavelengths) more truncated disk profiles, inverting the outer color gradient with respect to similar but unperturbed objects. The surface brightness of the disk, however, mildly decreases in H$\\alpha$ and in the UV bands while remaining mostly constant at longer wavelengths, even 5~Gyr after the interaction (Fig.~\\ref{figmuchange}) \\footnote{40\\% of  S0 galaxies in the Coma cluster underwent a star formation event in their centers in the last 5 Gyr, Poggianti et al. (2001).}. The differential evolution of the stellar disk due to gas stripping alone is thus not able to reproduce the $\\sim$ 0.65 mag increased central surface brightness of present-day lenticulars (Dressler 1980; Boselli \\& Gavazzi 2006). In the case of starvation stellar disks are not truncated ($R_e$ just slightly decreases) while surface brightnesses significantly decrease on long time scales.\\\\ Gravitational perturbations, such as tidal interactions with other galaxies (Merritt 1983), interactions with the cluster potential well (Byrd \\& Valtonen 1990) or a mixture of both %(called `galaxy harassment' by Moore et al$.$ 1996) must be invoked to reproduce the observed properties of nearby lenticulars.  As described in the introduction, similar conclusions have been obtained from the analysis of statistical, kinematical, structural and spectro-photometrical properties of nearby clusters.\\\\ This new study is consistent with the idea that the present evolution of late-type galaxies in clusters differs from that at earlier epochs, where late-type galaxies were mostly perturbed by dynamical interactions (pre-processing and/or galaxy harassment; Dressler 2004, Moore et al$.$ 1996) which were able to thicken the stellar disks thereby producing the present-day cluster lenticulars. \\\\ To conclude we can say that the combination of multifrequency observations of spatially resolved cluster galaxies combined with multi-zone spectro-photometric models of galaxy evolution provides an extremely useful tool to study the evolution of cluster galaxies. The results here presented for a typical HI-deficient, anemic Virgo cluster galaxy and extrapolated to the whole cluster galaxy population indicate that:\\\\ 1) The gas removal due to a ram-pressure stripping event can reproduce the observed radial truncation of the stellar disk stronger at shorter wavelengths. On the other hand,  starvation is not able to truncate gas or stellar disks.\\\\ 2) As a consequence, color gradients make cluster objects redder in the outer disk than in the inner regions, and are thus inverted with respect to normal, isolated late-type galaxies.\\\\ 3) This technique is useful to age-date the interaction. If NGC 4569 represents the typical HI-deficient cluster galaxy with truncated H$\\alpha$ and HI disk and normal intermediate age stellar disk, then our modeling suggests that gas stripping is a relatively recent event since it probably took place $\\leq$500 Myr ago in these systems.\\\\ 4) The effective surface brightness of the stripped galaxies remains constant or even mildly decreases after the interaction.\\\\ We hope to confirm this original result in the near future once multi-frequency data come available for a statistical significant sample of late-type cluster galaxies."
    },
    "0609/astro-ph0609181_arXiv.txt": {
        "abstract": "We studied the radio properties of very young massive regions of star formation in HII galaxies, with the aim of detecting episodes of recent star formation in an early phase of evolution where the first supernovae start to appear. Our sample consists of 31 HII galaxies, characterized by strong Hydrogen emission lines, for which low resolution VLA 3.5~cm and 6~cm observations were obtained. We complemented these observations with archival data at 20~cm. We found that the radio spectral energy distribution (SED) has a range of behaviours; 1) there are  galaxies where the SED is characterized by a synchrotron-type slope, 2) galaxies with a thermal slope, and, 3) galaxies with possible free-free absorption at long wavelengths. The latter SEDs were found in a few galaxies (e.g. UM\\,533, Tololo~1223-388) and represent a signature of heavily embedded massive star clusters closely related to the early stages of massive star formation.  Based on the comparison of the star formation rates (SFR) determined from the recombination lines and those determined from the radio emission we find that SFR(H$\\alpha$) is on average five times higher than SFR(1.4~GHz). We confirm this tendency by comparing the ratio between the observed flux at 20 cm and the expected one, calculated based on the H$\\alpha$ star formation rates, both for the galaxies in our sample and for normal ones. This analysis shows that this ratio is a factor of 2 smaller in our galaxies than in normal ones, indicating that they fall below the FIR/radio correlation. This result is further confirmed by the detection of high q-parameter values (the ratio of infrared to radio fluxes) in a few sources. These results suggest that the emission of these galaxies is dominated by a recent and massive star formation event in which the first supernovae (SN) just started to explode. This indicates that the radio emission is most likely dominated by free-free continuum, and that the emission at low frequencies may be optically thick, in agreement with the observed SEDs.  We combined the VLA data with age indicators based on optical observations (e.g.~equivalent width of \\Hb) together with the ratio between the far infrared and the radio continuum fluxes and proposed an evolutionary scenario to explain the observed trends.  We conclude that the systematic lack of synchrotron emission in those systems with the largest equivalent width of \\Hb\\ can only be explained if those are young starbursts of less than 3.5Myr of age, i.e.~before the first type II SNe start to explode. ",
        "introduction": "One of the outstanding problems in the study of starburst galaxies is understanding how massive stellar clusters (MSC) form and evolve. Most of what is known about MSC  comes from the study of relatively old objects, since they evolve fast and can be heavily embedded in their earlier stages, thus being difficult to detect (Johnson 2004). Recent attempts using radio observations have been successful in the detection of several of these embedded sources in nearby starburst galaxies. During the early stages of evolution, very young star clusters should either lack or have a deficit of synchrotron emission, being dominated by free-free radio emission. Given that the synchrotron emission observed in starbursts is due to supernova activity, it is a direct tracer of the end of the evolution of massive stars (M\\,$\\ge$ 8\\,M$_{\\odot}$). In the case of a coeval star formation burst, the first type II SNe appear after only 3.5\\,Myr and last until the 8\\,M$_{\\odot}$ stars explode at around 40\\,Myr. Consequently, the lack of synchrotron emission is a good indicator of very young clusters.  Recent radio observations (e.g. Turner, Beck, \\& Ho 2000; Johnson et al. 2001; Johnson et al. 2003; Cannon \\& Skillman 2004) have detected in a few nearby starburst galaxies, sources with compact flat radio spectrum which can be securely identified as young star clusters. In some of the cases these authors detect an inverted spectrum at low frequencies, due to free-free absorption, indicating that the cluster still is deeply embedded in the material from which it was formed. On a more global scale, Roussel et al. (2003) detected galaxies with abnormally high infrared to radio continuum ratios, identified as nascent starbursts.  One important question is how MSC do form? Are they the product of an instantaneous event or the star formation activity lasts several Myr instead? The situation is not clear, the excellent radio-FIR correlation (e.g. Yun et al. 2001) strongly suggests that star formation, at least at galactic scales, must proceed continuously; also in a study of the statistical properties of HII galaxies we found strong indications of continuous star formation (Terlevich et al.~2004). On the other hand, the presence of strong WR features in the spectra of starbursts and HII galaxies suggests a high degree of synchronization in the formation of massive stars. The question may be solved studying the properties of a complete sample of young systems.  Although these young starbursts are known, their numbers are relatively small. The detection of a larger number of young sources is an important step in the study of the early stages of massive star formation and evolution. This will allow us to better understand their properties and environment.  For this work we have taken a different approach, which consists in determining the radio properties of a sample of HII galaxies selected because their emission line spectra indicate extreme youth. HII galaxies are low mass objects whose emission and thus most observables are dominated by a massive burst of star formation. Their optical spectrum is identical to that observed in giant HII regions like 30-Doradus in the LMC. These properties make these galaxies ideal targets for a study of integrated properties in search for features related to MSCs. The best sources can then be studied in more detail with higher resolution observations. In this sense HII galaxies are considered ``young\". In fact they are probably the youngest stellar systems that can be studied in any detail. The youth scenario is supported not only by the strong emission lines from ionized gas, but also by their underlying stellar continuum properties, i.e. no stellar absorptions are detected in those HII galaxies with the strongest emission lines (EW(H$\\beta) > $150\\AA) and only weak hydrogen and helium and some very weak metal lines are detected on those more evolved HII galaxies (EW(H$\\beta) < $ 50\\AA). The lack of detection of stellar absorptions in the extreme HII galaxies is due to the fact that the optical spectrum of very young clusters is dominated by the light of massive blue supergiants close to the main sequence turn-off. The spectrum of these stars shows narrow and relatively weak absorptions of hydrogen and helium plus extremely weak metal lines. In the case of HII galaxies (and also HII regions) the weak and narrow stellar absorptions are filled with strong emission from the ionized gas of the associated HII region making it impossible to  detect directly their presence. This has been traditionally a problem for studies to derive highly precise helium abundances  as required for example for primordial helium determinations (e.g.~Skillman, Terlevich \\& Terlevich 1998). Thus very strong emission plus a ``featureless'' continuum is the signature of extremely young HII galaxies.  This ``youth\" of the HII galaxies does not imply that the whole galaxy is young, in fact the opposite is probably true given the colors of the extensions detected in some HII galaxies (Telles \\& Terlevich, 1997) and the simple fact that the metal content of all HII galaxies is typically around 0.1-0.2 of solar (e.g.~Hoyos \\& D\\'\\i az 2006) implying a substantial amount of chemical evolution or astration previous to the present dominant burst. Stellar population synthesis analysis (Raimann et al. 2000a, Raimann et al. 2000b) show that there is some evidence for older stellar populations in these sources, albeit at a very low level, again confirming that their emission is in fact dominated by recent star formation.  Thus ``young\" and ``age\" in this context refers to the present burst. HII galaxies could be, and probably are, much older.  Here we present the results of a study of integrated radio spectral energy distributions (SEDs) of nascent starburst candidates. We find several galaxies with flat spectrum, typical of thermal emission, or even with inverted spectra, indicating the presence of heavily embedded star clusters. With the data at hand we compare the radio and emission line (H$\\alpha$,H$\\beta$) properties of these galaxies and propose a qualitative evolutionary model.  This paper is organized in the following way. In Section~\\ref{Sec:Sample} we describe how the sample was selected. Section~\\ref{Sec:Observatios} describes the VLA observations at 6~cm (4.9~GHz) and 3.5~cm (8.4~GHz) and the data reduction. Section~\\ref{Sec:Radio} presents the radio SEDs, spectral indices and the proposed evolutionary model. In Section~\\ref{Sec:Opt} we study the optical properties of the sample galaxies and calculate the SFRs obtained from the optical emission lines. In Section~\\ref{Sec:OpticalRadio} we combined optical and radio data to discuss the nature of the selected galaxies and explore simple evolutionary scenarios. In Section~\\ref{Sec:FIR} we estimate the FIR to radio continuum flux ratios ($q$ parameter) and its relation to other observables to confirm the young nature of some of the observed galaxies. Finally, a summary of our results is presented in Section~\\ref{Sec:Discussion}. ",
        "conclusions": "\\label{Sec:Discussion}  The tight correlation between far infrared (FIR) and radio luminosity confirms that radio emission is also a good tracer of recent star formation (Yun et al. 2001). The origin of this correlation is the population of massive stars that are responsible for the dust heating, the FIR radiation and the supernova activity that generates the radio emission.  However, deviations from this correlation are expected both in the early phases of a starburst -- before the first Type II SNe explode -- or in the post-starburst phase when the synchrotron emission ages and is also contaminated by the emission from less massive stars (Bressan et al. 2002).  Based on the combination of optical, FIR and radio data, we have detected a large deviation from the FIR-radio  correlation within a sample of normal and starburst galaxies. Our sample of 31 galaxies is strongly biased to young systems which have strong EW(\\Hb) and relatively high SFR. Only a few galaxies in our sample follow the Yun et al. (2001) relation and the SFR calculated from radio continuum at 1.4 GHz is a factor of 4 below the SFR given by optical emission lines. We quantified this deficit introducing the $d$-parameter which is the ratio between the observed radio continuum flux at 1.4 GHz and the expected one calculated using the SFR provided by the optical emission lines (H$\\alpha$) and the empirical relation given by Yun et al. (2001) which relates the flux at 1.4GHz with the current SFR determined by the FIR emission. In most of the cases the $d$-parameter is well below 50\\% with only 2 galaxies having values above 75\\%.  The summary of our results is presented in Table~\\ref{Tab:Summary}. We tabulated the extreme values of the parameters which point to the existence of a recent star forming event in our sample of galaxies. The indicators with high values of the EW(\\Hb) and \\OIII /\\Hb\\ ratio, have a value of $\\alpha^{1.4}_5$ typical of a region dominated by thermal emission, a low $d$ parameter, and a high $q$ parameter. Only two galaxies, Mrk\\,1318 and Tol\\,1358-328 have the five indicators corresponding to very young objects while 5 galaxies present 4 of the 5 indicators. In the galaxies Tol\\,0226-390, Tol\\,1247-232 and UM\\,533 the only indicator not pointing to a very young system is the $q$ parameter which is in all cases very close to the canonical value. For the case of Tol\\,1304-386 we have no FIR information so we could not estimate the $q$ parameter. Mrk\\,711 has a relatively low EW(\\Hb) but a high \\OIII /\\Hb\\ ratio, pointing to the existence of a new burst on top of an older stellar population. In this galaxy all the other tracers of a young population are present. Another 10 galaxies show signs of young systems in 3 of the 5 tracers considered (see Table~\\ref{Tab:Summary}) and are also extremely good candidates to look for deeply embedded super stellar clusters where only thermal radio and mid infrared radiation can be detected.  For the early phase deviation to occur, a highly synchronized star formation is required, i.e. the massive star forming phase should be shorter than $\\sim$3-4Myr in order to prevent the appearance of the first supernovae associated to the more massive stars and the consequent synchrotron emission. This also suggests that there has been no large starforming event in the previous $10^8$yr. The results summarized above have important implications not only for low mass galaxy evolution but also for the study of young massive star clusters and for the understanding of the early stages of their evolution. We showed that by combining optical spectroscopy with radio continuum observations we can efficiently detect extremely young systems, before the first Type II SNe explode. These systems represent a significant enlargement of the current sample of nearby starbursts, like Henize\\,2-10 (Johnson \\& Kobulnicky 2003) and NGC\\,5253 (Beck et al. 1996, Turner et al. 2000, Turner \\& Beck 2004), known to have these properties. A detailed higher resolution study of these sources, to be published in a forthcoming paper, will be extremely valuable for the understanding of their properties and to study evolutionary scenarios. Furthermore, we hope that the study of a complete sample will allow the segregation between instantaneous and extended starformation scenarios."
    },
    "0609/astro-ph0609462_arXiv.txt": {
        "abstract": "In the latest in our series of papers on \\sat\\ and ground-based optical follow-up of new candidate magnetic cataclysmic variables (mCVs) found in the Sloan Digital Sky Survey, we report classifications of three systems: SDSS J144659.95+025330.3, SDSS J205017.84--053626.8, and  SDSS J210131.26+105251.5.  Both the X-ray and optical fluxes of \\ta\\ are modulated on a period of $48.7\\pm0.5$ min, with the X-ray modulation showing the characteristic energy dependence of  photo-electric absorption seen in many intermediate polars (IP). A longer period modulation and radial velocity variation is also seen at around 4 hrs, though neither dataset is long enough to constrain this longer, likely orbital, period well. \\tb\\ appears to be an example of the most highly magnetized class of mCV, a disk-less, stream-fed polar. Its 1.57 hr orbital period is well-constrained via optical eclipse timings; in the X-ray it shows both eclipses and an underlying strong, smooth modulation. In this case, broadly phase-resolved spectral fits indicate that this change in flux is the result of a varying normalization of the dominant component (a 41 keV \\mek), plus the addition of a partial covering absorber during the lower flux interval. \\tc\\ is a more perplexing system to categorize: its X-ray and optical fluxes exhibit no large periodic modulations; there are only barely detectable changes in the velocity structure of its optical emission lines; the X-ray spectra require only absorption by the interstellar medium; and the temperatures of the \\mek\\ fits are low, with maximum temperature components of either 10 or 25 keV.  We conclude that \\tc\\ can not be an IP, nor likely a polar, but is rather most likely a disc accretor-- a low inclination SW Sex star. ",
        "introduction": "The Sloan Digital Sky Survey \\citep[SDSS;][]{york00} has provided a wealth of new cataclysmic variables; the close binary systems with active accretion from a late main sequence star to a white dwarf \\citep[reviewed in][]{warn95}. Due to its sensitivity down to $\\sim$21st mag, SDSS is especially suited to the discovery of faint CVs with low accretion rates and short orbital periods. Included in the discoveries are dozens of systems with a noticeable emission line of \\ion{He}{2}, which is often a signature of a white dwarf with a high magnetic field. The highest field systems (polars) have no accretion disk and the white dwarf spin is synchronized to the orbit. The intermediate polars (IPs) usually have some outer disk, with magnetic curtains channelling the inner disk material to the magnetic poles of the white dwarf which is spinning faster than the orbital timescale. Polars can be identified by circular polarization, cyclotron harmonics and/or Zeeman splitting \\citep[review in][]{wick00}, while IPs are found through the detection of their spin and orbital periods. In addition, there is a class of high accretion rate disk systems with possible low magnetic field white dwarfs \\citep{rodr01}, termed SW Sex stars, which can also show strong \\ion{He}{2} lines in their optical spectra.  For the last few years, we have used \\sat\\ to identify the nature of the CVs with \\ion{He}{2} and to explore the nature of X-ray heating in these systems. Polars generally exhibit hard X-rays from the accretion shock and soft X-rays from the white dwarf surface heated by the hard X-rays. The ratio of soft to hard X-rays shows a dramatic change with magnetic field strength and accretion rate. The IPs are generally hard X-ray emitters with large absorption effects from the accretion curtains. Usually the spin period is much more pronounced in the X-ray than the optical.  The SW Sex stars usually show no X-ray variation and are hard X-ray emitters with low absorption.   The 3 sources in this study are all about 18th mag in the SDSS $g$ filter and were all identified as potential magnetic CVs in \\citet{szko03}. SDSS\\thinspace J144659.95+025330.3 (hereafter \\ta) showed no polarization and limited time-resolved spectra could not determine the orbital period. SDSS\\thinspace J205017.85--053626.8 (\\tb) has the strongest \\ion{He}{2} line of the 3 objects, with a peak flux comparable to H$\\beta$ and showed evidence of high and low states of mass transfer. A short span of time-resolved spectroscopy showed an orbital period near 2 hrs. Followup high speed photometry by \\citet{woud04} revealed brief (260s), deep (1.5 mag) eclipses and  determined a precise period of 1.57 hrs. The third system, SDSS\\thinspace J210131.26+105251.5 (\\tc) has very strong Balmer lines but did not show any velocity variation during 1.3 hr of time-resolved spectroscopy.  The {\\em XMM-Newton} data obtained on these 3 sources, combined with additional photometry, spectroscopy and spectropolarimetry has allowed us to determine that \\ta\\ is an IP, \\tb\\ is a likely polar, while \\tc\\ escapes an easy classification, but appears most likely to be an SW Sex star. Our observations and results are described below. ",
        "conclusions": "We have reported on observations allowing classification of three new SDSS CVs showing prominent \\ion{He}{2}: \\ta, \\tb\\ and \\tc.  Each represents a different class of CV, known to exhibit such high excitation emission lines.  \\ta\\ is a clear example of an IP.  We find two periodicities at about 4 hr and 49 min, arising respectively from the orbit and asynchronously spinning white dwarf.  The X-ray spectrum comprises a multi-temperature thermal plasma, with a maximum of 60 keV, modulated by phase-dependent local absorption. This likely originates from obscuration by the accretion curtains formed by the accreting material as it flows between the disrupted disc and the magnetic poles of the white dwarf.   \\tb\\ represents the most highly magnetized of the three, a fully spin-orbit synchronized polar, with an orbital period of 1.57 hr. In this case no disk forms, and all accretion is channeled along the field lines to a single dominant pole. Owing to its high inclination, there is a phase interval when the accretion stream obscures the X-ray emitting region, but most of the modulation on the spin/orbit appears consistent with simply its changing aspect.  The X-ray emission has two components: the dominant one is characteristic of the post-shock bremsstrahlung-cooled optically thin plasma of the accretion column, here with temperatures ranging from 1 keV to 40 keV; this also heats the surface of the white dwarf leading to additional soft blackbody (30 eV) emission.  In \\tc, we suggest that the combination of high accretion rate and lower magnetic field strength allows disk-mediated accretion, but also the formation of \\ion{He}{2} emission lines in the optical spectra, i.e. if it were at higher inclination it would appear as a classic SW Sex star.  Its X-ray spectra indicate a rather cool thermal plasma (\\til10 keV) more typical of disk systems, and an unobscured line of sight, inconsistent with an IP. From only weak modulation (unlike a polar) of its X-ray and optical lightcurves we suggest an orbital period in the 100 min range."
    },
    "0609/astro-ph0609148_arXiv.txt": {
        "abstract": "In order to complete our \\hi survey of galaxies in the Sculptor group area, VLA observations of NGC 24 and NGC 45 are presented. These two galaxies of similar magnitude $M_B \\sim -17.4$ lie in the background of the Sculptor group and are low surface brightness  galaxies, especially NGC 45. The \\hi distribution and kinematics are regular for NGC 24 while NGC 45 exhibits a kinematical twist of its major axis. A tilted-ring model shows that the position angle of the major axis changes by $\\sim 25\\degr$. A best-fit model of their mass distribution gives mass-to-light ratios for the stellar disk of 2.5 and 5.2 for NGC 24 and NGC 45  respectively. These values are higher than the ones expected from stellar population synthesis models. Despite the large dark matter contribution, the galaxy mass is still dominated by the stellar component in their very inner regions. These high mass-to-light ratios are typical of what is seen in low surface brightness galaxies and may indicate that, in those galaxies, disks are far from the maximum disk case. The halo parameters derived from the best-fit models are thus lower limits. ",
        "introduction": "\\label{sec:intro} At the beginning of the 90's, a series of papers was started on the \\hi properties of the Sculptor group galaxies (Puche \\& Carignan 1988; Carignan \\& Puche 1990a,b; Puche, Carignan \\& Bosma 1990; Puche \\& Carignan 1991; Puche, Carignan \\& van Gorkom 1991; Puche, Carignan \\& Wainscoat 1991). Since the Sculptor galaxies are all late-type spirals with very little or no bulge, they are ideal candidates for determining the basic dark halo parameters (Carignan \\& Freeman 1985). Their simple disk + halo stucture makes it easier to identify the contribution of each component to the rotation curve. Also, since the \\hi distribution of late-type spirals is usually extended (Huchtmeier, Seiradakis \\& Materne 1980), it is possible to measure their rotation curve out to many disk scale lengths, an essential condition to tie down the halo parameters (Lake \\& Feinswog 1989).   \\begin{figure*}\\centering \\includegraphics[width=\\columnwidth]{f1.eps}\\includegraphics[width=\\columnwidth]{f2.eps} \\caption{DSS blue images of NGC 24 (left) and NGC 45 (right). \\label{dssimages}} \\end{figure*}  As described in Puche \\& Carignan (1988), the Sculptor Group, the nearest group of galaxies from the Local Group, covers an area of $\\sim 20\\degr$ on the sky. In this area, besides the five main members studied so far (NGC 55, 247, 253, 300 and 7793), two more galaxies, NGC 24 and NGC 45 (Figure~\\ref{dssimages}), fall in the same region and have similar sizes and degree of resolution as well as $\\sim$ 20 known dwarf galaxies (C\\^ot\\'e, Freeman \\& Carignan 1997; C\\^ot\\'e, Carignan \\& Freeman 2000). However, when looking at the estimated distances, a mean distance of $2.4 \\pm 0.6$ Mpc is found for the main members (de Vaucouleurs 1978) but 6.8 Mpc (7$\\sigma$) and 5.9 Mpc (6$\\sigma$) for NGC 24 and NGC 45, respectively (based on a Hubble constant of 75 \\kms Mpc$^{-1}$, Tully 1988). Moreover, when doing dynamical mass estimates of the group, any combination of the main members gives consistently $\\sim3 \\times 10^{12}$ M$_{\\odot}$ for the group, while including NGC 24 and NGC 45 makes the estimate jump to $\\sim3 \\times10^{13}$ M$_{\\odot}$, clearly too large for a loose group like Sculptor. All this suggests strongly that NGC 24 and NGC 45 are not members of the group but background objects.  Even if they were not members of the Sculptor group, it was decided to study them because they belong to a very interesting subgroup of galaxies called the low surface brightness (LSB) spiral galaxies, especially NGC 45. The kinematical study of LSB galaxies currently feeds the ``cusp-core controversy'', (e.g. Swaters et al. 2003; Hayashi et al. 2004; de Blok 2005, and references therein). One of the many important questions which is brought to the forefront from the numerous studies of this type of object and which motivates part of the present work, is whether or not LSBs (and more generally low mass late-type spirals) can possibly form a distinct class in the Hubble sequence. The hypothesis emanating from studies on LSBs is: morphologically they belong to the spiral class, but dynamically they look more like the dwarf irregular class since they are dominated at all radii by their dark halo (Jobin \\& Carignan 1990; C\\^ot\\'e, Carignan \\& Sancisi 1991; Martimbeau, Carignan \\& Roy 1994).  Moreover, a study of optical rotation curves of spiral galaxies (Buchhorn 1992) showed convincingly that the optical rotation curves in the inner parts of these galaxies give \\rml ratios which are very high. The most likely explanation for this effect is that the dark matter component contributes a large fraction of the gravitational field even in the inner regions of LSBs galaxies (de Blok \\& Bosma 2002). Verheijen (1999) and Verheijen \\& Tully (1999) have shown convincingly the clear dichotomy between the rotation curves of HSB (high surface brightness) galaxies, which are most of the time close to the maximum disk situation, and those of LSB galaxies which are dominated at all radii by the dark matter component. Thus, in terms of their mass distribution, LSBs appear to be more closely related to dwarf irregular galaxies (e.g. DDO 154: Carignan \\& Freeman 1988; Carignan \\& Beaulieu 1989) where the luminous component is known to make a negligible contribution to the gravitational field everywhere in the disk, than to massive spirals (e.g. NGC 6946: Carignan et al 1990) where the luminous disk is the main contributor in the inner regions and the dark component only contributes significantly in the outer parts.   \\begin{figure*}\\centering \\includegraphics[width=\\columnwidth]{f3.eps}\\includegraphics[width=\\columnwidth]{f4.eps} \\caption{$I-$band luminosity profiles of NGC 24 (left) and $B-$band luminosity profile (Romanishin et al. 1983) of NGC 45 (right). \\label{lp}} \\end{figure*}  In order to get a better understanding of the mass distribution in NGC 24 and NGC 45, both optical and 21~cm radio observations were obtained and are presented in Section~\\ref{sec:obs}. In Section~\\ref{sec:photometry}, the luminosity profiles and the optical properties are discussed. After studying the \\hi content and its distribution in Section~\\ref{sec:neutrgas}, Section~\\ref{sec:kinematics} concentrates on the velocity field of these systems and presents their rotation curves. In Section~\\ref{sec:mass} the data are analyzed in terms of a two component model: a luminous (stellar$+$gaseous) disk and a dark halo represented by an isothermal sphere potential. Finally, the discussion in Sec.~\\ref{sec:discussion} is followed by a summary of the results in Sec.~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion} An optical and \\hi study of NGC 24 and NGC 45 has been presented. The main results are as follows:  (1) From the surface photometry of NGC 24 and NGC 45, it is found that while NGC 24 is on the faint side for normal galaxies, NGC 45 can be considered as a bona fide LSB galaxy. However, both galaxies have very similar absolute magnitudes $\\sim -17.4$.  (2) The \\hi distribution for NGC~24 and NGC~45 extends to $\\sim 1.3$ and 2.1$~{\\rm R_{Ho}}$ (respectively) at a level of $\\sim 10^{20}~{\\rm atoms\\,cm^{-2}}$. The \\hi extent of NGC~24 is very similar to the \\hi extent of the Sculptor late-type spiral galaxies NGC 247, NGC 300 and NGC 7793 (1.2, 1.5 and 1.4$~{\\rm R_{Ho}}$, respectively).  (3) The overall velocity fields of the two galaxies are very regular. NGC 45 exhibits a twist of the kinematical major axis, showing a $\\sim$25\\degr\\ variation of its major axis position angle as a function of radius. However, its disk does not have a classical warp since the inclination remains constant.  (4) The rotation curves derived from the velocity fields rise slowly and flatten at a velocity of $\\sim 110~{\\rm km\\,s^{-1}}$ and $\\sim 100~{\\rm km\\,s^{-1}}$ respectively. It extends out to $\\sim$11~kpc for NGC 24 and to $\\sim$17~kpc for NGC 45, which corresponds to $\\sim 7.5$ scale lengths in both cases.  The rotation curves, combined with the luminosity profiles, were used to study the mass distribution of NGC 24 and NGC 45. Using a best-fit model, the main results are :  (1) \\rml = 2.5, $r_c = 5.6~{\\rm kpc}$ and $\\sigma = 65~{\\rm km\\,s^{-1}}$ for NGC 24, and \\rml = 5.2, $r_c = 6.2~{\\rm kpc}$ and $\\sigma =55~{\\rm km\\,s^{-1}}$ for NGC 45.  (2) In both galaxies, the dark halo is the main contributor to the total mass of the galaxies with a contribution of more than $\\sim$80\\,\\% at the last measured point. Since the best-fit models are close to the maximum disk case, this can be considered a lower limit.  (3) The \\rml of 5.2 found for NGC 45 and 2.5 for NGC 24 are high when compared with the values predicted by stellar population synthesis models for galactic disks of same colors as NGC 24 and NGC 45. This result is similar to what is seen in other LSB galaxies. When a model for NGC 45 uses a \\rml of 2.0, which value is adopted from SPS models, it allows the galaxy mass to be entirely dominated by the dark component, but it also severely underestimates the velocity of the innermost points of the rotation curve.  This article is the first from a series that aims at measuring the shape of the mass density profile for the dark component of NGC 24 and NGC 45. When higher resolution optical kinematical \\ha\\ data obtained with Fabry-Perot interferometry become available, more accurate mass models of those galaxies will be presented."
    },
    "0609/astro-ph0609654_arXiv.txt": {
        "abstract": "{} {We present multiwavelength observations of the gamma-ray burst GRB 051028 detected by HETE-2 in order to derive its afterglow emission parameters and to determine the reason for its optical faintness when compared to other events.} {Observations were taken in the optical (2.0m Himalayan Chandra Telescope, 1.34m Tautenburg, 4.2m William Herschel Telescope) and in X-rays (Swift/XRT) between 2.7 hours and $\\sim$ 10 days after the onset of the event.} {The data can be interpreted by collimated emission in a jet with a typical value of $p$ = 2.4 which is moving in an homogeneous interstellar medium and with a cooling frequency $\\nu_{c}$ still above the X-rays at 0.5 days after the burst onset. GRB 051028 can be classified as a ``gray'' or ``potentially dark'' GRB. On the basis of the combined optical and {\\it Swift}/XRT data, we conclude that the reason for the optical dimness is not extra absorption in the host galaxy, but rather the GRB taking place at high-redshift. We also notice the very striking similarity with the optical lightcurve of GRB 050730, a burst with a spectroscopic redshift of 3.967, although GRB 051028 is $\\sim$ 3 mag fainter. We suggest that the bumps could be explained by multiple energy injection episodes and that the burst is intrinsically faint when compared to the average afterglows detected since 1997.  The non-detection of the host galaxy down to $R$ = 25.1 is also consistent with the burst arising at high redshift, compatible with the published pseudo-$z$ of 3.7 $\\pm$ 1.8.} {} ",
        "introduction": "The question whether a significant fraction of gamma ray bursts (GRBs) are intrinsically faint or true dark events remains unsolved (see Filliatre et al. 2005, Castro-Tirado et al. 2006 and references therein). For instance, \\object{GRB 000418} was detected in the near-IR (Klose et al. 2000) and it is one of the reddest (R-K = 4) together with  \\object{GRB 980329} (Reichart et al. 1999), \\object{GRB 030115} (Levan et al. 2006) and the recent \\object{GRB 050915A} (Bloom \\& Alatalo 2005). In most cases, it has been suggested that the cause of the reddening was dust extinction in the host galaxy. On the  other hand, \\object{GRB 021211} was found to be very dim at 24 hours, as a scaled-down version of \\object{GRB 990123} (Pandey et al. 2003).  With the launch of {\\it Swift} in Nov 2004, which has the ability to follow-up the events detected by the GRB detector onboard (BAT) or by other satellites like {\\it HETE-2} and {\\it INTEGRAL}, it is possible to zoom in on this population of optically faint events in order to disentangle their nature.  \\object{GRB 051028} was one of such event. It was discovered by {\\it HETE-2} on 28 Oct 2005, lying (90\\% confidence) on a 33\\arcmin $\\times$ 18\\arcmin error box centred at coordinates: RA (J2000) = 01$^{\\rm h}$48$^{\\rm m}$38\\fs6 Dec (J2000) = +47\\degr48\\arcmin30\\farcs0 (Hurley et al. 2005). The burst started at T$_{0}$ = 13:36:01.47 UT and a value of  T$_{90}$ = 16~s is derived, putting it in the ``long-duration'' class of GRBs. It had a fluence of 6 $\\times10^{-7}$ erg cm$^{-2}$ in the 2-30 keV range and 6 $\\times10^{-6}$ erg cm$^{-2}$ in the 30-400 keV range (Hurley et al. 2005). This event was also detected by {\\it Konus}/WIND in the 20 keV - 2 MeV range, with a duration of  $\\approx12$~s, a fluence of (6.78$_{-1.08}^{+0.61}$)   $ \\times10^{-6}$ erg cm$^{-2}$ in the 20 keV - 2 MeV  range and a peak energy E$_{p}$ =  298$_{-50}^{+73}$ keV (Golenetskii et al. 2005). {\\it Swift}/XRT started to observe the field $\\sim$ 7.1 hours after the event and detected the X-ray afterglow 5.2\\arcmin ~away from the center of the initial ${\\it HETE-2}$ error box (Racusin et al. 2005).  We report here results of multi-wavelength observations in optical and X-ray waveband and discuss the reasons for the apparent optical faintness of GRB 051028 in comparison with other bursts.    \\begin{table*} \\begin{center} \\caption{Journal of optical observations of the \\object{GRB 051028} field.} \\begin{tabular}{@{}lcccc@{}}  Date of 2005 UT & Telescope/ & Filter & Exposure Time & Magnitude\\\\ (mid exposure)  & Instrument &        &    (seconds)  &          \\\\ \\hline Oct 28,  16:18   & 2.0 HCT (HFOSC)  & $R_{c}$ &     300   &  20.63$\\pm$0.04 \\\\ Oct 28,  16:32   & 2.0 HCT (HFOSC)  & $R_{c}$ &     300   &  20.72$\\pm$0.05 \\\\ Oct 28,  16:47   & 2.0 HCT (HFOSC)  & $R_{c}$ &     300   &  21.14$\\pm$0.07 \\\\ Oct 28,  17:03   & 2.0 HCT (HFOSC)  & $R_{c}$ &     300   &  21.27$\\pm$0.07 \\\\ Oct 28,  17:43   & 1.34 Taut (CCD)  & $R_{c}$ &   1\\,080  &  21.23$\\pm$0.13 \\\\ Oct 28,  17:47   & 2.0 HCT (HFOSC)  & $R_{c}$ &     300   &  21.17$\\pm$0.08 \\\\ Oct 28,  21:42   & 4.2 WHT (PFC)    & $R$     &     300   &  21.97$\\pm$0.05 \\\\ Oct 29,  05:47   & 4.2 WHT (PFC)    & $R$     &     120   &  22.8$\\pm$0.3   \\\\ Oct 29,  20:15   & 4.2 WHT (PFC)    & $R$     &     720   &    $>$23.7      \\\\ Oct 31,  22:14   & 4.2 WHT (PFC)    & $R$     &   2\\,700  &    $>$25.1      \\\\ \\hline Oct 28,  16:25   & 2.0 HCT (HFOSC)  & $I_{c}$ &     300   &  19.79$\\pm$0.11 \\\\ Oct 28,  16:39   & 2.0 HCT (HFOSC)  & $I_{c}$ &     300   &  19.94$\\pm$0.06 \\\\ Oct 28,  16:55   & 2.0 HCT (HFOSC)  & $I_{c}$ &     300   &  20.29$\\pm$0.09 \\\\ Oct 28,  17:09   & 1.34 Taut (CCD)  & $I_{c}$ &   1\\,080  &  20.5 $\\pm$ 0.3 \\\\ Oct 28,  17:10   & 2.0 HCT (HFOSC)  & $I_{c}$ &     300   &  20.38$\\pm$0.08 \\\\ Oct 28,  17:55   & 2.0 HCT (HFOSC)  & $I_{c}$ &     300   &  20.35$\\pm$0.09 \\\\ Oct 28,  19:12   & 1.34 Taut (CCD)  & $I_{c}$ &   1\\,800  &  20.67$\\pm$0.23 \\\\ Oct 28,  20:33   & 1.34 Taut (CCD)  & $I_{c}$ &   3\\,600  &  20.75$\\pm$0.13 \\\\ Oct 28,  22:50   & 1.34 Taut (CCD)  & $I_{c}$ &   5\\,400  &  21.16$\\pm$0.16 \\\\ \\hline Oct 28,  18:28   & 1.34 Taut (CCD)  & $V$     &   1\\,080  &  22.08$\\pm$0.20 \\\\ \\hline \\label{tabla1} \\end{tabular} \\end{center} \\end{table*} ",
        "conclusions": "\\label{conclusiones}  We have presented multiwavelength observations of the long duration GRB 051028 detected by {\\it HETE-2} between 2.7 hours and $\\sim$ 10 days after the event. The X-ray afterglow of GRB 051028 can be compared to other GRB afterglows in the sense that its flux at 11 hours is typical, i.e., one can assume that the burst has occurred on a classical  n $\\sim$ 1 cm$^{-3}$ environment. The optical afterglow, on the other hand, is dim at a similar epoch (and comparable for instance to {\\object {GRB 030227}}, Castro-Tirado et al. 2003). We also noticed the remarkable similarity to the optical afterglow of {\\object {GRB 050730}}, a burst lasting $\\sim$ 10 times longer with comparable gamma-ray fluence\\footnote{The scarcity of the available X-ray data for GRB 051028 does not allow to make a straigth comparison with respect to the GRB 050730 X-ray afterglow.} at $z$ = 3.967 (see Pandey et al. 2006 and references therein). This indicates that the faintness of the optical emission is not due to a low-density environment as in the case of some short GRBs, such as GRB 050509b (Castro-Tirado et al. 2005). Instead, we propose that GRB 051028 occurred in a faint galaxy (with $R$ $>$25.1) at a high redshift consistent with the pseudo-$z$ = 3.7 $\\pm$ 1.8.  Thanks to the extraordinary repointing capabilities of {\\it Swift}, the accurate localisations for future events and the correspon\\-ding multiwavelength follow-up will shed more light on the origin of this faint optical afterglow population."
    },
    "0609/astro-ph0609524_arXiv.txt": {
        "abstract": "We discuss a set of tests which confront observations of cooling compact objects and theories of their thermal evolution. As an example we apply the recently developed $\\mathrm{Log\\, N}$-$\\mathrm{Log\\, S}$ test of compact star cooling theories to hybrid stars with a color superconducting quark matter core, we also apply and discuss other existing tests. While there is not yet a microscopically founded superconducting quark matter phase which would fulfill constraints from cooling phenomenology, we explore the hypothetical 2SC+X phase and show that the magnitude and density-dependence of the X-gap can be chosen to satisfy a set of tests: the temperature~--~age ($\\mathrm{T}$-$\\mathrm{t}$) test, the $\\mathrm{Log\\, N}$-$\\mathrm{Log\\, S}$ test, the brightness constraint, and the mass spectrum constraint. Some recent modifications of the population synthesis model used to obtain the $\\mathrm{Log\\, N}$-$\\mathrm{Log\\, S}$ distribution are briefly discussed. In addition, we propose to use the age-distance diagram as a new tool to study the local population of young isolated neutron stars. ",
        "introduction": "Compact objects known under the common name {\\it neutron stars} belong to the most facinating places in the Universe as the interest of different disciplines of science is focussed on them. In the present contribution we will elucidate how astronomical observations of cooling neutron stars (NSs) provide a tool to study the behaviour of matter at high density. Here astronomy meets QCD.  Sometimes the term {\\it neutron star} is used just as a common name for a large class of objects invented by the fantasy of theorists. They can be ``normal'' NSs, i.e. compact objects made mainly of hadrons even at the center -- {\\it hadron stars}. But there are other interesting possibilities. Central parts of NSs could contain more exotic forms of dense matter like a pion condensate or hyperonic matter. Finally, as a result of the deconfinement transition hadrons get dissolved into their quark constituents due to the Pauli principle: compact stars develop a quark matter core. In this case the NS is called a {\\it hybrid star}. Eventually, the whole body (maybe except a tiny envelope) is converted into quark matter. In this case we speak of a {\\it quark star}. So, a large variety of interesting possibilities appears. One of the very few ways to test which particular scenario for the internal structure  is realized in Nature is to study the thermal evolution of compact objects. This paper is devoted to the problem of testing theories of NS cooling by confronting results of computations with observational data.  This contribution is mostly based on the paper \\cite{pgb2006}. Still, some new results and discussions are given. At first, we describe different tests of cooling curves. Then we briefly describe the model for the cooling of hybrid stars to which we intend to apply the tests. In some details we dwell on the population synthesis scenario used here to produce the $\\mathrm{Log\\, N}$-$\\mathrm{Log\\, S}$ distributions. After this we apply the test to a set of cooling curves of hybrid stars. Finally, we discuss further improvements in the scenario and present our conclusions. ",
        "conclusions": "In conclusion, we want to underline that the idea of the conferences like QUARKS-2006 is a very fruitful one, as these meetings are the places where theoretical physicists and astrophysicists can have joint discussions: QCD meets astronomy again. Now it is obvious that there are theoretical models which can be tested only in ``celestial laboratories''. Neutron stars definitely are such labs.  Observations of cooling NSs provide a rare opportunity to study particular parts of the QCD phase diagram. We hope that the development of new tools to confront theory of cooling of compact objects with observations can help to achieve progress in both: astrophysics and  theoretical physics.  \\bigskip  \\noindent {\\bf Acknowledgments} The work of SP is supported by the RFBR grant 06-02-16025, by the Cariplo foundation and by the ``Dynasty'' Foundation. HG was supported in part by DFG under grant 436 ARM 17/3/05. SP thanks the Organizers for partial support."
    },
    "0609/astro-ph0609712_arXiv.txt": {
        "abstract": "We study the gravitational Faraday rotation, on linearly polarized light rays emitted by a pulsar, orbiting another compact object. We relate the rotation angle to the orbital phase of the emitting pulsar, as well as to other parameters describing its orbit and the orientation of the angular momentum of the binary companion. We give numerical estimates of the effect for the double-pulsar system PSR J0737-3039, and we note that the expected magnitude is exceedingly small, making the effect unlikely to be observed with present technology. It is however interesting per se, since in this phenomenon, gravito-magnetism plays a leading role, unlike what happens, for instance, when studying light bending or gravitational time delay, where it appears as a correction to the gravito-electric contribution. Also, we envisage the possibility that this effect could be relevant, at least in principle, for a pulsar orbiting a non charged black-hole. ",
        "introduction": "\\label{sec:intro}  The physical properties of the recently discovered double-pulsar system PSR J0737-3039 (\\cite{burgay03}, \\cite{lyne04}) make this system a rare laboratory for testing relativistic gravity. Namely, the system is composed of a 22-ms pulsar, PSR J0737-3039 A, in a slightly eccentric ($e \\simeq 0.088 $) $2.4 h$ orbit, together with PSR J0737-3039 B, whose period is 2.7 s: thanks to the high mean orbital speed ($v \\simeq 0.001 c$) this system is highly relativistic, and allows the detection of four post-Keplerian parameters, thus accurately constraining the masses of the two stars and providing also stringent tests of General Relativity (GR) (\\cite{kramer05}).  Actually, excluding cosmological issues, GR has passed all observational tests with excellent results (\\cite{will05}). However, we must remember that most of the tests of GR come from Solar System experiments, where the gravitational field is in the \"weak\" regime. On the other hand, it is expected that deviations from GR can occur in the \"strong\" field regime: hence, the Solar System experiments are inadequate to this end. On the contrary, the strong gravitational field is best tested by means of pulsars. Among the post-Newtonian effects that should be relevant in strong field regime (\\cite{will}), there are the so called \\textit{gravito-magnetic} (GM) effects, which arise from the rotation of the sources of the gravitational field: these effects are generally small, since the gravitational coupling with the angular momentum of the source is much weaker than the coupling with mass alone (the so called \\textit{gravito-electric} interaction), and this makes their detection very difficult (see \\cite{mashh1,ruggiero02}).  In previous papers (\\cite{tartaglia05,ruggiero05}), we pointed out that binary pulsar systems in general and, in particular, the double-pulsar PSR J0737-3039, could be very interesting for testing the GM effects. The relevance of GM effects in binary pulsar systems was investigated in the past and also more recently by many authors, who also pointed out the difficulties that the detection of these effects must face (\\cite% {doro95,laguna97,wex99,kopeikin02,rafikov06a,rafikov06b}), mainly due to their smallness, which is overwhelmed by other dominant effects.  Here, we would like to study a physical phenomenon where the gravito-magnetic interaction has a leading role, and which can be tested, at least in principle, in binary pulsar systems, i.e. systems where a pulsar is orbiting a compact object, such as another pulsar, as in PSR J0737-3039, or, hopefully, a black hole. The phenomenon we are going to deal with is the gravitational Faraday rotation.  The polarization vector of an electromagnetic (e.m.) wave propagating in a gravitational field undergoes a change of direction, as a result of the deflection of e.m. signals, and, in addition a rotation around the propagation vector. As for the latter, in particular, it has been studied and discussed by many authors, in the context of GR (see, for early works the paper by \\cite{plebanski60} and, for more recent references the paper by \\cite{sereno05a}). Actually this gravitational rotation of the plane of polarization, for linearly polarized light rays, is the gravitational analogue of the electro-magnetic Faraday effect, that is the rotation which a light ray undergoes when it propagates in a plasma under the action of a magnetic field, the role of the magnetic field being played, in this case, by the gravito-magnetic field. It is possible to show that, for astrophysical interesting cases, a spherical and non rotating source of gravitational field does not cause a net rotation of the polarization vector (\\cite{plebanski60,dyer92,faraoni93}), while things change when rotating sources come into play. Indeed, the gravitational Faraday rotation is a purely gravito-magnetic phenomenon (\\cite{sereno05a}). Actually, also the case of rotation is somewhat subtle, since in the lowest order of the weak field approximation, \\cite{plebanski60} concluded that the rotation of the polarization plane occurs only when the light rays penetrate into a rotating distribution of matter. More accurate calculations, that take into account the deflection of the light rays and higher-order terms in the weak field approximation, show that rotation occurs outside the horizon of a rotating (Kerr) black hole (\\cite{ishi88,nouri99,sereno05a}).  The phenomena we try and study here could be particularly relevant in the case of X-ray emission from accretion disks or other thick material envelopes surrounding rotating magnetized compact objects and specifically black holes. Indeed the intrinsic birefringence of the plasma in the magnetic field, together with the general relativistic effects due to the curvature in proximity of horizons, can be  the primary source of  polarization and of the rotation of the polarization plane of the emitted short wavelength radiation (\\cite{connors,broderick}).  Under the hypothesis that, to a reasonable order of approximation, the space-time around a rotating compact object (i.e. a pulsar, or a black hole) can be described by the Kerr metric, here we investigate the gravitational Faraday rotation induced on the light rays emitted by a pulsar orbiting around it. We study the effect and give estimates of its magnitude, and evaluate the possibility of detecting it.  The paper is organized as follows: after briefly reviewing the basics of propagation of e.m. signals in a gravitational field that lead to the gravitational Faraday rotation in Sec. \\ref{sec:em}, we apply the formalism to the case of the signals emitted by a pulsar orbiting a compact object in Sec. \\ref{sec:gravfar}. The estimate of the magnitude of the effect, together with the evaluation of the possibility of detecting it, is discussed in Sec. \\ref{sec:disc}, while the conclusions are summarized in Sec. \\ref{sec:conc} ",
        "conclusions": "\\label{sec:conc}  We have studied the gravitational Faraday rotation on light rays emitted by a pulsar, orbiting another compact object, such as a pulsar or a black-hole, whose gravitational field, to a reasonable order of approximation, can be described by the Kerr metric, of which we have considered the suitable weak-field limit.  In this phenomenon, gravito-magnetism plays a leading role, unlike what happens, for instance, when studying light bending or gravitational time delay, where it appears as a correction to the gravito-electric contribution.  The effect is a purely geometric one, contrary to the electromagnetic Faraday rotation, which depends on the frequency of the light rays; furthermore, it depends on the component of the angular momentum of the source of the gravitational field along the line of sight. So, when the latter is not null, a net rotation of the polarization plane is expected.  We have obtained a formula which relates the rotation angle to the orbital phase of the emitting pulsar, as well as to the parameters describing its orbit and to the orientation of the angular momentum of the binary companion.   The numerical estimates of the magnitude of this effect, in the context of the double-pulsar system PSR J0737-3039, show that at radio wavelengths it is overwhelmingly dominated by the electromagnetic Faraday rotation. Nevertheless, as a result of it's achromatic nature, gravitational Faraday rotation will dominate electromagnetic Faraday rotation at X-ray energies.  However, in practice, due too its exceedingly small nature, it is unlikely to be observed in pulsar systems.  Anyway, this effect could be of some relevance for  pulsar-black hole binaries, or accreting black holes (\\cite{connors}).  In such  systems, peculiar modulation of polarization orientation during the orbital motion may be evidence that an unseen companion is a black hole."
    },
    "0609/astro-ph0609238_arXiv.txt": {
        "abstract": "When faced with the task of constraining a galaxy's potential given limited stellar kinematical information, what is the best way of treating the galaxy's unknown distribution function (DF)?  Using the example of estimating black hole (BH) masses, I argue that the correct approach is to consider all possible DFs for each trial potential, marginalizing the DF using an infinitely divisible prior. Alternative approaches, such as the widely used maximum penalized likelihood method, neglect the huge degeneracies inherent in the problem and simply identify a single, special DF for each trial potential.  Using simulated observations of toy galaxies with realistic amounts of noise, I find that this marginalization procedure yields significantly tighter constraints on BH masses than the conventional maximum-likelihood method, although it does pose a computational challenge which might be solved with the development of a suitable algorithm for massively parallel machines.  I show that in practice the conventional maximum-likelihood method yields reliable BH masses with well-defined minima in their $\\chi^2$ distributions, contrary to claims made by Valluri, Merritt \\& Emsellem. ",
        "introduction": "A fundamental application of stellar dynamics is using observations of a galaxy's kinematics to constrain its potential~$\\psi(\\b x)$.  The galaxy is normally assumed to be collisionless and in a steady state, so that the dynamics of any population of stars are completely described by its phase space distribution function (DF), $f(\\b x, \\b v)$, which is the probability density of finding a star in the small volume of phase space around $(\\b x,\\b v)$.  By Jeans' theorem, this DF can depend on $(\\b x,\\b v)$ only through the integrals of motion of the (unknown) potential.  Most approaches to this task begin by considering the simpler problem of constraining $f$ given the observed data and some trial $\\psi$. The infinite-dimensional DF is parametrized by a finite sum of delta functions \\citep[e.g.,][]{schw79} or a truncated basis function expansion \\citep[e.g.,][]{dejonghe89,saglia00}, and the DF parameters are adjusted to optimize the fit to the observations.  If no set of parameters yields an acceptable fit while simultaneously representing a DF that is everywhere non negative, then the assumed potential can be ruled out.  This basic idea can be refined further.  \\citet{RT88} pointed out that naive application of this method will yield unrealistically spiky DFs, leading them to advocate the use of entropy (or something similar) as a regularizer.  This idea of regularizing the resulting DFs was made more explicit by \\citet{merritt93}, who cast the problem as one of finding the maximum penalized log-likelihood, ${\\cal L}'\\equiv-{1\\over2}\\chi^2+\\lambda P[f]$, which allows a trade off between goodness of fit, as measured by $\\chi^2$, and smoothness, as measured by the penalty function $P$. The choice of penalty function and the value to use for the tradeoff parameter~$\\lambda$ are arbitrary and subjective. Given a range of trial potentials, one finds the maximum penalized likelihood for each and then uses normal statistical methods to make statements about how well constrained the potential is.  This general approach has become the method of choice in stellar-dynamical searches for supermassive black holes (hereafter BHs) in galaxy centres, with choices of penalty function ranging from the mean-square second derivative of the DF \\citep[e.g.,][]{vdm98,capp02} or entropy \\citep[e.g.,][]{g03,sil05} through to models in which no regularization whatsoever has been applied \\citep[e.g.,][]{vdm98,ryan06}.  I refer to these as ``maximum-likelihood'' or ``maximum-penalized likelihood'' methods.    Another way to look at the problem of constraining potentials is to treat it as a straightforward mathematical inverse problem. \\citet{dm92} considered the case of constraining the potential of a spherical galaxy given perfect knowledge of its projected DF in the form of its line-of-sight velocity profiles (hereafter VPs) $\\L(R;\\vp)$.  By using a set of higher-order Jeans equations they showed that, given the potential, the DF $f(\\E,J^2)$ is completely determined by its projected VPs.  Constraining the potential, however, proved less amenable to their methods, but they found that choosing a potential too far from the true one would yield moments (and therefore DFs) that became negative, ruling out that potential.  Of course, one never has perfect knowledge of the full projected DF. More recently, \\citet[][hereafter VME04]{vme04} investigated the slightly less idealized problem of constraining the BH mass in a toy axisymmetric galaxy given noiseless measurements of a restricted number of (modified) moments of its VPs averaged over a number of spatial bins on the sky.  They showed that even when the potential has just one free parameter -- the BH mass -- there are many different potentials that can fit the available kinematics almost perfectly, even when the central spatial resolution of the kinematics is much finer than the BH's sphere of influence.  One thing that all these methods have in common is that they consider just one DF for each trial potential.  This is fine for the idealized case considered by \\citet{dm92}: given perfectly resolved, noiseless projected VPs of a spherical galaxy there is a unique DF for any assumed potential, even though this DF may not be non-negative everywhere.  But when the available data have finite spatial and velocity resolution, there will in general be many perfectly sensible, non-negative DFs that yield equally good fits to the data, and even more DFs producing fits that are only slightly worse.  Intuitively, one might expect that the more such DFs a potential admits, the more likely it is.  The purpose of the present paper is to revisit the the problem of constraining BH masses from a thoroughly Bayesian perspective, showing how it naturally incorporates this intuitive notion of counting up DFs.  To illustrate the ideas, I use simulated observations of some idealized spherical toy galaxies described in Section~2, modelling them under the assumptions given in section~3.  In section~4 I test how well the conventional maximum likelihood method recovers BH masses and counter some of the more pessimistic conclusions of VME04. Section~5 presents a Bayesian approach to the problem, which overcomes some of the inconsistencies of the maximum likelihood method. Finally, section~6 sums up and discusses the implications for BH masses in real galaxies. ",
        "conclusions": "Most BH mass estimates come from the conventional maximum-likelihood method.  They fit observations using models with specially constructed, unrealistically spiky DFs, yielding implausibly good fits to the observations.  Results from the toy galaxies considered here suggest that the (unpenalized) maximum-likelihood method nevertheless does yield reliable BH masses, but with overly pessimistic error estimates.  I confirm VME04's result that a wide range of BH masses can yield perfect fits to finite-resolution, noiseless observations.  Contrary to their somewhat speculative arguments, however, I show that one does not expect to find flat-bottomed $\\chi^2$ distributions in practice, unless one is blessed with very high signal-to-noise observations (figure~\\ref{fig:varynoise}) and fits only the low-order shapes of the VPs.   Although the maximum-likelihood method yields reliable BH masses, it is flawed because it considers only one DF for each BH mass.  The remedy is to consider all possible DFs for each potential, weighting each one by a suitably chosen prior.  This significantly improves the constraints on the BH mass, since the closer a trial potential comes to the true potential, the greater the number of non-negative DFs that are consistent with the observations.    \\subsection{Open questions and future work}  \\subsubsection{Choice of prior} An open question is to what extent these results depend on the choice of prior weights (equ.~\\ref{eq:mu}), and more generally on the use of the prior~(\\ref{eq:incontinent}), which is just one of many possible infinitely divisible distributions.  A very general argument \\citep{kingman} shows that drawing random realizations from most ID priors will yield spiky, uncorrelated distributions.  This is just what one expects if dealing with galaxies at the level of individual stars, but the present models are  far from this level of detail and one might plausibly expect some degree of correlation among neighbouring cells in phase space.  Making this idea quantitative is difficult, however.  \\subsubsection{ Computational scheme} It would be straightforward in principle to apply the ideas presented here to axisymmetric or triaxial galaxy models.  One could simply adopt the prior~(\\ref{eq:incontinent}), using the scheme described by \\citet{jens} to calculate the volumes used to assign the prior weights $\\mu$.  In practice, however, calculating the marginalized likelihood $\\pr(D|\\psi)$ will probably be very difficult.  The DFs of axisymmetric galaxies are three-integral, which means that, in order to ensure a fine enough discretization, one has to use many more DF components than for two-integral spherical models, with the Markov Chain Monte Carlo procedure used in section~\\ref{sec:marginalization} taking correspondingly more iterations to converge.  On the other hand, it is very likely that there are much more efficient methods than the combination of Gibbs sampling and thermodynamic integration used here.  \\subsubsection{More immediate problems} Before applying this method to real galaxies though, it is probably worth addressing the following more tractable problems first: \\begin{enumerate} \\item VPs are extracted from spectra, which suffer from poorly understood systematic errors \\citep{ryan06}. \\item For real observations, neither individual VP velocity bins nor (surprisingly) Gauss--Hermite coefficients are independent \\citep{ryan06}. \\item Most models of axisymmetric galaxies make an ad hoc assumption about the galaxy's three-dimensional light distribution $j(R,z)$, despite the fact that there are many $j(R,z)$ consistent with a given surface brightness distribution \\citep[e.g.,][]{kochanekrybicki}. Neglect of this degeneracy can lead to incorrect inferences about the galaxy's orbit structure \\citep{depropaper} which are likely to affect BH mass estimates. \\item The widely used $\\Delta\\chi^2$ criteria for obtaining uncertainties on BH masses are based on the assumption that $\\chi^2(M_\\bullet,\\Upsilon)$ is close to quadratic, which does not necessarily hold in practice. \\end{enumerate}"
    },
    "0609/astro-ph0609395.txt": {
        "abstract": "Orbital angular momentum ($J_{o}$), systemic mass ($M$) and orbital period ($P$) distributions of chromospherically active binaries (CAB) and W Ursae Majoris (W UMa) systems were investigated. The diagrams of $\\log J_{o}-\\log P$, $\\log M-\\log P$ and $\\log J_{o}-\\log M$ were formed from 119 CAB and 102 W UMa stars. The $\\log J_{o}\u0096-\\log M$ diagram is found to be most meaningful in demonstrating dynamical evolution  of binary star orbits. A slightly curved borderline (contact border) separating the detached and the contact systems was discovered on the $\\log J_{o}\u0096-\\log M$ diagram. Since orbital size ($a$) and period ($P$) of binaries are determined by their current $J_{o}$, $M$ and mass ratio $q$, the rates of orbital angular momentum loss ($d\\log J_{o} /dt$) and mass loss ($d\\log M/dt$) are primary parameters to determine the direction and the speed of the dynamical evolution. A detached system becomes a contact system if its own dynamical evolution enables it to pass the contact border on the $\\log J_{o}-\\log M$ diagram. Evolution of $q$ for a mass loosing detached system is unknown unless mass loss rate for each component is known. Assuming $q$ is constant in the first approximation and using the mean decreasing rates of $J_{o}$ and $M$ from the kinematical ages of CAB stars, it has been predicted that 11, 23 and 39 cent of current CAB stars would transform to W UMa systems if their nuclear evolution permits them to live 2, 4 and 6 Gyrs respectively. ",
        "introduction": "The angular momentum loss in close binaries is known to be replenished from the reservoir of orbital angular momentum (OAM) of the system by tidal locking. Since tidal locking operates as a mechanism to draw angular momentum from the orbit, the orbits of spin-orbit coupled binaries are then forced to shrink although actual angular momentum loss occurs at one or both components due to magnetically driven winds, also called magnetic braking \\citep{Sch59,K67, M68}. Shrinking orbits, then, require orbital periods to be decreasing. Authors such, \\cite{Huang66, OS70, van79, VR80, R82, M84, GB88, MV91,  MV96, S95, D99} all believed this mechanism is a main route to form W UMa  systems from the binaries initially detached with comparable periods. Showing high level of chromospheric and magnetic activity, RS CVn like objects including binary BY Dra type systems, shortly called CAB, are primary candidates to be the progenitors of the contact systems.  \\cite{K04} recently presented observational evidences of mass loss and orbital period decrease by comparing the total mass and period histograms between the kinematically young (age 0.95 Gyr) and older (age 3.86 Gyr) CAB sub samples. Secular orbital period decrease among CAB is indicated further by the fact that the shorter orbital period systems are older than the longer orbital period systems \\citep[see Table 5 of ][]{K04}. Moreover, 5.47 Gyr kinematical age for the field W UMa stars determined by \\cite{B05} appears consistent with the scenario. Consequently, the mean difference of 1.61\u0096-Gyr age between field W UMa and field CAB could be interpreted as a mean lifetime of contact stages. Not all contact systems have to be formed from detached progenitors. This is because \\cite{B05} have found that about 20 per cent of contact binary sample have ages less than 0.6 Gyr. Such a young age, being much less than the mean life time (1.61--Gyr) of contact stages, does not permit pre-contact detached phases. A Pre--contact detached stage usually has a duration up to several Gyr, e.g. 3.86 Gyr is an average. Therefore, those young W UMa must have been formed directly in the beginning of the main sequence or during the pre--main sequence contraction phase. Apparently, both mechanisms with and without pre-contact phases actively operate.  Recently, \\cite{D06} have determined a first order relative decreasing rates of OAM, systemic mass and orbital period for the detached CAB as $\\dot J_{o}/J_{o}=-3.48\\times 10^{-10}~yr^{-1}$, $\\dot M/M=-1.30 \\times 10^{-10}~yr^{-1}$ and $\\dot P/P=-3.96\\times 10^{-10}~yr^{-1}$ by using kinematical ages of sub samples formed according to OAM ranges. Mass loss, OAM loss and orbital period decrease driving one another appear to be confirmed by observational data. Nevertheless, it has been known that an isotropic mass loss from the surface(s) of one/or both component(s), despite it means OAM loss, forces the orbit to be bigger and makes the orbital period larger \\citep{P85}. It has also been shown by \\cite{D06} that not all rates of OAM loss would cause an orbit to shrink. Only if there is an amplification mechanism, such as magnetic loops forcing magnetized plasma to co--rotate, so mass loss at the Alfen radius carries away more angular momentum. If the amplification parameter, $\\bar A=(dJ_{o}/dM)/(J_{o}/M)$, is bigger than 5/3, then  the binary orbit would shrink and the orbital period would decrease. With the decreasing parameters given above, \\cite{D06} found $\\bar A=2.68$ as a first order observational mean.  Although, shrinking orbits and decreasing orbital periods are confirmed observationally and mean decreasing rates are available to support dynamical evolution, the details of the actual process are still unknown. Therefore, in this study, we investigate orbital angular momentum ($J_{o}$), systemic mass ($M$) and orbital period ($P$) distributions of CAB and W UMa systems together. ",
        "conclusions": "Total mass $M$ and orbital angular momentum OAM ($J_{o}$) are two basic physical quantities which determine sizes and periods of binary orbits. OAM has several different analytical formulae established with different combinations of mass ratio $q$, semi--major axis $a$, orbital period $P$ and systemic mass $M$. Because $M$ \\& $J_{o}$ are well known basic quantities and  mass loss ($dM$) \\& OAM loss ($dJ_{o}$) are again free parameters, from which the speed and direction of dynamical evolution can be inferred, the $\\log J_{o}-\\log M$ diagram is the best for studying the dynamical evolution of binary orbits.  There are limits in orbital dynamics. $d\\log J_{o}/d\\log M>3/2$ is for decreasing orbital sizes. $d\\log J_{o}/d\\log M>5/3$ is for decreasing orbital periods. if $d\\log J_{o}/d\\log M$ takes a value between 3/2 and 5/3, orbital sizes shrink despite corresponding orbital periods increase. According to mean dynamical evolution with $\\bar \\delta=d\\log J_{o}/d\\log M=2.68$, as provided by \\cite{D06}, the orbital sizes and periods of detached CAB systems are decreasing.  A contact border, which separates W UMa stars from CAB, is discovered on the $\\log J_{o}-\\log M$ diagram. So, One can predict OAM of a system on the contact border, called $J_{lim}$, which is determined only by $M$. Any OAM less than $J_{lim}$ implies a full contact for a binary. Similar to $J_{lim}$, the limiting period $P_{lim}$, and limiting orbital size $a_{lim}$ are computable using empirical formulae (11) and (12). Unlike $J_{lim}$, which is independent of $q$, $P_{lim}$ and $a_{lim}$ have a strong dependence on $q$.  Rather than computing the time left for each detached system to reach at the contact border, limiting borders corresponding 2, 4, and 6 Gyrs were computed by subtracting mean decreases ($\\Delta J$ and $\\Delta M$) of mean dynamical evolution from $J$ and $M$ values of the contact border. The CAB systems located between the contact border and those borders are the ones which are expected to be contact binaries within the time scales indicated.  A detached system heading towards the contact region must go trough mass transfer stages before becoming a contact binary. At any mass transfer stage, even if it is conservative, OAM $J_{o}$ is not conserved. For shrinking orbits, $J_{o}$ must decrease according to eq. (1). For conservative cases, the OAM difference between the two different size orbits of same system are to be compensated by spin AM. That is, $J_{o}$ still decreases while the total spin AM is increasing. If a detach system is sufficiently close to the contact border, during the first fast mass transfer stage, it has a changes to jump over into the region of contacts and to be one of them.  The traditional view of forming W UMa stars assumes that they only come from detached binaries of comparable periods. But, statistics available is convincing enough that there are very few such progenitors even to account for the low space density (0.2 per cent) of contact binaries in the solar neighbourhood \\citep{R02, R06}. Righteously \\cite{P06} says ``at this time the contact systems seem to appear OUT OF NOWHERE'' in order to emphasize the insufficiency of the traditional view. \\cite{P06}, therefore, consider mechanical tree--body orbital evolution, which is a similar formation model of close binaries in globular clusters \\citep{P03}, where the inner binary has relatively longer orbital period and has a better chance to evolve into a contact system after all \\cite{PR06} found that up to 50 per cent of W UMa binaries have companions.  The mean dynamical parameter $\\bar \\delta$ used in this study naturally contains affects of all mechanisms of OAM loss since its value is derived from stellar kinematics. Perhaps, it is not possible now to make a firm conclusion which mechanism is more efficient before a complete statistics and their analysis are available. Nevertheless, \\cite{B05} provided observational evidence that near contact detached binaries of few days orbital period are not the only source to form W UMa stars. They may also be born directly as contact systems. Moreover, it has also been found in this study that  considerable fraction of current detached CAB sample (39 per cent within 6 Gyrs) to have a potential to jump over the contact border on the $\\log J_{o}-\\log M$ diagram and become contact systems within a time scale of main-sequence life time. So, we encourage statistical studies to include these new findings after all contact binaries are rare with the local space density of just 0.2 per cent of the main sequence stars \\citep{R02, R06}."
    },
    "0609/astro-ph0609706_arXiv.txt": {
        "abstract": "We present results based on Spitzer Space Telescope mid-infrared (3.6-30\\,$\\mu$m) observations of the nearby IIP supernova 2005af. We report the first ever detection of the SiO molecule in a Type IIP supernova. Together with the detection of the CO fundamental, this is an exciting finding as it may signal the onset of dust condensation in the ejecta. From a wealth of fine-structure lines we provide abundance estimates for stable Ni, Ar, and Ne which, via spectral synthesis, may be used to constrain nucleosynthesis models. ",
        "introduction": "Mid-infrared (mid-IR) studies of supernovae (SNe) have been hampered by a combination of several factors: their intrinsic faintness at epochs of interest, the high background at mid-IR wavelengths, and strong terrestrial atmospheric absorption. With the sole exception of the extremely nearby (50\\,kpc) SN~1987A \\citep{roche:93, wooden:93}, this wavelength region has remained out of reach for SN studies. Yet, access to the mid-IR region is important for at least two reasons. First, a large number of fine-structure lines that are sensitive to the details of the explosion physics lie in the 5--40\\,$\\mu$m region.  These provide robust measurements of element abundances permitting direct comparisons with explosion models. Second, core-collapse SNe (CCSNe) have been proposed to be major sources of dust at high redshifts \\citep[e.g.,][]{tf:01}, but even in the local Universe, this hypothesis rests on little direct observational evidence. By monitoring the mid-IR evolution, where warm dust emits most strongly, we have a powerful means by which we can ascertain the epoch of dust condensation, and estimate the amount of fresh dust produced in the ejecta. Alternatively, via mid-IR measurements, we can study the development of an IR echo arising in the pre-existing dusty circumstellar medium, thereby constraining the pre-supernova evolution of the progenitor \\citep[e.g.,][]{meikle:06}.  The improvements in terms of sensitivity and spatial resolution afforded by the Spitzer Space Telescope ({\\it Spitzer}) finally allow us to probe SN behaviour using mid-IR diagnostics. To this end, we are pursuing a multi-epoch spectroscopic and photometric campaign of a significant sample of SNe of all types, within the framework of the Mid-IR Supernova Consortium (MISC). Fortuitously, since the launch of {\\it Spitzer}, at least half a dozen CCSNe have exploded within $\\sim$10\\,Mpc. Recently, we presented a first mid-IR study of the nearest (3.13\\,Mpc) of these, SN~2004dj \\citep{kotak:05}.  Here, we describe results for the Type IIP SN~2005af that highlight the diversity of this subgroup.  SN~2005af was discovered in the nearby edge-on spiral galaxy, NGC~4945, on 2005 Feb. 8.22 (UT) at V$=+12.8$ by \\citet{jp:05}. On the basis of a spectrum obtained four days later, \\citet{ff:05} classified it as a Type~IIP (plateau) supernova, at an epoch of about one month post-explosion. In what follows, we assume that the supernova exploded on 2005 Jan. 9, but note that the explosion epoch is probably uncertain by up to a few weeks. NGC~4945 is associated with the Centaurus group of galaxies, and estimates of its distance have ranged from 3.9 \\citep{dv:81} to 8.1\\,Mpc \\citep{baan:85}. In recent years, these values have tended to converge near the lower end of this range. Following the reasoning in \\citet[][and references therein]{bergman:92}, we adopt a distance of 3.9\\,Mpc. ",
        "conclusions": ""
    },
    "0609/astro-ph0609530_arXiv.txt": {
        "abstract": "Although damped Lyman alpha (DLA) systems are usually considered metal-poor, it has been suggested that this could be due to observational bias against metal-enriched absorbers. I review recent surveys to quantify the particular issue of dust obscuration bias and demonstrate that there is currently no compelling observational evidence to support a widespread effect due to extinction. On the other hand, a small sub-set of DLAs may be metal-rich and I review some recent observations of these metal-rich absorbers and the detection of diffuse interstellar bands in one DLA at $z \\sim 0.5$. ",
        "introduction": "At this conference on the metal-rich universe, a talk on damped Lyman alpha (DLA) systems may seem misplaced, since there is now a considerable literature on these absorption selected galaxies that describes them as metal-poor (e.g. Pettini 2004 and references therein). Faced with the almost universally sub-solar abundances of DLAs over all redshifts (the average metallicity of DLAs is $\\sim 1/30 Z_{\\odot}$ at $z \\sim 3$ and  $\\sim 1/10 Z_{\\odot}$ and $z \\sim 1$), questions have been raised concerning DLA selection techniques.  For example, are we missing a large fraction of metal-rich DLAs due to dust obscuration bias (Ostriker \\& Heisler 1984)?  Or are the bulk of the metals in absorbers below the canonical DLA column density threshold (P\\'eroux et al. 2006)?  In this contribution, I review the latest observations on the possibility of selection bias and discuss whether DLAs can ever be considered as metal-rich. ",
        "conclusions": ""
    },
    "0609/gr-qc0609094_arXiv.txt": {
        "abstract": "We study gravitationally collapsing models of pressureless dust, fluids with pressure, and the generalized Chaplygin gas (GCG) shell in (2+1)-dimensional spacetimes. Various collapse scenarios are investigated under a variety of the background configurations such as anti-de Sitter(AdS) black hole, de Sitter (dS) space, flat and AdS space with a conical deficit. As with the case of a disk of dust, we find that the collapse of a dust shell coincides with the Oppenheimer-Snyder type collapse to a black hole provided the initial density is sufficiently large.  We also find -- for all types of shell -- that collapse to a naked singularity is possible under a broad variety of initial conditions. For shells with pressure this singularity can occur for a finite radius of the shell. We also find that GCG shells exhibit diverse collapse scenarios, which can be easily demonstrated by an effective potential analysis. ",
        "introduction": "\\label{sec:intro}  Over the last few decades, general relativity in (2+1) dimensions has fascinated both field theorists and relativists because of its fertility as a test-bed for ideas about quantum gravity. One particular feature of interest is manifest when a negative cosmological constant is present.  Despite the fact that the spacetime geometry of this solution is an anti-de Sitter (AdS) spacetime, possessing negative constant curvature, a black hole can be present under a suitable choice of topological identifications \\cite{btz}. This solution has drawn much attention since its inception from a wide variety of perspectives \\cite{3dgravity}.  Shortly after the black hole solution was obtained, it was shown that it can be formed from a disk of pressureless dust undergoing gravitational collapse \\cite{mr} (the three-dimensional analogue of Oppenheimer-Snyder type collapse), generalizing earlier results that suggested matter could collapse to form conical singularities \\cite{gak}. Further study on this subject has been carried out from several viewpoints, including the formation of a black hole from colliding point particles \\cite{mat} and the more recent demonstration of critical phenomena in the context of collapse \\cite{3dcollapse}. These results are consistent with other results in four dimensions as well as results in two dimensions \\cite{mr2}.  Recently, a cosmological model of a (generalized) Chaplygin gas (GCG) was introduced as a possibile explanation of the present acceleration of the universe, the existence of dark energy, and the unification of dark energy and dark matter \\cite{kmp,btv,bbs}. Historically its original motivation was to account for the lifting force on a plane wing in aerodynamics \\cite{chap}. Afterwards, the same equation of state was rediscovered in the context of aerodynamics \\cite{tsien,karman}. A more interesting feature of this gas was recently renewed in an intriguing connection with string theory, insofar as its equation of state can be obtained from the Nambu-Goto action for $d$-branes moving in a $(d+2)$-dimensional spacetime in the light-cone frame \\cite{bh}. In addition, it has been shown that the Chaplygin gas is, to date, the only fluid that admits a supersymmetric generalization \\cite{hjp}; the relevant symmetry group was described in ref. \\cite{bj}. Moreover, further theoretical developments of the GCG were given in terms of cosmology and astrophysics \\cite{coas}. Inspired by the fact that the Chaplygin gas has a negative pressure, violating the energy conditions (in particular the null energy condition (NEC)), traversable wormhole solutions were found in four dimensions \\cite% {lobo}.  It is natural to ask whether or not a black hole can be formed from gravitational collapse of this gas in a finite collapse time. Much of the work on black hole formation deals with pressureless dust collapse; collapse of  this kind of exotic fluid to black holes so far has not received much treatment. Recent work \\cite{dc} \\ involved investigation of spherically symmetric clouds of a collapsing modified Chaplygin gas in four dimensions, where it was shown that it always leads to the formation of a black hole.  In this paper, we investigate some gravitational collapse scenarios of shells with a variety of equations of state, including the GCG shell. To set the stage we first consider the collapse of a shell of pressureless dust. In  dust collapse scenarios the evolution of the system is obtained by matching the inside and outside geometries using the junction conditions \\cite{isr,chase,poi}, \\begin{equation} \\label{eq:junctions} [g_{ij}]=0,~~[K_{ij}]=0, \\end{equation} where $[h] \\equiv h_{+} - h_{-}$ and ($+$) and ($-$) represent exterior and interior spacetimes, respectively. However for shells with pressure the junction condition for the extrinsic curvature in eq. (\\ref{eq:junctions}) is no longer valid, since there is a nonvanishing surface stress-energy on the boundary of the shell to take into account.  The main result of our investigation is that gravitational collapse in (2+1) dimensions does not necessarily lead to black hole formation for any of the fluid sources we study. The end points of collapse depend on the initial conditions, and can lead to either a black hole or the formation of a singularity and a Cauchy horizon\\footnote{This latter point has been overlooked in previous studies \\cite{gak}.}. This singularity is characterized by the onset of a divergent stress energy in the shell, whose intrinsic Ricci scalar also diverges in finite proper time for observers comoving with the shell.  For pressureless dust the singularity develops when the shell collapses to zero size. However for shells with pressure the singularity develops at some nonzero size characterized by the equation of state.  A similar scenario holds for the GCG shell. We also find that collapse is not the only possibility, but that shells can also expand out to infinity, possibly with a bounce depending on the initial conditions. Our results are consistent with earlier work on shell collapse in (2+1) dimensions \\cite{Olea,Steif}, generalizing them to include a more detailed analysis of collapse to naked singularities, and to situations in which a more general relationship between density and pressure is assumed.  The outline of our paper is as follows. In section \\ref{sec:shell}, we briefly present a formulation of the shell collapse and obtain the evolution equation for the dust shell radius. In section \\ref{sec:dustshell}, the gravitational collapses of pressureless dust shell are studied and compared to the result of dust cloud collapse in \\cite{mr}. In section \\ref{sec:pressure}, we study a collapse of a shell with an arbitrary pressure with no loss of generality. In section \\ref{sec:GCG}, the collapse of GCG shell is studied and some possible collapse conditions are found. Finally, we shall summarize and discuss our results in section \\ref{sec:discussions}. We consider the construction of some relevant Penrose diagrams and some basic properties of Jacobian elliptic functions in appendices. ",
        "conclusions": "\\label{sec:discussions}  Perhaps the most intriguing result of this paper is that shell collapse in $(2+1)$-dimensional gravity can violate --albeit somewhat mildly -- cosmic censorship for a broad range of initial conditions, whether we have pressureless dust shells, shells with pressure, or GCG shells. The situation is markedly different from that of a scalar field \\cite{scalarcollapse}, in which either a black hole is formed or the scalar field oscillates indefinitely without collapse. Here, as the shell collapses its density (and pressure, if any) diverge in finite proper time.  Although the exterior spacetime develops no curvature singularities (since spacetime in (2+1) dimensions has constant curvature outside of all matter sources), the stress-energy of the shell diverges in finite proper time and so the Einstein equations (and the second junction condition in eq. (\\ref{eq:junctions})) break down. Consequently the equation of motion describing the time-evolution of the shell is no longer valid. This yields a mild violation of cosmic censorship in that, strictly speaking, there is no definite manner in which to continue the spacetime beyond this event and its future light cone\\footnote{For a discussion of possible violations of cosmic censorship in fluid collapse in higher dimensions, see \\cite{hdfluid}.}. A Cauchy horizon forms if this endstate is not cloaked by an event horizon.  We find that dust shells can collapse to zero size in an AdS background, displaying similar behavior to that of a pressureless disk of dust \\cite{mr}. The endstate of collapse in both cases is one in which the shell/disk has finite velocity when it achieves zero size, with a diverging intrinsic Ricci scalar.  Unlike the situation in higher dimensions, this state need not be cloaked by an event horizon, in which case a Cauchy horizon is present.  However one might take the viewpoint that it is natural to consider matching this spacetime to one in which the shell bounces repeatedly from zero size to a maximal value and back again; the exterior space will always be that corresponding to an AdS point mass, with the interior space being one of several possibilities as outlined in the discussion in section 3. The collapse time is always finite. If a black hole is formed, the edge of the shell and the event horizon coincide in finite proper time, $t_{h}$.  For shells with pressure, the situation is more intriguing. In this case the endstate of shell collapse has a finite radius, since the material of the shell is no longer infinitely compressible. The intrinsic Ricci scalar becomes singular in finite proper time and the effective potential diverges. In other words, the shell collapses to a singular ring with a finite size within finite proper time. If the external geometry is initially that of a black hole, that singular ring will be screened by an event horizon. However this need not be the case, and a naked singular ring with a finite size can be formed if the exterior metric is that of a point mass. The role of pressure is that of shifting the singular point at $x=0$ to some finite radial position, sustaining the shell with a finite size.  Despite the qualitatively different physics of the GCG shell, we find that it can exhibit similar behaviour to the other two cases. It also presents a variety of scenarios (such as a collapse to a black hole or an indefinitely expanding shell) which depend upon the initial velocity and the shape of the effective potential. However there are some qualitative differences. It is possible for a GCG shell to collapse to a shell of finite radius in which the density is finite but the pressure diverges.  There is still a curvature singularity on the surface of the shell at the radius $x=x_s$. Even if the initial external geometry is a black hole, the singular ring need not be cloaked by an event horizon, and scenarios similar to the previous cases ensue.  While collapsing shells in $(2+1)$ dimensions are not easily translatable into realistic scenarios in $(3+1)$ dimensions, our study is of more than passing interest.  From a general relativistic viewpoint shell collapse highlights the importance of understanding what limits there may be to cosmic censorship. Indeed, since there are two possible endstates for collapse (either a black hole or not) then there should be some kind of critical phenomenon associated with this scenario as with the scalar field \\cite{scalarcollapse}. In this context it would be interesting to extend our results to rotating black holes, where it has been shown that cosmic censorship holds in dust collapse for the addition of a small amount of angular momentum \\cite{Olea}. From a string-theoretic viewpoint it would be interesting to understand the implications of this work for the AdS/CFT correspondence.  \\vspace{1cm}  \\textbf"
    },
    "0609/astro-ph0609689_arXiv.txt": {
        "abstract": "Large fluctuations in the electron column density can occur during the reionization process.  We investigate the possibility of deriving the electron density fluctuations through detailed mapping of the redshifted 21-cm emission from the neutral medium during reionization. We find that the electron-scattering optical depth and 21-cm differential brightness temperature are strongly anti-correlated, allowing optical depth estimates based entirely on redshifted 21-cm measurements. This should help to isolate the cosmic microwave background (CMB) polarization fluctuations due to reionization, allowing the removal of the patchy reionization polarization signal from the other polarization signals and a measurement of the primordial CMB quadrupole at various locations in the universe at the epoch of reionization. This latter application in principle allows three-dimensional mapping of the primordial density field at $z\\sim 1100$ over a large fraction of the Hubble volume. ",
        "introduction": "Two new frontiers in astrophysics and cosmology are the search for small angle polarization anisotropies in the cosmic microwave background and the efforts to map the redshifted 21-cm emission from neutral hydrogen at the epoch of reionization. On small angular scales, a significant source of CMB polarization comes from Thomson scattering of the primordial quadrupole anisotropy by compact regions of ionized gas. These ionized regions would appear in maps of the neutral gas as holes in a relatively uniform map. Thus we expect that where there are strong sources of Thomson scattering there should be missing 21-cm emission and where there is strong 21-cm emission there should be low Thomson optical depth, $\\tau_{\\rm es}$, i.e. the two should be anti-correlated. In this paper we investigate the relation between $\\tau_{\\rm es}$ and neutral hydrogen optical depth % in a quantitative way that properly captures the relevant physics using numerical simulations of structure formation with radiative transfer.  This topic is also related to the idea of Kamionkowski \\& Loeb\\ (1997) \\nocite{kamionkowski97a} for using polarized CMB anisotropy due to Thomson scattering by galaxy clusters of the CMB temperature quadrupole as a means to extract more information about the density field at $z\\sim 1100$ than we can measure in the CMB \\citep{sazonov99,seto00,baumann03,portsmouth04,seto05,bunn06,shimon06}. The information content of the quadrupole measurements has been carefully investigated by Bunn (2006) \\nocite{bunn06}, where it was found that higher redshift clusters are preferred for this, since there is less overlap with our measured CMB sky. Cluster $\\tau_{\\rm es}$ are typically around 0.005 \\citep{mason00}, and we will show that typical fluctuations of $\\tau_{\\rm es}$ during reionization can be of comparable size.  The quadrupole at the time of reionization should not be affected by the integrated Sachs-Wolfe effect and thus provides a remarkably clean measure of very large scale structure at $z\\sim 1100$.  Reionization occurs in a patchy way: strongly-clustered sources create large ionized regions embedded in a mainly neutral medium, which generates CMB polarization anisotropies on small scales. Scattering of the primordial CMB temperature quadrupole leads to linearly polarized signals from ionized regions, where the size of the ionized bubble sets the scale of the anisotropy. On large scales, the relevant quantity is just the mean ionization as a function of redshift, since the small scale structure is averaged out. This is the signal that WMAP detected as the hallmark of reionization.  The amplitude of the expected signal is expected to be on the order of the fluctuations in the electron scattering optical depth times the amplitude of the relevant primordial quadrupole. The latter should be on the order of 15 $\\mu K$  but the former is not well constrained at present. Calculations done assuming that the ionization fluctuations can be treated as a Gaussian random field find very small fluctuations, since modes along the line of sight (LOS) tend to average down \\citep{hu00}. However, recent simulations \\citep{2006MNRAS.369.1625I,21cmreionpaper} % show surprisingly large fluctuations in the optical depth, on the order of 0.01 coming just from the epoch of reionization, as we will show below. This optical depth is larger than that found in even the largest galaxy clusters in the local universe. The all-or-nothing nature of the ionized regions effectively breaks the Limber approximation, often leading to very little cancellation along the LOS. Recent work has suggested that patchy reionization scenarios can result in interesting levels of CMB polarization \\citep{mortonson06}.  In principle, one can use a H~I map to synthesize an effective $\\tau_{\\rm es}$ map and compare with a CMB polarization map to infer the CMB quadrupole at the location of the scattering. Naively, the fluctuation maps (after subtraction of the spatial means) should differ only by an overall amplitude scaling, provided the primordial quadrupole is slowly varying over the epoch of reionization.  We assume a flat $\\Lambda$CDM cosmology with parameters ($\\Omega_m,\\Omega_\\Lambda,\\Omega_b,h,\\sigma_8,n)=(0.27,0.73,0.044,0.7,0.9,1)$ \\citep{2003ApJS..148..175S}, where $\\Omega_m$, $\\Omega_\\Lambda$, and $\\Omega_b$ are the total matter, vacuum, and baryonic densities in units of the critical density, $\\sigma_8$ is the present rms linear density fluctuation on the scale of $8 h^{-1}{\\rm Mpc}$, and $n$ is the index of the primordial power spectrum of density fluctuations. ",
        "conclusions": "We have shown that fluctuations in 21-cm emission from the epoch of reionization and CMB polarization fluctuations should trace each other extremely well. High quality redshifted 21-cm observations will be able to reconstruct the optical depth to Thomson scattering to better than $20\\%$ accuracy; better accuracy can be achieved with a better understanding of reionization.  A map of the Thomson optical depth obtained from 21-cm emission will allow at least two important improvements in our knowledge. If reionization is a contributor to foreground ``B mode'' polarization anisotropy, this will allow cleaning of this foreground. Also, the polarization signal can be used to measure the primordial quadrupole at the time of reionization. A coarse grid of all-sky coverage will allow reconstruction of the large scale structure in the universe; this includes the volume of the Dark Ages, which is currently unobservable due to the lack of sources. With dense coverage a reconstruction of nearly the entire primordial density field at a snapshot of $z\\sim 1100$ within our Hubble volume would be possible; since each point only provides the quadrupole (a very large scale convolution) it is not clear that the reconstruction would be accurate in the presence of astrophysical contaminants, noise, and bulk flows that contribute velocity quadrupoles. However, this signal contains {\\em unique} information about the density fluctuations at $z \\sim 1100$ over a large fraction of our Hubble volume.  Thomson scattering of the primordial quadrupole at lower redshifts has been discussed in the context of using galaxy clusters as Thomson scatterers, but the 21-cm emission has several important advantages: 1) the Thomson optical depth can be accurately reconstructed from the 21-cm map, while galaxy cluster optical depths may not be measurable at the requisite sensitivity; the optical depth derived from Sunyaev-Zel'dovich measurements is weighted by the Compton $y$-parameter and is not appropriate for this purpose \\citep{knox04}; 2) the overlap in information of the CMB quadrupole at $z\\sim 10$ and  our measured CMB sky is much reduced when compared with clusters at $z \\sim 0.2-0.5$; local clusters are better as probes of the ISW effect; 3) the Sunyaev-Zel'dovich effect and gravitational lensing of the CMB will not be correlated with the CMB polarization signal and one does not expect radio halos, quasar and galaxy overdensities, or strongly lensed background objects.  With only a polarization map one could still measure the direction of the CMB quadrupole, which would allow tests of isotropy and homogeneity of the CMB at $z\\sim 10$, provided the reionization process is relatively sharp in redshift. Large ionized regions around quasars would provide interesting targets for ALMA to try to measure this signal, although imaging a field to $\\mu K$ sensitivity with a quasar at the center will require some care. While the signals will be difficult to image with high signal to noise, these observations offer tremendous opportunities for gaining information about our universe that is not available by any other means. The hurdles are not insurmountable and the observations required (high signal to noise neutral hydrogen maps and high resolution imaging of the CMB polarization fluctuations) are already on the ultimate wish list for physical cosmology."
    },
    "0609/astro-ph0609476_arXiv.txt": {
        "abstract": "We re-examine the gravitational acceleration (dipole) induced on the Local Group of galaxies by the Infrared Astronomical Satellite (IRAS) galaxy distribution of the Point Source Catalogue redshift survey (PSCz). We treat the cirrus-affected low galactic latitudes by utilizing a spherical harmonic expansion of the galaxy surface density field up to the octapole order. We find strong indications for significant contributions to the Local Group motion from depths up to $\\sim 185$ $h^{-1}$ Mpc and possible contribution even from $\\sim 210$ $h^{-1}$ Mpc, in agreement with the recent analysis of Kocevski \\& Ebeling of a whole sky X-ray cluster survey. What changes with respect to the previous PSCz dipole analyses is: (a) the large-scale dipole contributions and (b) an increase of the overall dipole amplitude due to the important contribution of the local volume ($\\mincir 4 \\;h^{-1}$ Mpc), which we now take into account. This results in a lower value of the $\\beta (\\equiv \\Omega_{\\rm m}^{0.6}/b)$ parameter, which we find to be $\\beta_{\\rm IRAS} \\simeq 0.49$ in real space. Therefore, for the concordance cosmological model ($\\Omega_{\\rm m}=1-\\Omega_{\\Lambda}=0.3$) the IRAS galaxies bias factor is $b_{\\rm IRAS}\\simeq 1$ which means that IRAS galaxies are good traces of the underlying matter distribution.  \\vspace{0.25cm}  \\noindent {\\bf Keywords}: galaxies: clusters: general - large-scale structure of universe -  infrared: galaxies ",
        "introduction": "It is well known that the peculiar velocity of the Local Group of galaxies is accurately determined from the CMB temperature dipole (Kogut et al. 1996; Bennett et al. 2003), being $u_{LG}=$622 km/sec towards $(l,b)=(277^{\\circ},30^{\\circ})$. Under the framework of linear theory, the most probable cause for this motion as well as for the observed peculiar motions of other galaxies and clusters is gravitational instability (see Peebles 1980). The latter is supported by the fact that the gravitational acceleration of the Local Group of galaxies, as traced by many different samples of extra-galactic mass tracers is well aligned with the general direction of the CMB dipole (eg. Yahil, Walker \\& Rowan-Robinson 1986; Lahav 1987; Plionis 1988; Lynden-Bell et al. 1988; Miyaji \\& Boldt 1990; Rowan-Robinson et al. 1990; Strauss et al. 1992; Hudson 1993; Scaramella et al. 1991; Plionis \\& Valdarnini 1991; Branchini \\& Plionis 1996; Basilakos \\& Plionis 1998; Schmoldt et al. 1999; Rowan-Robinson et al. 2000; Kocevski, Mullis \\& Ebeling 2004; Erdogd\\v{u} et al. 2006; Kocevski \\& Ebeling 2006).  However, what has been debated over the last decades, is the largest depth from which density fluctuations contribute significantly to the gravitational field that shapes the Local Group motion, which is called dipole convergence depth, $D_{\\rm conv}$.%  The outcome of many studies in the past, using different flux or magnitude limited galaxy and cluster samples, was  that the apparent value of the dipole convergence depth differed from sample to sample, in the range from 40 to 160 $h^{-1}$ Mpc, with a strong dependence to the sample's characteristic depth. For example, the dipole of the IRAS Point Source Catalogue for redshifts (hereafter PSCz) was claimed to mostly converge within a depth of $\\sim 100$ $h^{-1}$ Mpc with negligible, if any, contributions beyond $\\sim 140$ $h^{-1}$ Mpc (Rowan-Robinson et al. 2000; Schmoldt et al. 1999), while that of the optical Abell/ACO cluster sample, which is volume limited out to a large depth ($\\magcir 200 \\; h^{-1}$ Mpc), was claimed to converge at $\\magcir 160\\; h^{-1}$ Mpc but with significant contributions beyond $\\sim 140$ $h^{-1}$ Mpc (Scaramella et al. 1991; Plionis \\& Valdarnini 1991; Branchini \\& Plionis 1996). The latter result has been confirmed using X-ray cluster samples, which are free of the various systematic effects from which the optical catalogues suffer (Plionis \\& Kolokotronis 1998; Kocevski et al. 2004; Kocevski \\& Ebeling 2006). Of course, under the assumption that there is a linear bias relation between the cluster, the galaxy and the underlying matter density fluctuations, as proposed by Kaiser (1987), then the galaxy dipole should also show similarly deep contributions. This hinted that the shallower galaxy dipole convergence was spurious, and it could be attributed to the lack of adequately sampling of the distant density fluctuations.  This discussion was recently re-opened by two important papers, that of Erdogd\\v{u} et al. (2006) who analysed the 2 Micron All-Sky Redshift Survey (hereafter {\\it 2MRS}) and found $D_{\\rm conv} \\sim 60 \\; h^{-1}$ Mpc, and of Kocevski \\& Ebeling (2006) who analysed a complete whole-sky survey of X-ray clusters and found $D_{\\rm conv} \\magcir 200 \\;h^{-1}$ Mpc.  Due to the fact that the Erdogd\\v{u} et al. (2006) analysis accounted for the dipole contribution of local galaxies with known true distances and due to the potential importance of such contributions, especially on the light of the apparent difference between the original PSCz and {\\em 2MRS} dipole results (see section 5 in Erdogd\\v{u} et al. 2006), we decided to re-analyse the PSCz catalogue along the same lines. In particular, we estimate the PSCz dipole using an approach which is based on (a) a spherical harmonic expansion of the galaxy surface density  up to the octapole order in order to mask the ``zone of avoidance'' and (b) accounting for contributions to the PSCz dipole from the nearby galaxies. ",
        "conclusions": "We have revisited the PSCz dipole using a spherical harmonics expansion of the galaxy density field up to the octapole order in order to model the excluded low-galatic latitudes. We have also used measured distances from the literature for the nearby galaxies and we find that the amplitude of the dipole increases with respect to the previous PSCz dipole analysis. We also find indications for significant contributions to the gravitational field that shapes the Local Group motion from very large distances in agreement with a recent analysis of a deep all-sky X-ray cluster survey. Finally, within the linear biasing ansatz we find $\\beta_{\\rm IRAS} \\simeq 0.44$ and 0.49 in redshift and real space, respectively. This implies that within the framework of the concordance cosmological model ($\\Omega_{\\rm m}=1-\\Omega_{\\Lambda}=0.3$) the linear biasing factor of the IRAS galaxies is $b_{\\rm IRAS}\\simeq 1$.    {\\small"
    },
    "0609/astro-ph0609195_arXiv.txt": {
        "abstract": "{The detection of exolife is one of the goals of very ambitious future space missions or extremely large ground-based telescopes that aim to take direct images of Earth-like planets. While associations of simple molecules present in the planet's \\textit{atmosphere} ($O_2$, $O_3$, $CO_2$ etc.) have been identified as possible global biomarkers, we analyse here the detectability of vegetation on a global scale on Earth's surface. } {Considering its specific reflectance spectrum showing a sharp edge around 700 nm, vegetation can be considered as a potential global biomarker. This work, based on observational data, aims to characterise and to quantify this signature in the disk-averaged Earth's spectrum.} {Earthshine spectra have been used to test the detectability of the \"Vegetation Red Edge\" (VRE) in the Earth spectrum. We obtained reflectance spectra from near UV (320 nm) to near IR (1020 nm) for different Earth phases (continents or oceans seen from the Moon) with EMMI on the NTT at ESO/La Silla, Chile. We accurately correct the sky background and take into account the phase-dependent colour of the Moon. VRE measurements require a correction of the ozone Chappuis absorption band and Rayleigh plus aerosol scattering.} {The near-UV spectrum shows a dark Earth below 350 nm due to the ozone absorption. The Vegetation Red Edge is observed when forests are present ($4.0\\%$ for Africa and Europe), and is lower when clouds and oceans are mainly visible ($1.3\\%$ for the Pacific Ocean). Errors are typically $\\pm0.5$, and $\\pm1.5$ in the worst case. We discuss the different sources of errors and bias and suggest possible improvements.} {We showed that measuring the VRE or an analog on an Earth-like planet remains very difficult (photometric relative accuracy of 1\\% or better). It remains a small feature compared to atmospheric absorption lines. A direct monitoring from space of the global (disk-averaged) Earth's spectrum would provide the best VRE follow-up. } ",
        "introduction": "Future space missions like Darwin/TPF will provide the first images and low-resolution spectra of Earth-like extrasolar planets, with the final goal to extract a ''biosignature'' from the spectra. The TPF-C NASA mission will observe exo-terrestrial planets around the nearby stars with a large off-axis telescope equipped with a visible coronagraphic device. To prove the interest of the visible part of the spectrum in spite of the huge contrast ($10^{10}$) compared to thermal infrared ($10^7$), we study the presence of biomarker features. The spectrum of the light reflected by a planet, when normalized to the parent star spectrum, gives the planet reflectance spectrum, revealing planet atmospheric and ground colour, when the latter is visible through a partially transparent atmosphere. The planet will remain unresolved and its spectrum will be spatially integrated (i.e. disk-averaged) at the observed orbital phase of the planet. In this work, we estimate how the Earth's spectrum would look like from $\\sim$ 10 pc. This can be done - in principle at least - by integrating spatially-resolved spectra from low-orbit satellites. It can also be done from a space probe traveling into the  Solar System and looking back the Earth, as Mars Express did in 2003\\footnote{http://www.esa.int/esaCP/Pr\\_ 44\\_ 2003\\_ p\\_ EN.html}. An alternative method to obtain the Earth's disk-averaged spectrum consists of taking a spectrum of the Moon Earthshine, i.e. Earth light backscattered by the non-sunlit Moon. A spectrum of the Moon Earthshine directly gives a disk-averaged spectrum of the Earth at a given phase as seen from the Moon (since the Moon surface roughness \"washes out\" any spatial information on the Earth's colour).  Biosignatures are of two types: out-of-equilibrium molecules in the planet atmosphere (like oxygen and ozone) or ground colours characteristic of biological complex molecules (like pigments in vegetation). The visible and near infra-red Earthshine spectra published to date (\\cite{arnold02,woolf02,seager05,montanes05}) clearly show the atmospheric signatures and, at least, tentative signs of ground vegetation which thus appears as an interesting potential global biomarker. Vegetation indeed has a high reflectivity in the near-IR, higher than in the visible by a factor of $\\approx5$ (\\cite{clark99}). This produces a sharp edge around $\\approx 700\\rm nm$, the so-called Vegetation Red Edge (VRE). An Earth disk-averaged reflectance spectrum can thus rise by a significant fraction around this wavelength if  vegetation is in view from the Moon when the Earthshine is observed. Arcichovsky suggested in 1912 to look for chlorophyll absorption in the Earthshine spectrum, to calibrate chlorophyll in the spectrum of other planets (\\cite{arci1912}). This approach re-emerged again only in 1999 with preliminaryt tests at ESO by several among us (P.F., J.S. and D.B.) and at OHP (\\cite{arnold02}). The position of the VRE may not lie at 725 nm on an exo-Earth. Its possible confusion with mineral spectral features is discussed by Schneider (2004) and Seager et al. (2005).  The red side [600:1000 nm] of the Earth reflectance spectrum shows the presence of $O_2$ and $H_2O$ absorption bands and of the VRE, while the blue side [320:600 nm] clearly shows the Huggins and Chappuis ozone ($O_3$) absorption bands. The higher reflectance in the blue shows that our planet is blue due to Rayleigh scattering in the atmosphere, as detected by Tikhoff (1914) and Very (1915), and confirmed later with accurate Earthshine photometry by Danjon (1936).  The present paper presents new observations and emphasizes the difficulties of data reduction  and describes how it has been improved with respect to our first work (\\cite{arnold02}). In Section \\ref{background_sub}, we explain how we adapt Qiu's method (\\cite{qiu03}) to carefully subtract the sky background from Earthshine spectra. Section \\ref{sec:moon} describes how we correct the phase-dependent colour effect of the Moon, following a new processing of Lane \\& Irvine (1973) data. This correction was not done in Arnold et al. (2002). Section \\ref{discussion} presents the first - as fas as we know - near-UV integrated Earth spectra obtained from Earthshine observations that reveal a dark Earth below 350 nm due to the strong $O_3$ absorption (ozone Huggins absorption bands). While this work presents smaller VRE than Arnold et al. (2002) and Woolf et al. (2002), it nevertheless shows different values depending whether ocean or land are seen from the Moon at the epoch of observations. ",
        "conclusions": "We have presented spectra of the Earth reflectance from near-UV to near-IR (320 to 1020 nm) for four different nights. They show significant variations in Rayleigh scattering depending on the cloud cover (the Earth 'blue dot' can be almost white). One of the spectra was taken with mainly light from the Pacific Ocean while the others included parts of or the whole Africa and Europe. The Vegetation Red Edge is observed when land with forests are present, with values between 3 and 4\\% ($\\pm\\approx0.5\\%$), but remains very low otherwise ($\\approx1\\%$). A spectro-photometric accuracy better than 1\\% is therefore required for future ground based observations (with Extremely Large Telescopes) or space missions aiming to detect an analog to the terrestrial VRE. These instruments will all be equipped with an optimized coronagraph allowing starlight rejection better than $10^5$ to be able to provide the image of the planet itself.  A survey over one year or more with monthly observations would allow one to follow the VRE seasonal variations and undoubtedly improve our knowledge of the behavior of this biomarker.  In addition, it would be desirable to have a direct measurement and monitoring from space of the global (i.e. disk-averaged) spectrum of the Earth. One solution would be to observe Moon Earthshine: from space, the Earth reflectance would still be obtained with Eq.\\ref{albedo}, but would avoid the source of noise from one path through the atmosphere for both ES and MS light. Other solutions would be to integrate spatially-resolved spectra from low-orbit satellites, or ideally, to have a satellite far enough from the Earth to see it entirely and able to easily obtain disk-averaged spectra of the Earth radiance. The Earth reflectance would then simply be obtained after a division by the solar radiance, avoiding all sources of noise and bias due to the Moon.  Our work showed that measuring the VRE remains very difficult, and, although measurable, the VRE is a small feature when compared to $O_2$, $O_3$ and water vapor absorption lines. But $O_2$ present in the Earth atmosphere is 100\\% of photosynthetic origin and its deep absorption line at 760 nm is, for the Earth, a signature of photosynthesis. This suggests that if $O_2$ is observed at 760 nm on an exo-Earth, it is relevant to look for a VRE-like (probably small) feature in the spectrum, keeping in mind that this signature can be significantly different (spectral shape and wavelength) on an exo-Earth than on our planet. Clearly a sharp feature may be much more detectable than a smoother signature. An exhaustive analysis of possible artifacts remains necessary to consolidate the work done on this subject (\\cite{schneider04,seager05}). Phylogenetic analysis of plants and research on alternative photosynthesis chemistries are also of prime interest to better constrain this plausible biomarker."
    },
    "0609/astro-ph0609829.txt": {
        "abstract": "A number of trends among the properties of exoplanets have become evident in the years since the first one was announced in 1995. One particularly interesting trend began to emerge in 1997 -- the incidence of giant planets correlates with the metallicity of the host star. This has since been established with a high degree of statistical significance by several research groups. Other, more subtle, trends are beginning to appear as the sample size continues to grow and the statistics improve. I review the state of our knowledge concerning the observed compositional trends and their possible causes and suggest several research directions. ",
        "introduction": "The announcement of the first planetary mass object orbiting another sun-like star by \\citet{mayq95} opened a new field of empirical study to astronomers. At the time of this writing (August 2006), nearly 200 exoplanets have been reported. Nearly all have been found with the Doppler method, whereby the reflex motion of the host star about the system's center-of-mass is measured over at least one orbital period. A few have also been found with the transit and microlensing methods.  Several global trends have been found in the database of extrasolar planets around nearby FGK stars \\citep{marcy05, but06}: 1.2 per cent of the planets are found in orbits with semi-major axis, $\\it a$, less than 0.1 AU; the distribution of planet mass is a strongly decreasing function of planet minimum mass, $\\it M$,  ($dN/dM \\propto M^{-1.1}$) with a sharp cutoff above 12 M$_{\\rm J}$; $> 7$ percent of surveyed stars have planets with semimajor axis, $\\it a$, $< 5$ AU; eccentric orbits are common, with a median eccentricity value of 0.25; there is a paucity of high-mass planets with short-period orbits.  These discoveries have spurred much theoretical work. Some have focused on explaining the high eccentricities (see review by \\citet{trem04}), while others have explored mechanisms of planet migration to smaller orbits (see review by \\citet{thom05}). These studies are based on one of two broad categories of planet formation theories: core instability accretion (CIA, \\citet{poll96}); disk gravitational instability (DGI, \\citet{boss04} and references therein). The CIA model is more mature, and it has had some success in explaining some of the observed trends (e.g., \\citet{idalin05}).  The topic of the present review, however, concerns the compositions of the planet host stars. \\citet{gg97} first showed a link between the metallicities of the host stars and the presence of planets. Many studies have since been published on this topic, the most recent ones confirming the link with a high degree of statistical significance (e.g., \\citet{sant05,fv05}). In this review I will describe research on the composition of stars with planets (SWPs), focusing on studies published since my \\citet{gon03} review. In \\S~2 I summarize the ongoing exoplanet surveys and the spectroscopic observations and analyses of SWPs. I briefly review in \\S~3 survey statistics, including the preparation of control samples, biases and the incidence of exoplanets. In \\S~4 I summarize the results of recent spectroscopic chemical abundance analyses of SWPs, and in \\S~5 I discuss the proposed explanations for the observed abundance anomalies among SWPs. I end in \\S~6 by suggesting several research projects that may help advance our understanding of planet formation and evolution by studying the chemical compositions of SWPs. ",
        "conclusions": "The following list of suggested research topics may help advance our understanding of planet formation and evolution by studying compositional anomalies among SWPs. While most are not new suggestions, they need to be pursued more fully than they have been. They are arranged roughly in order of decreasing priority/increasing difficulty.  The most firmly established compositional anomaly among SWPs is their high mean metallicity relative to control samples. The search is now on for more subtle anomalies. As shown by several recent studies (e.g., \\citet{luck06,bond06,gil06}), mean compositional differences between SWPs and control samples, expressed as [X/Fe], must be smaller than about 0.15 dex. It is possible to detect compositional anomalies using current techniques even at this level, but it will be necessary to control carefully for systematic errors.  Apart from the controversies surrounding Li abundances, only \\citet{rob06} have claimed to discover statistically significant compositional anomalies among SWPs. Their findings need to be confirmed by independent analyses. Continuing spectroscopic study of nearby SWPs stars may confirm the other compositional differences suggested by some studies (e.g., Al/Fe). Addition of more metal-rich ([Fe/H] $> 0.3$) comparison stars will also be helpful in this regard. It is especially valuable to have several independent research groups testing each other's findings, given the subtleties of the claimed anomalies and the possibility of unaccounted for systematic errors (e.g., \\citet{schul06}). The recent history of research on this topic has shown on multiple occasions that findings that originally looked convincing were not confirmed following additional research.  On the theoretical front, published abundances of SWPs and control samples should be corrected for departures from LTE. Presently, only SWP O abundances derived from the O I 777 nm triplet lines include such corrections. The differential (relative to the Sun) corrections for most published abundances of these stars are expected to be considerably less than 0.1 dex. Studies could be modeled after \\citet{idiart00}, for example, which calculated non-LTE corrections to Ca and Mg abundances using published EW values. Another way to improve the abundances determinations is to replace the 1D stellar model atmospheres with 3D model atmospheres \\citep{asplund05}. These corrections are especially needed now that planet hunters are expanding their searches to stars outside the original dwarf FGK range.  Application of these corrections to spectroscopic analyses of FGKM dwarfs can be facilitated by setting up a grid of standard stars spanning the same range of spectral types. A small number of standard stars could be analyzed with the most recent 3D atmospheres and non-LTE calculations using time-consuming methods. The new standards could then be employed to correct chemical abundances of stars analyzed using the simpler traditional methods. For best results, the analyzed stars would be compared to standards closest in spectral type and metallicity. The standard stars could be analyzed again with every new generation of model atmospheres and analysis methods.  Several CPM pairs have been reported to display significant differences in [Fe/H]. These need to be confirmed with independent analyses. Proper interpretation of the results will require consideration not only of self-enrichment but also of diffusion. Both diffusion and self-enrichment are expected to become more important for hotter dwarfs. The two processes can be separated if a sufficient number of elements is included in the analyses \\citep{ash05}.  It would also be worthwhile to continue the search for signatures of self-enrichment among early type stars. In particular, the strong evidence for self-enrichment in J37 in NGC\\,6633 serves as motivation to conduct a general spectroscopic survey of similar stars in clusters of comparable age and metallicity. More generally, studies similar to \\citet{dot03}, which searched for trends between [Fe/H] and spectral type and between the scatter in [Fe/H] and spectral type in Hyades A and F stars, should continue to be refined and applied to additional clusters.  As noted recently by \\citet{jura06}, the atmospheres of cool helium-rich white dwarfs are sensitive to small amounts of accretion. Determination of the abundances of additional elements in their atmospheres should permit more stringent tests of the two competing explanations for the anomalous cases: interstellar and circumstellar accretion.  While attempts to detect the signature of self-enrichment via abundance-T$_{\\rm c}$ trends have not proven successful, this test should continue to be applied periodically as the database of SWP and control sample abundance determinations continues to grow. Although the test can be refined by any improvements to abundance determinations, the greatest benefits will come from improvements for elements with low values of T$_{\\rm c}$, such as C and O. O abundances, in particular, require careful attention, given recent controversies \\citep{ecu06b}. If self-enrichment is a rare process, then evidence for it will be seen in a small number of extreme stars, like J37. Otherwise, it could be detected as a statistically significant trend in a large sample. Regardless of which is the case, any analysis of this type should include as many F dwarf stars as possible.  Comparison of the metallicities of dwarfs and subgiants in a cluster will allow us to test self-enrichment much more effectively than similar studies of field stars. An analysis similar to that of \\citep{murr01} could be conducted on an old, solar metallicity open cluster with a well-populated turn-off and subgiant branch (e.g., M67). Stars that had been enriched on the main-sequence should display reduced surface metallicity on the subgiant branch. Such a study would avoid the weakness of \\citet{murr01}'s less homogeneous sample. An even better approach would be to compare enriched stellar evolutionary tracks \\citep{cody05} to star cluster color-magnitude diagrams.  As noted above, some planet formation models proposed to explain the metallicity-planet link also predict certain trends with stellar mass. Together, these two stellar parameters more tightly constrain such models than either of them alone. Planet surveys that target stars outside the main sequence FGK spectral type range, such as \\citet{gal06}, \\citet{endl06} and \\citet{john06}, will be especially helpful in testing the possible stellar mass-planet link.  To date, only one SWP has been targeted for asteroseismological analysis, $\\mu$ Arae \\citep{baz04,baz05}. The analysis results tend to support the self-enrichment model for this star, but additional data are needed to reach a firm conclusion. \\citet{baz05} suggest that interferometric radius measurements of $\\mu$ Arae accurate to better than 1 per cent would provide a significant new constraint on the models. Significant improvements in p-mode identifications would require a substantial new Doppler observing campaign.  Doppler planet searches in open clusters have not yet proven successful, but only one cluster, the Hyades, has been targeted. Focusing on an open cluster removes age and initial metallicity as variables. The Hyades is the closest cluster, and it is metal-rich, but its stars are apparently too young to allow detection of planets with confidence \\citep{paul04}. It will be necessary to target clusters comparable in age and metallicity to the Sun, such as M67. Doppler searches for planets in clusters beyond the Hyades will require at least 8-m class telescopes. Allocation of the needed large telescope time for these searches may have to await better estimates of the expected detection efficiencies from open cluster transit surveys.  \\subsection{Concluding Comments}  The composition of the typical nearby SWP differs from that of the typical nearby star. Several recent studies have established with a high degree of statistical confidence that, as a group, SWPs are more metal-rich than stars without known planets. The data are also suggestive of differences in the Al/Fe, Si/Fe and Ni/Fe ratios. Progress in testing the reality of these differences can be made by including more nearby stars without planets having [Fe/H] $> 0.3$ in the control samples.  The primordial hypothesis is the best single explanation for the anomalously high metallicities of nearby SWPs. The data are also suggestive of an orbital period bias, implying a metallicity-dependent mechanism that affects the final orbital position of a planet. However, an updated analysis of its statistical significance is required. Evidence for self-enrichment among SWPs is still weak, but there is the expectation that convincing evidence for it will be found among early F dwarf SWPs discovered in ongoing Doppler planet surveys.  Although the number of planets detected with the transit method is still small, it is already evident that their host stars are metal-rich. There is also evidence for a correlation between the metallicities of the host stars and the planets of transiting systems.  Taken together, these results favor the CIA model of giant planet formation. While some progress has been made in explaining the composition of SWPs within the framework of the CIA model, much remains to be done. The competing DGI model may still account for a small fraction of the known planets. Its contribution could be quantified once the CIA model is better constrained and the number statistics for metal-poor SWPs is improved."
    },
    "0609/astro-ph0609310_arXiv.txt": {
        "abstract": "The radio complex Sgr B region is observed with the X-Ray Imaging Spectrometers (XIS) on board Suzaku. This region exhibits diffuse iron lines at 6.4, 6.7 and 6.9~keV, which are K$\\alpha$ lines of Fe \\emissiontype{I} (neutral iron), Fe\\emissiontype{XXV} (He-like iron) and Fe\\emissiontype{XXVI} (H-like iron), respectively. The high energy resolving power of the XIS provides the separate maps of the K-shell transition lines from Fe\\emissiontype{I} (6.4~keV) and Fe\\emissiontype{XXV} (6.7~keV). Although the 6.7~keV line is smoothly distributed over the Sgr B region, a local excess is found near at $(l, b) = (\\timeform {0D.61}, \\timeform{0D.01})$, possibly a new SNR. The plasma temperature is \\textit{kT} $\\sim$3~keV and the age is estimated to be around several$\\times10^{3}$ years. The 6.4~keV image is clumpy with local excesses nearby Sgr B2 and at $(l, b) = (\\timeform{0D.74}, -\\timeform{0D.09})$. Like Sgr B2, this excess may be another candidate of an X-ray reflection nebula (XRN). ",
        "introduction": "% Sgr B is a molecular complex consisting of Sgr B1 and B2. In particular, Sgr B2 is the most massive giant molecular cloud with ultra compact H\\emissiontype{II} regions (UCH\\emissiontype{II}) \\citep{Gaume1995} and many maser sources near the cloud center \\citep{Mehringer1997}. These are the hints of high mass (HM) zero age main sequence (ZAMS) stars or young stellar objects (YSOs). However, extremely high absorption toward the cloud center ($N_{\\rm H} \\ge 10^{24}~{\\rm cm}^{-2}$, $A_{\\rm V}$= a few 100) prevents to detect any stars in optical or even in the infrared bands. \\citet{Takagi2002} have discovered many compact X-ray sources in the cloud center with Chandra. The X-ray fluxes and spectra indicate that these are likely HM YSOs. Since HM stars evolve very rapidly and finally undergo  supernova explosions, it will be reasonable to expect young SNRs near the Sgr B2 cloud. In facts, \\citet{Senda2002} have discovered a peculiar SNR candidate with Chandra.  Sgr B2 is also a strong 6.4~keV line source. \\citet{Koyama1996} and \\citet{Murakami2001} proposed that Sgr B2 is an X-ray reflection nebula (XRN), irradiated by the Galactic center (GC) source Sgr A$^{*}$. Sgr A$^*$ was  then  thought to be X-ray bright about 300 years ago, the light traveling time between Sgr B2 and Sgr A$^{*}$. In this XRN scenario, it is likely that other XRNe will be found in the molecular complex Sgr B. The Suzaku observation on the Sgr B region is intended to discover new SNRs and XRNe. This paper reports the first results of the Suzaku observation. ",
        "conclusions": "% \\subsection{The Overall Features} %  The X-ray spectrum of all the Sgr B region is given in figure \\ref{fig:sgrb-overall}. The spectra of the four XISs (XIS0-XIS3) are co-added and the night earth spectrum (non X-ray background; here, NXBG) is subtracted. With the superior energy resolution of the XIS for diffuse sources, we can clearly resolve the 6.4~keV, 6.7~keV and 6.9~keV lines. These are K$\\alpha$ lines from  neutral Fe\\emissiontype{I}, He-like Fe\\emissiontype{XXV} and hydrogenic Fe\\emissiontype{XXVI}. The 6.9~keV line may contain a small fraction of K$\\beta$ line of Fe\\emissiontype{I} (7.07 keV). Weak lines seen above ~7 keV are K$\\alpha$ of Ni\\emissiontype{I}(at $\\sim$ 7.5 keV), K$\\alpha$ of Ni\\emissiontype{XXVII} + K$\\beta$ of Fe\\emissiontype{XXV} (at $\\sim$7.8--7.9 keV) and K$\\beta$ of Fe\\emissiontype{XXVI} + K$\\gamma$ of Fe\\emissiontype{XXV} (at $\\sim$8.2--8.3 keV).  \\begin{figure} \\begin{center} \\FigureFile(80mm,80mm){gc_sgrb_overall_xisall_spec.eps} \\end{center} \\caption{The X-ray spectra from the full FOV of the XIS, but the CCD corners irradiated by the build-in calibration sources are excluded. All the four XIS data are co-added} \\label{fig:sgrb-overall} \\end{figure}  \\subsection{Discovery of a new SNR} % \\vspace*{1 cm} \\begin{figure} \\begin{center} \\FigureFile(80mm, 10mm){gc_sgrb_6700_publish.eps} \\end{center} \\caption{The 6.7~keV line map (the 6.58--6.74 keV band map) showing bright spot at the northwest corner. The fluxes are normalized by the 6.7~keV flat-field image. The source and background regions are shown by the solid and dotted ellipses, respectively.} \\label{fig:sgrb-6700img} \\end{figure} We made a narrow band image of 6.7~keV (the 6.58--6.74 keV band) in figure \\ref{fig:sgrb-6700img}. We see a clear 6.7~keV flux excess at the northwest corner. To confirm the 6.7 keV excess, we referred the archive data of Chandra (OBSID:944, effective exposure time was 99 ksec) and XMM (OBSID: 0203930101, effective exposure time was 42~ksec). Since the energy resolution of the Chandra ACIS is limited to separate the 6.4 and 6.7 keV lines, it is unclear whether the 6.7 keV source is present or not. The XMM image in the 6.7 keV band shows a clear elongated structure near the same position. Accordingly the presence of the 6.7 keV line source is no doubt. On the other hand, the continuum band image (e.g. 2-5 keV, or 2-8 keV bands) shows only a hint of enhancement, and show no clear structure.  This source is therefore very peculiar, which is dominant only in the 6.7 keV line. Since the dominance of the 6.7 keV line suggests that the excess is a new SNR (see also below for the discussion), we designate this source as Suzaku~J1747.0$-$2824.5 (G0.61+0.01) from its center position.  We made the NXBG-subtracted spectra from the solid ellipse for both the three FI CCDs (XIS0, XIS2 and XIS3 are co-added) and BI CCD (XIS1). In these spectra, the cosmic X-ray  background (CXB) and GCDX  are still included.  We therefore made the NXBG-subtracted spectra from the dotted ellipse in figure \\ref{fig:sgrb-6700img}, and subtract this local background (CXB + GCDX) from the source spectrum. All the spectra  have  been corrected for the vignetting at 6.7~keV. The results are given in figure \\ref{fig:sgrb-g061001-spec}, for the FIs and  BI, separately.  We see a pronounced peak at 6.7~keV, but no 6.9 keV line. The 6.7 line shape is asymmetric with a tail at lower energy. In order to verify the line structure, we drive fluxes of the 6.4, 6.7 and 6.9 keV lines ( K$\\alpha$ line of Fe \\emissiontype{I}, \\emissiontype{XXV} and  \\emissiontype{XXVI}) in the source and the background regions, applying a phenomenological model (a bremsstrahlung continuum and many Gaussian lines) in the raw data (no background subtraction). The resulting  6.4,  6.7 and 6.9 keV line fluxes are  2.22, 5.17 and 0.48 for the G0.61+0.01 (source) region, and  0.61, 0.68  and 0.30 for the background region, where the flux unit is  10$^{-6}$ photons~cm$^{-2}$~s$^{-1}$ ~arcmin$^{-2}$. In contrast to the 6.7 keV line, we see no large excess in the 6.9 keV line from the source region compared to the background region. Thus we confirm that G0.61+0.01 emits strong  6.7 keV line, but very weak 6.9 keV line. The  small excess of the  6.4 keV line makes the low energy  tail in the 6.7 keV line.  \\begin{figure} \\begin{center} \\FigureFile(80mm,80mm){gc_sgrb_blob_g_spec_vpshock_fi.eps} \\FigureFile(80mm,80mm){gc_sgrb_blob_g_spec_vpshock_bi.eps} \\end{center} \\caption{Left: the X-ray spectrum of the sum of 3 FI CCDs (XIS0, 2 and 3) for a new SNR (G0.61+0.01) with the best-fit VPSHOCK model. Right: same as the left, but of the BI CCD (XIS1)}  \\label{fig:sgrb-g061001-spec} \\end{figure} \\begin{table} \\caption{The best-fit parameters for G0.61+0.01 with the VPSHOCK model plus two emission lines} \\label{tab:g_fit} \\begin{center} \\begin{tabular}{lc} \\hline\\hline Parameter & Value \\\\ \\hline $N_{\\rm H}$ ($10^{23}$~H~cm$^{-2}$) & $1.6_{-0.4}^{+0.7}$ \\\\ $kT$ (keV) & $3.2_{-0.9}^{+2.3}$ \\\\ $n_{\\rm e}t$ ($10^{11}$ cm$^{-3}$ s) & $1.9_{-0.8}^{+4.7}$ \\\\ \\multicolumn{2}{l}{Abundances\\footnotemark[$*$]}\\\\ Ca & $3.5_{-2.4}^{+3.1}$ \\\\ Fe & $5.1_{-1.1}^{+1.2}$ \\\\ \\multicolumn{2}{l}{Neutral iron lines\\footnotemark[$a$]}\\\\ $I_{6.40}$ ($10^{-6}$ photons cm$^{-2}$ s$^{-1}$) & $5.1_{-2.5}^{+ 2.4}$ \\\\ $I_{7.06}$ ($10^{-6}$ photons cm$^{-2}$ s$^{-1}$) & 0.6 \\\\ \\multicolumn{2}{l}{Flux and Luminosity}\\\\ $F_{\\rm 2-10}$\\footnotemark[$\\dagger$] ($10^{-13}$ ergs cm$^{-2}$ s$^{-1}$) & $7.5_{-2.2}^{+1.1}$ \\\\ $L_{\\rm 2-10}$\\footnotemark[$\\ddagger$] ($10^{34}$ ergs s$^{-1}$) & $1.5_{-0.2}^{+0.1}$ \\\\ \\hline $\\chi^2$/dof & 99.0/78\\\\ \\hline \\multicolumn{2}{@{}l@{}}{\\hbox to 0pt{\\parbox{80mm}{\\footnotesize Note---The uncertainties indicate the 90\\% confidence limit. \\par\\noindent \\footnotemark[$*$] The elements which are not listed below are fixed at 1.0 (solar ratio). \\par\\noindent \\footnotemark[$a$] The line energy of K$\\alpha$ and K$\\beta$ is fixed at the theoretical value (6.40 and 7.06~keV, respectively; Kaastra and Mewe 1993) and the intensity of K$\\beta$ is fixed at 12.5\\% (Kaastra and Mewe 1993) of that of K$\\alpha$. \\par\\noindent \\footnotemark[$\\dagger$] Observed flux in the 2.0--10.0~keV band. \\par\\noindent \\footnotemark[$\\ddagger$] Absorption corrected luminosity in the 2.0--10.0~keV band. }\\hss}} \\end{tabular} \\end{center} \\end{table}  The FI and BI spectra are simultaneously fitted with a plane parallel shock model (VPSHOCK in the XSPEC package) adding two Gaussian lines at 6.4~keV and 7.06~keV. These two lines represent the K$\\alpha$ and K$\\beta$ lines of Fe \\emissiontype{I}, where the flux of latter line is fixed at 12.5\\% of the former (Kaastra and Mewe 1993). The best-fit results and parameters are shown in figure \\ref{fig:sgrb-g061001-spec} and table \\ref{tab:g_fit}. Although we detected the 6.4 keV line from G0.61+0.01,  it is very difficult to judge whether this line is really attributable to  G0.61+0.01, due to spilled-over flux from the adjacent source Sgr B2,  or due to a fluctuation of a larger scale structure in the 6.4~keV line.  As for the last possibility,  we see a large scale 6.4~keV enhancement in the northwest compared to the background region in the southeast (see figure \\ref{fig:sgrb-6400img}).  In any case,  we ignore this line in the discussion of G0.61+0.01 because the 6.4~keV line flux is only 3\\% of that in Sgr B2 (see tables 1 and  3). Since the Suzaku spatial resolution is not good enough, there may be possible contamination of unresolved point sources. To check this problem, we searched for point sources using the Chandra archive data (OBSID: 944, ~99 ksec exposure time) and found no point source in the source region. On the other hand,  in the background region, there are 48 point sources. The total flux in the 2-10 keV band is  5$\\times10^{-13}$ ergs~cm$^{-2}$~s$^{-1}$, which is only $\\sim$2\\% of the CXB + GCDX flux of 2.6$\\times10^{-11}$ ergs~cm$^{-2}$~s$^{-1}$, hence can be ignored in the present data analysis and discussion.  The best-fit temperature of $\\sim$3~keV and overabundance of Fe are consistent with an ejecta of  an SNR and are similar to those found in the central region of Sgr A East, a young  SNR near at the GC. The high temperature component of Sgr A  East is $kT\\sim$4--6 keV (Sakano et al. 2004, Park et al. 2005, Koyama et al. 2006c),  and iron is overabundant by factor of 4--5 (Maeda et al. 2002, Sakano et al. 2004, Park et al. 2005, Koyama et al. 2006c).  Thus G0.61+0.01 is likely an ejecta dominant central region of an SNR. We  note that Sgr A East has low temperature component of about 1 keV, while not in  G0.61+0.01. The absence of softer plasma may be due to the large absorption. The $N_{\\rm H}$ value of $1.6\\times 10^{23}$~H~cm$^{-2}$ is larger than that of typical value to the GC ($6\\times 10^{22}$~H~cm$^{-2}$ ) \\citep{Sakano2002}. Therefore, G0.61+0.01 would be located behind or in the rim of the Sgr B2 cloud. Since G0.61+0.01 is located in the south of an expanding radio shell \\citep{Oka1998}, which is probably interacting with the Sgr B2 cloud rim, we assume that the distance of G0.61+0.01 is the same as Sgr B2 and to be 8.5~kpc \\citep{Reid1988}. Then the 2--10 keV band luminosity is estimated to be $1.5\\times10^{34}$ ergs~s$^{-1}$, which is typical for an ejecta plasma of an  SNR. The size of G0.61+0.01 (the solid ellipse in figure \\ref{fig:sgrb-6700img}) is ~2.2$^\\prime\\times$4.8$^\\prime$, which corresponds to ~5.5 pc~$\\times$~12 pc  at a distance of 8.5 kpc. Assuming the plasma emission is due to a uniform density ellipsoid with the 3-axis radii of 2.7 pc, 2.7 pc and 6 pc, we estimate physical parameters of G0.61$+$0.01 ( table \\ref{tab:g_physpar}). Although  the iron abundance is 3-4 times of the solar, total number and mass of protons are fur larger  than those  of irons.  We therefore assume that electron density ($n_{\\rm e}$) is equal to the proton density ($ n_{\\rm p}$), and that protons carry most of the plasma mass ($m_{\\rm p}$). Dividing the radius of the major axis (6 pc)  by the sound velocity of the 3.2~keV plasma ($v=1.4\\times$10$^8$ cm~s$^{-1}$), we obtain the dynamical time scale ($t_{dyn}$) of $\\sim$4$\\times10^{3}$ years. If, instead, we  use the ionization parameter ($n_{\\rm e}t$) and electron density ($n_{\\rm e}$), then the ionization time scale ($t_{ioni}$) is estimated to be $\\sim7\\times10^{3}$ years. Since the source size of ~2.2$^\\prime\\times$4.8$^\\prime$ is comparable to that of the half power diameter ($\\sim2^\\prime$), the real size of G0.61+0.01 must be smaller. Therefore the quoted value of $t_{dyn}\\sim4\\times$10$^{3}$ years should be an upper limit, while that of  $t_{ioni}\\sim7\\times$10$^{3}$ years is a lower limit, because $n_{\\rm e}$ is inversely proportional to the root square of the plasma volume. Thus the age of G0.61+0.01 is probably around several$\\times$10$^{3}$ years.   Another possibility is that  G0.61+0.01 comprises  a part of a larger SNR. Since G0.61+0.01 is found at the edge of the XIS field, other parts of a candidate  SNR may be out of the XIS field. In this scenario, G0.61+0.01 may be  a part of the expanding radio shell discovered by \\citet{Oka1998}. The kinetic energy of the radio shell is a few of $10^{52}$ erg s$^{-1}$, within the range of single or multiple supernova explosions. Thus follow-up X-ray observations including this expanding radio shell is highly required. \\begin{table} \\caption{The physical parameters of G0.61+0.01} \\label{tab:g_physpar} \\begin{center} \\begin{tabular}{lc} \\hline\\hline Parameter & Value \\\\ \\hline EM\\footnotemark[$a$] (cm$^{-3}$) & $1.4\\times 10^{57}$ \\\\ $n_{\\rm e}$\\footnotemark[$b$] (cm$^{-3}$) & $0.9$ \\\\ $M$\\footnotemark[$c$] (\\MO) & 1.3 \\\\ $E$\\footnotemark[$d$] (ergs) & $2.4\\times 10^{49}$ \\\\ $t_{\\rm dyn}$\\footnotemark[$e$] (s) & $1.3\\times10^{11}$ \\\\ $t_{\\rm ioni}$\\footnotemark[$f$] (s) & $2.1\\times10^{11}$ \\\\ \\hline \\multicolumn{2}{@{}l@{}}{\\hbox to 0pt{\\parbox{60mm}{\\footnotesize Note---The plasma is assumed to be a uniform density ellipsoid with the 3-axis radii of 2.7 pc, 2.7 pc and 6 pc (see text). \\par\\noindent \\footnotemark[$a$] Emission measure (EM) = $n_{\\rm e} n_{\\rm H}V$ = $n_{\\rm e}^{2}V$, where $n_{\\rm e}$ and $n_{\\rm H}$ are  the electron and hydrogen  density and are assumed to be equal. \\par\\noindent \\footnotemark[$b$] The electron density. \\par\\noindent \\footnotemark[$c$] Total mass $(M) = n_{\\rm e}m_{\\rm p}V$, where $m_{\\rm p}$ is the proton mass and $V$ is the plasma volume. \\par\\noindent \\footnotemark[$d$] Thermal energy $(E) = 3n_{\\rm e}kTV$. \\par\\noindent \\footnotemark[$e$] The dynamical time scale:  the radius of the major axis of the plasma ellipsoid divided by the sound velocity of the $\\sim$3 keV plasma. \\par\\noindent \\footnotemark[$f$] The ionization time scale: the ionization parameter (see table 1) divided by the electron density } \\hss}} \\end{tabular} \\end{center} \\end{table}  \\subsection{Discovery of a New XRN}  We have made a narrow band image at 6.4~keV (the 6.33--6.46 keV band) in figure \\ref{fig:sgrb-6400img}. We see two bright spots in the north. One is Sgr B2 which has been already found as a strong 6.4 keV source \\citep{Koyama1996}, and the other is a newly discovered source. We again referred the same archive data of Chandra and XMM as the case of G0.61+0.01.  In the Chandra data, this excess is found near the edge of the ACIS  FOV. The XMM data show a clear excess near this source. The presence of the 6.4~keV line supports the presence of cool and dense gas clouds. We therefore designate this new source as Suzaku~J1747.7$-$2821.2 (M0.74$-$0.09) from its peak position. We made X-ray spectra of Sgr B2 and M0.74$-$0.09 from the circles given in figure \\ref{fig:sgrb-6400img}. The former is for comparison to the latter new source. The background spectrum is made from the dotted ellipse and subtracted in the same procedure as the case of G0.61+0.01.  \\begin{figure} \\begin{center} \\FigureFile(80mm,80mm){gc_sgrb_6400_publish.eps} \\end{center} \\caption{The 6.4~keV line map (the 6.33--6.46 keV band map) showing bright spots at Sgr B2 and M0.74-0.09. The fluxes are normalized by the 6.4~keV flat-field image. The sources and background regions are shown by the solid circles and dotted ellipse, respectively.}\\label{fig:sgrb-6400img} \\end{figure}  The background-subtracted spectra are shown is figures \\ref{fig:sgrb-B2-spec} and \\ref{fig:sgrb-m074009-spec}. We simultaneously fit the FIs and BI spectra with a model of absorbed power-law plus two Gaussians near at 6.4 and 7.06~keV, which are for the K$\\alpha$ and K$\\beta$ lines of  Fe\\emissiontype{I}. The best-fit parameters are shown in table \\ref{tab:m_fit}. This model nicely fits the data except an excess near the 6.7~keV line in the Sgr B2 spectra. In fact, the 6.7~keV line map (figure \\ref{fig:sgrb-6700img}) shows a weak enhancement at the position of Sgr B2. One possibility is that the 6.7~keV enhancement is a part of the new SNR candidate G0.61+0.01, because it is located in the close vicinity of Sgr B2. The other possibility is that the 6.7~keV enhancement is due to YSOs embedded in the center of Sgr B2. In fact, the Sgr B2 region is relatively crowded  with the Chandra point sources (13 point  sources), and at least some of them are  YSOs  with  a hint of the 6.7~keV line emission \\citep{Takagi2002}.  The total flux (in the 2--10 keV band) of the point sources is $\\sim$10$^{-13}$~ergs~cm$^{-2}$~s$^{-1}$, which is $\\sim$6\\% of the Sgr B2 flux (see table 3).  \\begin{figure} \\begin{center} \\FigureFile(80mm,80mm){gc_sgrb_B2_spec_fi.eps} \\FigureFile(80mm,80mm){gc_sgrb_B2_spec_bi.eps} \\end{center} \\caption{Left: the X-ray spectrum of the sum of the 3 FI CCDs (XIS0, 2 and 3) for Sgr B2 with an absorbed power-law model and two Gaussian lines. Right: same as the right but of the BI CCD (XIS1)} \\label{fig:sgrb-B2-spec} \\end{figure}  The Sgr B2 cloud has been studied extensively with ASCA and Chandra.  \\citet{Koyama1996} and \\citet{Murakami2001} concluded that the 6.4~keV emission is due to  fluorescence by strong X-rays coming from Sgr A$^{*}$, hence named the X-ray reflection nebula (XRN). In this paper, we found a clear K$\\beta$ line at 7.06~keV with consistent flux ratio to the K$\\alpha$ line (6.4~keV) in the fluorescent X-ray origin and deep Fe edge at 7.1~keV. These discoveries provide additional supports for the XRN scenario of Sgr B2. Further details on the Sgr B2 results with Suzaku will be presented in a separate paper.  \\begin{figure} \\begin{center} \\FigureFile(80mm,80mm){gc_sgrb_M074-009_spec_fi.eps} \\FigureFile(80mm,80mm){gc_sgrb_M074-009_spec_bi.eps} \\end{center} \\caption{Same as figure \\ref{fig:sgrb-B2-spec}, but for a new source M0.74-0.09} \\label{fig:sgrb-m074009-spec} \\end{figure}  In the M0.74$-$0.09 region, Miyazaki and Tsuboi (2000) reported flux peaks of the CS (J=1-0) line emission at $(l, b) = (\\timeform{0D.761}, -\\timeform{0D.117})$ and $(\\timeform{0D.764}, -\\timeform{0D.064})$, clear evidence for the presence of a molecular cloud. The XIS spectrum of this region exhibits a strong 6.4~keV line with an equivalent width of 1.6~keV, a 7.06~keV line and an Fe edge structure at 7.1~keV (see table 3). All these features are consistent with being from K$\\alpha$, K$\\beta$ and K-edge from Fe\\emissiontype{I}. The flux of the 7.06~keV line is about 10\\% of that of the 6.4~keV line, which is also consistent with the fluorescent X-ray origin (Kaastra and Mewe 1993). We note here that the background region is the same as the case of G0.61+0.01, hence with the same argument in section 3.2, possible point source contribution can be ignored.  Unlike Sgr B2, no hint of HM YSO is found so far.  No bright point source is found in the Chandra image. Therefore, the X-rays can not be the scattering and fluorescence by embedded YSOs. If the X-rays from Sgr B2 and M0.74$-$0.09 are due to the Thomson scattering and fluorescence of the same irradiating external source like Sgr A$^*$, then the $N_{\\rm H}$ ratio between these sources should be similar to the 6.4~keV line flux ratio. The observed $N_{\\rm H}$ ratio is 0.42, while that of the 6.4~keV line flux is 0.36, in good agreement of the fluorescence scenario by a single irradiation source.  Therefore the XRN scenario by the past activity of Sgr A$^*$, which was successfully applied for Sgr B2 may also be applied for M0.74$-$0.09.  The counter scenario against the XRN is that the 6.4~keV line emission is produced by the collision of electrons. Since the cross section of iron K-shell ionization is maximum at the electron energy of a few 10~keV \\citep{Tatischeff2002}, the most probable source is low energy electrons (LEE) as proposed for the origin of the Galactic Ridge iron K-shell emission \\citep{Valinia2000}. Since a few 10~keV electrons are absorbed in less than $10^{22}$~H~cm$^{-2}$ of depth \\citep{Tatischeff2002} , the produced X-ray spectrum should have no large absorption edge. Our observation, however, shows a clear absorption of (4.0--9.6) $\\times 10^{23}$~H~cm$^{-2}$, in far excess to the Galactic interstellar absorption \\citep{Sakano2002}. Thus the LEE origin is unlikely, unless we assume a special geometry such that the 6. 4keV source is deep in or behind the dense cloud.  \\begin{table*} \\caption{The result of spectral fittings of Sgr B2 and M0.74$-$0.09 with a power-law and two Gaussian models}\\label{tab:m_fit} \\label{tab:m_fit} \\begin{center} \\begin{tabular}{lcc} \\hline\\hline Parameter & Sgr B2 & M0.74$-$0.09 \\\\ \\hline Absorbed power-law model: & & \\\\ Column density $N_{\\rm H}$ ($10^{23}$ cm$^{-2}$) & $9.6_{-0.8}^{+2.5}$ & $4.0_{-1.1}^{+1.4}$ \\\\ Photon index $\\Gamma$ & $3.2_{-0.6}^{+0.9}$ & $1.4_{-0.7}^{+0.4}$ \\\\ Gaussian 1 (Fe\\emissiontype{I} K\\emissiontype{$\\alpha$}): & & \\\\ Line energy (eV) & $6399_{-5}^{+5}$ & $6406_{-6}^{+6}$ \\\\ Intensity ($10^{-5}$ photons cm$^{-2}$ s$^{-1}$) & $16.5_{-0.3}^{+0.8}$ & $5.9_{-1.0}^{+1.4}$ \\\\ Equivalent Width (keV) & 1.13 & 1.55 \\\\ Gaussian 2 (Fe\\emissiontype{I} K\\emissiontype{$\\beta$}): & & \\\\ Line energy (eV)\\footnotemark[$a$] & 7058 & 7065 \\\\ Intensity ($10^{-5}$ photons cm$^{-2}$ s$^{-1}$) & $1.4_{-0.5}^{+0.5}$ & $0.6_{-0.3}^{+0.3}$ \\\\ Equivalent Width (keV) & 0.13 & 0.18 \\\\ \\hline Observed flux\\footnotemark[$\\dagger$] ($10^{-12}$ ergs cm$^{-2}$ s$^{-1}$) & $1.5_{-0.9}^{+0.1}$ & $1.3_{-0.8}^{+0.2}$ \\\\ Luminosity\\footnotemark[$\\ddagger$] ($10^{34}$ ergs s$^{-1}$) & $9.7_{-5.1}^{+0.1}$ & $2.6_{-0.9}^{+0.4}$ \\\\ $\\chi^2$/dof & 154.7/89 & 54.4/66 \\\\ \\hline \\multicolumn{3}{@{}l@{}}{\\hbox to 0pt{\\parbox{180mm}{\\footnotesize Note---The uncertainties indicate the 90\\% confidence limit. \\par\\noindent \\footnotemark[$a$] The energy gap between K$\\alpha$ and K$\\beta$ is fixed at the theoretical value (+659~eV) (Kaastra and Mewe 1993). \\par\\noindent \\footnotemark[$\\dagger$] Observed flux in the 4.0--10.0~keV band. \\par\\noindent \\footnotemark[$\\ddagger$] Absorption corrected luminosity in the 4.0--10.0~keV band. }\\hss}} \\end{tabular} \\end{center} \\end{table*}"
    },
    "0609/astro-ph0609126_arXiv.txt": {
        "abstract": "We outline the WIMP dark matter parameter space in the Constrained MSSM by performing a comprehensive statistical analysis that compares with experimental data predicted superpartner masses and other collider observables as well as a cold dark matter abundance. We include uncertainties arising from theoretical approximations as well as from residual experimental errors on relevant Standard Model parameters.  We present high--probability regions for neutralino dark matter direct detection cross section, and we find that $ 10^{-10}\\pb \\lsim \\sigsip \\lsim 10^{-8}\\pb$ for direct WIMP detection (with details slightly dependent on the assumptions made). We highlight a complementarity between LHC and WIMP dark matter searches in exploring the CMSSM parameter space. We conclude that most of the $95\\%$ probability region for the cross section will be explored by future one--tonne detectors, that will therefore cover most of the currently favoured region of parameter space. ",
        "introduction": "\\label{sec:intro}  Two of the most challenging questions facing particle physics today are the instability of the Higgs mass against radiative corrections (known as the ``fine--tuning problem'') and the nature of dark matter. Unlike the Standard Model (SM), weak scale softly broken supersymmetry (SUSY) provides solutions to both of them. Firstly, the fine--tuning problem is addressed via the cancellation of quadratic divergences in the radiative corrections to the Higgs mass.  Secondly, assuming $R$--parity, the lightest supersymmetric particle (LSP) is a leading weakly interactive massive particle (WIMP) candidate for cold dark matter (CDM). Despite these and other attractive features, without a reference to grand unified theories (GUTs), low energy SUSY models suffer from the lack of predictivity due to a large number of free parameters (\\eg, over 120 in the Minimal Supersymmetric Standard Model (MSSM)), most of which arise from the SUSY breaking sector.  The MSSM with one particularly popular choice of universal boundary conditions at the grand unification scale is called the Constrained Minimal Supersymmetric Standard Model (CMSSM)~\\cite{kkrw94}. The CMSSM is defined in terms of five free parameters: common scalar ($\\mzero$), gaugino ($\\mhalf$) and tri--linear ($\\azero$) mass parameters (all specified at the GUT scale) plus the ratio of Higgs vacuum expectation values $\\tanb$ and $\\text{sign}(\\mu)$, where $\\mu$ is the Higgs/higgsino mass parameter whose square is computed from the conditions of radiative electroweak symmetry breaking (EWSB). The economy of parameters in this scheme makes it a useful tool for exploring SUSY phenomenology.  Many studies have explored the phenomenology of the CMSSM or other SUSY models, mostly by evaluating the goodness--of--fit of points scanned using fixed grids in parameter space, see \\eg~\\cite{grid1,grid2,grid3,grid4,grid5,grid6}. However, this approach has several severe limitations.  Firstly, the number of points required scales as $k^N$, where $N$ is the number of the model's parameters and $k$ the number of points for each of them. Therefore this approach becomes highly inefficient for exploring with sufficient resolution  parameter spaces of even modest dimensionality, say $N>3$. Secondly, narrow ``wedges'' and similar features of parameter space can easily be missed by not setting a fine enough resolution (which, on the other hand, may be completely unnecessary outside such special regions). Thirdly, extra sources of uncertainties (\\eg, those due to the lack of precise knowledge of SM parameter values) and relevant external information (\\eg, about the parameter range) are difficult to accommodate in this scheme.  In this work we report results from a Bayesian exploration of the CMSSM parameter space, obtained through the use of Markov Chain Monte Carlo methods. In particular we focus on the prospects for direct neutralino dark matter detection with the next generation of dark matter searches. We refer the reader to \\cite{rrt1} for full details. The Bayesian approach has several technical and statistical advantages over the more traditional fixed--grid scan technique, the most important being perhaps the ability to incorporate all relevant sources of uncertainties, \\eg the residual uncertainty in the value of SM parameters. This means that the inferred high probability regions in terms of \\eg neutralino mass and scattering cross section take fully into account all sources of uncertainty relevant to the problem. For other works applying a similar approach to the CMSSM, see~\\cite{bg04} (with some relevant differences) and more recently \\cite{al05,allanach06}. ",
        "conclusions": "We have presented a detailed investigation of the prospects for dark matter detection in the framework of the CMSSM parameter using state--of--the--art Bayesian methods. The power and flexibility of the approach allows to probe many previously unexplored ranges of parameters and to fully incorporate the effects of remaining uncertainties in relevant SM parameters and other theoretical uncertainties in computing mass spectra and observables.  We have shown that the WIMP dark matter direct detection elastic scattering cross section $\\sigsip$ presents a wide spread of values (below today's limits) at around $10^{-9\\pm1}\\pb$ and a strong anticorrelation with $\\mchi$. In addition, a region at relatively large $\\sigsip\\simeq 1.6\\times10^{-8}\\pb$, and fairly independent of $\\mchi$, appears to be a feature of the focus point region (despite large theoretical uncertainties) and will be the first to be tested in direct detection experiments.  Despite the remaining uncertainties arising from the local dark matter density profile and the large errors on nuclear form factors, it seems that a large fraction of the high--probability parameter space of neutralino dark matter in the CMSSM will be within reach of currently running upgraded dark matter detector, and most of it will be explored by future one--tonne detectors. We have highlighted the complementarity of these searches with direct SUSY searches at colliders and in particular showed that the focus point region, mostly outside the reach of the LHC, will be among the first to be probed by direct dark matter detection experiments. Those results are largely robust with respect to changes in the {\\em a priori} allowed range for the parameters or to the exclusion of the anomalous magnetic moment of the muon measurement from the analysis.  {\\em Acknowledgements:} R.T.\\ is supported by the Royal Astronomical Society through the Sir Norman Lockyer Fellowship. R.Rda is supported by the program ``Juan de la Cierva'' of the Ministerio de Educaci\\'{o}n y Ciencia of Spain. R.Rda and R.T. would like to thank the European Network of Theoretical Astroparticle Physics ILIAS/N6 under contract number RII3-CT-2004-506222 for financial support."
    },
    "0609/astro-ph0609256_arXiv.txt": {
        "abstract": "The nonlinear mean-field dynamo due to a shear-current effect in a nonhelical homogeneous turbulence with a mean velocity shear is discussed. The transport of magnetic helicity as a dynamical nonlinearity is taken into account. The shear-current effect is associated with the ${\\bf W} {\\bf \\times} {\\bf J}$ term in the mean electromotive force, where ${\\bf W}$ is the mean vorticity due to the large-scale shear motions and ${\\bf J}$ is the mean electric current. This effect causes the generation of large-scale magnetic field in a turbulence with large hydrodynamic and magnetic Reynolds numbers. The dynamo action due to the shear-current effect depends on the spatial scaling of the correlation time $\\tau(k)$ of the background turbulence, where $k$ is the wave number. For Kolmogorov scaling, $\\tau(k) \\propto k^{-2/3}$, the dynamo instability occurs, while when $\\tau(k) \\propto k^{-2}$ (small hydrodynamic and magnetic Reynolds numbers) there is no the dynamo action in a sheared nonhelical turbulence. The magnetic helicity flux strongly affects the magnetic field dynamics in the nonlinear stage of the dynamo action. Numerical solutions of the nonlinear mean-field dynamo equations which take into account the shear-current effect, show that if the magnetic helicity flux is not small, the saturated level of the mean magnetic field is of the order of the equipartition field determined by the turbulent kinetic energy. Turbulence with a large-scale velocity shear is a universal feature in astrophysics, and the obtained results can be important for elucidation of origin of the large-scale magnetic fields in astrophysical sheared turbulence.  \\bigskip  {\\it Keywords:} Nonlinear shear-current dynamo; Sheared turbulent flow; Magnetic helicity transport ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609060_arXiv.txt": {
        "abstract": "We review in these notes the status of the construction of the Pierre Auger Observatory and present the first Physics results, based on the data collected during the first year and a half of operation. These results are preliminary, once the work to understand the systematics of the detectors are still underway. We discuss the cosmic ray spectrum above 3 EeV, based on the measurement done using the Surface Detector and the Fluorescence Detector, both, components of the observatory.  We discuss, as well, the search for anisotropy near the Galactic Center and the limit on the photon fraction at the highest energies. ",
        "introduction": "The cosmic ray spectrum, measured at the top of the atmosphere covers a huge range in energy, going from 10~MeV to energies above $10^{20}$ eV, with a differential flux that spans 31 decades \\cite{nag00}. The techniques to survey this spectrum goes from instruments aboard satellite flights, balloon borne detectors, to counters that monitor the fluxes of neutrons and muons at the Earth surface, and at higher energies to wide area arrays of particle detectors. The spectrum can be divided roughly into four regions with very distinct behavior. The first one, with energies below 1~GeV, has a very distinctive character from the rest of it. Its shape and cut-off is strongly dependent on the phase of the solar cycle, a phenomenon known as {\\it solar modulation}. Actually, there is an inverse correlation between the intensity of cosmic rays at the top of the atmosphere and the level of the solar activity \\cite{lon92,she85}.  The region above 1~GeV show a spectrum with a power law dependence, $N(E)$~d$E=K~E^{-x}$~d$E$, where the spectral index $x$ varies as $2.7~<~x~<~3.2$. The region between 1~GeV and the {\\it knee} region at $4\\times10^{15}$~eV,  is characterized by an index $x~\\simeq~2.7$. These cosmic rays most likely are produced at supernova explosions and their remnants  \\cite{aha04}.  At the knee ($4\\times10^{15}$~eV) the power law index steepens to 3.2 until the so called {\\em ankle}, at $5\\times10^{18}$~eV. The origin of the cosmic rays in this region is less clear and subject of much conjecture. Above the ankle the spectrum flattens again to an index $x~\\simeq~2.8$ and this is interpreted by many authors as a cross over from the steeper galactic component to a harder extra galactic source for the cosmic rays  \\cite{cro99,oli00}.  The existence of cosmic rays with energies above $10^{20}$~eV presents a puzzling problem in high energy astrophysics. The first event of this class of phenomena was observed at the beginning of the 60's by John Linsley at the Volcano Ranch experiment, in New Mexico  \\cite{lin63,lin63a}. Since then, many events were detected, at different sites, using quite distinct techniques \\cite{law91,efi91,bir93,yos97,tak98}. If those cosmic rays are common matter, that is, protons, nuclei or even photons, they undergo  well known nuclear and electromagnetic processes, during their propagation through space. Their energies are degraded by the interaction with the cosmic radiation background ({\\sc cmb}), way before they reached the Earth  \\cite{gre66,zat66}. After travelling distances at the scale of 50~Mpc, their energies should be under $10^{20}$~eV, thus restricting the possible conventional astrophysical objects, which could be sources of them, and those should be easily localized through astronomical instruments. On the other hand, it is very difficult to explain the acceleration of charged particles, with energies up to $3~\\times~10^{20}$~eV or even greater, on known astrophysical objects, by means of electromagnetic forces, the only conventional ones capable of long range and long periods of acceleration. ",
        "conclusions": "The Pierre Auger Observatory is still under construction but has already the largest integrated exposure to high energy cosmic rays. The combination of fluorescence and the surface detector measurements allow for the reconstruction of the shower geometry and its energy with much greater quality than what could be achieved with either detector standing alone. Each of the detectors have different systematics, allowing for valuable information for cross-checking the results from each of them.  The observatory should be finished by mid 2007, accumulating by then a much larger exposure than what was used for the preliminary results presented here. That will allow for the search of anisotropies in the southern sky, as well as the test of the predicted GZK suppression."
    },
    "0609/astro-ph0609511_arXiv.txt": {
        "abstract": "Even quiescent solar prominences may become active and sometimes erupt. These events are occasionally linked to coronal mass ejections. However we know very little about the plasma properties during the activation and eruption processes. We present new computations of the helium line profiles emitted by an eruptive prominence. The prominence is modelled as a plane-parallel slab standing vertically above the solar surface and moving upward as a solid body. The helium spectrum is computed with a non local thermodynamic equilibrium radiative transfer code. The effect of Doppler dimming / brightening is investigated in the resonance lines of \\ion{He}{i} and \\ion{He}{ii} formed in the EUV, as well as on the \\ion{He}{i} $\\lambda\\lambda$\\,10830 \\AA\\ and 5876 \\AA\\ lines. We focus on the line profile properties and the resulting integrated intensities. It is shown that the helium lines are very sensitive to Doppler dimming effects.  We also study the effect of frequency redistribution in the formation mechanisms of the resonance lines and find that it is necessary to use partial redistribution in frequency for the resonance lines. ",
        "introduction": "We use our non local thermodynamic equilibrium radiative transfer code where the prominence is modelled as a plane-parallel slab standing vertically above the solar surface and moving upward as a solid body ({Fig.~\\ref{protu}}). We consider isothermal and isobaric prominence atmospheres only. Other computational details are given in \\cite{lgv06} (astro-ph/0608221).  \\begin{figure} \\centering \\resizebox{0.9\\hsize}{!}{\\includegraphics{protu-dop.eps}} \\caption{Representation of our 1D plane-parallel prominence. The incident radiation coming from the solar disk depends on the plasma radial velocity $\\vec{V}$ perpendicular to the line-of-sight.} \\label{protu} \\end{figure} ",
        "conclusions": "{Partial redistribution in frequency is necessary} to compute the profiles of the helium resonance lines emitted by moving material in prominences. {Velocity effects are more visible when thermal emission is low.} Together with hydrogen lines the helium lines offer the possibility of a powerful diagnostic of the eruptive prominence plasma. In a future work, we will infer {the full velocity vector} by combining {Doppler dimming / brightening effects} on hydrogen and helium lines with the {apparent motion} of the prominence material brought by SOHO or future imagers (e.g. SOLAR-B, STEREO)."
    },
    "0609/astro-ph0609727_arXiv.txt": {
        "abstract": "Using the {\\em Hubble Space Telescope}, we have obtained the first color-magnitude diagram (CMD) to reach the main-sequence turnoff of the Galactic globular cluster NGC~6388.  From a comparison between the cluster CMD and 47~Tucanae's, we find that the bulk of the stars in these two clusters have nearly the same age and chemical composition. On the other hand, our results indicate that the blue horizontal branch and RR Lyrae components in NGC~6388 are intrinsically overluminous, which must be due to one or more, still undetermined, non-canonical second parameter(s) affecting a relatively minor fraction of the stars in NGC~6388. ",
        "introduction": "NGC~6388 and NGC~6441 are two of the most intriguing Galactic globular clusters (GC's). The integrated-light study by \\citet*{mrea93} revealed a strong far-UV flux for these metal-rich \\citep*[${\\rm [Fe/H]} \\simeq -0.60$ and $-0.53$, respectively;][]{h96} bulge GC's. The far-UV flux in resolved GC's is dominated by hot horizontal branch (HB) stars \\citep*[e.g.,][]{bdea95}, especially when rare UV-bright post-asymptotic giant branch stars are not present. Accordingly, the most likely explanation for the far-UV flux in NGC~6388 and NGC~6441 was immediately recognized to be hot HB stars. However, while the UV-upturn phenomenon in elliptical galaxies is often attributed to blue HB stars \\citep[e.g.,][]{bdea95}, no resolved metal-rich GC had been known with a blue HB morphology. The tendency for metal-rich GC's to have red HB's while metal-poor GC's have predominantly blue HB's reflects the ``first parameter'' of HB morphology. As a consequence, the presence of blue HB stars in NGC~6388 and NGC~6441 would represent an example of the so-called ``second-parameter (2$^{\\rm nd}$P) phenomenon.''  The presence of blue HB stars extending almost as faint in $V$ as the turnoff (TO) point in both NGC~6388 and NGC~6441 was confirmed by \\citet{mrea97}, who presented {\\em Hubble Space Telescope} ({\\em HST}) photometry for both these GC's from the survey by \\citet{gpea02}. A remarkable feature of the published diagrams is the presence of a {\\em strongly sloped} HB at colors where other GC's have a much more nearly ``horizontal'' HB. As emphasized by \\citet[][, hereafter SC98]{sc98}, such a sloped HB cannot be simply the result of an older age or of enhanced mass loss along the red giant branch (RGB): while these are able to move a star horizontally along the HB, neither is able to increase the luminosity of a blue HB star compared to the red HB or RR Lyrae stars. Likewise, while strong differential reddening might explain the sloping HB of a red HB cluster, it obviously cannot produce RR Lyrae and blue HB stars. SC98 conclude therefore that {\\em non-canonical} 2$^{\\rm nd}$P candidates must be at play in NGC~6388 and NGC~6441.  However, these conclusions were challenged by \\citet[][, hereafter R02]{grea02}, who computed models with non-standard values of the chemical abundance and mixing length parameter. Some of their models did reveal sloped HB's, but only as a consequence of an anomalously {\\em faint red HB} (in $V$), together with a blue HB having a $V$-band luminosity consistent with the canonical models (see, e.g., their Fig.~2).  Additional insights are provided by stellar variability and spectroscopic studies. \\citet{nsea94}, \\citet{alea99}, \\citet{pea00,pea01,pea02,pea03}, and \\citet{mcea06} have shown that the RR Lyrae variable stars in these GC's, which occupy the normally ``horizontal'' part of the HB, have much longer periods than field RR Lyrae of similar metallicity, thus strongly suggesting that they are intrinsically more luminous (SC98). Moreover, theoretical calculations by \\citet{pea02} have shown that, contrary to the suggestions by \\citet{crea02}, the RR Lyrae components in both clusters cannot be explained in terms of evolution away from a position on the blue HB---and neither can the sloping nature of the HB be reproduced in this way. On the other hand, the first spectroscopic measurements of the gravities of blue HB stars in both NGC~6388 and NGC~6441 \\citep*{smea99} revealed surface gravities that are higher than predicted by even the canonical models, thus arguing against an anomalously bright blue HB~+~RR Lyrae component in these clusters. However, a recent reassessment of the spectroscopic gravities of blue HB stars in NGC~6388 by \\citet{ms06} indicates that the actual gravities should, in fact, be lower than the canonical values. It appears that the 1999 values must have been in error by a substantial amount, probably due to unresolved blends in the crowded inner regions of these massive \\citep[$M_V \\sim -9.5$;][]{h96} GC's.  In an effort to shed light on this puzzling situation, we have made use of the data obtained for NGC~6388 in the course of our snapshot {\\em HST} program to study stellar variability in the cluster, and also of archival data, to produce its deepest-ever color-magnitude diagram (CMD). In \\S2 we describe this dataset and reduction procedures. In \\S3 we compare our CMD with 47~Tucanae's (NGC~104). We close in \\S4 by discussing the implications of our results for our understanding of the origin of the peculiar HB morphology of NGC~6388.    \\begin{figure}[t] \\figurenum{1} \\begin{center} \\includegraphics*[width=3.4in]{f1.eps} \\caption{The first deep CMD of NGC~6388 clearly revealing the TO point of the cluster ({\\em red $\\times$ symbols}) is overplotted on the 47~Tuc CMD ({\\em plus signs}) in the $V$, $B\\!-\\!I$ plane. To produce this plot, the 47~Tuc CMD was shifted by +3.2 in $V$ and by +0.85 in $B\\!-\\!I$, so as to align their red HB components. } \\end{center} \\label{fig:cmd} \\end{figure} ",
        "conclusions": "In the present {\\em Letter}, we have shown that, apart from the blue HB and RR Lyrae components, the CMD's of NGC~6388 and 47~Tuc are very similar, thus strongly suggesting that the bulk of the stars in these two clusters are very similar. Differences in age, metallicity, $[\\alpha/{\\rm Fe}]$, $Y$, and $\\alpha_{\\rm MLT}$ between the two GC's, if present, should be small. Our results suggest that the red HB component of NGC~6388 is neither underluminous (as in the ``canonical tilt'' scenario of R02) nor overluminous (as in the high-$Y$ scenario of SC98), except for its bluest stars. On the other hand, both the RR Lyrae and blue HB stars in NGC~6388 {\\em are} significantly overluminous compared to field RR Lyrae stars of similar [Fe/H], thus explaining NGC~6388's anomalously long RR Lyrae periods. The lack of a sizeable luminosity difference between the red HB's of NGC~6388 and 47~Tuc indicates that {\\em the 2$^{\\it nd}$P that leads to the production of an overluminous RR Lyrae~+~blue HB component in NGC~6388 must be non-canonical in nature}---that is, it must be neither age nor RGB mass loss, or else a sloped HB would not result (SC98). However, since only a relatively small fraction ($\\simeq 17\\%$; \\citeauthor{mc06} \\citeyear{mc06}) of the HB stars are in the RR Lyrae strip or on the blue HB, only a minor fraction of the cluster stars should be affected by this non-canonical 2$^{\\rm nd}$P, which could be either $Y$ or/and the helium-core mass at the He-flash. (A fraction of the red HB stars, especially the bluer 25\\% or so, could also be affected, though to a much smaller extent.) Additional, detailed studies of the blue HB, RR Lyrae, and main sequence components will be required before we are in a position to conclusively decide what 2$^{\\rm nd}$P(s) is (are) responsible \\citep[see][ for a recent review]{mc06}."
    },
    "0609/astro-ph0609661_arXiv.txt": {
        "abstract": "We report observations of gaseous methanol in an edge-on torus surrounding the young stellar object L1551 IRS5. The peaks in the torus are separated by $\\sim$ 10,000 AU from L1551 IRS5, and contain $\\sim$ 0.03 $M_{\\rm \\earth}$ of cold CH$_3$OH. We infer that the CH$_3$OH abundance increases in the outer part of the torus, probably as a result of methanol evaporation from dust grain surfaces heated by the shock luminosity associated with the shocks associated with the jets of an externally located x-ray source. Any methanol released in such a cold environment will rapidly freeze again, spreading CH$_3$OH throughout the circumbinary torus to nascent dust grains, planitesimals, and primitive bodies. These observations probe the initial chemical conditions of matter infalling onto the disk. ",
        "introduction": "In the earliest stages of their formation, low-mass young stellar objects are embedded in  flattened gaseous envelopes. Surrounding these nascent protostars, planetary systems, primitive bodies and their attendant dust grains will condense from the infalling material. In this paper, we study the distribution of the organic molecule methanol around a protostellar core, as it is an important constituent of the young material in the disk. The L1551 molecular core (distance 140 pc) contains a young binary system (seen as an infrared and radio source L1551 IRS-5) hidden inside a dense envelope providing $\\sim$ 150 magnitudes of visual extinction (Stocke \\al 1988, White \\al 2000, Fridlund \\al 2005). A molecular outflow emanates from the center of the disk (Snell \\al 1980, Kaifu \\al 1984, Fridlund \\& White 1989, Rainey \\al 1987, Parker \\al 1991), with atomic jet(s) associated with an x-ray emitting region being observed to the SW (Fridlund \\& Liseau 1998, Favata \\al 2002, 2003, Bally \\al 2003, Fridlund \\al 2005). The protostellar disk is surrounded by a massive (radius $\\sim$ 20,000 AU) cool envelope that exhibits both rotational and infall motions (Takakuwa \\al 2004, Moriarty-Schieven \\al 2006). Fridlund \\al (2002) used HCO$^+$, H$^{13}$CO$^+$ and $^{13}$CO \\j10 observations to estimate that the mass of the disk is $\\sim$ 2.5$\\pm$1.5 \\Msun. ",
        "conclusions": ""
    },
    "0609/astro-ph0609382_arXiv.txt": {
        "abstract": "{} {In this work we investigate the dynamic behavior of inter-network regions of the solar chromosphere.} {We observed the chromosphere of the quiet Sun using a narrow-band Lyot filter centered at the \\mbox{{\\ion{Ca}{ii}\\,K$_\\mathrm{2v}$}} emission peak with a bandpass of 0.3\\,\\AA. We achieved a spatial resolution of on average 0\\farcs 7 at a cadence of 10\\,s.} {In the inter-network we find a mesh-like pattern that features bright grains at the vertices. The pattern has a typical spatial scale of 1\\farcs 95 and a mean evolution time scale of \\mbox{53\\,s} with a standard deviation of \\mbox{10\\,s}. A comparison of our results with a recent three-dimensional radiation hydrodynamical model implies that the observed pattern is of chromospheric origin. The measured time scales are not compatible with those of reversed granulation in the photosphere although the appearance is similar. A direct comparison between network and inter-network structure shows that their typical time scales differ by at least a factor of two.} {The existence of a rapidly evolving small-scale pattern in the inter-network regions supports the picture of the lower chromosphere as a highly dynamical and intermittent phenomenon.} ",
        "introduction": "Deducing the thermal structure of the solar chromosphere is complicated by the limited number of diagnostics and the difficulties to interpret them. An example are the core regions of the calcium resonance lines (\\mbox{\\ion{Ca}{ii}\\,H and K}) that are commonly assumed to originate from the chromosphere \\citep[cf.][]{val81}. Unfortunately the simplifying assumption of local thermodynamic equilibrium (LTE) is not valid for the emission reversal peaks and the cores of these lines. Rather, non-local effects must be taken into account for a realistic description of chromospheric plasma and radiation field. This is not only a problem for the interpretation of observations of these spectral features but also for the numerical modeling of the chromosphere. The recent three-dimensional radiation hydrodynamics model by \\citet[][hereafter W04]{wedemeyer04a} shows a mesh-like pattern at chromospheric heights in non-magnetic inter-network regions, consisting of cool regions and hot filaments that are a product of propagation and interaction of oblique shock waves and adiabatic expansion of post-shock regions. The authors argue that the amplitude of the temperature variations in the model chromosphere is somewhat uncertain due to the too simple radiative transfer (grey, LTE). But the pattern is nevertheless representative as it is a product of acoustic waves that are excited in the lower well-modeled layers. The spatial scales of the pattern are comparable to the underlying granulation but the dynamical time scales are much shorter. Mapping the chromospheric small-scale structure of inter-network regions thus requires both high spatial and temporal resolution. This might be the reason why most earlier observations failed to detect this phenomenon (Table~\\ref{tbl:cadata}).  Here we present new observations taken with a Lyot filter of very narrow transmitted wavelength range and a spatial resolution high enough to resolve the chromospheric small-scale pattern of inter-network regions.  In Sect.~\\ref{sec:obs} we describe the observations and data reduction, followed by results and discussion in Sects.~\\ref{sec:result} and \\ref{sec:discuss}, respectively.   \\begin{figure*}[htp] \\centering \\includegraphics[width=18cm]{woeger_1.eps} \\caption{(A) Calcium image with superimposed network mask and closeup regions (numbered squares), (B) the closeups of the network region and (C) the inter-network region, (D) the MDI magnetogram, (E) a  2-D power map of the temporal frequency band at \\mbox{5.5\\,mHz} with a width of \\mbox{0.9\\,mHz} and (F) the map of \\mbox{1/e} time scales. All axes are in arcseconds. } \\label{fig:caimage} \\end{figure*}  \\begin{table}[pb] \\centering \\caption{Recent observations of the \\ion{Ca}{ii} resonance lines. The properties of the data show that the presented data set is unique regarding the combination of spectral and spatial resolution (\\mbox{$\\Delta\\lambda$} is FWHM). } \\begin{tabular}{lrrrr} \\hline Author & $\\lambda$ & $\\Delta\\lambda$ & scale        & $\\Delta$t\\\\ & $[\\AA]$   & $[\\AA]$         & [\\arcsec/px] & [s]    \\\\ \\hline this work & 3933.5 & 0.30 & 0.1485 & 10\\\\ \\cite{rutten04dot2} & 3968.8 & 1.35 & 0.071 & 30\\\\ \\cite{tritschler02} & 3933.2 & 0.60 & 0.350 & 6\\\\ \\cite{depontieu02} & 3933.0 & 3.00 & 0.083 & $\\sim$20\\\\ \\cite{kneer93} & 3933.0 & 0.60 & 0.720 & 15\\\\ \\hline \\end{tabular} \\label{tbl:cadata} \\end{table} ",
        "conclusions": "\\label{sec:discuss}  We provide a new method to discriminate network and inter-network structures in filtergrams which show chromospheric structure when the data are not accompanied by information about the underlying magnetic field.  Figure \\ref{fig:time_int} demonstrates that the pattern evolution time scale provides an excellent criterion to distinguish between network and inter-network regions: all inter-network points are grouped in the time interval below \\mbox{100\\,s}, whereas the network has time scales above \\mbox{200\\,s}.  Temporal thresholding delivers an accurate mask, M2, for network pixels.  Structures in the inter-network similar to those found in our observations have already been detected in various other works. \\citet[][hereafter LW05]{leenaarts05} investigate a similar pattern at higher spatial resolution -- but lower spectral resolution -- and identified it with reversed granulation.  This result is also supported by the much longer time scale of 2\\,min they determine for the pattern.  In Fig.~1 of \\citet{kneer93} the pattern can be surmised. \\citet{tritschler02} present observational data with time scales of inter-network structures close to the ones obtained in this work. Their spatial resolution did not suffice to resolve the small-scale pattern and they used a broader bandwidth of 0.6\\,\\AA.  This work is the first one to observe the small-scale pattern with a mean evolution time scale much shorter than \\mbox{100\\,s}.  In order to verify the chromospheric origin of the detected inter-network pattern, we compare its spatial and evolution time scale to the recent 3D model by W04.  The temperature amplitudes of the model chromosphere might be somewhat uncertain due to the simplified treatment (grey, LTE), whereas the excitation of the waves occurs in realistically modelled lower layers, making the dynamics and related time scales more reliable.  A direct comparison with synthetic intensity maps for \\mbox{\\ion{Ca}{ii}\\,K$_\\mathrm{2v}$} is not possible yet as the model chromosphere is dominated by strong shocks that cause numerical problems. We can, however, perform a qualitative comparison with horizontal slices of gas temperature from the model chromosphere. The cuts exhibit a pattern of granular size at a geometrical height around \\mbox{1000\\,km} above \\mbox{$\\tau$\\,=\\,1} which is rapidly evolving with time scales of around \\mbox{$\\tpe{} = 20$ - 30\\,s}. The pattern is quite similar to the one we observe, but the evolution time scale is still a factor of two shorter than the observational value. This can be explained as follows: \\begin{figure}[tp] \\centering \\includegraphics[width=8.5cm]{woeger_4.eps} \\caption{Time scale versus intensity. The gap between 100 and \\mbox{200\\,s} indicated with a shaded horizontal band shows the different behavior of network and inter-network structures.  } \\label{fig:time_int} \\end{figure} First of all, the correlation between gas temperature at \\mbox{$z = 1000$\\,km} and intensity in the \\mbox{\\ion{Ca}{ii}\\,K$_\\mathrm{2v}$} spectral feature might be poor as the assumption of LTE is not valid in this strongly scattering line.  In LTE local changes in temperature are instantaneously translated into changes of the source function and thus intensity changes although the intensity response is non-linear in amplitude.  In contrast to the temperature slices at a well-defined geometrical height, the emergent intensity (at a given wavelength) is a non-local quantity that has contributions from an extended height range.  For the line core, the main contribution of intensity originates from chromospheric heights around \\mbox{1000 - 1500\\,km} above \\mbox{$\\tau$\\,=\\,1} \\citep[cf.][]{val81}, but already at \\mbox{0.15\\,\\AA} off-core considerable contributions to the filtergram intensity from lower geometrical heights have to be expected (Rammacher, priv. comm.).  The emergent intensity is thus integrated along the line of sight over many layers which individually would show different pattern evolution times \\citep[decreasing with height, see][]{wedemeyerphd}.  The integrated pattern is consequently smoothed so that the relative intensity variations become smaller and result in longer evolution time scales. This effect is even enhanced due to the fact that the observed intensity refers to corrugated surfaces of optical depth instead of plane horizontal cuts as used for the model.  The intensity pattern will be smoothed even more compared to the fine structure in the temperature cuts.  W04 calculate a time scale of $\\tpe{, z} >$\\,120\\,s for changes in gas temperature in a plane horizontal cut at the bottom of the photosphere, whereas the grey emergent intensity (under assumption of LTE, which reflects conditions at the bottom of the photosphere), yields an increase of the typical time scale to $\\tpe{, \\tau} = 200$\\,s.  This is due to the fact that variations tend to be smaller on surfaces of optical depth because these are shaped by inhomogeneities themselves. In addition, seeing and instrumental effects lead to a spatial smearing of the intensity signal and thus to longer time scales. We found that smoothing the temperature cuts to the spatial resolution encountered in our data increases the time scales by about 25\\%. The simulations of W04 do not include magnetic field. The existence of inter-network fields (with intrinsic field strengths varying from \\mbox{1\\,kG} down to \\mbox{$\\la$100\\,G}) is undisputed as they have been recently measured with high polarimetric precision \\citep[see ][and references therein]{socas+lites+pillet2004, lites+socas2004}. The contribution of possibly unresolved magnetic elements in inter-network regions tends to lengthen the effective evolution time scale. A preliminary analysis of temperature cuts from the 3D magnetohydrodynamics model by \\citet{schaffenberger05} results in a slightly higher value of $\\tpe{,MHD} = 27$\\,s at the same height ($z = 1000$\\,km) compared to 24\\,s in the purely hydrodynamical model by W04. \\begin{figure}[tp] \\centering \\includegraphics[width=8.5cm]{woeger_5.eps} \\caption{Temporal power for network (dashed) and inter-network (solid).} \\label{fig:power1d} \\end{figure} Consideration of all these effects leads to a reasonable agreement between observation and numerical model. An interpretation of the pattern as reversed granulation is inappropriate as the corresponding time scale is much longer than the 53\\,s found here.  LW05 state 120\\,s and 90\\,s for observed and synthetic \\mbox{\\ion{Ca}{ii}\\,H} line wing intensities, respectively, which clearly originate from the mid-photosphere \\citep[LW05,][]{wedemeyerphd}.  These values are in line with the time scale of order of 90\\,s for the reversed granulation pattern observed in the Fe\\,I line core at 709.0\\,nm by \\citet{janssen06}. We conclude that the small-scale inter-network pattern reported here evolves too fast to be due to reversed granulation and is thus of chromospheric origin.  The pattern is very likely the intensity signature of the mesh-like, fast evolving pattern found in recent 3D hydrodynamical models of the solar chromosphere.  It is presumably a product of propagation and interaction of shock waves (W04) that are excited below the observed layer. The same is true for the \\ion{Ca}{ii} grains \\citep[see also][]{carlsson97a} that we consider an integral part of the pattern.  Collision of two neighbouring wave fronts results in compression and heating of the gas at the contact region which shows up as a band of enhanced intensity, making up the ``strings'' of the mesh-like pattern.  At a mesh vertex, however, more than two wave fronts meet, resulting in stronger compression and thus in an intensity higher than the remaining pattern.  Due to the limitations of earlier observations (see Table~\\ref{tbl:cadata}) only the brightest components have been detected as grains whereas the pattern itself remained undetected. A more detailed study including the simultaneous G-band data is in preparation."
    },
    "0609/astro-ph0609107_arXiv.txt": {
        "abstract": "X-ray observations of hot, intergalactic gas in galaxy groups provide a useful means of characterizing the global properties of groups.  However, X-ray studies of large group samples have typically involved very shallow X-ray exposures or have been based on rather heterogeneous samples. Here we present the first results of the {\\em XI} (XMM/IMACS) Groups Project, a study targeting, for the first time, a redshift-selected, statistically unbiased sample of galaxy groups using deep X-ray data.  Combining this with radio observations of cold gas and optical imaging and spectroscopy of the galaxy population, the project aims to advance the understanding of how the properties and dynamics of group galaxies relate to global group properties.  Here, X-ray and optical data of the first four galaxy groups observed as part of the project are presented.  In two of the groups we detect diffuse emission with a luminosity of $L_{\\rm X} \\approx 10^{41}$~erg~s$^{-1}$, among the lowest found for any X-ray detected group thus far, with a comparable upper limit for the other two.  Compared to typical X-ray selected groups of similar velocity dispersion, these four systems are all surprisingly X-ray faint.  We discuss possible explanations for the lack of significant X-ray emission in the groups, concluding that these systems are most likely collapsing for the first time.  Our results strongly suggest that, unlike our current optically selected sample, previous X-ray selected group samples represented a biased picture of the group population. This underlines the necessity of a study of this kind, if one is to reach an unbiased census of the properties of galaxy groups and the distribution of baryons in the Universe. ",
        "introduction": "In hierarchical theories of structure formation, structures of progressively increasing size separate out from the Hubble expansion, recollapse and virialise. Groups of galaxies are the characteristic structures which have formed by the present epoch, and are believed to contain the bulk of the matter in the Universe \\citep*{fuku98}. As such, an understanding of the Universe requires an understanding of galaxy groups.  The depth of the potential wells of groups is similar to that of individual galaxies, and galaxy velocities within groups are only a few hundred km~s$^{-1}$. Under these circumstances, galaxies can interact strongly with one another, and with the group potential (e.g.\\ \\citealt{sers74,meno92,verd98,mend03,mile04}).  Hence, not only are groups the most common environment for galaxies, but they provide an environment which has a strong effect on galaxy properties, and which is itself evolving, as groups collapse, virialise and grow through mergers and accretion.  Despite the crucial role played by groups in cosmic structure formation and evolution, they have received relatively little attention compared to larger clusters. This is at least partly due to the fact that typical groups may contain only a handful of bright galaxies in their inner regions. In optical data, groups are therefore often difficult to detect with confidence in projection.  X-ray emission from a hot intragroup medium (IGM) provides a much more reliable method of detecting the potential well of a virialised group. Some earlier works based on pointed {\\em ROSAT} X-ray observations utilized this to study the global properties of groups and those of their member galaxies in a coherent manner \\citep{mulc96,mulc98,zabl98,hels00,hels03,osmo04}.  However, although care had been taken to include groups spanning a wide range of properties, as in the study of \\citet{osmo04}, from a statistical viewpoint these works all suffered from the fact that {\\em ROSAT} only targeted a heterogeneous sample of hand-picked groups.  X-ray investigations of optically selected group samples include those of \\citet{mahd97}, and later \\citet{mahd00}, who invoked data from the {\\em ROSAT} All-Sky Survey (RASS) to study a redshift-selected, unbiased sample of groups in X-rays.  However, no detailed information on the galaxy dynamics within these groups was extracted. \\citet{burn96} also employed RASS data to investigate the X-ray properties of a sample of groups, drawn from the optically selected and statistically complete poor cluster catalogue of \\citet{ledl96} and \\cite{whit99}. This work furthermore included a detailed investigation of the galaxy dynamics within the groups, though the systems were originally selected by photometric enhancements and were not based on velocity information.  RASS data also formed the basis for the search for X-ray gas in the optically selected \\citet{hick82} catalogue of groups conducted by \\citet*{ebel94}, a work later augmented by \\cite{ponm96}, who included pointed ROSAT data.  Even in the latter study, however, the majority (roughly two-thirds) of the groups only had RASS coverage.  Due to the shallow RASS exposures of typically a few hundred seconds, the X-ray properties of groups with $L_{\\rm X} \\la 10^{42}$~erg~$^{-1}$ could generally not be investigated in these studies. Moreover, in many cases the data only allowed for a detection of hot intragroup gas, whereas no detailed characterization of gas properties could be carried out.  A further concern for RASS--based results for the diffuse X-ray luminosity $L_{\\rm X}$ from groups is the potential danger of contamination from active galactic nuclei due to the broad point spread function (PSF) of the {\\em ROSAT} PSPC at large off-axis angles.  With the exception of dedicated large-scale structure surveys (e.g., \\citealt{pier04}), more recent observations with {\\em XMM-Newton} and {\\em Chandra} have largely been targeting groups already studied by {\\em ROSAT} in pointed observations or serendipitously detected in X-rays as part of the RASS. It is not clear to what extent these groups, most of which are X-ray selected, can be seen as representative of the overall group population at low redshift. Studies of clusters of galaxies indicate that X-ray selection could incorporate a serious bias, with X-ray and optically selected samples showing highly disparate X-ray properties and in some cases \\citep{dona02} being largely non-overlapping.  \\citet{gilb04} found that a significant fraction of optically selected clusters do not have a clear X-ray counterpart in {\\em ROSAT} data, even though the physical reality of the X-ray faint systems were confirmed by spectroscopic follow-up. A similar conclusion, although based on optical multi-colour imaging rather than spectroscopy, was reached by \\citet{bark06} using {\\em Chandra} data.  The optically selected high-redshift clusters of \\citet*{lubi04} were found to be X-ray underluminous for their velocity dispersion $\\sigma_{\\rm v}$, deviating from the $L_{\\rm X}$--$\\sigma_{\\rm v}$ relation derived for rich, nearby X-ray clusters.  A similar result was shown for colour-selected clusters by \\citet{hick04}, while \\citet{pope05} found that a subset of their Abell clusters were X-ray underluminous with respect to their virial mass. These results clearly suggest that X-ray selection could provide a biased picture of the cluster population, and there is little reason to assert that the situation for groups should be significantly different. Given the lower X-ray luminosities of groups, such a bias is likely to be even stronger for these systems.  As a consequence of the fact that most X-ray studies of groups have either been concentrating on X-ray selected targets picked from a variety of different group catalogues, or have employed the shallow RASS data to investigate group X-ray emission, we currently have no {\\em unbiased} census of the properties of hot gas and their relation to the dynamics of galaxies within galaxy groups.  The advent of large redshift surveys, coupled with the high sensitivity, spatial resolution, and spectral capability of {\\em XMM-Newton} and {\\em Chandra}, has made it possible to remedy this situation. Groups can be selected in redshift space, the properties of any hot IGM can be assessed using deep X-ray data, and group dynamics can be studied with multi-slit spectrographs.  We have launched a project to take advantage of this situation, by targeting a statistically representative sample of groups with both {\\em XMM-Newton} and the IMACS multi-object camera and spectrograph \\citep{bige03}, installed on the 6.5-m Baade/Magellan telescope at Las Campanas.  The primary aim of the {\\em XI} Project is to understand the nature and evolution of the galaxy group population, and to study the way in which the properties of groups are related to those of their member galaxies. By performing a sensitive test for the presence of any hot IGM in the selected groups, one of the key outcomes of this project will be a reliable estimate of the fraction of optically selected groups which actually contain a hot IGM. Another issue to be addressed is whether the properties of any hot X-ray gas can be used as an evolutionary tracer of the dynamical state of a group.  The thrust of the present paper is to provide an initial attack on these questions, using the results obtained for the first four groups observed by {\\em XMM} as part of the project. The large collecting area and field of view of {\\em XMM} makes it the ideal X-ray instrument for this study. IMACS has a field of view perfectly matched to that of {\\em XMM} ($\\sim 30$~arcmin in both cases), and allows us to obtain several hundred spectra over this entire field in a single exposure, down to $M_B \\approx -15$ at the selected sample redshift of $z\\approx 0.06$.  We assume $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$, and a Hubble constant of $H_0=70$~km~s$^{-1}$~Mpc. The virial radius $r_{\\rm vir}$ of each group can then be identified with $r_{100}$, the radius enclosing a mean density of 100 times the critical density $\\rho_c$ \\citep*{eke96}. If the dark matter density of the groups follows an 'NFW' profile \\citep*{nava95} with concentration parameters in the plausible range $c=5-20$ \\citep{bull01}, we then have $r_{500}\\approx 0.5r_{\\rm vir}$.  In the adopted cosmology, the characteristic redshift of our sample of $z=0.06$ corresponds to a luminosity distance $D\\approx 275$~Mpc, and 1~arcmin to $\\approx 70$~kpc. ",
        "conclusions": "As part of an ongoing effort to investigate the X-ray and optical properties of a substantial, statistically unbiased, and kinematically selected sample of galaxy groups, for the first time using deep X-ray data, we have performed {\\em XMM-Newton} observations of the first four groups in this sample. In two of the groups, we detect an X-ray emitting intragroup medium, with luminosities among the lowest found for X-ray detected groups so far. The two other groups observed here remain undetected in the X-ray, and all four groups are found to be X-ray underluminous for their velocity dispersions when compared to expectations from X-ray bright groups. Furthermore, none of the groups hosts a dominant elliptical galaxy at the centre of the X-ray emission.  Our results therefore suggest that the nature of the IGM in these optically selected groups may be very different from that seen in standard X-ray selected group samples. The fact that we are finding some of the faintest IGM yet seen in groups in our first four observations clearly demonstrates that the {\\em XI} Project is exploring new parameter space. It also establishes that the predominantly X-ray selected group samples studied in the past with e.g.\\ {\\em ROSAT} were not representative of the overall group population.  In contrast, the unbiased redshift selection employed in the {\\em XI} Project appears to be targeting a different class of groups not previously studied in much detail.  The low levels of IGM emission in the groups could be an indication that i) the groups are in the process of collapsing for the first time, ii) the gravitational potentials are too shallow for the gas to emit much X-ray emission (i.e.\\ the gas is essentially too cool to produce X-rays), or iii) there is simply little or no intragroup gas. We find the first explanation to be the more attractive, at least until the validity of the other scenarios can be more firmly addressed by subjecting a larger sample of optically selected groups to deep X-ray observations.  The idea that many groups are collapsing for the first time at the present epoch is also consistent with expectations from other optical group studies and with results of cosmological simulations.  In contrast to the more well-studied class of X-ray bright groups, the type of groups studied in this paper truly represents the most common galaxy environments in the Universe. Investigating the detailed properties of a larger sample of such groups is therefore crucial if one is to obtain an unbiased understanding of the nature of the group population and a census of the distribution of baryons in the Universe."
    },
    "0609/astro-ph0609331_arXiv.txt": {
        "abstract": "We present a new formalism to describe the ratios and profiles of emission lines from hydrogen in Balmer-dominated shocks.  We use this model to interpret the measured widths and ratios of broad and narrow H$\\alpha$, H$\\beta$ and Ly$\\alpha$ emission lines in supernova remnants (SNRs).  Our model results agree fairly well with those obtained previously by Chevalier, Kirshner \\& Raymond (1980) and are consistent with observations of several SNRs.  The same model fails to account for the ratio of broad to narrow line emission from the reverse shock in SNR 1987A as observed by Heng et al. (2006).  We suggest that this discrepancy between theory and observation results from a faulty assumption that Balmer-dominated shocks can be treated as sharp discontinuities.  If the spatial structure of the shock transition zone is taken into account, the predicted ratios of broad to narrow line emission in most SNRs will change by modest factors, but the ratio in SNR 1987A will increase substantially.  Significantly greater shock velocities will be required to account for the observed full widths at half-maximum of the broad emission lines in most SNRs. ",
        "introduction": "\\label{sect:intro} When a fast astrophysical shock enters neutral interstellar gas, we see an optical spectrum that is dominated by H$\\alpha$ and other Balmer lines and is nearly devoid of the forbidden lines typically seen in emission line nebulae.  The Balmer emission is caused by impact excitation of neutral hydrogen atoms by fast ions and electrons in the shocked gas.  The emission lines in such ``Balmer-dominated'' or ``non-radiative'' shocks have two components.  The first is a narrow component with a line width characteristic of the cold interstellar gas, which results from direct excitation of the neutral hydrogen atoms.  The second is a broad component with a line width characteristic of the thermal velocity broadening of the shocked protons.  The broad component results from charge transfer reactions, which pass electrons from the nearly stationary hydrogen atoms to the shocked protons.  The resulting fast hydrogen atoms may be created in excited states or may be excited by subsequent collisions with other fast ions and electrons.  The resulting ratio of the broad to narrow components and the profiles of the broad Balmer lines depend on the equilibration of electron and ion temperatures in the shocked gas, on resonance trapping of the Lyman lines, and on the velocity and the inclination of the shock surface.  The theory to interpret the emission spectra of such Balmer-dominated shocks was developed originally by Chevalier \\& Raymond (1978, hereafter CR78) and Chevalier, Kirshner \\& Raymond (1980, hereafter CKR80) and augmented by Smith et al. (1991, hereafter S91) and Ghavamian et al. (2001, hereafter G01).  It has been used to interpret the emission spectra of several supernova remnants (SNRs), including: SN 1006 (CKR80; Kirshner, Winkler \\& Chevalier 1987, hereafter KWC87; S91; Winkler \\& Long 1997; Ghavamian et al. 2002, hereafter G02); Kepler (Fesen et al. 1989, hereafter F89; Blair, Long \\& Vancura 1991, hereafter BLV91; Sollerman et al. 2003; Sankrit et al. 2005); Tycho/SN1572 (KWC87; S91; Ghavamian et al. 2000; G01), RCW 86 (Long \\& Blair 1990, Smith 1997; G01; S03), portions of the Cygnus Loop (Raymond et al. 1983; Fesen \\& Itoh 1985; Hester, Raymond \\& Blair 1994; G01), and four remnants in the Large Magellanic Cloud (Tuohy et al. 1982; S91; Smith et al. 1994).  Generally, the theoretical models fit the observations fairly well, but significant discrepancies (by factors $\\sim 2$) between the theoretical and observed broad-to-narrow line ratios persist in some cases, as we shall discuss.  Supernova 1987A has provided a dramatic new example of a Balmer-dominated shock.  In contrast to the SNRs cited above, broad Balmer emission from SNR 1987A comes from hydrogen atoms in the rapidly expanding ($\\sim 12,000$ km~s$^{-1}$) supernova debris rather than the stationary interstellar gas.  In this case, the highest velocity component of the emission lines, called ``surface emission,'' results from the excitation of freely-streaming hydrogen atoms that cross the reverse shock. Observations of the profiles of H$\\alpha$ and Ly$\\alpha$ emission from the reverse shock with the Space Telescope Imaging Spectrograph (STIS) have enabled us to measure the rapidly evolving flux of hydrogen atoms across the reverse shock and to reconstruct partially the three-dimensional geometry of the shock surface (Michael et al. 1998a, b, 2003, with the last hereafter M03; Heng et al. 2006, hereafter H06). In addition to the surface emission, H06 identified a lower velocity component of the H$\\alpha$ and Ly$\\alpha$ emission lines, which they called ``interior emission.'' Evidently, this interior emission results from charge transfer of electrons from the freely-streaming hydrogen atoms to protons in the shocked gas and is analogous to the broad component seen in the Balmer lines of other SNRs.  We also note that Balmer-dominated bow-shocks have been observed in four pulsar nebulae (Bucciantini 2002).  Here, we revisit the theory of Balmer-dominated shocks.  We hope to resolve the discrepancies between observations and theory, to understand better how to interpret the line profiles, and to extend the existing theoretical models to the very high shock velocities seen in SNR 1987A.  In \\S\\ref{sect:assumptions}, we state the assumptions of our model and mention their limitations.  In \\S\\ref{sect:rates}, we discuss the cross sections used and present the resulting velocity distribution functions and rate coefficients.  In \\S\\ref{sect:tree}, we present a new formalism for deriving the intensities and the profiles of broad and narrow emission lines.  In \\S\\ref{sect:results}, we compare our results to previous models and to existing data on Balmer-dominated remnants and SNR 1987A.  Finally, in \\S\\ref{sect:discussion}, we discuss the limitations of the model and future development that is needed. ",
        "conclusions": "\\label{sect:discussion}  For most SNRs, the models presented here agree fairly well with previous models and seem to account for the observed data of both broad line FWHM and ratio $\\Re_{b,n}(\\mbox{H}\\alpha) \\approx 1$ of broad-to-narrow line intensities.  However, these same models fail to account for the observed ratio $\\Re_{b,n}(\\mbox{H}\\alpha)$ of interior-to-surface emission for the conditions of the reverse shock in SNR 1987A.  Either the model is wrong, or the interior H$\\alpha$ emission from SNR~1987A is dominated by some excitation mechanism other than charge transfer at the reverse shock.  As we mentioned in \\S\\ref{sect:assumptions}, the most questionable assumption of the model is that the hydrogen atoms enter a proton plasma which has been suddenly decelerated according to shock jump conditions.  This assumption is flawed --- the approximation is good only when the neutral fraction is small (see Lim \\& Raga 1996 for a discussion on the effects of larger neutral fractions).  Indeed, the reverse shock has a transition zone of finite thickness $\\Delta z \\sim n_p \\sigma_{I,p}$, within which the ionization, excitation and charge transfer reactions occur.  Consider the first hydrogen atoms to enter this zone from upstream.  When they become ionized, the resulting protons must have streaming velocities comparable to that of the hydrogen atoms.  Within the transition zone, the newly created protons must be decelerated and compressed so that they emerge on the downstream side of the zone with velocity $v_p = v_s/4$.  But most of the ionizations must occur when the relative streaming velocity of the hydrogen atoms and the protons is substantially less than the relative velocity of $3v_s/4$ that is assumed in the present model.  Given the fact that $\\Re_{b,n}(\\mbox{H}\\alpha)$ decreases precipitously for $v_s \\gtrsim 2000$ km s$^{-1}$ (Fig. \\ref{fig:87a}), we see that our model might grossly underestimate the actual value of this ratio in the case of SNR 1987A.  To properly interpret the H$\\alpha$ and Ly$\\alpha$ emission from SNR 1987A, we must construct a model for this transition zone similar to that of Whitney \\& Skalafuris (1963) and Skalafuris (1965, 1968 and 1969).  In a sense, the analysis presented here is a warm-up to that task.  Once we have such a model, we can use the results of the present analysis to calculate the emissivity and profiles of broad and narrow lines throughout the transition zone.  The assumption that a Balmer-dominated shock can be treated as a discontinuous jump is faulty for all SNRs, not just SNR~1987A.  Accordingly, we may ask why the model seems to work for most SNRs, for which $v_s \\lesssim 2000$ km~s$^{-1}$.  If one allowed for the transition zone in constructing a model for actual SNRs, one would construct a weighted average of $\\Re_{b,n}(\\mbox{H}\\alpha)$ over relative velocities ranging up to $3v_s/4$.  But, as Fig. \\ref{fig:halpha} shows, the ratio $\\Re_{b,n}(\\mbox{H}\\alpha)$ does not vary greatly for $200 \\lesssim v_s \\lesssim 2000$ km~s$^{-1}$.  Therefore, the value of $\\Re_{b,n}(\\mbox{H}\\alpha)$ appropriately averaged over a range of relative velocities would still agree fairly well with the observed values.  On the other hand, these considerations bring into question the validity of the relationship between the broad line FWHM and the shock velocity predicted by the current models.  In a more realistic model that takes into account the finite thickness of the transition zone, a greater fraction of the H$\\alpha$ emission would take place where the relative velocity is less than $3v_s/4$.  The result would be a reduced value of the FWHM of the broad emission for a given value of $v_s$ --- or, equivalently, a greater value of $v_s$ required to account for a given FWHM.  This consideration suggests that the current models may be under-estimating significantly the shock velocities of SNRs.  The present model is also inadequate to interpret current observations of the ratio $\\Re_{b,n}(Ly\\alpha)$ in SNR 1987A.  The observed brightness of Ly$\\alpha$ is modified by resonant scattering in the supernova debris.  The resonance scattering can double the brightness of Ly$\\alpha$ on the near side of the supernova debris and greatly suppress it on the far side, as pointed out by H06. However, since the Ly$\\alpha$ photons are produced in the transition zone, and their mean free paths are less than $\\Delta z$, we need a model of the transition zone to account properly for this resonant scattering."
    },
    "0609/astro-ph0609041_arXiv.txt": {
        "abstract": "We identify a group of $\\sim$20 co-moving, mostly southern hemisphere, $\\sim$200\\,Myr old stars near Earth.  Of the stars likely to be members of this Carina-Near Moving Group, in either its nucleus ($\\sim$30\\,pc from Earth) or surrounding stream, all but 3 are plausible members of a multiple star system.  The nucleus is (coincidentally) located quite close to the nucleus of the AB Doradus moving group notwithstanding that the two groups have substantially different ages and Galactic space motions, UVW. ",
        "introduction": "During the past decade various young stellar associations and clusters have been identified within 100\\,pc of the Sun \\citep[see][for a review]{ZS2004}.  These comoving groups can illuminate our understanding of the properties of low and intermediate mass stars at times that correspond to the long-ago era when planet building and heavy bombardment were occurring in our Solar System.  The Spitzer Infrared Observatory is targeting virtually all these young stars to determine if they are surrounded by dusty debris disks.  And the stars are being observed at near-infrared wavelengths with the Hubble Space Telescope and with adaptive optics systems on most of the world's largest telescopes; the goal is to detect thermal emission from young massive planets.  In the present Letter we add another example -- the Carina-Near Moving Group -- to the growing menagerie of nearby young associations.  The nucleus -- consisting of stars in the relatively narrow range of spectral type F through K0 of which HR\\,3070 (= HIP~38160) is the earliest spectral type and the brightest -- is only $\\sim$30\\,pc from Earth (Table~1).  Some of the surrounding stream stars, whose group membership is therefore less secure than those in the nucleus, are even closer to Earth; for example, GJ\\,358 is within 10\\,pc.  Based primarily on lithium abundance and X-ray flux (see Section~\\ref{sec3}), we estimate the group to be somewhat older than the Pleiades.  Some nuclear stars in our proposed group were initially suggested as co-moving by \\citet{Eggen}. Later, \\citet[hereafter MU2000]{MU2000} argued that they are part of a spatially very extensive, X-ray bright, moving group they dubbed ``Carina-Vela''. However, other than X-ray luminosity, their choice of group membership was predicated entirely on kinematics, while lacking even radial velocity determinations.  Our past experience \\citep[][]{SBB2002} with stars proposed \\citep{MF01} as members of the TW Hydrae Association, demonstrated that a purely convergent-point kinematic approach is insufficient to reliably pinpoint members of nearby young associations.  Nonetheless, in the 5 years following publication of their paper, neither \\citeauthor{MU2000} nor any other astronomer has verified or denied their kinematic model.  For example, both \\citet{C2000} and \\citet{W2003} studied group member HIP\\,37563 without referencing the MU2000 paper.  Indeed, unware of the MU2000 paper, we independently stumbled on the nearer (to Earth) members of their proposed association while pursuing our research program of identification of young nearby stars. Apparently, in their young star program, \\citet{Jensen} also find some overlap with the MU2000 sample. However, based on the few stars explicitly mentioned in the Jensen~et~al abstract, the stars in their program appear, on average, to be about 3 times more distant from Earth than stars in our Table~1.  Our investigations of radial velocities and youth indicators such as lithium abundance and H$\\alpha$ intensity are described in the following sections.  The essence of our results may be summarized as follows:  we confirm some of the MU2000 suggested members, deny others, and add some of our own.  In addition, we think it is not helpful, and may even be wrong, to associate very nearby stars such as those listed in Table~1 with the much more distant stars that MU2000 place together into a single moving group. Specifically, as is evident from their Fig. 5, the nearest members of their ``Carina-Vela'' group are three times closer to Earth than their proposed more distant members, with essentially no stars at intermediate distances.  In view of these large spatial separations, lack of evidence that Galactic space velocities $UVW$ are all similar, and lack of data on lithium abundances and activity indicators, at present we see no compelling reason to suppose the Carina-Vela group stars shared a common origin.  Table~1 of MU2000, ``Candidate members of the Carina-Vela moving group'', contains 58 entries most of which are not particularly near Earth.  Of these 58, only 5 are included in our Table~1 as candidate members of the moving group of stars we are considering and one MU2000 candidate is a definite non-member of our proposed moving group.  Rather than being spread out to distances $>100$\\,pc and with generally unknown $UVW$ (i.e.  MU2000's Carina-Vela Group), our Table~1A and B stars are characterized by proximity to Earth and rather tight constraints (a few km/s) on measured $UVW$ velocities.  In addition, stars listed in our Table~1B are not in the Carina-Vela direction. Finally, based on our Figures 1 and 2 and the discussion in Section 3, we believe that the Carina-Near stars must be older than the Pleiades, whereas IC\\,2391 which belongs to the Carina-Vela moving group proposed by MU2000 is generally regarded as younger than the Pleiades. In view of these many differences, we designate the stars we are considering to be the ``Carina-Near Moving Group''. ",
        "conclusions": "We have identified a group of comoving $\\sim$200\\,Myr old stars that partially surround the Sun.  At least four of these stars are close enough to Earth to have been included in the original Gliese catalog of nearby stars (1969):  GJ\\,140, 358, 900, and 907.1. Like all previously discovered and younger associations of stars within 100\\,pc of Earth -- TW~Hydrae, Tucana/Horologium, $\\beta$~Pictoris, AB~Doradus, $\\eta$~Cha, Cha-Near -- the nucleus of the Carina-Near Moving Group is located deep in the southern hemisphere.  This nucleus, is (coincidentally) located quite close to the nucleus of the AB~Doradus Moving Group notwithstanding that the two groups have substantially different ages and Galactic space motions, UVW: AB~Dor $(-8,-27,-14)$; Carina-Near $(-26,-18,-2)$. Indeed, the U velocity of the Carina-Near group is substantially more negative than the U velocities of all the other groups listed here.  In any event, identification of this nearby group of stars, many of which have masses similar to the Sun, extends the time-line covered by these other groups to older ages. As instrumentation capabilities improve, astronomers will be able to follow the planetary formation process over the age interval from 8 to 200\\,Myr."
    },
    "0609/astro-ph0609088_arXiv.txt": {
        "abstract": "The explosion of a carbon-oxygen white dwarf as a Type Ia supernova is known to be sensitive to the manner in which the burning is ignited. Studies of the pre-supernova evolution suggest asymmetric, off-center ignition, and here we explore its consequences in two- and three-dimensional simulations.  Compared with centrally ignited models, one-sided ignitions initially burn less and release less energy. For the distributions of ignition points studied, ignition within two hemispheres typically leads to the unbinding of the white dwarf, while ignition within a small fraction of one hemisphere does not.  We also examine the spreading of the blast over the surface of the white dwarf that occurs as the first plumes of burning erupt from the star. In particular, our studies test whether the collision of strong compressional waves can trigger a detonation on the far side of the star as has been suggested by \\citet{plewa2004a}. The maximum temperature reached in these collisions is sensitive to how much burning and expansion has already gone on, and to the dimensionality of the calculation. Though detonations are sometimes observed in 2D models, none ever happens in the corresponding 3D calculations. Collisions between the expansion fronts of multiple bubbles also seem, in the usual case, unable to ignite a detonation. ``Gravitationally confined detonation'' is therefore not a robust mechanism for the explosion. Detonation may still be possible in these models however, either following a pulsation or by spontaneous detonation if the turbulent energy is high enough. ",
        "introduction": "\\label{intro_sect}  In the currently favored model for Type Ia supernovae (SNe~Ia), a carbon-oxygen white dwarf (WD) grows to almost the Chandrasekhar mass, then explodes due to a thermonuclear instability \\citep[e.g.,][]{hillebrandt2000a}. While the modeling of the explosion itself has reached a high level of sophistication, with multi-dimensional studies being routinely carried out by several groups \\citep[e.g.][]{reinecke2002d,gamezo2003a,plewa2004a,roepke2004c, roepke2005b, garcia2005a, roepke2006b}, the initial conditions of this process remain largely unknown. This is unfortunate since the geometry of the flame ignition has a large effect on the explosion strength \\citep{niemeyer1996a,livne2005a,roepke2006a,schmidt2006a}.  Runaway commences once the WD has accreted sufficient matter from a binary companion to approach a central density $\\sim 3 \\times 10^9$ g cm$^{-3}$ where plasma neutrino losses are exceeded by energy generation from a highly screened carbon fusion reaction.  The stage for flame ignition is set by a century of convective carbon burning in the progenitor WD. It remains unclear, however, if the first sparks to develop a nearly discontinuous temperature gradient on their perimeters (the ``flame'') are concentrated in the center of the star \\citep{hoeflich2002a} or spread around by the convective flow in which they are embedded. Since this convective flow may have a dipole character, one natural possibility is lopsided ignition displaced somewhat off-center \\citep[e.g.,][]{woosley2004a}.  It is known that, once born, the flame experiences an extended period of subsonic propagation---a ``deflagration'' \\citep{nomoto1976a}. Prompt detonation is excluded on the grounds that it would produce spectroscopy, nucleosynthesis, and a light curve very different from observations \\citep[see][for a review]{filippenko1997a}. It would also require a degree of isothermality in the core that would be very difficult to achieve \\citep{woosley1990a}. The deflagration poses a computational challenge since the ashes of the burning are buoyant, and that leads to instabilities and turbulence that can only be followed with any accuracy in a multi-dimensional calculation. The difficulty is compounded by the large range of spatial scales---sub-millimeter for the flame width and Kolmogorov scale to 2,000 km for the WD---and the high Reynolds number, $Re \\sim 10^{14}$.  Calculations of the same WD differ in outcome because of the assumptions about ignition, various techniques used to treat flame instabilities, and turbulence. It is not practical to resolve both the flame and the star, so full-star models, such as the ones presented here, rely upon an effective flame model and a subgrid scale model for the turbulence. Qualitatively, the flame advances radially at a speed given by the flotation of the largest plumes, but the lateral spreading of each plume and the overall efficiency of the explosion can vary, depending upon the way turbulence is handled, and on the dimensionality and resolution of the model.  Results of different approaches have been published in a variety of studies with the general conclusion that a pure deflagration can give a viable explosion \\citep{reinecke2002d,gamezo2003a,roepke2005b}, not too different from what is observed \\citep{travaglio2004a,blinnikov2006a}. It remains controversial, however, if these models can give light curves as bright as some observations indicate, or can explain all of the spectroscopic features \\citep{gamezo2003a, kozma2005a}.  Moreover, these successful models all have in common the assumption of nearly isotropic central ignition. It may be that nature provides only anisotropic ignition conditions, or it may be that the the observational constraints on a pure deflagration will ultimately prove too stringent. In these cases, a transition to detonation may need to occur \\citep{plewa2004a,livne2005a}.  The idea of a delayed detonation has been around for some time \\citep{khokhlov1991a,woosley1994a}, but the physics of that transition, if it happens, is still uncertain \\citep{niemeyer1997b, niemeyer1999a}. Recent two-dimensional calculations have suggested that burned material may quickly ascend to the surface of the still gravitationally bound star, sweep around it and, by collision and compression on the opposite side, trigger a detonation in the unburned material \\citep{plewa2004a}. This is called by its proponents ``gravitationally confined detonation'', or GCD.  Here, we follow the evolution of one-sided ignitions for a range of assumptions regarding the initial conditions in two-dimensional (2D) and three-dimensional (3D) simulations. The results are sensitive to the fuel consumption and energy release during the rising stage of the plume. Therefore a correct description of the turbulent deflagration flame is as crucial here as in other models. We describe our approach in Sect.~\\ref{sect:mod}.  Before that, we give a brief motivation of the ignition scenarios explored here. Although \\citet{roepke2006a} showed that ignition conditions cannot be explored reliably in 2D simulations, we can and do use surveys in cylindrical symmetry to get a feeling for the parameter range to be explored, as well as the dependence of the results on numerical resolution (Sect.~\\ref{sect:2d}). Full-star 3D simulations are presented in Sect.~\\ref{sect:3d}, and the consequences regarding the possibility of triggering a detonation are discussed in Sect.~\\ref{sect:sdet}. Conclusions are drawn in Sect.~\\ref{sect:concl}. ",
        "conclusions": "\\label{sect:concl}   The evolution of thermonuclear supernova models that ignite asymmetrically has been followed in two and three dimensions. Parameters of the setup, such as the numerical resolution, the displacement of the center of flame ignition from the center of the WD, and the ignition shape were explored in a systematic way in 2D, and the more interesting cases were explored in 3D.  For all 2D simulations, the energy release in the nuclear burning falls short of unbinding the star. In the cases with one-sided ignition, the flame floats rapidly to the surface \\citep[see also,][]{livne2005a}, spreading laterally as it goes due to burning and instabilities. Since the star is still gravitationally bound, the emerging ashes sweep around the core of the WD and collide on the opposite side, in agreement with \\citet{plewa2004a}. The collision strength, and thus the maximum temperature reached in the compressed fuel correlate inversely with the nuclear energy released on the flame's way to the surface.  The question of whether the compression of unburned material in the collision region is adequate to trigger a detonation was explored in detail, and the necessary criteria were set out. A detonation can only occur if the fuel temperature exceeds approximately $1.9 \\times 10^9 \\, \\mathrm{K}$ at a density above $10^7\\, \\mathrm{g} \\, \\mathrm{cm}^{-3}$. At lower densities, detonation requires higher temperatures and eventually becomes impossible, for any temperature, for densities less than about $10^6 \\, \\mathrm{g} \\, \\mathrm{cm}^{-3}$.  The conditions for initiating a detonation were met in several 2D calculations in which the flame ignited in a spherical bubble more than $200 \\, \\mathrm{km}$ off-center.  Less efficient prescriptions for the burning might have found similar conditions in simulations that ignited closer in. For the initial conditions and flame propagation model assumed, the results of \\citet{plewa2004a}, may be reasonable. However, since the flame model applied there is not based on a consistent treatment of the flame's interaction with the turbulent cascade, it is difficult to judge its validity.  \\emph{In three dimensions, all simulations fell far short of initiating a detonation}. Some even released sufficient energy to unbind the star. There are several reasons for this difference. Lacking the artificial symmetry of 2D simulations, the focusing of the collision in 3D models can be weaker. Indications for this were found in our simulations. At least two of the 3D simulations gave significantly lower collision temperatures than predicted by 2D models that burned similar amounts of fuel (see Fig.~\\ref{fig:et1}). A second, probably dominant effect is that 3D models release more energy than their 2D analogs. This is mostly due to the additional degree of freedom in developing flame surface area due to instabilities. Another contribution to the difference may be the improved subgrid-scale turbulence model applied in the 3D simulations. However, this effect is expected to be minor \\citep{schmidt2006c}. An asymmetrically ignited 3D model based on a different flame implementation \\citep{calder2004a} burned $\\sim$$0.075\\, M_\\odot$ of material. This corresponds to a nuclear energy release of $\\sim$$1.2 \\times 10^{50} \\, \\mathrm{erg}$---a value that falls in the range spanned by our parameter study. Therefore, although the Calder et al. model was not followed beyond the breakout of the ashes from the surface, one expects that that model would also have failed to trigger a detonation.  While we found no example of a successful detonation in our 3D simulations, this possibility cannot be completely ruled out since the exploration of the parameter space was incomplete. An interesting possibility is that of double-sided ignitions, which may be possible in a teardrop-shaped ignition overshooting through the WD's center (or, in a simpler configuration, as two opposed bubbles). Such a configuration shortens the way the material has to travel towards the collision spot once ashes break out of the star's surface on both sides. Therefore the expansion of the star may not be as advanced as in the collision on the far side of a single-bubble breakout. In this case, higher temperatures and densities are expected in the compressed fuel. On the other hand, burning on both sides of the star releases more energy while the flames propagate towards the surface. This increases the expansion prior to collision. In our simulations, the latter effect dominated and the collision was weak in a 2D model. A similar 3D model became unbound and no detonation was found in either. Thus, if a multiple surface breakout scenario is to work at all, no more than two widely separated ignition kernels are admissible or there will be too much expansion. An open question is whether a special placement of two or three bubbles spanning a smaller angle than $180^{\\circ}$ with the WD's center might favor the first effect, increasing the collision strength by shortening the path of the surface material. However, the bubbles cannot be too close or they would merge quickly due to burning, without significantly compressing the unburned material between them, and there cannot be very many, or they will prematurely unbind the star.  Keeping in mind the uncertainties of the flame model and the incompleteness of the parameter space explored in 3D simulations, we conclude, that although a detonation due to the colliding surface material may, in principle, occur for certain---possibly artificial---ignition configurations, it cannot serve as a robust model for SNe~Ia. The simulations presented here indicate that it may not be realized in nature at all.  For models that remain gravitationally bound, failure to initiate a detonation will lead to pulsations of the WD star \\citep{nomoto1976a}. This may be a second chance for triggering a detonation \\citep{arnett1994a,bravo2006a} and this occurrence will be addressed in a follow-up study."
    },
    "0609/astro-ph0609277_arXiv.txt": {
        "abstract": "{Structural and dynamical parameters of faint open clusters are probed with quality 2MASS-photometry and analytical procedures developed for bright clusters.} {We derive fundamental parameters of the faint open clusters Lyng\\aa\\,2, BH\\,63, Lyng\\aa\\,12 and King\\,20, the last three of which have no prior determinations. We also focus on the structure and dynamical state of these clusters.} {\\jj, \\hh\\ and \\ks\\ 2MASS photometry with errors smaller than 0.2\\,mag are used to build CMDs, radial density profiles, colour-colour diagrams, luminosity and mass functions. Colour-magnitude filters are used to isolate probable member stars. Field-star decontamination is applied to Lyng\\aa\\,2, Lyng\\aa\\,12 and King\\,20.} {Reddening values are in the range $0.22\\leq\\ebv\\leq1.9$, with BH\\,63 the most reddened object. Ages of Lyng\\aa\\,2, King\\,20, Lyng\\aa\\,12 and BH\\,63 are $\\approx90$, $\\approx200$, $\\approx560$ and $\\approx700$\\,Myr, respectively. The radial density distributions of Lyng\\aa\\,12 and King\\,20 are well-represented by King profiles. Lyng\\aa\\,2 and BH\\,63 are very small with core and limiting radii of $\\approx0.12$\\,pc and $\\approx1.5$\\,pc. Yet, they fit in the small-radii tail of the open cluster size distribution. Lyng\\aa\\,12 and King\\,20 have $\\rc\\approx0.43$\\,pc and $\\rl\\approx3.9$\\,pc. Lyng\\aa\\,2 and Lyng\\aa\\,12 are inside the Solar circle. Total stellar masses (extrapolating the MFs to stars with 0.08\\,\\ms) range from $\\approx340$\\,\\ms\\ (BH\\,63) to $\\approx2300$\\,\\ms\\ (Lyng\\aa\\,12). Observed masses are $\\sim1/4$ of these values. In all clusters the core mass function is flatter than the halo's.} {Faint open clusters can be probed with 2MASS when associated with colour-magnitude filters and field-star decontamination. BH\\,63 appears to be in an advanced dynamical state, both in the core and halo. To a lesser degree the same applies to King\\,20. Marginal evidence of dynamical evolution is present in the cores of Lyng\\aa\\,2 and Lyng\\aa\\,12.} ",
        "introduction": "\\label{intro}  Because it is relatively simple to estimate ages and distances of open clusters (OCs), they have become fundamental probes of Galactic disc properties (Lyng\\aa~\\cite{Ly82}; Janes \\& Phelps \\cite{JP94}; Friel \\cite{Friel95}; Bonatto et al. \\cite{BKBS06}; Piskunov et al. \\cite{Pi2006}). However, the proximity of most of the OCs to the plane and the corresponding high values of reddening and field-star contamination (Bonatto et al. \\cite{BKBS06} and references therein) usually restrict this analysis to the more populous OCs and/or those located a few kpc from the Sun.  To probe disc structure Bonatto et al. (\\cite{BKBS06}) employed 654 OCs with fundamental parameters (such as age, distance from the Sun, and reddening) mostly from the WEBDA\\footnote{\\em http://obswww.unige.ch/webda} database (Mermilliod \\cite{Merm1996}). They found that a large fraction of the faint and/or poorly populated OCs must be overwhelmed in the field, particularly in bulge/disc directions. Piskunov et al. (\\cite{Pi2006}) used 650 OCs from ASCC-2.5 data to infer kinematical similarities among OC groups. Both works estimate a total population of $\\sim10^5$ OCs in the Galaxy. Dias et al. (\\cite{Dias2002}) reported 1756 optically visible OCs and candidates in their catalogue. These works suggest that besides those not yet observed, there is a large number of OCs and candidates whose properties have not been explored.  \\begin{figure*} \\begin{minipage}[b]{0.50\\linewidth} \\includegraphics[width=\\textwidth]{fig1a.ps} \\end{minipage}\\hfill \\begin{minipage}[b]{0.50\\linewidth} \\includegraphics[width=\\textwidth]{fig1b.ps} \\end{minipage}\\hfill \\caption[]{Left panel: $4\\arcmin\\times4\\arcmin$ XDSS R image of BH\\,63. Right panel: $15\\arcmin\\times15\\arcmin$ XDSS R image of Lyng\\aa\\,2. Images centered on the optimised coordinates (cols.~5 and 6 of Table~\\ref{tab1}).} \\label{fig1} \\end{figure*}   \\begin{figure*} \\begin{minipage}[b]{0.50\\linewidth} \\includegraphics[width=\\textwidth]{fig2a.ps} \\end{minipage}\\hfill \\begin{minipage}[b]{0.50\\linewidth} \\includegraphics[width=\\textwidth]{fig2b.ps} \\end{minipage}\\hfill \\caption[]{Same as Fig.~\\ref{fig1} for the $11\\arcmin\\times11\\arcmin$ field around Lyng\\aa\\,12 (left panel) and the $9\\arcmin\\times9\\arcmin$ field around King\\,20 (right panel).} \\label{fig2} \\end{figure*}  Efforts to analyse unstudied OCs and derive their fundamental parameters will contribute to future disc studies by unveiling the properties of individual OCs themselves and constraining theories of molecular cloud fragmentation, star formation and stellar and dynamical evolution. These processes will be better understood as the number of studied clusters, faint ones in particular, increases. A side-product of more constrained parameters will be better statistics of OCs as a class.  As a first step in this direction we studied a series of faint OCs (Bica, Bonatto \\& Dutra \\cite{BBD2003}; Bonatto, Bica \\& Dutra \\cite{BBD2004}; Bica \\& Bonatto \\cite{BiBo2005}) using near-IR photometry. For spatial and photometric uniformity we work with \\jj, \\hh\\ and \\ks\\ 2MASS\\footnote{The Two Micron All Sky Survey, All Sky data release (Skrutskie et al. \\cite{2mass1997}), available at {\\em http://www.ipac.caltech.edu/2mass/releases/allsky/}} Point Source Catalogue (PSC) photometry. We have developed methods to disentangle cluster and field stars in colour-magnitude diagrams (CMDs), as well as in radial stellar density distributions. The use of field-star decontamination and colour-magnitude filters (hereafter C-MF) have produced more robust parameters (e.g. Bonatto \\& Bica \\cite{BB2005}; Bonatto et al. \\cite{BBOB06}). In particular, field-star decontamination constrains more the age, especially for low-latitude OCs (Bonatto et al. \\cite{BKBS06}). These procedures have proven useful also in the analysis of faint and/or distant OCs (Bica, Bonatto \\& Dutra \\cite{BBD2003}; Bonatto, Bica \\& Dutra \\cite{BBD2004}; Bica \\& Bonatto \\cite{BiBo2005}). Recently we discussed the advantages of using C-MFs in the construction of radial density profiles for M\\,52 and NGC\\,3960, whose fields are affected by different amounts of differential reddening (Bonatto \\& Bica \\cite{BB2006}). As by-products we have been deriving more constrained structural parameters such as core and limiting radii, and mass function slopes, that allow inferences on dynamical evolution.  In the present work we focus on the faint OCs Lyng\\aa\\,2, Lyng\\aa\\,12, King\\,20 (Alter et al. \\cite{Alter70} and references therein), and BH\\,63 that was catalogued by van den Bergh \\& Hagen (\\cite{BH75}).  We intend to derive photometric, cluster structure and dynamic-related parameters working with a set of methods which provide more accurate results. The target clusters are Lyng\\aa\\,2 (for comparison purposes with the literature), and the as yet unstudied OCs Lyng\\aa\\,12, King\\,20 and BH\\,63. Since BH\\,63 is very reddened, and consequently differential reddening may affect its field, it required an approach slightly different than those for the other OCs. Throughout this paper errors correspond to $1\\sigma$ Poisson fluctuations.  \\begin{table*} \\caption[]{General data on the clusters} \\label{tab1} \\tiny \\renewcommand{\\tabcolsep}{0.6mm} \\renewcommand{\\arraystretch}{1.5} \\begin{tabular}{lcccccccccccccc} \\hline\\hline &\\multicolumn{3}{c}{Derived from XDSS}&\\multicolumn{9}{c}{Derived from 2MASS}\\\\ \\cline{2-4}\\cline{6-14} Cluster&$\\alpha(2000)$&$\\delta(2000)$&D&&$\\alpha(2000)$&$\\delta(2000)$&$\\ell$&$b$&Age&$\\ebv$&\\ds&\\dgc&\\dgc&Alternative \\\\ &(hms)&($^\\circ\\arcmin\\arcsec$)&(\\arcmin)&&(hms)&($^\\circ\\arcmin\\arcsec$)&($^\\circ$)&($^\\circ$)&(Myr)&&(kpc)&(kpc)&(kpc)&Names\\\\ (1)&(2)&(3)&(4)&&(5)&(6)&(7)&(8)&(9)&(10)&(11)&(12)&(13)&(14)\\\\ \\hline BH\\,63&09:20:39&$-$49:13:23&2.5&&09:20:39.5&$-$49:13:20.5&271.63&$+$0.40&$700\\pm100$&$1.9\\pm0.6$&$2.3\\pm0.3$ &$8.3\\pm0.3$&$7.5\\pm0.3$ &ESO\\,212SC2\\\\  Lyng\\aa\\,2&14:24:35&$-$61:19:50&13.0&&14:24:21.1&$-$61:19:21.0&313.84&$-0.43$&$90\\pm10$&$0.22\\pm0.04$ &$0.9\\pm0.1$&$7.4\\pm0.1$&$6.6\\pm0.1$  &OCl$-$916,\\,ESO\\,134SC2,\\,BH\\,157\\\\  Lyng\\aa\\,12&16:46:04&$-$50:45:50&8.0&&16:46:04.4&$-$50:45:40.1&335.70&$-3.46$&$560\\pm100$&$0.22\\pm0.04$ &$1.0\\pm0.2$&$7.1\\pm0.2$&$6.3\\pm0.2$    &OCl$-$973,\\,ESO\\,226SC22\\\\  King\\,20&23:33:15&$+$58:28:00&6.0&&23:33:17.0&$+$58:28:33.1&112.85&$-$2.86&$200\\pm20$&$0.65\\pm0.03$ &$1.9\\pm0.2$&$8.9\\pm0.2$&$8.1\\pm0.2$   &OCl$-$262 \\\\ \\hline \\end{tabular} \\begin{list}{Table Notes.} \\item Cols.~2 and 3: measured by us on XDSS images. Col.~4: angular diameter estimated on XDSS images. Col.~10: reddening in the cluster's central region (Sect.~\\ref{age}). Col.~12: \\dgc\\ calculated using $R_O=8.0$\\,kpc (Reid \\cite{Reid93}) as the distance from the Sun to the Galactic center. Col.~13: \\dgc\\ using $R_O=7.2$\\,kpc (Bica et al. \\cite{BBBO06}). \\end{list} \\end{table*}  This paper is organized as follows. In Sect.~\\ref{Target_OCs} we discuss basic properties and review literature data (when available) on the 4 clusters. In Sect.~\\ref{2mass} we present the 2MASS data, subtract the field-star contamination (except for BH\\,63), derive fundamental cluster parameters, and analyse the radial density profiles. In Sect.~\\ref{MF} we derive luminosity and mass functions (LFs and MFs), and compute stellar content properties. In Sect.~\\ref{CWODS} we discuss dynamical states. Concluding remarks are given in Sect.~\\ref{Conclu}. ",
        "conclusions": "\\label{Conclu}  In this paper, we derived photometric and structural parameters of the faint open clusters BH\\,63, Lyng\\aa\\,2, Lyng\\aa\\,12, and King\\,20 with \\jj, \\hh\\ and \\ks\\ 2MASS photometry, restricted to stars with observational uncertainties $\\epsilon_{J,H,K_S}<0.2$\\,mag. BH\\,63 is significantly more reddened than the others, with $\\ebv\\approx1.9$. Distance from the Sun ranges from $\\ds\\approx0.9$\\,kpc (Lyng\\aa\\,2) to $\\ds\\approx2.3$\\,kpc (BH\\,63). Lyng\\aa\\,2 and Lyng\\aa\\,12 are located inside the Solar circle.  The use of C-MFs allowed us to obtain accurate parameters for such faint clusters including mass function slopes for the core and halo. In particular, radial density profiles built with C-MF photometry are better defined and yield more constrained structural parameters than the observed ones. The RDPs of Lyng\\aa\\,12 and King\\,20 are well-represented by a King profile. Lyng\\aa\\,2 and BH\\,63 appear to be very small, with core and limiting radii of $\\approx0.12$\\,pc and $\\approx1.5$\\,pc. However, they fit in the small core and limiting radii tail of the OC distribution. Lyng\\aa\\,12 and King\\,20 have $\\rc\\approx0.43$\\,pc and $\\rl\\approx3.9$\\,pc.  The open clusters of the present sample have flat core MFs, suggesting dynamically evolved systems, particularly BH\\,63, which has a flat MF in  its halo as well. The total mass locked up in stars (extrapolated down to $0.08$\\,\\ms) ranges from $\\approx340$\\,\\ms\\ (BH\\,63) to $\\approx2300$\\,\\ms\\ (Lyng\\aa\\,12).  We also analysed the clusters with dynamical and structural diagnostic diagrams. BH\\,63 appears to be an extreme case of OC dynamical evolution both in  its core and halo, perhaps due to interaction with the nebula FeSt\\,1-108. King\\,20 shows signs of dynamical evolution as well, but to a  lesser degree. The cores of Lyng\\aa\\,2 and Lyng\\aa\\,12 present marginal evidence of dynamical evolution.  The present work shows that it is possible to explore faint OCs with 2MASS photometry, provided care is taken to identify probable member stars (to better define the cluster sequence on CMDs) and discard stars with discordant colour (for more intrinsic cluster RDPs and MFs). The use of field-star decontamination and colour-magnitude filters on the fields of BH\\,63, Lyng\\aa\\,2, Lyng\\aa\\,12 and King\\,20 allowed us to obtain  good results for the fundamental, structural and dynamical-evolution related parameters."
    },
    "0609/astro-ph0609794_arXiv.txt": {
        "abstract": "Technical libraries are currently experiencing very rapid change.  In the near future their mission will change, their physical nature will change, and the skills of their employees will change.  While some will not be able to make these changes, and will fail, others will lead us into a new era. ",
        "introduction": "We are at the end of the paper era.  For the past two centuries libraries have been primarily devoted to building and maintaining vast collections of paper, documenting the myriad intellectual activities of mankind.  These activities are nearing the end of their useful lives.  Many existing libraries will not be able to meet the challenges, and they will close.  Others will provide the intellectual leadership necessary to support the new information seeking and storing paradigm, and will thrive. ",
        "conclusions": ""
    },
    "0609/astro-ph0609749_arXiv.txt": {
        "abstract": "{} {We use catalogues of superclusters of galaxies from the 2dF Galaxy Redshift Survey to study the properties of galaxies in superclusters. } {We compare the properties of galaxies in high and low density regions of rich superclusters, in poor superclusters and in the field, as well as in groups, and of isolated galaxies in superclusters of various richness.} {We show that in rich superclusters the values of the luminosity density  smoothed on a scale of 8 \\Mpc\\   are higher than in poor superclusters: the median density in rich superclusters is $\\delta \\approx 7.5$, in poor superclusters $\\delta \\approx 6.0$. Rich superclusters contain high density cores with densities $\\delta > 10$ while in poor superclusters such high density cores are absent. The properties of galaxies in rich and poor superclusters and in the field are different: the fraction of early type,  passive galaxies in rich superclusters is slightly larger than in poor superclusters, and is the smallest among the field galaxies. Most importantly, in high density cores of rich superclusters ($\\delta > 10$) there is an excess of early type, passive galaxies in groups and clusters, as well as among those which do not belong to groups or clusters. The main galaxies of superclusters have a rather limited range of absolute magnitudes.  The main galaxies of rich superclusters have larger  luminosities  than those of poor superclusters and of groups in the field (the median values  are correspondingly $M_{bj} = -21.02$, $M_{bj} = -20.9$  and $M_{bj} = -19.7$ for rich and poor superclusters and  groups in the field).} {Our results show that both the local (group/cluster) environments and global (supercluster) environments influence galaxy morphologies and their star formation activity. ",
        "introduction": "It is presently well established that galaxies belong to various systems from groups and clusters to superclusters, forming the supercluster-void network. Early studies of superclusters of galaxies were reviewed by Oort (\\cite{oort83}). These studies were based on observational data about galaxies, as well as on data about nearby groups and clusters of galaxies. Classical cluster catalogues were constructed by Abell (\\cite{abell}) and Abell \\etal (\\cite{aco}) by visual inspection of Palomar plates. The first relatively deep all-sky catalogues of superclusters of galaxies were complied by Zucca \\etal (\\cite{z93}) and Einasto \\etal (\\cite{e1994}, \\cite{e1997}, \\cite{e2001}) using data about Abell clusters.   The modern era of the study of various systems of galaxies began when new deep redshift surveys of galaxies were published. These surveys cover large regions of sky and allow to investigate the distribution of galaxies up to fairly large distance from us. These surveys formed the basis for new catalogues of groups, clusters and superclusters of galaxies.  The first of such catalogues was the Las Campanas catalogue of groups by Tucker \\etal (\\cite{Tucker00}).  The Las Campanas Galaxy Redshift Survey and the Sloan Digital Sky Survey were also used to compile catalogues of groups, clusters and superclusters by Einasto et al. (\\cite{e03a}, \\cite{e03b}, hereafter E03a and E03b) and Basilakos (\\cite{bas03}).  Group and supercluster catalogues based on the 2 degree Field Galaxy Redshift Survey (2dFGRS) were published by Eke \\etal (\\cite{eke04}), Yang \\etal (\\cite{yang04}) and Tago et al. (\\cite{tago06}, hereafter T06), and by Erdogdu et al. (\\cite{erd04}) and Porter and Raychaudhury (\\cite{pr05}).   The pioneering studies of the properties of galaxies in clusters by Davis \\& Geller (\\cite{dg76}) and Dressler (\\cite {d80}) showed that there exists a correlation between the spatial density of galaxies and their morphology -- early type galaxies are located preferentially in the central regions of clusters, where the local densities are high, while late type galaxies are located mostly in outer regions of clusters, having lower local densities around them. Einasto (\\cite{e91}) showed that clustering of galaxies depends on both their luminosity and morphology.  Already early studies of the morphological segregation of galaxies at supercluster scales demonstrated that this segregation extends to scales of 10--15 \\Mpc\\ (Giovanelli, Haynes and Chincarini \\cite{gio86}, Einasto and Einasto \\cite{e87} and Mo et al.  \\cite{mo92}).  The data about galaxies in the Las Campanas Galaxy Redshift Survey, the Sloan Digital Sky Survey and the 2 degree Field Galaxy Redshift Survey (2dFGRS) enable us to study the properties and the spatial distribution of galaxies in detail. Numerous papers have demonstrated segregation of galaxies by their spectral type, luminosity and colour index (Norberg et al. \\cite{n01}, \\cite{n02}, Zehavi et al. \\cite{zehavi02}, Goto et al. \\cite{go03}, Hogg et al. \\cite{hogg03}, \\cite{hogg04}, Balogh et al. \\cite{balogh04}, De Propris et al. \\cite{depr03}, Madgwick et al. \\cite{ma03b}, Croton et al. \\cite{cr05} and Blanton et al. \\cite{blant04}, \\cite{blant06} among others).  Blanton et al. (\\cite{blant06}) come to the conclusion that the blue galaxy fraction and the recent star formation history in general, depend mostly on the local environment of galaxies.  In this study the local environment was defined as the spatial density on the 0.5--1~\\Mpc\\ scale, and the global density on the 5--10~\\Mpc\\ scale.  Croton et al. (\\cite{cr05}) showed that galaxy populations depend also on the large scale environment. Balogh et al. (\\cite{balogh04}) compared the populations of star-forming and quiescent galaxies in groups from the 2dFGRS and SDSS surveys and in small (1.1~\\Mpc) and large scales (5.5~\\Mpc) and showed that the relative numbers of these galaxies depend both on the local and global environments. Even low density environments contain a large fraction of non-star-forming galaxies.  On the basis of the 2dFGRS we recently compiled a new catalogues of superclusters of the 2dF galaxies (Einasto et al. \\cite{e06a}, hereafter Paper I). This catalogue  is available electronically at the web-site \\texttt{http://www.aai.ee/$\\sim$maret/2dfscl.html}.  In (Einasto et al. \\cite{e06b}, hereafter Paper II) we studied various properties of these superclusters: their multiplicity, geometry, luminosity functions and other properties. We also compared the properties of real superclusters with simulated superclusters from the Millennium Simulations (Springel et al. \\cite{springel05}), and from the semianalytical mock catalog by Croton et al. (\\cite{cr06}).  In the present paper we continue our study of the properties of superclusters. We study the properties of galaxies: their  luminosities, spectral types and colors in rich and poor superclusters  and for comparison also in the field. These data enable to analyse populations of galaxies of different luminosity, morphology and star formation rate in various environments: in rich and poor superclusters, as well as in groups located in superclusters and in the field. The use of a large catalogue of superclusters enables us for the first time to study the properties of galaxies in a large number of superclusters of various richness.    The paper is composed as follows.  In the next Section we shall describe the supercluster data. Then we study  the properties of galaxies in superclusters, the density distribution in superclusters of various richness, and the properties of galaxies in groups located in regions of different large scale density in superclusters. Then we compare the luminosities of main galaxies in superclusters and in groups located in the field. In the last Sections we discuss the results and list our conclusions. ",
        "conclusions": "A detailed study of luminosity functions of galaxies from the 2dF survey in regions of different density of the large-scale environment was made by Croton et al.  (\\cite{cr05}), and in clusters by De Propris et al. (\\cite{depr03}). De Propris et al. (\\cite{depr03}) found that the luminosity functions of early type galaxies in clusters are brighter and steeper than those in the field, and that clustering of passive galaxies is stronger than clustering of actively star forming galaxies (Madgwick et al. \\cite{ma03b}). Using densities smoothed on a scale of 8~\\Mpc\\, Croton et al.  (\\cite{cr05}) divided the volume under study into 7 regions of various environmental densities, from extreme voids to cluster populations and found that the brightest galaxies in voids are approximately 5 times fainter than those in clusters. Even larger differences between luminosities of galaxies in high and low density regions were found in Einasto et al. (\\cite{e05b}). In the present paper we showed, in accordance to these results, that the luminosity-density dependence in important for galaxies of all types both in superclusters and among field galaxies at all densities.  Several recent studies address the problem whether the properties of galaxies, their star formation activity in particular, correlates with the local and/or global environment of galaxies, defined, for example, as the clustercentric distance. The well-known morphology-density relation is an example of such correlation (Dressler (\\cite {d80}). One question asked in these studies is whether there exists a critical density so that star formation is supressed in all groups/clusters where the density exceeds this critical value. Lewis et al. (\\cite{lew02}) used the data about star-forming galaxies in the 2dFGRS to show that there exists a correlation between the star formation activity and the local galaxy density which holds for galaxies at distances at least two virial radii from the cluster/group centre.  Balogh et al. (\\cite{balogh04}) compared the populations of star-forming and quiescent galaxies in small (1.1~\\Mpc) and large scales (5.5~\\Mpc) and showed that the relative numbers of these galaxies depend both on the local and global environments. Even low density environments contain a large fraction of non-star-forming galaxies. They concluded that the galaxy population must be only indirectly related to their present-day environment. Possible physical mechanisms must have been more effective in the past, and perhaps affeced the star formation rate in very short (less than 1 Gyr) timescales, like starbursts induced by galaxy interactions in close pairs of galaxies (see Balogh et al. \\cite{balogh04} and references therein).   Gray et al. (\\cite{gray04}) found, using data about the supercluster A901/902, strong evidence that the highest density regions in clusters are populated mostly with quiescent galaxies, while star forming galaxies dominate in outer/lower density regions of clusters. Similarly, Haines et al. (\\cite{hai06}) demonstrated that the colours of galaxies in the core region of the Shapley supercluster depend on environment, redder galaxies being located in cluster cores. They also found large concentrations of faint blue galaxies between clusters.  Our results are in accordance with those;in addition we showed that even at supercluster scales, the properties of galaxies and their environmental densities are correlated. In high density cores of rich superclusters the fraction of quiescent (red) galaxies is higher than this fraction in lower density regions even for those galaxies which do not belong to groups or clusters.    Porter and Raychaudhury (\\cite{pr05}) investigated the star formation rate in groups of galaxies from the Pisces-Cetus supercluster, according to their spectral index $\\eta$. They concluded that galaxies in rich clusters have lower star formation rates than galaxies in poor groups.  This is in accordance with our results, showing that galaxies from a higher density environment have lower star formation rates than galaxies from a lower density environment.   Our previous analysis in Paper II demonstrated that geometrical properties of rich and poor superclusters are different, rich superclusters have larger sizes, their shapes and compactness differ from those of poor superclusters. The mean density of superclusters  increases gradually with increasing total luminosity or richness of superclusters. This demonstrates that rich superclusters are physical systems with properties different from those of poor systems, they do not represent just percolations of several loose systems.  Our present study reveals additional differences between rich and poor superclusters. Rich superclusters contain high density cores which are absent in poor supercluters. The fraction of early type and passive galaxies in groups and clusters and among isolated galaxies in high density cores of rich superclusters is higher than in groups in poor superclusters.  We shall analyse the detailed properties of high density regions in rich superclusters in another study (Einasto et al., in preparation).   Hilton et al. (\\cite{hil05}) found that the fraction of early spectral type galaxies is significantly higher in clusters with a high X-ray flux. Many of these clusters belong to rich superclusters (Einasto et al. \\cite{e2001}, Belsole et al. \\cite{bel04}), so this result is in accordance with the present paper.  Our analysis  shows that main galaxies of superclusters form a specific class of galaxies with a very limited range of luminosities. Main galaxies of rich superclusters are more luminous than main galaxies of poor superclusters. A similar conclusion has been reached by Einasto \\& Einasto (\\cite{e87}) using data on nearby superclusters. This is in accordance with the result of De Propris et al. (\\cite{depr03}) that there is an excess of very bright galaxies in cores of clusters. Main galaxies of superclusters are formed by multiple merger processes, as indicated by direct observations and numerical simulations (see Laine et al. \\cite{laine03} and Gao et al. \\cite{gao05}). It has been known already for a long time that first-ranked cluster galaxies have a small dispersion of absolute magnitudes (Hubble \\& Humason \\cite{hubble31}, Hubble \\cite{hubble36}, Sandage \\cite{sandage76}).  More recent studies by Postman \\& Lauer (\\cite{postman95}) and Laine et al. (\\cite{laine03}) have shown that absolute magnitudes of brightest cluster galaxies have a scatter of about 0.24 - 0.33 mag. The scatter of luminosities of supercluster main galaxies is larger than the scatter of brightest cluster galaxies.  There may be several reasons for this.  One possibility is that we have found the main group and its main galaxy by an automated search routine, and supplementary information on the morphological type, colour etc. has not been used. For this reason our sample of main galaxies is probably not as homogeneous as samples of first-ranked cluster galaxies investigated by Hubble, Sandage, Postman and others.  Our analysis shows a large scatter of the properties of galaxies in poor superclusters.  In rich superclusters this scatter is small.   Numerical simulations show that dynamical evolution in high-density regions is determined by a high overall mean density that speeds up  clustering of particles (Einasto et al. \\cite{e05b} and references therein, Gao, Springel and White \\cite{gsw05}).  In high density regions clustering starts early and continues until the present.  The haloes that populate high density regions are themselves also richer, more massive and have larger velocities than the haloes in low density regions.  In low density filaments that cross voids, as well as in the outer low density regions of high density systems, the mean density is low and thus the evolution is slow, and in these regions haloes themselves are also poor, less massive and have small velocities. These differences affect the evolution and properties of galaxies in various environments.   In this paper we have used a catalogue of superclusters of galaxies from the 2dF galaxy redshift survey to study the properties of galaxies in superclusters and the properties of the richest superclusters.  Our main conclusions are the following.  \\begin{itemize}  \\item{} The density distributions in rich and poor superclusters are different. The densities in rich superclusters are higher than in poor superclusters, and rich superclusters contain high density cores which are absent in poor superclusters.  \\item{} Rich superclusters contain a higher fraction of early type, passive, red galaxies than poor superclusters.  \\item{} The properties of galaxies are correlated with the values of the luminosity density field smoothed on a scale of 8 Mpc/h: early type,   passive, non-starforming  galaxies have higher environmental  densities while late type, active, star/forming galaxies have  lower environmental densities. This  trend extends to field galaxies and to the lowest densities in our sample.  \\item{}  The fraction of early type, passive galaxies is the highest in rich groups/clusters in high density regions of rich superclusters. In these high density regions even among isolated galaxies the fraction of star forming galaxies is smaller than the fraction of star forming galaxies among isolated galaxies in poor superclusters and in the field.    \\item{} Main galaxies of rich superclusters have larger luminosities (median value $-$21.2) than main galaxies of poor superclusters (median value $-$20.9) and main galaxies of groups in the field (median value $-$19.7).    \\end{itemize}  Our results show that both local (group/cluster) environments and global (supercluster) environments are  important in ifluencing galaxy morphologies and their star formation activity. This indicates the importance of the role of superclusters, and specially rich superclusters as a high density environment, which affects the properties of their member galaxies and groups/clusters of galaxies."
    },
    "0609/astro-ph0609055_arXiv.txt": {
        "abstract": "We present a perturbation theory for studying the instabilities of non-axisymmetric gaseous discs. We perturb the dynamical equations of self-gravitating fluids in the vicinity of a non-axisymmetric equilibrium, and expand the perturbed physical quantities in terms of a complete basis set and a small non-axisymmetry parameter $\\epsilon$. We then derive a linear eigenvalue problem in matrix form, and determine the pattern speed, growth rate and mode shapes of the first three unstable modes. In non-axisymmetric discs the amplitude and phase angle of travelling waves are functions of both the radius $R$ and the azimuthal angle $\\phi$. This is due to the interaction of different wave components in the response spectrum. We demonstrate that wave interaction in unstable discs with small initial asymmetries, can develop dense clumps during the phase of exponential growth. Local clumps, which occur on the major spiral arms, can constitute seeds of gas giant planets in accretion discs. ",
        "introduction": "Unstable density waves inside self-gravitating gaseous media play an important role in structuring of spiral galaxies and accretion discs. Prior works in the literature deal with the instability of differentially rotating discs, which are initially circular. Aoki et al. (1979) carried out a global analysis of spiral modes in the soft-centred, gaseous Kuzmin (1956) disc using Hunter's (1963) device. Lemos et al. (1991) studied the unstable axisymmetric modes of the more puzzling scale-free discs and solved the dispersion relation to determine the sufficient conditions for ring formation and for the existence of breathing modes. Recently Goodman \\& Evans (1999, hereafter GE) constructed the normal modes of the scale-free Mestel (1963) disc in the Cowling approximation. For self-gravitating modes they computed the ratio of growth rate to pattern speed by solving a recurrence relation in the Mellin-transform space. Density waves in the presence of an external potential field, like that of a massive central object, have also been extensively studied using semi-analytical and finite difference techniques (e.g., Papaloizou \\& Lin 1989; Savonije \\& Heemskerk 1990; Papaloizou \\& Savonije 1991).  The assumption of axisymmetry of the equilibrium (or initial) state, however, imposes significant simplifications on the solution of dynamical equations. In fact, unstable modes of circular discs are decoupled because the only $\\phi$-dependent term, $\\exp({\\rm i}m\\phi)$ (${\\rm i}=\\sqrt{-1}$), is factored out of the linearized Poisson and hydrodynamic equations. Here $\\phi$ is the azimuthal angle and $m$ is the wave number. Such a factorization is impossible if the unperturbed state is non-axisymmetric. Azimuthal modes of perturbed non-axisymmetric discs can interact with the Fourier components of the equilibrium state and result in more complex patterns.  In this paper we attempt to formulate the instability problem of non-axisymmetric discs based on a perturbation theory, and present a mechanism for the generation of dense clumps through {\\it wave interaction} in {\\it linear regime}. That occurs when a non-axisymmetric disc is destabilized. We apply our formulation to the simple discs of Jalali \\& Abolghasemi (2002, hereafter JA), which reduce to Mestel's disc in axisymmetric limit. We have chosen JA models as our unperturbed state because of two reasons. Firstly, closed form (exact) expressions are available for physical quantities like the potential--density pair, pressure, sound speed and velocity components. Secondly, both the potential and the surface density have a single Fourier component. This minimizes the number of interacting terms in our subsequent analysis.  We review steady-state, non-axisymmetric, scale-free discs of JA in section \\ref{sec:non-axi-symmetric-discs} and introduce a cutout function that helps us handle the singularity at the centre. We present a first-order perturbation theory for non-axisymmetric gaseous discs in section \\ref{sec:perturbation-formulation}, where we also introduce a basis set by which we expand the perturbed physical quantities. We then derive an eigenvalue problem for determining the normal modes of non-axisymmetric discs and discuss on the algorithms used for solving the resulting eigensystem. In sections \\ref{sec::modes-of-circular-discs} and \\ref{sec:modes-nonaxi-discs}, we study the unstable modes of circular and non-axisymmetric discs for the fundamental wave numbers $m=0$, 1 and 2, and discuss on the results and their physical implications in section \\ref{sec:discussion}.  \\begin{figure} \\centerline{\\hbox{\\epsfxsize=8.5truecm\\epsfbox{fig1.eps}}} \\caption{The region under the solid line is the allowable zone in the parameter space of a bisymmetric disc with $n=2$.} \\label{fig:fig1} \\end{figure} ",
        "conclusions": "\\label{sec:discussion} The classical procedure in the instability analysis of gaseous (and stellar) discs is to start with an axisymmetric equilibrium and then perturb the governing equations in the vicinity of this equilibrium state. The subsequent eigenvalue problem [e.g., equation (\\ref{eq:axisymmetric-eigenvalue-equation})] results in a complex frequency $\\omega_m =\\Omega_m+{\\rm i}s_m$, which implies an $m$-fold barred/spiral pattern $\\Sigma_m=A_m(R)\\cos[m\\phi+\\Theta_m(R)]$ that rotates by the angular velocity $\\Omega_m/m$ and exponentially grows by the rate $s_m$. Nevertheless, most events of physical significance happen in nonlinear regime when the modes begin to interact. In such a circumstance, more than one azimuthal wave component is present in the response spectrum. The most general form of a density wave in nonlinear regime is \\begin{equation} \\Sigma_1(R,\\phi,t)=\\sum_{m=-\\infty}^{+\\infty}A_m(R,t)e^{{\\rm i}[m\\phi+\\Theta_m(R,t)]}, \\label{eq::general-density-wave} \\end{equation} which means that the amplitude and phase angle of $\\Sigma_1$ are functions of all spatial and time variables. A primary result of (\\ref{eq::general-density-wave}) is the generation of dense clumps at the locations of the maxima of $\\Sigma_1$. Clumps may undergo gravitational collapse, merge and give birth to gas giants. This is a widely accepted formation scenario of protoplanets (e.g., Boss 1997; Mayer et al. 2004) because it can take place in short time scales. Asymmetries like the presence of a double star (Boss 2006) can also generate dense clumps.  According to (\\ref{eq::mode-nonaxi}) and calculations of \\S\\ref{sec:modes-nonaxi-discs}, unstable density waves of non-axisymmetric discs have a rich structure (in linear regime) even for very small non-axisymmetry parameter $\\epsilon$. Although the temporal evolution of the $m$th mode is governed by the simple function $\\exp({\\rm i}\\omega_m t)$, the eigenfunction associated with $\\omega_m$ is by no means simple and it contains infinite number of azimuthal wave components, which in turn are capable of generating local clumps. Since realistic accretion discs have not initially the {\\it nice} axisymmetry feature, their instability can give birth to protoplanets sooner than what nonlinear hydrodynamical simulations predict. A protoplanet, as a possible product of the gravitational collapse of a dense clump, is separated from the rest of the evolving gas and migrates until it is captured by a resonant island in the phase space.  Our perturbation theory can be applied to any non-axisymmetric configuration regardless of the nature of the centre. Cutouts, which are special treatments designed for scale-free models, are not needed in cored systems with finite total mass. The most important factor in success of our perturbation theory is the smallness of the non-axisymmetry parameter $\\epsilon$, which prohibits the interaction of azimuthal wave components in the zeroth order terms. Our equilibrium density has a single azimuthal Fourier component [see equation (\\ref{eq:simple-disc})]. Accordingly, the harmonic numbers of wave components differ in steps of $nl$ ($l=0,\\pm1,\\pm2,\\cdots$). For bisymmetric discs studied in this paper, this means that wave components have only even or odd harmonic numbers. We will have the most complete spectrum if the equilibrium fields include Fourier terms with both even and odd harmonic numbers, or if the disc is lopsided with $n=1$.  We presented three length scales to the perturbed equation (\\ref{eq::perturbations-for-all}). The first two are the cutout radii, $R_{\\rm in}$ and $R_{\\rm out}$, and the third one is the length scale of the basis functions, $R_C$. At a first look, such a treatment of scale-free systems seems unfavorable because in this way we destroy their scale-invariance and cuspy features. However, with appropriate selection of the exponents of the cutout functions, $N_{\\rm in}$=$N_{\\rm out}$=2, and the weight function of the velocity components, $w_2(\\xi)$, the coefficient matrix $\\textbf{A}$ of our eigenvalue problem became a function of the two dimensionless parameters $\\alpha$ and $\\gamma$. We easily decreased $\\gamma$ and noticed that unstable waves do not extend to distant regions although the mass of our equilibrium models (with flat rotation curves) is infinite at $R\\rightarrow \\infty$. We then let $\\alpha$ tend to zero with the aim of removing the inner cutout. This needed more terms of the series of associated Legendre functions, but the qualitative features of our results remained unchanged as long as mode shapes and ${\\rm arg}(\\omega_m)$ were concerned. There was an excellent match between our results and those of Lemos et al. (1991) who used logarithmic spirals to model their perturbed quantities. Logarithmic spirals do not introduce any length scale to variational equations.  The critical value of $\\Mach$, below which our bisymmetric discs are stable, is a function of the fundamental wave number $m$. Numerical results show that the models with $\\alpha =0.5$ are stable for $m=0$ excitations when $\\Mach \\lta 2.7$. This limit approximately becomes $1.6$ and $1.4$ for $m=1$ and 2, respectively. The most interesting property that we have found for unstable non-axisymmetric discs is that the $m=0$ and $m=1$ modes can be stabilized by increasing the non-axisymmetry parameter $\\epsilon$. This happens for $m=0$ when $2.7 \\lta \\Mach \\lta 2.91$ and for $m=1$ when $1.6 \\lta \\Mach \\lta 2.3$. There are some remarkable differences between the stabilized $m=0$ and $m=1$ modes. The stabilized $m=0$ mode has the form of a ring, which rotates according to the results displayed in Fig. \\ref{fig:fig5}{\\em a}. It occurs when the wave components with the azimuthal harmonic numbers $m-nl=-nl$ and $m+nl=nl$ cancel each other for $l>0$. The stabilized $m=1$ mode, however, is a non-spiral, equiangular neutral mode. It emerges for a sufficiently large $\\epsilon$, when the wave component with the harmonic number $m-n=-1$ ($l=-1$) cancels the fundamental component with $m+n\\times 0=1$ ($l=0$). It is worth noting that the $m=2$ mode cannot be stabilized by increasing $\\epsilon$, because the interacting wave components of ${\\cal O}(\\epsilon)$ have the harmonic numbers $m-n=0$ ($l=-1$) and $m+n=4$ ($l=1$). Neither of these first-order components is capable of suppressing the effect of the fundamental component with the harmonic number $m+n\\times 0=2$. On the other hand, the wave component with the harmonic number $m-n\\times 2=-2$ is of ${\\cal O}(\\epsilon ^2)$, which is too weak to compensate the zeroth order $m=2$ component. Based on the above reasonings, we speculate that the $m=2$ mode can be stabilized in discs whose surface density includes the fourth harmonic of the azimuthal angle or when $n=4$ in the simple discs of JA. Furthermore, the unstable $m=1$ and $m=2$ modes of lopsided discs (with $n=1$) cannot be stabilized by increasing the non-axisymmetry parameter."
    },
    "0609/astro-ph0609325_arXiv.txt": {
        "abstract": "We present raytracing computations for light emitted from the surface of a rapidly-rotating neutron star in order to construct light curves for X-ray pulsars and bursters. These calculations are for realistic models of rapidly-rotating neutron stars which take into account both the correct exterior metric and the oblate shape of the star. We find that the most important effect arising from rotation comes from the oblate shape of the rotating star. We find that approximating a rotating neutron star as a sphere introduces serious errors in fitted values of the star's radius and mass if the rotation rate is very large. However, in most cases acceptable fits to the ratio $M/R$ can be obtained with the spherical approximation. ",
        "introduction": "\\label{s:intro}  One of the most fundamental problems in neutron star astrophysics is the determination of the mass-radius relation through observations. This would allow the observational determination of the equation of state (EOS) of cold supernuclear density material. While the measurement of mass is possible if the star is in a binary, neutron stars are too small to allow a direct measurement of their radii. One promising indirect method for inferring the radius is through observations and modeling of light curves for X-rays emitted from X-ray pulsars and neutron stars that exhibit type I X-ray bursts. This method requires the raytracing of photons emitted from the star's surface in order to predict the signal detected by the observer. In this paper we present the first raytracing calculations for rapidly-rotating neutron stars that include the correct metric and correct shapes of the star, as well as allowing for arbitrary emission and detection directions.  Rapidly-rotating neutron stars are a promising target for an effort to constrain the EOS since observations of their light curves could potentially allow the determination of a star's mass and radius. In the case of a slowly rotating neutron star, the light curve characteristics are controlled by the ratio $M/R$, but the value of $M$ or $R$ for the star is not observable through the light curve. However, if the star is rapidly-rotating, the Doppler boost resulting when the star's equatorial velocity is of the order $v/c \\sim 0.1$ or larger creates an asymmetry in the light curve which makes it possible to measure $v$. Since the spin frequency is always known, the star's radius can be extracted. Knowledge of $R$ and $M/R$ could then potentially be used to constrain the EOS. The purpose of this paper is to determine how accurately these two parameters can be determined through light curve fitting.  In this paper we investigate the effect of rapid rotation on pulse shapes. In order to isolate these effects, we have chosen to simplify our treatment of the shape, size, location and emissivity of the emitting region. It should be understood that the effects due to rapid rotation discussed in this paper are only one ingredient in the modeling of pulse shapes and it will be necessary to combine these effects with more realistic treatments of the emitting regions in order to model real data.  In order to proceed with a program of fitting light curves in order to infer the radius of a neutron star a number of assumptions must be made. The most fundamental assumption made in this paper is that the light originates from the surface of the star and that no material lying between the star and observer causes scattering. We make this assumption so that we can focus on the effects due to rotation on the pulse shape, but it is possible that light is emitted from an extended region off of the star's surface and that the light interacts with the matter surrounding the star. For example, in the case of the slowly rotating X-ray pulsars complicated accretion columns are needed in order to fit the observations \\citep{lea04}. Pulse shape models of real data should allow for the possibility of extended emission regions.     Another important input to light curve models is the spectrum and emissivity of the emitting region. In the case of the 2.5 ms X-ray pulsar \\saxj, \\cite{PG03} have shown that a hybrid spectrum including blackbody and Comptonized emission is needed. In their models, thermal emission is directed towards the normal to the surface and the Comptonized emission is directed away from the normal, creating a strong anisotropy. Other models for the pulse shapes observed in type I X-ray bursts suggest  anisotropic emission as well (\\cite{ML98}, \\cite{Bhat05}). Any anisotropy in the emission is known to have a strong effect on the light curve, so realistic light curve models must include the effects of anisotropy. In our present work, we have simplified our models to only allow isotropic emission. We expect that the magnitude of the rotational effects discussed in this paper are similar to the magnitude of the effects caused by anisotropic emission.   The assumed shape of the emitting region also strongly affects the pulse profile. Most studies have focused on either infinitesimal spots or simple Gaussian profiles. Recently \\cite{KR05} have produced MHD simulations of accretion flows onto neutron stars that show the hot spots on X-ray pulsars could have more complicated shapes, confirming predictions made by \\cite{GL78}. For the case of type I X-ray bursts it is possible for normal modes to be excited \\citep{Heyl04}, which would produce patterns associated with the modes \\citep{Heyl05,LS05}. In our calculations we have chosen the simplest possible pattern (an infinitesimal uniform brightness spot) in order to focus on the effects due to rotation. Models of real data must allow for more complicated emission shapes and brightness patterns.  The gravitational field outside of the star is important, for it gravitationally redshifts the photons and bends their direction of propagation. For non-rotating stars the Schwarzschild metric is all that is required to describe the photon trajectories \\citep{PFC83}. The Schwarzschild + Doppler (S+D) approximation \\citep{ML98,PG03} was introduced as a simple method for calculating the light curves of slowly rotating stars. In the S+D approximation the gravitational field is modeled by the static Schwarzschild solution and the rotational effects are approximated by using the special relativistic Doppler transformations. One of the goals of this paper is to find the conditions under which the S+D approximation will fail. For rapidly-rotating neutron stars, the metric must be computed numerically using one of many known algorithms, such as that of \\cite{CST94}. The Kerr black hole metric can be made to agree to first order in angular velocity with the metric of a rotating neutron star metric by choosing the Kerr parameter $a$ to be $J/M$ where $J$ is neutron star's angular momentum. The Kerr black hole is a poor approximation to a rotating neutron star for spin frequencies above 400 Hz (for an example involving the innermost stable circular orbit, see Miller et al 1998). However, some calculations (Braje et al 2000) have suggested that the Kerr metric may suffice for raytracing applications. This motivates \\citep{BRR00,Bhat05} the use of the Spherical Kerr approximation (SK) where a spherical surface is embedded in the Kerr spacetime. In this paper we will show that when the S+D approximation fails, the SK approximation also fails, so that the SK approximation is not necessarily a better choice than S+D. In this paper we compute the metric of rapidly-rotating neutron stars using realistic equations of state and perform raytracing on the computed numerical spacetime. In a previous paper \\citep{CLM05} (CLM) we performed similar calculations that were restricted to the equatorial plane. In this paper we generalize to allow arbitrary initial spot locations and arbitrary observer locations.  The final important input to a raytracing program is the shape of the star's surface, which for a rotating fluid star is an oblate spheroid. In almost all previous calculations it has been assumed that the surface of the star is a sphere, irrespective if the exterior gravitational field was assumed to be Schwarzschild or Kerr. We are aware of only one previous calculation, by \\citet{BR01} that takes the oblate surface into account. In that paper they consider further scatterings of photons within the magnetosphere, so that the effect of oblateness is not clearly separated from other effects. In order to motivate the choice of a spherical surface, it is often stated that the ratio of polar axis to equatorial axis is very close to unity, even for rapid rotation. However, this ratio is neither coordinate invariant, nor directly observable. The oblate shape of the star has two main effects. (I) The gravitational field near the the poles is stronger than near the equator, so the redshift and photon deflection are larger for photons emitted near the poles. (II) Light is emitted at different angles with respect to the normal to the surface if the surface is an oblate spheroid instead of a sphere.  As a result some parts of a spherical star that an observer can't see will be visible if the star is actually oblate (the converse is also true). Of these effects, the second is the most important, and leads to failure of the S+D and SK approximations for rapid rotation.  Another application which makes use of raytracing is the calculation of absorption line profiles. Absorption line profiles have been computed in the past using the S+D approximation \\citep{OP03,VS04} and using the SK approximation \\citep{BML06}. Interestingly, \\citet{BML06} have identified signs of frame-dragging in the line profiles computed using the Kerr metric for very rapidly-rotating neutron stars. \\citet{CMBW06} compared line profiles using oblate stars (using the same type of method used in this paper) with the line profiles calculated in the spherical approximation and found that the differences are insignificant for line profile calculations. However, it should be noted that the calculations done by \\citet{CMBW06} assumed that light from all of the star contributes to the line profile, which would tend to average out differences arising from the shape of the star.   In this paper we wish to explore the rapid rotation effects of the metric and oblateness on the calculated light curves of neutron stars and on the fitting of a star's mass and radius from the light curves. In order to isolate these effects, we consider the simplest possible assumptions for all other aspects. We assume isotropic emission from an infinitesimal spot, but allow for spots and observers at any location. As discussed in the previous paragraphs, there is enormous uncertainty in the shape, size and emissivity of the emitting regions which will have the effect of making it difficult to discern the relativistic effects discussed in this paper. However, we will show that for very rapidly rotating neutron stars, the relativistic rotational effects are large and should be incorporated in models of real data. In section \\ref{s:light} we describe the method used to compute the light curves for rotating neutron stars. In section \\ref{s:approx} we outline the four types of approximation schemes that could be used to produce light curves. In section \\ref{s:results} we compare the light curves from the exact and approximate methods. We show how programs that attempt to fit light curves using either spherical approximation will converge to incorrect values of mass and radius at spin frequencies larger than about 300 Hz, although they may converge to the correct value of the ratio $M/R$. We conclude with a discussion of these results in section \\ref{s:conclusions}. ",
        "conclusions": "\\label{s:conclusions}  We have investigated the effect of rapid rotation on the light curves produced by a small spot on a neutron star's surface. Our calculations are exact, in that we numerically compute the spacetime metric for the rotating star, include the correct shape of the star and solve for the geodesics in the numerical spacetime. We find that the most important effect on the light curve arising from rotation comes from the oblate shape of the star's surface.   All past work involving fits of X-ray data have assumed that the surface of the star is spherical. Our present computations show that the assumption of a spherical surface leads to large deviations from the correct light curve for spin frequencies larger than 300 Hz. However, our test runs show that the assumption of a spherical surface can still fit the correct value of $M/R$ at spin frequencies $\\le 500$ Hz, even if the individually fitted values of mass and radius are incorrect.   In our light curve calculations, we find that if the star's surface is assumed to be a sphere, the light curves resulting from the use of the Kerr metric are almost the same as those resulting from the use of the Schwarzschild metric. The difference in these light curves is smaller than the size of the error bars in any present data, so it seems that there is no significant advantage in the use of the SK approximation over the S+D approximation.  We calculated light curves in an approximation where an oblate surface is embedded in either the Schwarzschild (OS) or Kerr (OK) spacetime. In both cases, the resulting light curves are very close to the exact light curves computed using the numerical spacetime of the rotating neutron star. This suggests that an approximation similar to the Schwarzschild + Doppler approximation that includes an oblate surface could be a useful way to fit data. The construction of a fitting program that extracts the mass and radius of a star from a light curve using an oblate surface is non-trivial. In future work, we plan to develop a code that will allow an oblate shape to be incorporated into a S+D fitting program that includes finite spot sizes and anisotropic emission. We anticipate that such a program will allow the accurate extraction of the masses and radii of very rapidly-rotating neutron stars."
    },
    "0609/astro-ph0609113_arXiv.txt": {
        "abstract": "We have analyzed Chandra observations of 18 low-luminosity early-type galaxies with $L_B \\lax 3 \\times 10^{10} L_{\\odot B}$.  Thermal emission from hot gas with temperatures between 0.2 and 0.8~keV comprises 5-70\\% of the total 0.5-2.0~keV emission from these galaxies.   We find that the total X-ray luminosity from LMXBs (resolved plus the power-law component of the unresolved emission) scales roughly linearly with the K-band luminosity of the galaxies with a normalization comparable to that found in more luminous early-type galaxies. All of the galaxies in our sample are gas poor with gas masses much less than that expected from the accumulation of stellar mass loss over the life time of the galaxies. The average ratio of gas mass to stellar mass in our sample is $M_{gas}/M_*=0.001$, compared to more luminous early-type galaxies which typically have $M_{gas}/M_*=0.01$. The time required to accumulate the observed gas mass from stellar mass loss in these galaxies is typically $3 \\times 10^8$~yr.  Since the cooling time of the gas is longer than the replenishment time, the gas cannot be condensing out of the hot phase and forming stars, implying that the gas is most likely being expelled from these galaxies in a wind. The one exception to this is NGC4552, which is the most optically luminous galaxy in our sample and has the highest gas content.  Using recent estimates of the Type Ia supernova rate and AGN heating rate in early-type galaxies, we find that, on average, heating by Type Ia supernovae should exceed AGN heating in galaxies with $L_B \\lax 3 \\times 10^{10} L_{\\odot B}$. We also find that heating by Type Ia supernovae is energetically sufficient to drive winds in these galaxies, even if the present Type Ia supernova rate is overestimated by a factor of two or the present stellar mass loss rate is underestimated by a factor of two.  There is significant scatter in the gas properties of galaxies with comparable optical luminosities. Nearly continuous heating by Type Ia supernovae alone cannot account for the large scatter in observed gas mass.  We suggest that the scatter in gas properties could arise from periodic AGN outbursts, during which time AGN heating dominates over Type Ia supernova heating, or environmental factors, such as a high pressure environment which would suppress the formation of strong galactic winds. ",
        "introduction": "Heating by supernovae and active galactic nuclei (AGN) are thought to play important roles in galaxy formation, generating the observed correlation between the bulge mass of a galaxy and the mass of the central supermassive black hole, the upper mass cutoff of galaxies and reheating the gas in galaxies, groups and clusters. In order to reconcile the predictions of the cold dark matter hierarchical clustering scenario with the observed luminosity function of galaxies and their mass-to-light ratios, it is thought that supernova feedback regulates star formation in galaxies less luminous than about $3 \\times 10^{10} L_{\\odot B}$, while AGN feedback regulates star formation in more massive galaxies (e.g., White \\& Frank 1991; Bower et al. 2005; Croton et al. 2006).  Only Type II supernovae provide a natural feedback mechanism due to the short time delay between the onset of star formation and the first supernova. While Type Ia supernovae (SN Ia) cannot provide a feedback mechanism that regulates concurrent star formation, SN Ia can play a significant role in the fate of gas shed by evolving stars in low-mass galaxies long after star formation ceases.    Einstein observations showed that the bulk of the X-ray emission from optically luminous early-type galaxies arises from hot gas in hydrostatic equilibrium (Forman, Jones \\& Tucker 1985). A harder X-ray component, more prevalent among low-luminosity early-type galaxies, was detected by ASCA and assumed to arise from low-mass X-ray binaries (LMXBs) due to the old stellar population in these systems (Kim, Fabbiano \\& Trinchieri 1992). Based on the analysis of ASCA data, White, Sarazin \\& Kulkarni (2002) found that the luminosity of the hard X-ray component was more strongly correlated with the luminosity in globular clusters rather than the total optical luminosity of the galaxies. Chandra, with its superior angular resolution, has resolved populations of point sources in many early-type galaxies (e.g., Sarazin, Irwin \\& Bregman 2000; Angelini, Lowenstein \\& Mushotzky 2001; Kraft et al. 2001; Blanton, Sarazin \\& Irwin 2001) and found that 20 to 80\\% of the point sources reside in globular clusters (Sarazin et al. 2003).  While the hot gas content of X-ray luminous early-type galaxies has been well studied by Chandra and previous X-ray telescopes, little is known about the properties of the gas in less luminous early-type galaxies.  In low-luminosity early-type galaxies, the X-ray emission from LMXBs can exceed the thermal emission from hot gas, so it is imperative that the LMXBs be detected and excised from the analysis of the diffuse emission. Theoretical models concerning the evolution of early-type galaxies show that for typical SN Ia rates, early-type galaxies fainter than $M_B \\approx -20$ should possess SN Ia driven galactic winds (e.g., David, Forman \\& Jones 1990, 1991). In addition to heating by SN Ia, AGN heating may also expel a significant portion of the gas shed by evolving stars in the shallow potential well of low-luminosity galaxies.  Chandra has detected X-ray cavities filled with radio emitting plasma and AGN driven shocks in many elliptical galaxies and clusters (e.g., McNamara  \\begin{table*}[t] \\begin{center} \\caption{Low-Luminosity Early-Type Galaxy Sample} \\begin{tabular}{lccccccc} \\hline Name & $D_L$ & $M_B$ & Type & $L_K$ & $\\sigma_*$ & ObsIds & Exposure\\\\ & (Mpc) & & & ($L_{\\odot ,K}$) & (km~s$^{-1}$) & & (ksec) \\\\ \\hline\\hline ESO 0428-G014 & 24.3 & -19.96 & -1.6 &  $3.62 \\times 10^{10}$ & - & 4866 & 28.1 \\\\ NGC 821  & 24.1 & -20.19 & -5.0 &  $8.79 \\times 10^{10}$ &  199.0 & 4006 4408 & 22.9 \\\\ NGC 1023 & 11.4 & -20.21 & -3.0 &  $8.44 \\times 10^{10}$ &  204.0 & 4696 & 10.5 \\\\ NGC 1386 & 16.5 & -18.97 & -0.6 &  $3.47 \\times 10^{10}$ &  166.0 & 4076 & 19.6 \\\\ NGC 1389 & 21.7 & -19.24 & -3.3 &  $3.91 \\times 10^{10}$ &  132.0 & 4169 & 38.4 \\\\ NGC 2434 & 21.6 & -20.10 & -5.0 &  $7.38 \\times 10^{10}$ &  205.0 & 2923 & 19.2 \\\\ NGC 2787 & 7.48 & -17.76 & -1.0 &  $1.46 \\times 10^{10}$ &  194.0 & 4689 & 31.4 \\\\ NGC 3115 & 9.68 & -20.19 & -3.0 &  $8.95 \\times 10^{10}$ &  266.0 & 2040 & 34.6 \\\\ NGC 3245 & 20.9 & -19.95 & -2.0 &  $5.92 \\times 10^{10}$ &  220.0 & 2926 & 9.6  \\\\ NGC 3377 & 11.2 & -19.18 & -5.0 &  $2.75 \\times 10^{10}$ &  131.0 & 2934 & 38.7 \\\\ NGC 3379 & 10.6 & -19.94 & -5.0 &  $7.31 \\times 10^{10}$ &  201.0 & 1587 & 31.2  \\\\ NGC 3608 & 22.9 & -20.11 & -5.0 &  $5.24 \\times 10^{10}$ &  204.0 & 2073 & 38.1 \\\\ NGC 4251 & 19.6 & -20.03 & -2.0 &  $6.67 \\times 10^{10}$ &  119.0 & 4695 & 10.5 \\\\ NGC 4435 & 11.4 & -18.68 & -2.0 &  $2.38 \\times 10^{10}$ &  157.0 & 2883 & 24.6 \\\\ NGC 4459 & 16.1 & -19.83 & -1.0 &  $7.58 \\times 10^{10}$ &  172.0 & 2927 & 9.5 \\\\ NGC 4552 & 15.3 & -20.36 & -5.0 &  $1.05 \\times 10^{11}$ &  261.0 & 2072 & 53.3 \\\\ NGC 4697 & 11.7 & -20.28 & -5.0 &  $9.12 \\times 10^{10}$ &  165.0 & 4727 4728 & 72.6 \\\\ NGC 5866 & 15.3 & -20.10 & -1.0 &  $9.04 \\times 10^{10}$ &  159.0 & 2879 & 28.0 \\\\ \\hline \\end{tabular} \\end{center}  \\noindent Notes:  Galaxy name, luminosity distance ($D_L$), absolute blue magnitude ($M_B$), Galaxy Type in RC3, K-band luminosity ($L_K$) based on $M_{K\\odot}=3.33$ within the the same aperture as that used to extract the X-ray spectrum, stellar velocity dispersion ($\\sigma_*$), Chandra ObsIds used in the analysis, and the cleaned ACIS exposure time. \\end{table*}  \\smallskip  \\noindent et al. 2000, Finoguenov \\& Jones 2002, Fabian et al. 2003, Blanton et al. 2003, Nulsen et al. 2005a, Nulsen et al. 2005b; Forman et al. 2005).  Based on the analysis of cavities found in a sample of galaxies, groups and clusters, Birzan et al. (2004) determined that the mechanical power of AGNs can be up to $10^4$ times their radio power. Thus, the thermal and dynamic properties of the hot gas in low-luminosity early-type galaxies are sensitive probes of the present SN Ia rate and AGN activity in these galaxies.  This paper is organized in the following manner.  In $\\S$~2, we present our low-luminosity early-type galaxy sample.  Section 3 contains the details of our Chandra data reduction.  The spectroscopic results for the LMXBs and diffuse emission are discussed in $\\S$~4 and the observed scaling between the total X-ray luminosity of the LMXB population and the K-band luminosity of the galaxies is presented in $\\S$~5.  Sections 6 and 7 discuss the properties of the hot gas, the inferred characteristics of the galactic winds, and the implications regarding the Fe abundance in galaxies with galactic winds.  In $\\S$~8, we examine the relative importance of AGN and SN Ia heating in these galaxies.  Possible origins of the large scatter in the observed gas properties among the galaxies are discussed in $\\S$~9 and the main results of our paper are summarized in $\\S$~10. ",
        "conclusions": "If heating by SN Ia is the dominant mechanism for expelling the gas shed by evolving stars from low-luminosity early-type galaxies, it is a puzzle as to why there is so much scatter in $M_{gas} / M_*$ and $L_{X,gas}$ (see Figs. 7 and 11) given that supernova heating is essentially a continuous process. Since the gas masses are computed in different apertures (see Table 5), we checked to see if there is a correlation between $M_{gas} / M_*$ and aperture size and found none.  While the time-averaged SN Ia heating rate shown in Fig. 11 exceeds the time-averaged AGN heating rate, except for possibly the most luminous galaxies in our sample, the episodic nature of AGN outbursts could produce times during which mechanical AGN heating exceeds SN Ia heating. Thus, the large observed scatter could be a reflection of recent AGN activity.  The sound crossing time within the central 10~kpc in these galaxies is approximately $3 \\times 10^7$~yr. If AGN outbursts repeat on a shorter time scale, then the gas will not be able to establish a steady-state wind. In the very center of the galaxies, heating by SN Ia cannot be thought of as a continuous process.  For example, within an enclosed stellar mass of $10^6 M_{\\odot}$, there will only be one SN Ia every $3.0 \\times 10^7$~yr.  During this time, the stars within this region will shed a total of 50~$M_{\\odot}$ of gas.  If this gas is converted into mechanical energy by the central black hole with an efficiency of 10\\%, this will produce $1.0 \\times 10^{55}$~erg of energy at a rate of $1.0 \\times 10^{40}$~erg~s$^{-1}$, which is comparable to the SN Ia heating rate of an entire galaxy (see Fig. 10). Thus, it could be that the gas in the very center of the  \\begin{inlinefigure} \\center{\\includegraphics*[width=0.90\\linewidth,bb=10 142 570 700,clip]{f11.eps}} \\caption{Scatter plot of the bolometric X-ray luminosity of the gas vs. the K-band luminosity of the galaxies. The inverted arrows are 90\\% upper limits. The solid line is the estimated SN Ia heating rate and the dashed line is the estimated mechanical AGN heating rate.} \\end{inlinefigure}  \\smallskip  \\noindent galaxies does not partake in the global SN Ia wind, but periodically accretes into the central black hole.   Environment can also play a significant role in the gas content of early-type galaxies. XMM-Newton and Chandra observations of early-type galaxies in the Coma cluster have shown that these galaxies can be either under luminous or over luminous in X-rays relative to field galaxies with comparable optical luminosities due to the competing effects of adiabatic compression, ram pressure stripping and thermal evaporation (Vikhlinin et al. 2001; Finoguenov \\& Miniati 2004; Hornschemeier et al. 2005).  Due to the large distance of the Coma cluster, these observations were only sensitive to the X-ray properties of luminous early-type galaxies which have massive hydrostatic coronae in regions outside of dense groups or clusters. Low luminosity field galaxies, however, should have strong winds due to the lack of a confining ambient medium.    NGC4552 is located in the densest environment of any of the galaxies in our sample and clearly shows the effects of its environment. It is located $72^{\\prime}$ (a projected distance of 360~kpc) from the center of the Virgo cluster and has the highest gas content and hosts the most radio luminous AGN in our sample. The Chandra image shows that NGC 4552 has a leading cold front and a trailing wake of ram pressure stripped gas (Machacek et al. 2005).  Based on the pressure jump across the cold front, Machacek et al. estimated that NGC4552 is traveling at 1600~km~s$^{-1}$. The average gas density and pressure in NGC4552 derived from the values in Table 5 are $<n_e> = 1.7 \\times 10^{-3}$~cm$^{-3}$ and $<P>=1.9 \\times 10^{-3}$~keV~cm$^{-3}$.  Using the results in Machacek et al., the ram pressure of the Virgo gas at the leading cold front in NGC4552 is $P_{ram} = 2.2 \\times 10^{-3}$~keV~cm$^{-3}$, which is greater than the average thermal pressure of the gas, but less than the central pressure of $P_0 = 3.5 \\times 10^{-2}$~keV~cm$^{-3}$.  These calculations are in good agreement with the fact that some of the outer gas in NGC452 is being stripped as it falls into the Virgo cluster, but not the central gas. Based on the inferred velocity of NGC4552, it should have traveled approximately 2~Mpc during the time it has taken to accumulate its present gas mass.  The ROSAT all-sky survey data show that the gas in Virgo follows a $\\beta=0.47$ profile (Schindler et al 1999). Assuming the Virgo gas is isothermal and a radial trajectory for NGC4552, the thermal pressure of the Virgo gas surrounding NGC4552 would have increased by a factor of 20 during this time and the ram pressure by an even greater factor due to the increasing infall velocity. Thus, the ambient pressure during most of this time was sufficient to suppress the formation of a galactic wind, but insufficient to strip the gas from the galaxy. This would have lead to a substantial increase in gas mass which could have triggered the AGN outburst."
    },
    "0609/astro-ph0609439_arXiv.txt": {
        "abstract": "We present deep F606W, F814W  ACS photometry of the recently discovered globular cluster B514, the outermost known globular in the M31 galaxy. The cluster appears quite extended and member stars are unequivocally identified out to $\\sim 200$ pc from the center. The Color Magnitude Diagram reveals a steep Red Giant Branch (RGB), and a Horizontal Branch (HB) extending blue ward of the instability strip, indicating that B514 is a classical old metal-poor globular cluster. The RGB locus and the position of the RGB Bump are both consistent with a metallicity $[Fe/H]\\sim -1.8$, in excellent agreement with spectroscopic estimates. A preliminary estimate of the integrated absolute V magnitude ($M_V\\la -9.1$) suggests that B514 is among the brightest globulars of M31. ",
        "introduction": "Stars and stellar systems orbiting the most external regions of their parent galaxy are crucial tracers of the mass distribution and the formation history of the parent galaxy as a whole. For example, it is well known that substructures related to accretion events are best preserved in the outskirts of a galactic halo \\citep[see, for instance,][and references therein]{bj04}. In this context globular clusters may serve as excellent tracers of substructures in the outer region of their parent galaxy. For instance, \\citet{bfi} were able to identify the accretion signature of the  of the Sagittarius dwarf galaxy among the globular clusters in the outer halo of the Milky Way.  As large substructures are also identified  in the nearby spiral galaxy M31 \\citep{m31nat} it becomes increasingly important to have a complete census of the cluster population in the  galaxy outskirts - still lacking to date \\citep[see][]{G05,G06} - as well as to collect information about cluster ages, metal content, kinematics, etc. Several groups are currently working on this line \\citep{hux04,bat}, including ourselves \\citep[see][]{G04,G05,G06}. In particular, in \\citet[][hereafter G05]{G05} we reported on the discovery of the outermost known cluster of M31, B514 located at a projected distance of $R_p\\simeq 55$ Kpc from the center of M31, not far from the galaxy major axis. The integrated photometry and spectroscopy indicates that B514 is likely an old metal-poor globular cluster (G05), however direct imaging, at a spatial resolution sufficient to resolve the system into stars, is the only way to ultimately ascertain the nature of a candidate globular in M31 \\citep[see][]{G04,G06}.  Here we present the first results of follow-up observations of B514 taken with the Wide Field Camera (WFC) of the Advanced Camera for Surveys (ACS) on board of the Hubble Space Telescope (HST), confirming that the object is a genuine old and metal poor cluster lying in the remotest region of M31. ",
        "conclusions": ""
    },
    "0609/astro-ph0609396_arXiv.txt": {
        "abstract": "We report the discovery of twenty-one hitherto unknown bright southern ultracool dwarfs with spectral types in the range M7 to L5.5, together with new observations of a further three late M dwarfs previously confirmed. Three more objects are already identified in the literature as high proper motion stars; we derive their spectral types for the first time. All objects were selected from the 2MASS All Sky and SuperCOSMOS point source databases on the basis of their optical/near-infrared colours, $J$-band magnitudes and proper motions. Low resolution (R\\,$\\sim$\\,1000) $JH$ spectroscopy with the ESO/NTT SOFI spectrograph has confirmed the ultracool nature of 24 targets, out of a total of 25 candidates observed. Spectral types are derived by direct comparison with template objects and compared to results from H$_2$O and FeH indices. We also report the discovery of one binary, as revealed by SOFI acquisition imaging; spectra were taken for both components. The spectral types of the two components are L2 and L4 and the distance $\\sim$\\,19\\,pc. Spectroscopic distances and transverse velocities are derived for the sample. Two $\\sim$\\,L5 objects lie only $\\sim$\\,10\\,pc distant. Such nearby objects are excellent targets for further study to derive their parallaxes and to search for fainter, later companions with AO and/or methane imaging. ",
        "introduction": "In recent years, the search for faint dwarfs in the Solar neighbourhood has been the subject of much attention. Such studies have been motivated by our still highly incomplete knowledge of the Sun's coolest and least massive neighbours, as pointed out by, for example, \\citet{hen97}. M dwarfs account for around 70\\% of the stellar number density in the Solar neighbourhood, and yet \\citet{hen02} estimated as many as 63\\% of stellar systems within 25\\,pc remain undiscovered, while the figure for missing systems within 10\\,pc may be near 25\\% \\citep{rei03}, or even higher (30\\%; \\citealt{hen02}). Current and recent searches rely primarily on proper motion (\\citealt{sub05}; \\citealt{dea05}; \\citealt{pok03}), in conjunction with photometry and spectroscopic follow-up (\\citealt{hen04}; \\citealt{lod05}; \\citealt{sch05}; \\citealt{pha06}; \\citealt{pok05}).  Efforts in the field have been enormously aided by the digitization of wide-field survey plates by e.g. the SuperCOSMOS project (\\citealt{ham01}; \\citealt{hen04}), and by the advent of all-sky near-infrared surveys, most notably 2MASS \\citep{skr97} and DENIS \\citep{epc97}, together with the Sloan Digital Sky Survey (SDSS; \\citealt{yor00}).  Moreover, observations of the low-mass dwarf population in the Solar neighbourhood have required the establishment of two new spectral classes; the L and T dwarfs (\\citealt{kir99}; \\citealt{mar99}), with effective temperatures below that of the coolest M dwarfs. These discoveries have resulted in a large part from the exploitation of all-sky surveys: we may cite \\citet{bur04a} and \\cite{chi06} as two recent examples notable in for finding large numbers of both L and T dwarfs. In the near future, the field will be further boosted by next-generation infrared surveys, notably the UKIDSS Large Area Survey \\citep{law06}. New L and T dwarf discoveries are increasingly populating the immediate Solar neighbourhood.  In this paper, we report the discovery of twenty-one new ultracool dwarfs (later than M6; \\citealt{bou05}), the majority of which are estimated spectroscopically to lie closer than $\\sim$\\,30\\,pc, and which have spectral types later than M7. The two latest objects, with $\\sim$\\,L4-L6 spectral types, are good candidates to be very nearby; close to 10\\,pc, and with reference to theoretical isochrones \\citep{bar03}, are very likely substellar, i.e. with masses below the hydrogen burning limit of 0.072\\,M$_{\\odot}$ \\citep{leg01}. In the next Section, we overview the aims of this research and methods of target selection. The spectroscopic observations are detailed in Sect. 3. In Sect 4., we present our results and discuss the methods we have used to analyse our spectra and derive spectroscopic distances and transverse velocities for the targets. These results are discussed in Sect. 5. In Sect. 6 we show that one of the targets is a binary, with $\\sim$\\,L2 and $\\sim$\\,L4 components. A summary of our conclusions is presented in Sect. 7.   \\begin{table*} \\vbox to 220mm{\\vfil Landscape Table 1 to go here. \\caption{} \\vfil} \\label{t1} \\end{table*}  \\section[]{Survey aims and target selection}  \\subsection{Survey overview and goals}  The aim of this project is to locate the remaining undiscovered, yet apparently bright, southern ultracool dwarfs in the Solar neighbourhood. By primarily using a combination of their red optical and near-infrared colours, and secondarily {\\it via} selecting objects from this photometric sample which have the highest reduced proper motions, the survey presented here is designed to reveal bright, and therefore nearby dwarfs undiscovered by previous studies. By requiring candidates to have an $R$-band detection, the search is biased towards nearer and/or more intrinsically bright, earlier type dwarfs. A secondary bias is introduced by the use of the reduced proper motion, which as a combination of absolute magnitude and transverse velocity \\citep{pok03} may lead to the preferential discovery of high velocity objects with possible thick-disk or halo kinematics. In general, however, our aim is to find those ultracool dwarfs, optically bright enough to be detected in the SuperCOSMOS database, which lie above our reduced proper motion threshold. This corresponds to only a few $\\times$\\,10$^2$\\,mas\\,yr$^{-1}$ in most cases (see Table \\ref{t1}).  \\subsection{Target selection}  \\begin{figure} \\includegraphics[width=9cm]{test1.eps} \\caption{Illustration of the $R-K, J-K$ cut used to select candidate ultracool dwarfs. Known such objects (filled triangles) lie above and to the left of the line shown. Candidates (circles) in this region were retained.} \\label{z1} \\end{figure}  \\begin{figure} \\includegraphics[width=9cm]{test2.eps} \\caption{As Fig.\\,\\ref{z1}, for $J-K, I-J$. Candidates lie below and to the right of the line shown. } \\label{z2} \\end{figure}  \\begin{figure} \\includegraphics[width=9cm]{test3.eps} \\caption{Illustration of the removal of false matches with the SuperCOSMOS $R-I$ colour. True matches to known ultracool dwarfs lie above the horizontal line. } \\label{z3} \\end{figure}  \\begin{figure} \\includegraphics[width=9cm]{test.eps} \\caption{Efficiency of the $R-K$ colour to select ultracool dwarfs. The plot shows 90 out of 103 known dwarfs with spectral types M6 or later and with SuperCOSMOS $RI$ counterparts retained by our photometric selection methods (triangles). Circles are 317 candidates occupying the same region of the $R-K,I-J$ diagram, from which 61 high proper motion candidates were drawn.} \\label{f1} \\end{figure}  Targets were selected by interrogation of the current photometric and kinematic data provided by the 2MASS All-Sky point source and SuperCOSMOS\\footnote{{\\tt http://surveys.roe.ac.uk/ssa/}} surveys covering the complete Southern Hemisphere. Candidate ultracool dwarfs were drawn from 2MASS with $|b|$\\,$>$\\,10$^{\\circ}$, and with $JHK$ data passing cuts based on those of \\citet{cru03}. Preliminary cuts on the 2MASS sample were:  \\begin{description} \\item{11\\,$\\leq$\\,$J$\\,$\\leq$\\,17  \\item $J-K$\\,$>$\\,1  \\item $J$\\,$\\leq$\\,3\\,$\\times$\\,$(J-K)$\\,+\\,10.5} \\end{description}  \\noindent yielding an initial sample of 1055887 2MASS objects, which was immediately reduced to 109237 by removing any point with photometric quality flag 'U' (upper limit) on any magnitudes.  Excluding 52387 objects in regions of high stellar surface density (e.g. the LMC and SMC, star-forming and other regions at low galactic latitude defined by \\citet{cru03}, with very slight modifications), the search yielded 8101 objects simultaneously satisfying the following cuts:  \\begin{description} \\item{$(J-H)$\\,$\\leq$\\,1.75\\,$\\times$\\,$(H-K)$\\,+\\,0.1875  \\item$(J-H)$\\,$\\geq$\\,1.75\\,$\\times$\\,$(H-K)$\\,-\\,0.475} \\end{description}  and additionally:  \\begin{description} \\item{$(J-H)$\\,$\\leq$\\,0.8 where 0.3\\,$\\leq$\\,$(H-K)$\\,$\\leq$\\,0.35  \\item $(J-H)$\\,$>$\\,1\\,--\\,$(H-K)$} \\end{description}  \\noindent where the last cut is a translation of the initial $J-K$\\,$>$\\,1 criterion.  Objects with $\\delta$\\,$\\ga$\\,3$^{\\circ}$ were not considered further, so as to ensure SuperCOSMOS optical coverage of the sample,  keeping 4074 objects, and of these only 2MASS photometry better than flag\\,=\\,CCC (SNR\\,$\\geq$\\,5 and $\\Delta$$JHK$\\,$\\leq$\\,0.2), was kept, retaining 2317 objects.  A search radius of 18\\arcsec\\,was used to search for a SuperCOSMOS optical counterpart for each object, to allow for proper motion, and which,  given the typical difference in epoch between the two surveys ($\\sim$\\,20\\,yr) allowed retention of candidate field dwarfs with $\\mu_{\\rm tot}$ up to $\\sim$\\,1\\,\\arcsec\\,yr$^{-1}$. Although closer epoch differences could allow higher $\\mu$ objects through, with possibly less accurate measurements, this survey has not been designed to yield the very highest proper motion objects, which are in any case expected to be far less numerous. In contrast to, for example \\citet{pok05} and \\citet{dea05}, this research instead is primarily designed to find late-type ($\\ga$\\,M7), ultracool dwarfs {\\it via} a combination of their optical and near-infrared colours, using proper motion to remove more distant contaminants.  At the same time, a total of 135 known ultracool objects taken mostly from \\citet{cru03} were found in the SuperCOSMOS data (i.e. there was one or more optical counterpart with $R$ and $I$ photometry) and a  total of 1245 other counterparts to the 2MASS data were also found, but these included false matches and a high proportion of contaminating objects (e.g. AGB carbon stars). Further cuts were therefore applied to the dataset of optical counterparts to the 2317 2MASS sources, to retain as many known ultracool dwarfs and discard other data points, using the loci of known objects in the $R-K,J-K$ (Fig.\\,\\ref{z1}) and $J-K,I-J$ (Fig.\\,\\ref{z2}) diagrams as a guide. These cuts were:  \\begin{description}  \\item{$(R-K)$\\,$\\geq$\\,4$\\times$\\,$(J-K)$\\,+\\,0.625  \\item $(J-K)$\\,$\\leq$\\,0.5$\\times$\\,$(I-J)$\\,+\\,0.25}  \\end{description}  together with:  \\begin{description} \\item{5.5\\,$\\leq$\\,$(R-K)$\\,$\\leq$\\,9} \\end{description}  \\noindent to isolate L dwarfs.  \\begin{figure} \\includegraphics[width=9cm]{pmbw.eps} \\caption{Selection of targets by reduced proper motion (H$_{\\rm J}$\\,=\\,$J$\\,+\\,5log($\\mu_{\\rm tot}$) with $\\mu$ in mas\\,yr$^{-1}$) and $I-J$ colour. Small triangles are data for the 90 known L dwarfs retained by our selection methods (see text) {\\it after} use of optical vs. infrared colour criteria. The new object outside the selection criteria was originally mismatched in the 2MASS/SuperCOSMOS cross-correlation.} \\label{f2} \\end{figure}  Moreover, false matches were almost completely removed by rejecting points with $(R-I)$\\,$\\leq$\\,2.2 (Fig.\\,\\ref{z3}), although note that the SuperCosmos $R$ filter does not have the same characteristics as Cousins $R$, so this last cut is not necessarily applicable when comparing 2MASS magnitudes to optical magnitudes from other surveys. Ninety known dwarfs were retained out of 103 passing the first three near infrared-only colour cuts. New candidate L dwarfs with optical counterparts passing the above optical/near-infrared cuts numbered 317 in total. The efficiency of the $R-K$ colour to select ultracool dwarfs is shown in Fig.\\,\\ref{f1} where the 90 known objects and 317 potential candidate new dwarfs are shown to occupy the same region of the $R-K,I-J$ diagram. Indeed, a 2D Kolmogorov-Smirnov test shows this is the case with 99.3\\% probability. Our final southern sample consists of 61 of these 317 , selected to have reduced proper motion H$_{\\rm J}$\\,$\\geq$\\,25 (in milliarcsecond units) and $(I-J)$\\,$\\geq$\\,3.1, typical for L dwarfs \\citep{ken04}. The 90 known L dwarfs passing these photometric cuts are plotted in Fig.\\,\\ref{f2} together with the 21 new discoveries presented here, and three further late M/early L objects from \\citet{lod05} for which we present new spectroscopic data. In the last part of the selection process, candidates appearing in the SIMBAD database, or as unpublished objects listed in DwarfArchives.org\\footnote{{\\tt www.dwarfarchives.org}}, or rejected by visual inspection of the SuperCOSMOS $R$-band images, were removed. 25 targets remained after this last rejection process, and 24 out of 25 targets observed spectroscopically are found here to be ultracool dwarfs later than $\\sim$\\,M7, the one interloper being a reddened early-type object, as evidenced by the low stellar density in the surrounding 5\\arcmin$\\times$5\\arcmin~ SuperCOSMOS $I$-band field and despite the high galactic latitude $b$\\,=\\,$-$36$^{\\circ}$. Optical and infrared photometry are given in Table \\ref{t1} together with proper motions and other basic data. Derived spectral types, spectroscopic distances and transverse velocities are given in Table \\ref{t2}. We give details of the methods used to determine these in Sect. 4.  \\section[]{Spectroscopic observations and data reduction}  Observations were carried out in three runs at the ESO New Technology Telescope (NTT) with the SOFI spectrograph, during October 25--29, 2005, January 17--19, 2006 and April 4--7, 2006, allowing observation of the complete sample of southern targets at all RAs. The 0.6\\arcsec~ slit was employed with the blue $JH$ grism yielding a wavelength coverage of 9300--16500\\,\\AA~ at a spectral resolution R\\,=\\,$\\lambda/\\Delta\\lambda$\\,$\\sim$\\,1000. The detector is a Hawaii HgCdTe 1024$\\times$1024 array with 18.5\\,\\micron~pixels. Accurate sky subtraction was facilitated by nodding the telescope along the slit, typically by 30--45\\arcsec. Detector integration times (DIT) were in the range 60\\,sec to 120\\,sec, with between 4 and 12 integrations at each nod position, yielding total integration times for science targets ranging between 16 and 48\\,min, depending on target $J$ magnitude. The resulting spectra have signal-to-noise ratios of at least $\\sim$\\,100 and often better. Conditions were generally good with optical seeing in the range 0.6\\arcsec -- 1\\arcsec; however during the January 2006 run in particular, conditions of high humidity saturated the telluric H$_2$O band at 1.35\\,\\micron~ and prevented its accurate removal. A, F or late B-type telluric standards were observed before and after each target observation at airmasses within 0.1 and often within 0.05 of the target. Dome flats and xenon lamp arcs were taken each afternoon, and again at the end of the night. Observations of template M6.5, M9, L1, L3 and L5 dwarfs were taken with the same instrumental setup; multiple observations of some of these standard stars were obtained during different observing runs.  \\begin{figure*} \\includegraphics[bb=0 0 543 700,width=17cm]{types.eps} \\caption{Spectral types by direct comparison to SpeX standards \\citep{cus05} and known objects observed with the same instrumental setup as the targets. Unweighted means of adjacent ($\\la$\\,1 subclass) standards from both datasets were permitted to obtain the best fits to the spectra for new objects.} \\label{f3} \\end{figure*}  \\begin{table*} \\centering \\begin{minipage}{110mm} \\caption{Derived data for the target objects. Comparison of spectral types from direct comparison (col. 5, see Fig.\\,\\ref{f3}) with spectral types derived from the spectral indices of \\citet{sle04} (cols. 2 \\& 4) and \\citet{mcl03} (col. 3; see Sect. 4.2). Where the H$_2$O-l index is affected by poor telluric correction, no value is given. The FeH index is invalid for spectral types later than $\\sim$\\,L3. Col. 6 gives the spectroscopic distance with errors reflecting the stated error on the spectral type, itself derived by comparison to template objects. Errors on transverse velocities (Col. 7) reflect only the error on the distance. } \\begin{tabular}{@{}lllllll@{}} \\hline Name & H$_2$O-l & H$_2$OB & FeH & spec. & d/pc & v$_{\\rm trans}$/km\\,s$^{-1}$ \\\\ \\hline 2M0004$-$20 & L0.5 & L0 & M9 & M8\\,$\\pm$\\,1 & 17.7$^{+3.8}_{-2.4}$ &       69.3\\,$\\pm$\\,12.1\\\\ 2M0005$-$21 & L0 & M9.5 & M9 & M8\\,$\\pm$\\,1 & 26.5$^{+5.7}_{-3.5}$ &    \t83.0\\,$\\pm$\\,14.4\\\\ 2M0123$-$42 & M8 & M9.5 & M9 & M8\\,$\\pm$\\,1 & 25.1$^{+5.4}_{-3.3}$ &       31.8\\,$\\pm$\\,5.4\\\\ 2M0127$-$31 & M9 & M9 & M7 & M7\\,$\\pm$\\,0.5 & 24.3$^{+3.3}_{-2.4}$ &    \t39.4\\,$\\pm$\\,4.5\\\\ 2M0128$-$55 & L2.5 & L3 & L1 & L1\\,$\\pm$\\,0.5 & 22.7$^{+1.5}_{-1.5}$ &   \t31.5\\,$\\pm$\\,2.1\\\\ 2M0228$-$63 & - & L1.5 & L2 & L0\\,$\\pm$\\,1 & 23.2$^{+3.0}_{-2.7}$ & \t70.0\\,$\\pm$\\,8.5\\\\ 2M0352+02   & - & L0.5 & M9.5& L0\\,$\\pm$\\,1 & 18.6$^{+2.4}_{-2.2}$ &   \t40.2\\,$\\pm$\\,5.0\\\\ 2M0357$-$12 & - & M9 & M7.5 & M8\\,$\\pm$\\,0.5 & 28.9$^{+2.7}_{-2.1}$ &  \t18.5\\,$\\pm$\\,1.5\\\\ 2M0405$-$15 & - & M8.5 & M8 & M8\\,$\\pm$\\,0.5 & 30.1$^{+2.8}_{-2.2}$ &  \t40.1\\,$\\pm$\\,3.3\\\\ 2M0511$-$46 & - & M8.5 & M8 & M7\\,$\\pm$\\,1 & 42.8$^{+14.0}_{-7.5}$  &  \t35.1\\,$\\pm$\\,8.8\\\\ 2M0526$-$44 & - & L2 & M9.5 & L1\\,$\\pm$\\,0.5 & 26.1$^{+1.7}_{-1.7}$ &  \t20.8\\,$\\pm$\\,1.4\\\\ 2M0855$-$25 & L0 & L1 & M9.5 & M9\\,$\\pm$\\,0.5 & 22.8$^{+1.6}_{-1.4}$ &   \t41.8\\,$\\pm$\\,2.8\\\\ 2M1019$-$27 & M9 & L1 & L0 & M9.5\\,$\\pm$\\,0.5 & 24.3$^{+1.6}_{-1.4}$ &\t74.2\\,$\\pm$\\,4.6\\\\ 2M1146$-$39 & - & L1.5 & M9& M9.5\\,$\\pm$\\,0.5 & 25.6$^{+1.6}_{-1.5}$ &\t35.4\\,$\\pm$\\,2.1\\\\ 2M1347$-$76 & - & L2 & L1  & L0\\,$\\pm$\\,1 & 25.9$^{+3.3}_{-3.0}$ &\t47.3\\,$\\pm$\\,5.7\\\\ 2M1520$-$44 & L3 & L4 & - & L4.5\\,$\\pm$\\,1 & 9.6$^{+2.2}_{-1.9}$ &    \t33.7\\,$\\pm$\\,7.0\\\\ 2M1523$-$23 & L0 & L2.5 & L1 & L2.5\\,$\\pm$\\,0.5 & 22.1$^{+1.8}_{-1.8}$ &\t36.3\\,$\\pm$\\,3.0\\\\ 2M1548$-$16 & M8.5 & L1.5 & L0.5 & L2\\,$\\pm$\\,1 & 20.8$^{+3.2}_{-3.2}$ &\t18.6\\,$\\pm$\\,2.9\\\\ 2M1750$-$00 & L8 & L6 & - & L5.5\\,$\\pm$\\,0.5 & 8.0$^{+0.9}_{-0.8}$ &    18.6\\,$\\pm$\\,1.9\\\\ 2M1842$-$39 & - & M6 & M9  & M9\\,$\\pm$\\,1 & 29.7$^{+4.6}_{-3.4}$ &       88.3\\,$\\pm$\\,11.9\\\\ 2M1956$-$17 & M8 & L0 & M9 & L0\\,$\\pm$\\,1 & 25.4$^{+3.3}_{-3.0}$ &\t21.9\\,$\\pm$\\,2.7\\\\ 2M2243$-$59 & L1 & L1.5 & L0 & L0\\,$\\pm$\\,1 & 29.4$^{+3.8}_{-3.4}$ &\t34.0\\,$\\pm$\\,4.2\\\\ 2M2255$-$57 & L6 & L7.5 & - & L5.5\\,$\\pm$\\,0.5 & 11.5$^{+1.3}_{-1.2}$ &  \t85.8\\,$\\pm$\\,9.0\\\\ 2M2321$-13$ & L1 & L2 & L3 & L1\\,$\\pm$\\,0.5 & 31.7$^{+2.1}_{-2.1}$   &\t86.5\\,$\\pm$\\,5.7\\\\ \\hline \\end{tabular} \\label{t2} \\end{minipage} \\end{table*}  Data reduction was undertaken within the {\\sc iraf}\\footnote{{\\sc iraf} is distributed by the National Optical Astronomy Observatories, which is operated by the Association of Universities for Research in Astronomy, Inc. (AURA) under cooperative agreement with the National Science Foundation.} environment, using standard techniques for the reduction of infrared spectroscopic data. As the data covers the region between the $J$ and $H$-bands, the master flatfields contain regions where structure in the dispersion direction was present (owing to the water vapour column between the dome and detector). These were interpolated using the {\\it fixpix} routine. This means that in these regions there is no correction for pixel-to-pixel variation, but only regions of strong telluric absorption are affected. A normalised flatfield was obtained by dividing the master flat by the same image smoothed by an 11$\\times$11 pixel boxcar. All science and telluric standard frames were divided by the normalised flat.  After flatfielding, science frames taken in each nod position were subtracted pairwise and spectra extracted using the {\\it apall} routine. Optimal and non-optimal extractions were identical to within typically $\\la$\\,a few\\,\\%. Arc spectra were extracted using a profile of a relatively bright target star and calibrated with the {\\it identify} and {\\it reidentify} routines. The fit to the dispersion has typical RMS\\,$\\sim$\\,0.1\\,\\AA. Individual spectra were only summed (using the {\\it sarith} routine with no pixel rejection) after wavelength calibration. Telluric correction was again performed using {\\it sarith}; no attempts were made to scale or shift the telluric spectrum to improve the correction. Hydrogen lines at 0.954, 1.005, 1.092 and 1.282\\,\\micron~ were removed from the telluric spectra within the {\\it splot} routine, before division. The telluric standard giving the best correction was normally the one closest in airmass and time but different standards from the same night were employed if they gave a better, less noisy, correction. In about 30\\% of the cases, part of the 1.35\\,\\micron~ water band was interpolated over; in these cases there was too little flux in the region in either target or standard to provide a useful correction. Finally the spectra were multiplied by a blackbody spectrum (F$_{\\rm\\lambda}$) with the T$_{\\rm eff}$ of the chosen telluric standard star, and normalised to unity at the highest point of the final spectrum. ",
        "conclusions": "Using a combination of optical and near-infrared survey data from SuperCOSMOS and 2MASS, twenty-one hitherto undescribed sources are shown to be ultracool dwarfs with spectral types in the range M7 to L5.5 using low resolution spectroscopy from NTT/SOFI. The sample objects have $JHK$ magnitudes and colours comparable to known late M and L dwarfs and are bright enough to have optical detections in digitised plate data. They have proper motions typically a few $\\times$10$^2$ mas\\,yr$^{-1}$. Spectroscopic distances are derived, typically within 30\\,pc, showing these objects to be important additions to our knowledge of low-mass, low-luminosity objects in the Solar neighbourhood.  Two objects in particular, 2M1750$-$00 and 2M2255$-$57, are of particular interest, with spectral types L5-L6 and spectroscopic distances near 10\\,pc. The latter has a high proper motion, 1.5\\arcsec\\,yr$^{-1}$. These two objects are likely brown dwarfs. In common with the rest of the sample, they are of great interest for further study using AO, HST, or methane imaging, to uncover lower-mass, later type companions. There is circumstantial evidence for the existence of such companions from the spectroscopic data presented here. Furthermore, these two objects in particular merit a parallax determination.  We have found a further object, 2M1520$-$44, to be a resolved L dwarf binary consisting of L2 and L4 components, lying $\\sim$\\,19\\,pc distant. Further study should reveal this source as a common proper motion pair. As such, it may prove a useful test-bed for evolutionary models of very low-mass stars and brown dwarfs."
    },
    "0609/astro-ph0609675_arXiv.txt": {
        "abstract": "% Recent results with the Potsdam Wolf-Rayet (PoWR) models have shown that Wolf-Rayet mass loss can be explained by radiative wind driving.  An inspection of the galactic WR sample, however, reveals that a significant part of the observed WR stars are in conflict with our models. This group is chiefly formed by intermediate spectral subtypes. Among the population of late-type WN stars we find that the enigmatic WN\\,8 subtypes, which are well-known for their photometric variations, are in disagreement with our models. So are there other driving mechanisms working in these objects? ",
        "introduction": "The Potsdam Wolf-Rayet (PoWR) hydrodynamic model atmospheres combine fully line-blanketed non-LTE models with the equations of hydrodynamics \\citep[see][]{gra1:05,ham1:03,koe1:02,gra1:02}.  In these models the wind structure ($\\rho(r)$, $v(r)$, and $T(r)$) is computed in line with the full set of non-LTE populations and the radiation field in the co-moving frame (CMF). In this numerically rather expensive approach no simplifying assumptions for the computation of the radiative acceleration $a_{\\rm rad}$ are made. In contrast to all previous wind models $a_{\\rm rad}$ is obtained by direct integration instead of making use of the Sobolev approximation \\begin{equation} \\label{eq:arad} a_{\\rm rad} = \\frac{1}{c} \\int \\chi_\\nu F_\\nu {\\rm d}\\nu. \\end{equation} In this way, complex processes like strong line overlap or the redistribution of radiation are automatically taken into account.  The models thus describe the conditions in WR\\,atmospheres in a realistic manner, in particular also at large optical depth.  Moreover they provide synthetic spectra, i.e.\\ they allow for a direct comparison with observations.  Utilizing these models, we obtained the first fully self-consistent Wolf-Rayet wind model \\citep{gra1:05}, for the case of an early-type WC star.  Moreover we have recently examined the mass loss from late-type WN stars and its dependence on metallicity \\citep{gra2:06,gra1:06}.  To illustrate the complexity of the involved processes we show in Fig.\\,\\ref{fig:3d} the distribution of the radiative flux, the opacity, and their product (i.e., the radiative acceleration per frequency interval) within the WC model. It can be clearly seen that the initial FUV flux, originating from the bottom of the wind, is blocked by successively recombining continua.  This radiation is re-emitted at lower frequencies in the form of emission lines, forming the typical WR emission line spectrum.  The radiative acceleration predominantly arises from Fe-group line opacities absorbing this radiation. These lines partly form the observable iron forest short-ward of $1500\\,$\\AA.  They show a pronounced ionization structure and contribute by themselves to the complex 'background' flux.  This coexistence of line emission and absorption, in combination with strong overlap, is the most probable reason for substantial deviations from the standard theory of line-driven winds \\cite[][see also Bjorkman this volume]{cas1:75}.  In our WR models we tend to detect extremely small or even negative force multiplier parameters \\citep[see][]{gra1:05}, which can be explained by a 'statistical' line locking effect due to the strong overlap.  In addition to the complex radiative processes we find a strong influence of wind clumping on the dynamics of O and WR star winds \\citep{gra2:03,gra1:05}. Wind clumping particularly affects the recombination rates because these scale with $n_{\\rm e}^2$. A down-scaling of $\\dot M \\propto 1/\\sqrt D$ (where $D$ denotes the clumping factor) thus preserves the emission measure in the atmosphere.  Because WR atmospheres are dominated by recombination processes, the ionization structure as well as the emission line spectrum are, approximately, preserved under such a transformation \\citep{ham1:98}. The same holds for the radiative {\\em force}. The resulting radiative {\\em acceleration} thus scales with $\\sqrt D$ under this transformation. ",
        "conclusions": "Our hydrodynamic atmosphere calculations for WR\\,stars have shown that WR-type stellar winds can be driven alone by radiative acceleration.  In particular we found that the proximity to the Eddington limit is the primary trigger of the enhanced mass loss of WR\\,stars. Nevertheless, difficulties remain with H-free objects of intermediate spectral subtype, and late-type WN stars with strong stellar winds (mostly of spectral type WN\\,8). In both cases, problems are caused by the gap between the 'hot' and the 'cool' Fe opacity-peaks. For the intermediate subtypes this gap causes a deficiency of line opacities at the wind base. This problem could be overcome by an additional support due to fast stellar rotation. For the WN\\,8 stars, the inward increase of the radiative force due to the hot Fe opacity-peak causes problems in the hydrostatic layers, presumably leading to the observed variability. These oscillations most probably affect the mass loss of these stars.  \\newpage"
    },
    "0609/astro-ph0609169_arXiv.txt": {
        "abstract": "We study the evolution of two fundamental properties of galaxy clusters: the luminosity function (LF) and the scaling relations between the total galaxy number $N$ (or luminosity) and cluster mass $M$. Using a sample of 27 clusters ($0\\le z \\le 0.9$) with new near-IR observations and mass estimates derived from \\xray temperatures, in conjunction with data from the literature, we construct the largest sample for such studies to date.  The evolution of the characteristic luminosity of the LF can be described by a passively evolving population formed in a single burst at $z=1.5-2$. Under the assumption that the mass-temperature relation evolves self-similarly, and after the passive evolution is accounted for, the $N$--$M$ scaling shows no signs of evolution out to $z=0.9$. Our data provide direct constraints on halo occupation distribution models, and suggest that the way galaxies populate cluster-scale dark matter halos has not changed in the past 7 Gyr, in line with previous investigations. ",
        "introduction": "\\label{sec:intro}    The subject of cluster galaxy evolution has always been central to cluster studies.  For the early types that comprise the majority of the cluster galaxies, studies based on the fundamental plane and/or color-magnitude relations suggest that their stars are primarily old, and evolve passively with time \\citep[\\eg][]{bower92,stanford98,kelson00}.  Given the hierarchical nature of structure formation, changes in the galaxy population over cosmic time are expected.  Numerous investigations have reported such evidence, \\eg a higher fraction of blue galaxies in more distant clusters \\citep[\\eg][]{butcher78,ellingson01}, a decline in the S0 galaxy abundances towards higher-$z$ \\citep{dressler97}, and a population that appears to be in the post-starburst phase \\citep{poggianti99}.     A powerful method for quantifying galaxy populations is via the luminosity function (LF). The characteristic luminosity of the cluster galaxies as a whole is reflected by the ``knee'' of the LF ($M_*$); comparison of LFs at different redshifts thus will indicate the evolution of the mean galaxy population and constrain the formation epoch of cluster galaxies \\citep[e.g.][]{barger98,depropris99,strazzullo06}.   Here we present an analysis of the evolution of the cluster $K_s$-band LF from $z=0$ to $z=0.9$, which has several novelties and advantages over previous studies. Foremost, all our clusters have mass estimates accurate to $\\lesssim 30\\%$ (from the \\xray temperature $T_X$). We construct the LF within the same portion of the cluster virial radius, which facilitates comparisons of clusters at different redshifts and of different masses.  We also make use of the cluster mass information and construct the LF in terms of space density. Finally, our data generally are deeper, and cover a larger fraction of the cluster virial region.  These aspects allow us to integrate the LF to estimate the total galaxy luminosity ($L$) and number ($N$), and study their correlation with the cluster mass ($M$).  Based on a large sample of nearby clusters, we have found tight correlations between $L$ or $N$ with $M$ (\\citealt{lin03b,lin04}, hereafter \\citetalias{lin03b}, \\citetalias{lin04}, respectively), which already provide an interesting constraint on the regularity of the cluster galaxy formation process.   In this study we pay special attention to the evolution of the $N$--$M$ relation, which directly constrains simple clustering models such as the halo occupation distribution (HOD; e.g.~\\citealt{berlind02}) out to $z\\sim 1$. In such models, the way galaxies populate dark matter halos (e.g.~the number of cluster galaxies above some luminosity limit as a function of halo mass, known as the halo occupation number) are tuned until the observed galaxy clustering at both small and large scales is reproduced.  Physics of galaxy formation is studied by identifying mechanisms that reproduce the halo occupation statistics (e.g.~\\citealt{zheng05}). Furthermore, the evolution of the $N$--$M$ relation may reveal the mass function evolution of subhalos in cluster-size halos \\citep{cooray05a}.    In \\S\\ref{sec:data} we describe the sample selection, observations, data reduction, and our methods for the star-galaxy separation. We stack the cluster data to construct the composite LFs at different redshifts (\\S\\ref{sec:app}), which enable us to constrain the evolution of $M_*$ and the formation epoch of the cluster galaxies.  We then proceed to solve for the LF parameters of individual clusters, and examine the evolution of the $L$--$M$ and $N$--$M$ relations (\\S\\ref{sec:sc}).  We discuss our main findings in \\S\\ref{sec:sum}.  We assume the cosmological parameters to be (\\omegam, \\omegal, $H_0$) = (0.3, 0.7, $70\\,h_{70}$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$). Throughout the paper we denote $K_s$-band by $K$-band for brevity, and adopt the Vega magnitudes unless noted % \\citep[$K_{{\\rm AB}}-K_{{\\rm Vega}}=1.85$,][]{blanton06}. ",
        "conclusions": "\\label{sec:sum}   We can understand the no-evolution of the $N$--$M$ relation by looking at the integration of the LF: \\beq N(M,z) = V \\phi_* \\Gamma(\\alpha+1, L_{l}(z)/L_*(z)), \\eeq where $V$ is the cluster volume, $L_{l}$ is the luminosity corresponding to the lower limit of the integration, and $\\Gamma$ is the incomplete gamma function. Because $L_l$ is set to be a fixed fraction of $L_*$, it is the product of $V$ and $\\phi_*$ that determines $N$. For two clusters $i$ and $j$ which have the same mass but at different redshifts (say $z_j>z_i$), $V_j<V_i$ (assuming that $r_\\Delta \\propto 1/H(z)$, where $H(z)$ is the expansion rate of the Universe). One thus expects that on average, $\\phi_{*j}>\\phi_{*i}$. Indeed, $\\phi_*$ derived for our clusters agrees with this expectation. {\\it That the increase in the galaxy number density is offset by the decrease in the virial volume is the main reason for the lack of evolution of the $N$--$M$ relation.}    Previous studies have found suggestions of lack of HOD evolution toward $z\\sim 1$, based on clustering statistics \\citep{yan03,phleps06,coil06}. In a recent attempt to extract both cosmological information and cluster physics from a large sample of clusters, it is found that the optical observable--mass relation, although only weakly constrained, does not evolve with redshift \\citep{gladders06}. We note, however, passive evolution has not been accounted for in this analysis (S.~Majumdar 2006, private communication). Our study is complementary in the sense that we probe the (high mass end of) occupation number directly, with reliable cluster mass estimates from $T_X$.    It is possible that changes in both the red and blue galaxy populations with redshift are hidden from the evolution of the $N$--$M$ relation, if they conspire to make the total galaxy number unchanged. We are in the process of acquiring color information in order to constrain the evolution of these populations separately.   The fact that only when the large, nearby cluster sample of \\citetalias{lin04} is included can we obtain reasonably tight constraint on the $N$--$M$ evolution points to the need for a {\\it much} larger sample with reliable mass estimates at higher-$z$. This is also demanded for large optical/NIR cluster surveys which may rely on optical/NIR based cluster mass estimators (\\eg FLAMEX). Based on simple Monte Carlo simulations, we estimate that at least 20 clusters at $z=0.8-1.0$ are required to constrain $\\gamma$ to 0.1 level. While more clusters at intermediate redshifts may provide better constraints on the evolution of the slope of the scaling relation, $z>1$ clusters clearly will help pinpoint the formation epoch of the cluster galaxies."
    },
    "0609/astro-ph0609443_arXiv.txt": {
        "abstract": "Primordial stars are likely to be very massive $\\geq$30$\\Msun$, form in isolation, and will likely leave black holes as remnants in the centers of their host dark matter halos. We expect primordial stars to form in halos in the mass range $10^{6}-10^{10}\\Ms$. Some of these early black holes, formed at redshifts z$\\gtsim$10, could be the seed black hole for a significant fraction of the supermassive black holes found in galaxies in the local universe. If the black hole descendants of the primordial stars exist, their mergers with nearby supermassive black holes may be a prime candidate for long wavelength gravitational wave detectors. We simulate formation and evolution of dark matter halos in $\\Lambda$CDM universe. We seed high-redshift dark matter halos with early black holes, and explore the merger history of the host halos and the implications of black hole's kick velocities arising from their coalescence. The central concentration of low mass early black holes in present day galaxies is reduced if they experience even moderate kicks of tens of ${\\rm km \\,s^{-1}}$. Even such modest kicks allow the black holes to leave their parent halo, which consequently leads to dynamical friction being less effective on the low mass black holes that were ejected, compared to those still embedded in their parent halos. Therefore, merger rates with central supermassive black holes in the largest halos may be reduced by more than an order of magnitude.  Using analytical and illustrative cosmological N--body simulations, we quantify the role of kicks on the merger rates of black holes formed from massive metal free stars with supermassive black holes in present day galaxies. ",
        "introduction": "It is firmly established that most galaxies have super-massive black holes (SMBH) at their centers. Their observed masses are in the range 10$^6 \\, \\Msun \\ltsim$ M $\\ltsim 10^9 \\, \\Msun$ and it appears that there are correlations (Kormendy $\\&$ Gebhardt 2001, Gerhard 2001) between their masses and the bulk properties of the galactic bulges hosting them (Laor 2001, Merritt $\\&$ Ferrarese 2001, Tremaine et al. 2002, Gebhardt et al. 2001, Graham et al. 2001). These correlations point to a link between the formation of SMBHs and the evolution of their hosts. It also appears that SMBHs are linked to the properties of the host dark matter halos. If the SMBH precursors were present from very early on, then their mergers, together with growth by accretion, could account for the abundance of the SMBHs today (Schneider et al. 2002).   Ab initio numerical simulations of the formation of the first luminous objects within the standard structure formation framework, finds that metal free stars form in isolation, and may have masses $30\\,\\Msun \\ltsim$ m $\\ltsim 300 \\,\\Msun$ (Abel et al. 2000, Abel et al. 2002, Bromm et al. 2002). In current models of structure formation, dark matter initially dominates and pregalactic objects form from small initial density perturbations. As they assemble via hierarchical merging, the metal-free primordial gas cools slowly through the rotational lines of hydrogen molecules. As the gas cools, it sinks to the center of the dark matter potential well. However, as the cooling rates are small, the clumps fail to fragment as they collapse. This leads to a top-heavy initial stellar mass function (IMF) and to the production of very massive stars, unlike the modern stellar IMF which declines rapidly with increasing mass.   Metal-free, high mass Population III stars lose only a small fraction of their mass in their lifetime of $\\sim$3 Myr. This suggests that a large population of primordial massive black holes (MBH, with mass $\\ltsim 10^3 \\Msun$) could be an end product of such pregalactic star formation (Heger et al. 2003). Since they form in rare high-$\\sigma$ density peaks (Couchman $\\&$ Rees 1986, Madau $\\&$ Rees 2001), relic MBHs are predicted to cluster in the cores of more massive halos (Abel et al. 2002) which in turn are formed by subsequent mergers of the original parent dark matter halos. It has been suggested (Volonteri et al. 2003, Islam et al. 2003) that mergers of massive halos and clustering of these MBHs should start at redshifts as large as z$\\sim$20. Subsequent growth of MBHs would proceed through accretion of baryonic gas and coalescence, possibly leading to the formation of intermediate mass black holes (IMBH, with mass in range $10^3\\,\\Msun \\ltsim$ m $\\ltsim 10^6 \\,\\Msun$). In the most optimistic scenario (Haiman $\\&$ Loeb 2001), 10$^9\\Msun$ black holes could form as early as z$\\sim$10.  In this paper we will explore the evolution of the seed black hole in dark matter halos. Mergers of dark matter halos hosting at least one black hole (Bromm $\\&$ Loeb 2003) at high redshift form black hole binaries, and their subsequent coalescence under gravitational radiation may give them a significant kick velocity (Favata et al. 2004, Merritt et al. 2004). These recoils are caused by anisotropic emission of gravitational waves which carry away linear momentum. The recoil velocity is highly controversial requiring full general relativistic calculations. However, recent estimates place most recoils in the range 10-100 ${\\rm km \\,s^{-1}}$, although kicks of a few hundred ${\\rm km \\,s^{-1}}$ are not unexpected, and the largest recoils should not be above 500 ${\\rm km \\,s^{-1}}$ (Favata et al. 2004, Blanchet et al. 2005, see also Baker et al 2006). The amplitude of the kick determines if the black hole will be ejected from its host dark matter halo. Even the most massive dark matter halo at redshift z$\\geq$11 can not retain a black hole that receives $\\geq$150${\\rm km \\,s^{-1}}$ kick, while at redshift of z$\\geq$8, kicks of 300${\\rm km \\,s^{-1}}$ are required to cause ejection (Merritt et al. 2004, Fig 3.). This is valid for black holes with masses $10^3\\,\\Msun \\ltsim$ m $\\ltsim 10^8 \\,\\Msun$. Kicks of 40${\\rm km \\,s^{-1}}$ should be sufficient to: eject IMBHs from globular clusters; displace IMBHs from the centers of dwarf galaxies; and perhaps provide a sufficient population of IMBHs for merging at the centers of spiral galaxies (Merritt et al. 2004).   We use N-body TREEcode simulations (GADGET, Springel et al. 2001) to investigate the {\\it observational implications} of black hole kicks through the asymmetric emission of gravitational radiation during the merger. In this academic exercise, we assume that Population III stars collapse to form massive black holes at high redshifts. These black holes grow through gas accretion and mergers at the centers of dark matter halos. We use collisionless particles at the centers of dark matter halos as tracers of IMBHs. If kicks are important, the growth of IMBHs into supermassive black hole will be suppressed. We test two extreme scenarios for this behaviour and whether there are predictable consequences. In the first scenario, kicks do not occur in any dark matter halo at any redshift. In the second, kicks occur in every dark matter halo at an initial redshift and we introduce a prescription for assigning kicks to our tracer particles. As a result, these two scenarios set limits to the IMBHs' merger rate. The first scenario provides a maximum merger rate, and the second scenario provides a lower bound on the merger rate (the absolute lower bound is that no mergers occur, for example if black hole formation is suppressed and primordial stars leave no compact remnants). A better approach would be to track every individual merger, and assign kicks when mergers occur. Our approach is computationally less expensive and provides only constraints. We have chosen z=8 as an initial redshift for assigning kicks to our tracers for IMBHs. At redshift z$\\sim$12 supernovae rates for Pop III stars fall to zero (Wise $\\&$ Abel 2005). The massive black holes resulting from Pop III collapse have already settled in the centers of dark matter halos by redshift z$\\sim$8, the initial redshift for kicks. As mentioned before we are calculating constraints for the merger rates and not studying individual mergers. The MBH tracers are tracked from the initial to the final redshift and used to quantify the changes in distributions of IMBHs in these two scenarios. Through detailed analysis of their trajectories, we study how IMBHs respond to the dynamical evolution of dark matter halos, testing extreme ranges of outcome to check for the plausible scope for detection. Our results also apply to black holes ejected by the gravitational slingshot mechanism. The slingshot velocities fall in the recoil range (Aarseth 2005).   Our numerical simulations are followed by an analytical treatment of dynamical friction. As long as the dark matter halos have a bound core with density greater than the local galactic density and have \"many\" particles in them, the numerical dynamical friction will mimic \"real\" dynamical friction with reasonable fidelity (Vine $\\&$ Sigurdsson 1998). If the halo density becomes smaller than the local galactic density, or if the secondary halo structure is down to being resolved by only \"few\" bound particles, then the dynamical friction prescription breaks down and relaxation due to numerical noise and large scale potential fluctuations dominate; integration of the dynamical friction term becomes meaningless. This happens basically where we stop the simulations, and that is where we analytically estimate a conservative lifetime to coalescence.   We are  exploring the effects of kicks on IMBH merger rates and the formation of SMBHs that are detected in the nuclei of luminous galaxies. Knowledge of IMBH merger rates is essential for the future experiments such as LISA (Danzmann 2003) and BBO (Big Bang Observer) (Cornish $\\&$ Crowder 2005). LISA will be able to study much of the last year of inspiral of IMBH into SMBHs, as well as the waves from the final collision and coalescence of IMBH binaries. For binary black holes with masses in the range 100$\\Msun\\ltsim$ M $\\ltsim10^4\\Msun$, LISA can observe the last few years of inspiral, but not the final collision even at cosmological distances. Equal-mass black hole binaries enter LISA's frequency band roughly 1000 years before their final coalescence, independent of their masses, for the range 100$\\Msun\\ltsim$ M $\\ltsim10^6\\Msun$ (Thorne 1995, Cornish $\\&$ Levin 2002, Vecchio et al. 2004). Third generation gravitational wave observatories, such as BBO mission will have arcsecond-arcminute precision in identifying and localizing every merging black hole binary independent of redshift, anywhere in the universe over the years to months before their actual merger. In combination with ground-based detectors at higher frequencies, BBO could measure the mass, angular momentum, and dynamic spacetime structure of these compact objects with unprecedented precision.   In \\S 2 we present a straightforward analytical treatment of dynamical friction. In \\S 3 we describe the simulation setup, the codes used for analysis, and present results for 3 different high resolution simulations with varying kick parameters for early black holes. In \\S 4 we describe the post-merger evolution through dynamical friction based on the arguments in \\S 2, before we conclude in \\S 5. ",
        "conclusions": "We have performed and analyzed high-resolution collisionless simulations of the evolution of structure in a $\\Lambda$CDM model. We have followed the formation and evolution of dark matter halos, and by assuming that each halo hosts an IMBH, we have studied the formation of SMBHs. We focussed on two specific cases. First, IMBHs together with their host halos merge through dynamical friction. Then, we explored the effect of kicks on the IMBH distribution. We found that kicks lead to the ejection of IMBHs from their host halos in many cases. Analytically, it is clear that the dynamical friction will act more efficiently on the host halos than on the less massive black holes formed within them. Our illustrative calculations highlight some of the expected differences in the final density distribution of the black holes. They can be quite different, even in the presence of the modest kick velocities we have imposed. We note that, in general, IMBHs ejected from their parent sub-halos at high redshift still get bound to the most massive dark matter halo at later times but they form a ''hot population'' with a flatter density profile than the underlying dark matter.  In order for dynamical friction to work accurately in numerical simulations, the density of the background (in this case, the primary halo density) needs to be resolved well. Higher resolution gives smaller softening lengths for particles and smaller softening lengths give more realistic dynamical friction. With the resolution of our current simulations, dynamical friction can not be efficiently realized in the model after z$\\sim$1 and r$\\ltsim$r$_{vir}/30$, because the background density of the parent halos is not resolved well enough. However, our simulations are able to account for effects of dynamical friction on the sub-halos at high z (Vine $\\&$ Sigurdsson 1998). Subsequent papers will track the dynamical evolution of the sub-halos and IMBHs at late times, using higher resolution simulations and semi-analytic implementations of dynamical friction (Binney $\\&$ Tremaine 1987). A little over 4$\\%$ of IMBHs merge with the SMBH at the center in less than one Hubble time. If this is the dominant way of creating SMBHs, then it is efficient even if the IMBHs are ejected from their parent halos, with merger numbers only reduced by a factor of two or so for modest kicks. LISA observations should strongly constrain any kick on IMBHs formed from very massive Pop III stars in low mass proto-galactic sub-halos.  If these IMBHs accrete from the interstellar medium they may also be found by deep X--ray observations (Agol $\\&$ Kamionkowski 2002) in the Milky Way. The observational signature will depend sensitively on the spatial distribution of these black holes in our galaxy. Similarly, the next generation of micro-lensing searches may be able to constrain high mass black holes from the longest events (Agol et al. 2002)."
    },
    "0609/astro-ph0609733_arXiv.txt": {
        "abstract": "{ We investigate the radio spectral energy distributions (SEDs) of young star-forming galaxies and how they evolve with time. The duration and luminosity of the nonthermal radio emission from supernova remnants (SNRs) are constrained by using the observational radio SEDs of \\sbs\\ and \\izw, which are the two lowest-metallicity blue compact dwarf galaxies in the nearby universe. The typical radio ``fluence'' for SNRs in \\sbs, that is the radio energy emitted per SNR over its radiative lifetime, is estimated to be $\\sim 6$--$22\\times 10^{22}~{\\rm W~Hz^{-1}~yr}$ at 5 GHz. On the other hand, the radio fluence in \\izw\\ is $\\sim 1$--$3\\times 10^{22}~{\\rm W~Hz^{-1}~yr}$ at 5 GHz. We discuss the origin of this variation and propose scaling relations between synchrotron luminosity and gas density. We have also predicted the time dependence of the radio spectral index and of the spectrum itself, for both the ``active'' (\\sbs) and ``passive'' (\\izw) cases. These models enable us to roughly age date and classify radio spectra of star-forming galaxies into active/passive classes. Implications for high-$z$ galaxy evolution are also discussed. ",
        "introduction": "\\label{sec:intro}  Radio emission from galaxies is known to be connected with star formation activity. Two radiative processes are responsible: thermal free-free radiation from ionized gas in \\hii\\ regions and synchrotron radiation from relativistic electrons spiraling in magnetic fields (Condon \\cite{condon92}). The former originates from ionized gas around massive stars, while the latter comes from supernova remnants (SNRs) whose progenitors are massive stars. Since massive stars have short lifetimes, both thermal and nonthermal emission should trace the current star formation rate (SFR).  Nevertheless, the physical basis of the radio emission is not fully established. The observed radio flux in luminous evolved starbursts and normal spiral galaxies is dominated by nonthermal emission, but not from the short timescale radiation of discrete SNRs; more than $90$\\% of it is probably due to the diffusion of cosmic ray electrons in the galaxy disk (Condon \\cite{condon92}) over timescales of $10^7$--$10^8$\\,yr (Helou \\& Bicay \\cite{helou93}). Indeed, if the theoretically estimated adiabatic timescale (e.g., Woltjer \\cite{woltjer72}) and the observationally obtained $\\Sigma$--$D$ (surface brightness vs.\\ diameter) relation for discrete SNRs (e.g., Clark \\& Caswell \\cite{clark76}) are used to derive the nonthermal emission in galaxies (e.g., Biermann \\cite{biermann76}; Ulvestad \\cite{ulvestad82}), the emission is underpredicted by a factor of 10 or so.  In very young galaxies with ages $\\la 10$\\,Myr, however, the cosmic-ray diffusion mechanisms have not yet had a chance to dominate the radio emission. Nonthermal synchrotron radiation from discrete SNRs would be expected to dominate over the diffuse component. Moreover, in galaxies with starbursts younger than $\\sim 3$\\,Myr, there should be very little synchrotron emission since the SNRs have not yet exploded. Indeed, some galaxies with a deficit of nonthermal emission are observed (Roussel et al.\\ \\cite{roussel03}) and appear to be due to starbursts observed within a few Myrs of their onset, after a long period ($\\ga$ 100\\,Myr) of quiescence.  In this paper, we investigate theoretically the time evolution of the radio spectral energy distribution (SED) in young star-forming galaxies. In principle, this is a straightforward process because we know that the radio SED can be determined by the star formation history, for which we have already developed the formalism (Hirashita et al.\\ \\cite{hhf02}; Hirashita \\& Hunt \\cite{hirashita04}). In practice, it is a difficult exercise because of the unknown nature of the time evolution of the nonthermal radio component. Hence, we adopt nearby blue compact dwarf galaxies (BCDs) as an observational sample to constrain our theoretical model of radio emission. Most BCDs have a young age of {\\it current star formation} (although most of them had previous starburst episodes), and a low metallicity (e.g., Izotov \\& Thuan \\cite{izotov04}). This implies that they are relatively unevolved chemically and may provide a reasonable template to investigate star formation properties of young galaxies. The BCDs as a class have radio spectra with different properties than those of evolved luminous star-forming galaxies (Klein et al.\\ \\cite{klein91}); flatter radio spectra of BCDs indicate that they have a lower fraction of nonthermal emission than more luminous systems. Hence, they may lack the diffuse emission that characterizes the radio spectrum of larger disk galaxies, and thus are ideal targets for constraining our models.  Two classes of star-formation activity in BCDs have recently emerged observationally, as proposed by Hunt et al.\\ (\\cite{hunt-cozumel}) and Hunt \\& Hirashita (2006, in preparation). They argue that the star-formation modes in the two most metal-poor galaxies, \\object{SBS 0335$-$052} and \\object{I Zw 18}, are very different, in spite of their similar metallicities ($12+\\log ({\\rm O/H})=7.3$ and $12+\\log ({\\rm O/H})=7.2$, respectively; Skillman \\& Kennicutt \\cite{skillman93}; Izotov et al.\\ \\cite{izotovetal99}). The major star-forming region of \\sbs\\ is compact and dense (radius $r_{\\rm SF}\\la 40$ pc, number density $n\\ga 600$ cm$^{-3}$; Dale et al.\\ \\cite{dale01}; Izotov \\& Thuan \\cite{izotov99}). Moreover, \\sbs\\ hosts several super star clusters (SSCs), detectable H$_2$ emission lines in the near-infrared (NIR) (Vanzi et al.\\ \\cite{vanzi00}), a large dust extinction ($A_V\\sim 16$ mag; Thuan et al.\\ \\cite{thuan99}; Hunt et al.\\ \\cite{hunt01}; Plante \\& Sauvage \\cite{plante02}), and high dust temperature (Hunt et al.\\ \\cite{hunt01}; Dale et al.\\ \\cite{dale01}; Takeuchi et al.\\ \\cite{takeuchi03}). On the contrary, the star-forming regions in \\izw\\ are diffuse ($r_{\\rm SF}\\ga 100$ pc, $n\\la 100$ cm$^{-3}$), and contain no SSCs. NIR H$_2$ emission has not been detected (Hunt et al., private communication), and the dust extinction is moderate ($A_V\\sim 0.2$ mag; Cannon et al.\\ \\cite{cannon02}). We consider a region with such properties as ``passive'' following Hunt et al.\\ (\\cite{hunt-cozumel}), while a \\sbs-like star-forming property is considered ``active''. The similar metallicities of \\sbs\\ (active) and \\izw\\ (passive) imply that the chemical abundance is not a primary factor in determining the star-forming properties. We argue that the compactness of star-forming regions, which affect gas density, gas dynamics, and so on, is important in the dichotomy of active and passive modes.  In this paper, we extend this active-passive classification into the radio regime by focusing on \\sbs\\ and \\izw\\ as representatives of metal-poor BCDs, and also of two star-forming modes in BCDs: ``active'' and ``passive''. This paper is organized as follows. First, in Sect.\\ \\ref{sec:model} we explain the model that describes the evolution of radio SEDs. Some basic results are given in Sect.\\ \\ref{sec:constraint}, where observational data of \\sbs\\ and \\izw\\ are compared with the results for active and passive modes. The constraints obtained here are compared with theoretical descriptions of the physical processes of synchrotron radiation in Sect.\\ \\ref{sec:mag}. In Sect.\\ \\ref{sec:time} we generalize the results of our models with predictions for the time evolution of radio SEDs distinguishing between active and passive modes, and discuss implications for high-redshift star formation. Finally we give our conclusions in Sect.\\ \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  In order to investigate the time evolution of the radio SEDs of metal-poor star-forming galaxies, we have constructed a radio SED model by treating the thermal and nonthermal components consistently with the star formation history. In particular, the duration and luminosity of the nonthermal radio emission from supernova remnants (SNRs) has been constrained by using the observational radio SEDs of \\sbs\\ and \\izw, both of which are reasonable proxies for young galaxies in the nearby universe. In \\sbs, the typical radio energy emitted per SNR over its lifetime is estimated to be ${\\cal E}(0.1-100~{\\rm GHz})\\sim 8\\mbox{--}32 \\times 10^{47}~{\\rm erg}$. In \\izw, another representative of metal-poor star-forming galaxies, we find that ${\\cal E}(0.1-100~{\\rm GHz})\\sim 2\\mbox{--}6\\times 10^{47}$ erg. Both estimates are significantly larger than previous estimates based on the $\\Sigma-D$ relation. These values of the two ``template'' galaxies can be simultaneously explained by a simple density scaling relation of $B\\propto n^{0.35-0.73}$, indicating that the magnetic pressure scales with the gas pressure, or that the magnetic fields are amplified by SNe, or both.  We have also predicted the time dependence of the radio spectral index and of the radio SED itself, for both the active (\\sbs) and passive (\\izw) cases. These models enable us to roughly age date and classify radio spectra of star-forming galaxies into active/passive classes. Since the radio emission around $\\nu\\la 1$ GHz can be affected by the free-free absorption especially in the active class, the spectral index defined at $\\nu\\ga 5$ GHz should be used to estimate the contribution from the nonthermal component. On the other hand, the free-free absorption feature around $\\nu\\sim 1$ GHz could be used to select the active class."
    },
    "0609/astro-ph0609035.txt": {
        "abstract": "We develop a new theoretical model for the spectral formation process in accretion-powered X-ray pulsars based on a detailed treatment of the bulk and thermal Comptonization occurring in the accreting, shocked gas. A rigorous eigenfunction expansion method is employed to obtain the analytical solution for the Green's function describing the scattering of radiation injected into the column from a monochromatic source located at an arbitrary height above the stellar surface. The emergent spectrum is calculated by convolving the Green's function with source terms corresponding to bremsstrahlung, cyclotron, and blackbody emission. The energization of the photons in the shock, combined with cyclotron absorption, naturally produces an X-ray spectrum with a relatively flat continuum shape and a high-energy quasi-exponential cutoff. We demonstrate that the new theory successfully reproduces the phase-averaged spectra of the bright pulsars Her~X-1, LMC~X-4, and Cen~X-3. In these luminous sources, it is shown that the emergent spectra are dominated by Comptonized bremsstrahlung emission. ",
        "introduction": "More than 100 pulsating X-ray sources have been discovered in the Galaxy and the Magellanic Clouds since the seminal detection of Her X-1 and Cen X-3 more than three decades ago (Giacconi et al. 1971; Tananbaum et al. 1972). These sources display luminosities in the range $\\xlum \\sim 10^{34-38}{\\ \\rm ergs \\ s^{-1}}$ and pulsation periods $0.1 {\\ \\rm s} \\lapprox P \\lapprox 10^3 {\\ \\rm s}$, and comprise a variety of objects powered by rotation or accretion, as well as several anomalous X-ray pulsars whose fundamental energy generation mechanism is currently unclear. The brightest pulsars are accretion-powered sources located in binary systems, with emission fueled by mass transfer from the ``normal'' companion onto one or both of the neutron star's magnetic poles. The accretion flow is channeled by the strong magnetic field into a columnar geometry, and the resulting emission is powered by the conversion of gravitational potential energy into kinetic energy, which escapes from the column in the form of X-rays as the gas decelerates through a radiative shock and settles onto the stellar surface. The spectra of the accretion-powered sources are often well fitted using a combination of a power-law spectrum in the $5 - 20\\,$keV energy range plus a blackbody component with a temperature in the range $T \\sim 10^{6-7}\\,$K (e.g., Coburn et al. 2002; di Salvo et al. 1998; White et al. 1983) and a quasi-exponential cutoff at energy $E \\sim 20-30\\,$keV. There are also indications of cyclotron features and iron emission lines in a number of sources (e.g., Pottschmidt et al. 2005). The observations suggest typical magnetic field strengths $\\sim 10^{11-12}\\,$G.  Previous attempts to calculate the spectra of accretion-powered X-ray pulsars based on static or dynamic theoretical models have generally yielded results that do not agree very well with the observed profiles (e.g., M\\'esz\\'aros \\& Nagel 1985a,b; Nagel 1981; Yahel 1980; Klein et al. 1996). Due to the lack of a fundamental physical model, most X-ray pulsar spectral data have traditionally been fitted using multicomponent functions of energy that include absorbed power-laws, cyclotron features, iron emission lines, blackbody components, and high-energy exponential cutoffs. The resulting model parameters are often difficult to relate to the physical properties of the source. However, the situation improved recently with the development by Becker \\& Wolff (2005a,b) of a new model for the spectral formation process based on the ``bulk'' or ``dynamical'' Comptonization (i.e., first-order Fermi energization) of photons due to collisions with the rapidly compressing gas in the accretion column . In the ``pure'' bulk Comptonization model of Becker \\& Wolff, the transfer of the proton kinetic energy to the photons, essential for producing the emergent spectrum, is rigorously modeled using a transport equation that includes a first-order Fermi term that accounts for the energy gain experienced by photons as they scatter back and forth across the accretion shock. The seed photons for the upscattering process are provided by the blackbody radiation injected at the surface of the dense ``thermal mound,'' located at the bottom of the accretion column (Davidson 1973). By ignoring the effect of thermal Comptonization in the column, the model corresponds physically to the accretion of a gas with thermal velocity much less than the dynamical (bulk) velocity. This ``cold plasma'' criterion is well satisfied in X-ray pulsars because the typical inflow speed at the top of the column is $\\sim 0.5\\,c$ and the characteristic temperature is $\\sim 10^7\\,$K. The resulting X-ray spectrum is characterized by a power law with photon index $\\alpha_{\\rm X} > 2$ combined with a low-energy turnover representing unscattered Planck radiation. The X-ray spectrum computed using the pure bulk Comptonization model has no high-energy cutoff, and therefore the photon index must exceed two in order to avoid an infinite photon energy density. Despite this restriction, the bulk Comptonization model has successfully reproduced the observed spectra of several steep-spectrum sources, including X Persei and GX 304-1 (Becker \\& Wolff 2005a,b).  Many of the brightest X-ray pulsars, such as Her X-1 and Cen X-3, display spectra with photon indices $\\alpha_{\\rm X} \\le 2$ in the $5 - 20\\,$keV energy range, combined with high-energy quasi-exponential cutoffs at $20-30\\,$keV (e.g., White et al. 1983). This type of spectral shape cannot be explained using the pure bulk Comptonization model of Becker \\& Wolff (2005b). The presence of the exponential cutoffs at high energies, combined with the relatively flat shape at lower energies, suggests that thermal Comptonization is also playing an important role in luminous X-ray pulsars by transferring energy from high- to low-frequency photons via electron scattering. The two steps in the thermal process (Compton scattering followed by inverse-Compton scattering) are described mathematically by the Kompaneets (1957) equation (see also Becker 2003). Although it is clear that the majority of the photon energization in X-ray pulsars occurs via the first-order Fermi (i.e., bulk Comptonization process), we find that in bright sources such as Her X-1, the role of the thermal Comptonization process must also be considered in order to reproduce the observed spectra.  In this paper, we extend the pure bulk Comptonization model developed by Becker \\& Wolff (2005b) to include both bulk and thermal Comptonization by incorporating the full Kompaneets operator into the transport equation used to model the development of the emergent radiation spectrum. The exact solution to this equation yields the Green's function for the problem, which represents the contribution to the observed photon spectrum due to a monochromatic source at a fixed height in the accretion column. By exploiting the linearity of the mathematical problem, the Green's function can be used to compute the emergent spectrum due to any spatial-energetic distribution of photon sources in the column. We calculate X-ray pulsar spectra by convolving the Green's function with bremsstrahlung, cyclotron, and blackbody sources, with the first two distributed throughout the column, and the latter located at the surface of the thermal mound. The accretion/emission geometry is illustrated schematically in Figure~1. Seed photons produced inside the column by various emission mechanisms experience electron scattering as they diffuse through the column, eventually escaping through the walls to form the emergent X-ray spectrum. The escaping photons carry away the kinetic energy of the gas, thereby allowing the plasma to settle onto the surface of the star.  The remainder of the paper is organized as follows. In \\S~2 we briefly review the nature of the primary radiation transport mechanisms in the accretion column with a focus on dynamical, thermal, and magnetic effects. The transport equation governing the formation of the radiation spectrum is introduced and analyzed in \\S~3, and in \\S~4 the exact analytical solution for the Green's function describing the radiation distribution inside the accretion column is derived. The spectrum of the radiation escaping through the walls of the accretion column is developed in \\S~5, and the physical constraints for the various model parameters are considered in \\S~6. The nature of the source terms describing the injection of blackbody, cyclotron, and bremsstrahlung seed photons into the accretion column is discussed in \\S~7. Emergent X-ray spectra are computed in \\S~8, and the results are compared with the observational data for several luminous X-ray pulsars. The implications of our work for the production of X-ray spectra in accretion-powered pulsars are discussed in \\S~9. ",
        "conclusions": "\\begin{figure}[t] \\begin{center} \\epsfig{file=f8.eps,height=10.0cm} \\end{center} \\caption{Same as Fig.~6, except the data correspond to Cen X-3 and the theoretical spectra were computed based on the model~3 parameters reported in Table~1. The data were reported by Burderi et al. (2000; {\\it circles and crosses}). This case also includes interstellar absorption with hydrogen column density $N_{\\rm H}=2.0 \\times 10^{22}\\,{\\rm cm}^{-2}$.} \\end{figure}  The new theoretical model developed here represents the first complete calculation of the X-ray spectrum associated with the physical accretion scenario first suggested by Davidson (1973), including bulk and thermal Comptonization. For given values of the stellar mass $M_*$ and the stellar radius $R_*$, our model has six fundamental free parameters, namely the accretion rate $\\dot M$, column radius $r_0$, the electron temperature $T_e$, the magnetic field strength $B$, the photon diffusion parameter $\\xi$ (eq.~[\\ref{eq3.12}]), and the Comptonization parameter $\\deltapar$ (eq.~[\\ref{eq4.5}]). The accretion rate is constrained by the observed luminosity, which effectively removes this parameter from the set. A unique solution for the remaining free parameters can be found by comparing the model spectrum with the observations for a particular source. The analytical nature of the new model makes it amenable to incorporation into standard analysis packages such as XSPEC.  We have shown that the combination of bulk and thermal Comptonization naturally leads to the development of emergent spectra in accretion-powered X-ray pulsars that are in good agreement with the observational data for three bright sources. The main result of the derivation is the analytical solution for the column-integrated Green's function describing the escaping radiation spectrum, given by equation~(\\ref{eq5.10}). Based on this closed-form solution, we have developed expressions for the emergent spectra resulting from the reprocessing of seed photons injected as a result of cyclotron, blackbody, and bremsstrahlung emission, given by equations~(\\ref{eq7.5}), (\\ref{eq7.12}), and (\\ref{eq7.15}), respectively (we note that additional seed photons may also be supplied via external illumination of the column, which is not considered here). In each case, we confirm that the number of photons escaping from the column per unit time is exactly equal to the number injected, as expected in the steady-state scenario considered here. Our main results, summarized in Figures~6, 7, and 8, indicate that the spectra observed from the bright sources Her X-1, LMC X-4, and Cen X-3 are dominated by Comptonized bremsstrahlung.  The treatment of the full energy and angle dependence of the scattering cross sections for the two polarization modes computed using equations~(\\ref{eq2.2}) and (\\ref{eq2.4}) is not possible using an analytical approach. As an alternative, in this paper we have adopted a two-dimensional approximation by modeling the scattering of photons propagating either parallel or perpendicular to the magnetic field using the energy- and mode-averaged expressions given by equations~(\\ref{eq2.6}) and (\\ref{eq2.7}), respectively. The results we have obtained for the directional components of the cross section listed in Table~2 are not unreasonable, given the uncertainties in the detailed physical properties of the accretion column. In particular, we find that the cross section for photons propagating parallel to the field direction, $\\sigpar$, is substantially reduced relative to the Thomson value. This result is consistent with equation~(\\ref{eq2.6}), since $\\epsilon \\ll \\epsilon_c$ for most of the photons in typical X-ray pulsar spectra. Furthermore, the cross section values reported in Table~2 confirm that $\\sigpar \\ll \\sigbar \\ll \\sigperp$, as expected (see \\S~6.2 and Canuto et al. 1971). Compared to the behavior obtained in our model, we expect that utilization of the exact cross sections given by equations~(\\ref{eq2.2}) and (\\ref{eq2.4}) would tend to increase the efficiency of thermal Comptonization, relative to the bulk process, for photons with energy $\\epsilon$ close to the cyclotron energy $\\epsilon_c$. This effect is therefore expected to cause some modification of the parameter values required to fit the observed X-ray pulsar spectra. Since the radiative transfer problem including the exact cross sections is not tractable analytically, further study of this issue awaits the development of a complete numerical model, which we intend to pursue in future work.  The validity of the approximate velocity profile employed here (eq.~[\\ref{eq3.15}]) was established in Figure~5 by demonstrating excellent agreement between the bulk Comptonization spectrum computed using our new model in the limit $T_e \\to 0$ and the corresponding spectrum obtained using the pure bulk Comptonization model of Becker \\& Wolff (2005b), which is based on the exact velocity profile (eq.~[\\ref{eq3.16}]). The values obtained for the accretion column radius $r_0$ (reported in Table~1) comply with the corresponding upper limits computed using equation~(\\ref{eq6.6.3}), which was derived based on the detailed dynamics of neutron star accretion flows (Lamb et al. 1973; Harding et al. 1984). We find that $r_0$ is essentially equal to the upper limit for the two brightest sources considered here (Cen X-3 and LMC X-4), while it is about an order of magnitude smaller than this limit for Her X-1. Since the luminosity of Her X-1 is roughly an order of magnitude lower then the other two sources, this suggests a possible connection between the luminosity and the column size that should be further investigated by developing numerical models that simultaneously treat the dynamics and the spectral formation in X-ray pulsars.  The transfer of energy from the gas to the photons is generally dominated by bulk Comptonization in accretion-powered X-ray pulsars. However, our results confirm that unsaturated thermal Comptonization nonetheless plays a significant role in the spectral formation process. Specifically, we find that the thermal process effectively transfers energy from high to low frequency radiation, which contributes to the observed quasi-exponential cutoffs at high energies and at the same time causes flattening of the spectrum at lower energies. While it has been known for some time that pulsars generally have spectra that are well fitted using a combination of a power-law plus a quasi-exponential cutoff, until now this form has been adopted in a purely ad hoc manner. The new model described here finally provides a firm theoretical foundation for this empirical result, based on model parameters that are directly tied to the physical properties of the source. It therefore represents a significant step in the development of a comprehensive theory for the spectral formation process in accretion-powered X-ray pulsars.  The predictions of our analytical model can be tested both observationally and also by developing numerical models that allow the relaxation of several of the key assumptions made here. For example, within the context of a numerical model one can include the exact spatial variation of the shape and strength of the dipole magnetic field, as well as the full energy and angle dependence of the electron scattering cross section. Although these effects are beyond the scope of the present paper, we nonetheless believe that our main conclusions regarding the spectral shape will be verified by more detailed models.  The authors would like to thank Ken Wolfram, Paul Ray, Lev Titarchuk, and Kent Wood for a number of stimulating conversations. The authors are also grateful to the anonymous referee who provided a number of insightful comments that led to significant improvements in the manuscript. PAB would also like to acknowledge generous support provided by the Office of Naval Research.  \\clearpage"
    },
    "0609/astro-ph0609505_arXiv.txt": {
        "abstract": "We present results on the structure of the near-surface layers of the Sun obtained by inverting frequencies of high-degree solar modes from ``ring diagrams''. We have results for eight epochs between June 1996 and October 2003. The frequencies for each epoch were obtained from ring diagrams constructed from MDI Dopplergrams spanning complete Carrington rotations. We find that there is a substantial latitudinal variation of both sound speed and the adiabatic index $\\Gamma_1$ in the outer 2\\% of the Sun. We find that both the sound-speed and $\\Gamma_1$ profiles change with changes in the  level of solar activity. In addition, we also study differences between the northern and southern hemispheres of the Sun and find a small asymmetry that appears to reflect the difference in magnetic activity between the two hemispheres. ",
        "introduction": "\\label{sec:intro}  Seismic data have been successfully used to determine the solar interior structure (e.g., Gough {\\it et al.}~1996). Departures from spherical symmetry  have been less well studied. Among the more  detailed studies is one by Antia  {\\it et al.}~(2001). This study showed that the solar sound-speed profile is not spherically symmetric, and that it depends on both radius and latitude. Helioseismic determinations of solar structure --- both the spherically symmetric and the latitudinally dependent parts ---  are  restricted to the deeper layers of the Sun ($ r < 0.95 R_\\odot$) because of the limitations of the mode sets that are routinely determined by the Michelson Doppler Imager (MDI) on board the Solar and Heliospheric Observatory (SOHO) and from the ground-based Global Oscillations Network Group (GONG). These data sets provide frequencies for {\\it p}-modes with degree $\\ell \\le 190$ and {\\it f}-modes with degree  $\\ell\\le 300$ for the most part, and these low- and intermediate-degree modes are not very useful in determining the structure of the near-surface layers of  the Sun. However, the {\\it f}-modes, which are confined to the outermost layers of the Sun, are more sensitive to rotation than structure,  and these provide reliable determinations of solar rotation to fairly shallow depths. Solar rotation and other large-scale flows are known to have a strong latitudinal dependence (e.g., Thompson {\\it et al.}~1996; Schou {\\it et al.}~1998); this dependence is also known to show a time variation that is correlated with solar activity (e.g.,  Schou 1999; Basu \\& Antia 2000, 2002, 2003; Howe {\\it et al.}~2000, 2004a). Changes in the latitudinal variation of structure in near-surface regions are, however, difficult to study because of the lack of availability of frequencies for high-degree modes.  The changes in frequencies of intermediate-degree global modes that are routinely determined from MDI and GONG data show that there are no clearly observable changes in solar structure below $0.95R_\\odot$ (Eff-Darwich {\\it et al.}~2002; Basu 2002). There is tentative and indirect evidence that there are indeed solar-cycle related changes in the adiabatic index $\\Gamma_1$ in the shallower layers of the Sun, around the He II ionization zone, i.e. $\\approx 0.98R_\\odot$, (Basu \\& Mandel 2004; Verner {\\it et al.}~2006). Chou \\& Serebryanskiy (2005) find some marginal evidence for temporal variations near the base of the convection zone, which contradicts earlier negative results (Eff-Darwich {\\it et al.}~2002; Basu 2002). Similarly, inversions of the frequency splittings of these modes do not show any significant time variations in the latitudinal distribution of sound speed and density below $0.95R_\\odot$ (Antia {\\it et al.}~2001, 2003). However, the ``surface term'' from the inversion (see \\S~\\ref{sec:inv}) changes with time, implying that there could be changes in the structure closer to the solar surface.  The surface term also shows a distinct correlation with the distribution of magnetic fields at the solar surface (Fig.~8 of Antia {\\it et al.}~2001). This leads us to believe that we may be able to detect changes in the structure of the outer layers of the Sun associated with magnetic activity changes by using  frequencies of high-degree modes. It is already known that frequencies of high-degree modes vary with local magnetic field strength (e.g., Hindman {\\it et al.}~2000; Rajaguru {\\it et al.}~2001; Howe {\\it et al.}~2004b), and in principle, these frequency differences can be inverted to infer corresponding variations in solar structure or other properties. Earlier studies have shown that the profiles of sound speed and $\\Gamma_1$ in regions of high magnetic activity differ from those in quiet regions (Kosovichev {\\it et al.}~2000, 2001; Basu, Antia \\& Bogart~2004). These studies show that for magnetically active regions, sound speed and adiabatic index $\\Gamma_1$ are lower than that of quiet regions in the immediate sub-surface layers (about the outer seven Mm or so).  If there are indeed changes in the outer layers of the Sun, we need frequencies of high-degree modes to detect these changes. Unfortunately, it is very difficult to determine the frequencies of very high degree modes through global analysis (e.g., Korzennik {\\it et al.}~2004 and references therein) and as a result neither GONG nor MDI provide data on high degree modes as standard products, and frequency sets with high-degree modes are not available for different epochs. High-degree solar  modes ($l\\ge150$) which are trapped in the outer parts of the solar envelope have lifetimes that are much shorter than the sound travel time around the Sun and hence the characteristics of these modes are mainly determined by the average conditions in the local neighborhood rather than those of the entire sphere. The properties of these modes are better determined by local helioseismic techniques. Ring-diagram (plane-wave $k$--$\\nu$) and time-distance analysis (Duvall \\ea~1993) are examples of such local helioseismic techniques. Frequencies obtained from ring-diagram analysis have been successfully used to determine the difference in structure between active and quiet regions of the Sun (Basu, Antia \\& Bogart 2004).  Since high-degree global-mode data are not standard products of the GONG or MDI projects, we use a ring-diagram analysis to obtain solar oscillation frequencies of different latitudinal zones of the Sun at different times. This technique allows us to do a differential analysis of the structure between different regions. We use these frequencies to determine the differences in structure between the equator and higher latitudes. Since we are interested in the near-surface layers where ionization occurs, we need to invert for both the sound speed ($c$) and the adiabatic index ($\\Gamma_1$) in order to be able to interpret the results, since any temperature change in the ionization zone will also change $\\Gamma_1$ there. To study how the latitudinal dependence of structure changes with time and activity, we analyze several epochs corresponding to different solar activity levels. The differences in structure between the equator and higher latitudes are expected to be less than those found between strong active regions and quiet regions, because of the smaller differences in the zonally averaged magnetic fields, nevertheless we expect the differences to be detectable.  The rest of the paper is organized as follows: In \\S~\\ref{sec:data} we discuss the analysis technique and the regions studied; we give a brief description of the inversion process in \\S~\\ref{sec:inv}; we describe and discuss our results in \\S~\\ref{sec:res}, and state our conclusions in \\S~\\ref{sec:conclu}. ",
        "conclusions": "\\label{sec:conclu}  We have analyzed data from eight full Carrington rotations to determine the latitudinal distribution of the axisymmetric components of sound speed and $\\Gamma_1$, and to determine whether those distributions change with time and magnetic activity level. Ring-diagram analysis allows us to determine frequencies of high-degree modes, permitting us to determine solar structure in the near-surface layers, where most solar-cycle related changes are believed to occur.  Our results are valid in the radius range $0.975$--$0.996R_\\odot$ and most reliable in the range $0.985$--$0.995R_\\odot$. We have studied latitudes up to $\\pm45^\\circ$.  For all epochs studied, we find significant frequency differences between the equator and the higher latitudes.  These differences cannot be explained as being caused by projection effects alone; most of the differences are  believed to reflect physical causes. Assuming that the frequency differences are caused by differences in structure alone, we have inverted them to determine the sound speed and the adiabatic index differences between the equator and higher latitudes. We find significant differences in structure between the equator and the higher latitudes for all epochs. The differences are largest for latitudes $ > 30^\\circ$. These regions tend to have higher sound speed and $\\Gamma_1$ than the equator. The differences in $\\Gamma_1$ between the equator and higher latitudes imply temperature changes in these regions. Using global modes, Antia {\\it et al.}~(2001) also found a region of positive sound speed difference around a latitude of $60^\\circ$. Although the magnitude of the difference is much smaller in deeper layers, of the order of $10^{-4}$, the qualitative behavior is similar. We expect the relative magnitude of the effect to increase in the near surface layers, as the global analysis suggests that most of the temporal variations in the solar oscillation frequencies arise from the surface layers. Because of the low pressure near the surface, magnetic fields are expected to be more effective in modifying the structure.  We find that the surface term from the inversions changes with time. The surface term is reasonably correlated with the difference between the magnetic activity index of the higher latitude regions and the equator. The correlation improves if only $|\\delta {\\rm MAI} | > 5$~G regions are considered. The surface term is also correlated with the 10.7 cm flux, though the correlation is strongly dependent on the latitude. The latitudinal dependence of the correlation coefficient can be explained by changes in the latitudinal distribution of the magnetic field as solar activity changes.  The asphericity of the Sun changes with time, with the difference between the higher latitudes and the Equator increasing with increasing levels of activity. The differences can be explained by the change in the magnetic field distribution on the Sun as the activity level increases. The change with time and activity is in opposite directions for latitudes above and below $30^\\circ$. This is consistent with the evolution of the magnetic field distribution. As with the surface term, the differences are correlated both with local and global magnetic field indices. The magnitude of changes between the active and quiet periods seen in this study is about a factor of 3--10 less than what was observed between individual active and quiet regions (Basu {\\it et al.}~2004). This is expected, as the average magnetic field strength over a whole latitude band is much smaller than that in active regions alone. Although it is difficult to separate the effects of magnetic field and structural variations, one may expect the magnetic field to play a significant role near the surface. The ratio of magnetic to gas pressure is expected to be comparable, to order of magnitude, to the variations found in sound speed. At $r=0.99R_\\odot$ a variation of $10^{-3}$ may correspond to a magnetic field of 4000 G, which is considerably larger than the typical magnetic field observed at the surface.  We have also looked at temporal variations of structure by inverting the frequency differences between times when the Sun was active (2000--2002) and when it was relatively quiet (1996--1998). Assuming  that there are no time-dependent systematic errors in the data, these differences were inverted to look at how the structure of the outer layers of the Sun  differed during the two epochs. We find that there are significant differences in the near-surface structure of the Sun during quiet and active periods. These results, however, may be affected by systematic observational variations. It is also possible that the results are systematically biased in some subtle way by the different projection effects involved in tracking or not tracking the data.  We have also studied the north-south asymmetries in the structure of the outer layers of the Sun. We find that when the Sun was active during cycle 23, the sound speed in the northern hemisphere was higher than that in the southern hemisphere, at least at high latitudes. Somewhat similar behavior is seen for $\\Gamma_1$. This can be explained by differing amounts of asymmetry in the magnetic fields of the northern and southern hemispheres when the Sun was quiet and when the Sun was active."
    },
    "0609/astro-ph0609068_arXiv.txt": {
        "abstract": "{} {  % The goal of this work is to measure the evolution of the Galaxy Stellar Mass Function and of the resulting Stellar Mass Density up to redshift $\\simeq 4$, in order to study the assembly of massive galaxies in the high redshift Universe.  } {We have used the GOODS-MUSIC catalog, containing $\\sim$3000 Ks-selected galaxies with multi-wavelength coverage extending from the U band to the Spitzer $8 \\mu$m band, of which 27\\% have spectroscopic redshifts and the remaining fraction have accurate photometric redshifts. On this sample we have applied a standard fitting procedure to measure stellar masses. We compute the Galaxy Stellar Mass Function and the resulting Stellar Mass Density up to redshift $\\simeq 4$, taking into proper account the biases and incompleteness effects. } {Within the well known trend of global decline of the Stellar Mass Density with redshift, we show that the decline of the more massive galaxies may be described by an exponential timescale of $\\simeq 6$~Gyrs up to $z\\simeq 1.5$, and proceeds much faster thereafter, with an exponential timescale of $\\simeq 0.6$~Gyrs.  We also show that there is some evidence for a differential evolution of the Galaxy Stellar Mass Function, with low mass galaxies evolving faster than more massive ones up to $z\\simeq 1-1.5$ and that the Galaxy Stellar Mass Function remains remarkably flat (i.e. with a slope close to the local one) up to $z\\simeq 1-1.3$. } {The observed behaviour of the Galaxy Stellar Mass Function is consistent with a scenario where about 50\\% of present--day massive galaxies formed at a vigorous rate in the epoch between redshift 4 and 1.5, followed by a milder evolution until the present-day epoch.} ",
        "introduction": "The observational evidence of the continuous increase of the stellar content of galaxies over cosmic times has emerged only recently from a series of observations and surveys which have made use of new sensitive IR instrumentation.  Following  the first pioneering studies (\\cite{giallongo98}, \\cite{brinchmann00}, \\cite{papovich01}) that set-up the technique for estimating stellar masses in high redshift galaxies, several surveys pointed out that the global stellar content of the Universe, as measured by the average stellar mass density, increases with cosmic time (\\cite{brinchmann00}, \\cite{dickinson03}, \\cite{fontana03}, \\cite{rudnick03}, \\cite{glazebrook04}, \\cite{drory04}, \\cite{fontana04}, F04 hereafter, \\cite{rudnick06}).  While the average stellar mass density provides a global picture of the process of stellar assembly, the Galaxy Stellar Mass Function (GSMF in the following) provides a more detailed view on how this process evolves as a function of the galaxy mass itself. At low redshift, accurate GSMF have been obtained from the 2dF (\\cite{Cole01}) or 2MASS--SDSS (\\cite{bell03}) surveys.  Because of the large area and of the accuracy in the redshift estimate that is needed to compile a GSMF, very few GSMFs have been obtained so far. The MUNICS survey (\\cite{drory04}) and the K20 survey (F04) first explored the evolution of the GSMF up to $z\\simeq 1-2$ in a somewhat complementary fashion. The former adopted a wide, relatively deep sample on a wide area ($\\simeq 5000$ objects at $K\\leq19$), mostly relying on photometric redshifts, sampling the GSMF up to $z\\simeq 1.5$, while the latter adopted a deeper but smaller sample ($\\simeq 500$ objects at $K\\leq20$), with excellent spectroscopic coverage, extending up to $z=2$. In the overlapping redshift ranges, the two GSMF agree quite well, and both analyses suggested a decline in the density of massive galaxies at $z\\simeq 1$, of the order of 50-70\\% with respect to the local value. Recently, \\cite{drory05} extended such analysis to higher ($z\\simeq 5$) redshift, adopting a combination of $I$ and $K$--selected samples.  F04 also pointed out a tentative evidence of a differential evolution of the GSMF, such that massive galaxies appear to evolve less than low mass galaxies, at least up to $z\\simeq 1.5$. As we shall show in the following, this trend is also shown by our new data.  Along a different line, evidence has emerged that the GSMF for galaxies of different spectral or morphological types shows a significant evolution, with an increase in the fraction of stellar mass residing in late type galaxies at $z\\simeq 1-1.5$ with respect to low redshift (F04, \\cite{bundy05}, \\cite{franceschini06}).  Although these surveys have already provided a first picture of the evolution of the GSMF, the uncertainties involved in this exercise are still large.  In particular, the exploration of the high redshift Universe is still very limited. In several cases, the lack of long--wavelength data hampered the estimate of the stellar mass at high redshift, for which an adequate sampling of the rest frame optical--near infrared part of the spectrum is essential. In addition, the collection of a large sample of high redshift galaxies requires a combination of large areas and deep near--IR observations. For different reasons, all the surveys mentioned above suffer of these limitations, although to a different extent.  In this context, the GOODS-South survey provides an excellent opportunity of improving in a significant way our knowledge of the high--$z$ GSMF. The combination of deep, wide IR observations with Spitzer (in its four channels from 3.5 to 8~$\\mu$m) and VLT-ISAAC (in $J$, $H$ and $Ks$), coupled with high quality imaging data in the optical domain (both ACS and VLT--VIMOS), and of a large, extended spectroscopic coverage make the GOODS-South field ideal for this investigation. From this public data set we have obtained a multicolour catalog of faint galaxies, that we named GOODS--MUSIC (GOODS MUltiwavelength Southern Infrared Catalog), that we describe in \\cite{grazian06a}.  In this work we use the GOODS--MUSIC sample to improve the previous estimates of the GSMF in several ways. The most important is that our analysis includes the $3.5-8 \\mu$m Spitzer observations of the complete $Ks$--selected data set. In addition, it extends to $Ks$ magnitudes deeper than any previous survey, enabling us to obtain a complete sample of galaxies up to $z\\simeq 4$, and to measure the slope of the low mass side of the Galaxy Stellar Mass Function up to $z\\simeq 1.3$. Compared to our previous analysis of the K20 survey data, the present data set provides a final sample that is 6$\\times$ larger than the K20 one, on which we adopt a more sophisticated technique to parametrize the evolution of the Galaxy Stellar Mass Function with a redshift-dependent Schechter fit, that provides interesting clues on the evolution of massive galaxies.  The paper is organized as follows. In Sect.2, after reminding the basic feature of our dataset and of the procedure that we adopt to extract the stellar masses from each galaxy, we discuss how the inclusion of the Spitzer bands affects the mass estimates and highlight the selection effects that will affect our analysis. In Sect.3, we present the basic results of our analysis, namely the Galaxy Stellar Mass Function and the resulting mass density, both in a binned as well as in a parametric fashion. In Sect.4, we compare our basic findings with the prediction of recent theoretical models. Finally, in Sect.5, we focus on the highest redshift range, to discuss the reliability of the photometric redshifts on a few, intriguing objects that might be at very high redshift. The discussion and conclusions are summarized in Sect.6.  All the magnitudes cited below are in the AB system. To scale luminosities and compute volumes we have adopted a ``concordance'' cosmological scenario with $\\Omega_\\Lambda = 0.7$ and $H_0=70$~km~s$^{-1}$~Mpc$^{-1}$. ",
        "conclusions": "In this work we have presented an analysis of the evolution of the Galaxy Stellar Mass Function (GSMF) and of the corresponding stellar mass density up to $z\\simeq 4$.  It has been obtained from the GOODS-MUSIC sample, a $Ks$--selected catalog of 2931 galaxies with 14 bands photometry, extending from $0.35$ to $8\\mu$m, extracted from the public data of the GOODS-South survey.  We have derived accurate stellar masses for this large sample of galaxies, adopting the standard technique of fitting spectral synthesis models (Bruzual \\& Charlot 2003) on the multiwavelength data. To compare with previous works, we have used the Salpeter IMF, ranging over a set of metallicities (from $Z=0.02 Z_\\odot$ to $Z=2.5 Z_\\odot$ and dust extinction ($0<E(B-V)<1.1$, with a Calzetti extinction curve).  On this data set, we have computed the Galaxy Stellar Mass Function and the resulting stellar mass density with the usual $1/V_{max}$ formalism. To provide a visualization of the GSMF free of the well known biases of the $1/V_{max}$ approach (namely, the sensitivity to large scale structures and the arbitrarity of the binning), we also provide a global fit to the data with a Schechter function with smooth, redshift--dependent free parameters (``STY approach'').  We have carefully discussed the selection effects that play a role in our $Ks$-selected sample, and that must be taken into account in the analysis.  We have shown that, while we can detect massive ($M_*>10^{11}M_\\odot$) passively evolving galaxies up to $z\\simeq 4$, our sample becomes progressively biased against star--forming, dust-enshrouded objects with $E(B-V)>0.5$ already at $z>2$. Since there are evidences that these objects do exist at $z<2$, it is possible that our census of galaxies at $z>2$ is incomplete, and that our estimates in this redshift range (as in any $Ks$--selected sample) should be considered as a lower limit.  The major results that we have found with these data are the following:  - We show that the inclusion of the Spitzer data (from 3.5 to 8 $\\mu$m) significantly improves the reliability of the mass estimate at $z>2$, as expected. At lower redshift, there is no significant improvement and the observed scatter is consistent with that induced by the model degeneracy.  - We confirm the well known trend of global decline of the stellar mass density with redshift. The total mass density is about a factor 2$\\times$ lower than local at $z\\simeq 1$, and about 10\\% of the local at $z\\simeq 3$.  - We compute the GSMF in several redshift bins, from $z=0.4$ to $z=4$, and we show that a simple scaling of the Schechter parameters is able to provide a smoothly evolving rendition of the GSMF.  - Restricting only to galaxies with $M_*>10^{11}M_\\odot$, their mass density evolves relatively mildly up to $z\\simeq 1-1.5$. At $z\\simeq 1$, their integrated mass density is about 50\\% of the local value.  At $z>1.5$, massive galaxies become to evolve much faster, such that at $z\\simeq 3$ they provide at most one tenth of the local density. This trend may be described by two exponential e--folding times, i.e. $\\rho(z)\\propto e^{-t/\\tau}$, with $\\tau\\simeq 6$Gyrs at $z<1.5$ and $\\tau\\simeq 0.6$Gyrs at $z>1.5$, although this number may be biased low because of the selection effects described above.  - As far as low mass galaxies are concerned, we show that the GSMF remains remarkably flat (it steepens by less than 0.1 for each unit redshift) up to $z\\simeq 1-1.3$, the highest redshift bin where it can be reliably measured. At the same time, a sensible decrease occurs in the number of low mass galaxies: the density of galaxies with $M\\simeq 10^{10}M_\\odot$ at $z\\simeq 1$ is 4 times lower than the local one.  - In this context, we finally show that there is a clear evidence for a differential evolution of the Galaxy Stellar Mass Function up to $z\\simeq 1.5$, with less massive galaxies evolving more than massive ones. Such a trend is evident in Fig.\\ref{massfun}, that shows that the GSMF in the redshift bins up to $z\\simeq 1.3$ remains progressively closer to the local one for increasing stellar masses, and it is also substantiated by two quantitative evidences. First, from $z=0$ to $z\\simeq 1.5$ the total stellar mass density decreases more than that due to massive galaxies only (i.e. those with $M^*>10^{11}M_\\odot$). More intriguingly, the STY fit to the overall evolution of the GSMF shows that the characteristic mass $M^*$ increases up to $z\\simeq 1.3$, reflecting the change in the shape of the GSMF, that from $z\\simeq 0$ to $z\\simeq 1.3$ is progressively more skewed toward higher masses.  - After $z\\simeq 1.5$, however, the evolution of the mass density contained in massive galaxies is much faster, as shown by Fig.\\ref{massdens}, showing that this is their main epoch of formation.  It is straightforward to see that such differential evolution of the GSMF is a natural consequence of the ``downsizing'' scenario for galaxy evolution, that has been found in many independent surveys. Following the original definition by Cowie et al. 1996, it has been shown by several authors (e.g. \\cite{brinchmann00}, \\cite{fontana03}, \\cite{perez05}, \\cite{feulner05}) that the specific star--formation rate increases with redshift, showing that massive galaxies become progressively more actively star--forming as redshift increases. By looking at the GSMF, we observe the consequences of such a trend: more massive systems form in a vigorous phase at high redshift, that is largely complete at $z\\simeq 1-1.5$, such that the corresponding section of the GSMF is already close to the local value. Lower mass systems, on the contrary, continue to grow their stellar content at even lower redshifts, such that the increase of the faint side of the GSMF is large also at low $z$.   We remark that this picture is not plagued by the biases against the detection of dusty massive galaxies at high $z$ that exist in our sample, as in any $Ks$ complete one. As we discussed in the text, we might be missing high redshift, star--forming dusty galaxies with large extinction, but we are essentially complete with respect to passively evolving, dust-free galaxies of $M_*>10^{11}M_\\odot$ up to $z\\simeq 4$.  The low density of high redshift massive systems provides decisive evidence that the formation and assembly of local, massive bulges and ellipticals did not form in a single phase at very high redshift.  We have explored whether such a picture is qualitatively consistent with the predictions of most recent theoretical models for galaxy formation. We have compared our data with a set of theoretical models, including semianalytic models (\\cite{bower06}, \\cite{menci06}, \\cite{monaco06}) as well as hydrodynamical ones (\\cite{nagamine05a}, \\cite{nagamine05b}). All these models predict a relatively mild evolution of the stellar mass density contained in massive galaxies from $z=0$ to $z\\simeq 1.5$, that is broadly consistent with the observed data, and a much faster evolution at higher $z$, again in agreement with the data.  For all the models, however, at intermediate redshifts the match with the slope and the normalization of the GSMF at intermediate or low masses ($M_*\\simeq 10^{10} M_\\odot$) is still critical, when not even poor. Interestingly, this implies that the ``downsizing'' scenario (that is based on the evolution ratio between massive and low mass galaxies) is not reproduced by any model because of the overprediction of low mass galaxies, and not because  of an underprediction of the massive ones.   At high redshift, the detailed agreement with the observed data is still sensitive to the description of the physical processes inserted in the models. Hydrodynamical models tend to overpredict the observed mass density, as already noted by the authors (\\cite{nagamine04}), while the semianalytic models that include feedback from AGNs are closer to the data.  In conclusion, the new data presented in this work provide an overall description of the rise of the stellar mass, and in particular of that residing in massive galaxies, in which about half of such stellar mass appear to have been assembled during the first 2-4 Gyrs after recombination, followed by a milder increase over the remaining cosmic time. Although encouraging, the comparison with the theoretical expectations provides evidences that some fundamental physical processes, likely affecting both low and high mass galaxies, are still incorrectly represented, and that the density of massive galaxies at high redshift is indeed a very useful test for these models.  At this purpose, more effort is needed to improve the reliability of the observational estimates at intermediate and high redshift. From the observational point of view, larger spectrophotometric surveys on independent fields are definitely necessary to smear out the effects of sample variance. These might in particular affect two of our major results, namely the increase of the characteristic mass $M^*$ up to $z\\simeq 1$ and the amount of massive galaxies at $z>2$.  In addition to such (obvious) caveats, several systematic effects are still to be properly minimized. To mention a few, those related to the choice of the IMF and to the differences that may arise from the treatment of post--AGB stars in spectral synthesis models (\\cite{maraston05}), and the possible contribution of dust--enshrouded galaxies to the overall mass budget.  In addition, these findings rise important questions about the physical processes that led to the rise of the stellar mass density in massive galaxies: what is the physical nature of the galaxies that contribute to the stellar mass density at high redshifts, and what is the physical mechanism that drove this rise, i.e. how much it is related to star--formation occurring within the observed galaxies as compared to the contribution from merging processes.  Although present--day surveys are starting to explore these issues, these will remain among the more challenging questions of the next years."
    },
    "0609/astro-ph0609542_arXiv.txt": {
        "abstract": "We have developed a new software tool, \\textsc{Lensview}, for modelling resolved gravitational lens images. Based on the \\textsc{LensMEM} algorithm, the software finds the best fitting lens mass model and source brightness distribution using a maximum entropy constraint. The method can be used with any point spread function or lens model. We review the algorithm and introduce some significant improvements. We also investigate and discuss issues associated with the statistical uncertainties of models and model parameters and the issues of source plane size and source pixel size.  We test the software on simulated optical and radio data to evaluate how well lens models can be recovered and with what accuracy. For optical data, lens model parameters can typically be recovered with better than 1\\% accuracy and the degeneracy between mass ellipticity and power law is reduced. For radio data, we find that systematic errors associated with using processed radio maps, rather than the visibilities, are of similar magnitude to the random errors. Hence analysing radio data in image space is still useful and meaningful.  The software is applied to the optical arc HST J15433+5352 and the radio ring MG1549+3047 using a simple elliptical isothermal lens model. For HST J15433+5352, the Einstein radius is 0\\farcs525 $\\pm 0.015$ which probably includes a substantial convergence contribution from a neighbouring galaxy. For MG1549+3047, the model has Einstein radius 1\\farcs105 $\\pm 0.005$ and core radius 0\\farcs16 $\\pm 0.03$. The total mass enclosed in the critical radius is $7.06 \\times 10^{10}M_{\\odot}$ for our best model.  We finish by discussing issues relating to modelling of resolved lens images for this method and some alternatives. ",
        "introduction": "The discovery of the first Einstein rings at both radio \\citep{1988Natur.333..537H} and optical \\citep{Warren1996MN} wavelengths brought with them the promise of a new age of precision galaxy mass measurements. Strong lensing is a powerful astrophysical tool with applications ranging from measuring galaxy structural properties and evolution to cosmological parameter measurements using time delays (for $H_0$) or the statistical properties of lens samples.  Some applications of lensing, such as determining Hubble's constant ($H_0$) from time delays of multiple-image quasars, rely on accurate knowledge of both the radial mass profile of the lens mass and the location of the lens centre. Quadruple images of lensed QSOs commonly have a range of solutions for mass slope and ellipticity \\citep[e.g.][]{1994AJ....108.1156W} which cannot be resolved without additional assumptions or information. Einstein rings can, in principle, help to solve this problem because they roughly trace the shape of the tangential critical line, hence the ellipticity of the mass distribution is well determined. In addition, the many lines-of-sight provided by Einstein rings sample the gravitational potential much more thoroughly than QSO lenses. Hence, there is a great deal of information contained in Einstein ring images.  Modelling resolved gravitationally lensed images, such as Einstein rings, is a much more complicated problem than modelling lensed QSO images. The lensed image of a resolved source is determined both by the properties of the lensing galaxy and the source brightness profile. Consequently, algorithms for modelling resolved images have gone through several stages of development. The Ring Cycle algorithm \\citep{1988A&A...191...39K,1989MNRAS.238...43K} reconstructed a source based on the assumption that the data is a true representation of the lensed surface brightness. \\textsc{LensCLEAN} \\citep{1992ApJ...401..461K} uses a CLEAN style algorithm to build a non-parametric source from a processed radio map. This was subsequently improved \\citep{1996ApJ...464..556E} to work directly on radio visibilities rather than CLEANed radio maps. More recently, it was improved again \\citep{2004MNRAS.349....1W} and applied to the lens B0218+357 \\citep{2004MNRAS.349...14W} to make an accurate $H_0$ measurement. Also recently, the `semilinear' method was developed by \\citet{2003ApJ...590..673W} who noted that for a fixed lens model, determining the source model that best describes the data can be posed as a linear inversion problem and can therefore be tackled with standard matrix inversion techniques. \\citet{2005PASA...22..128B,2006ApJ...637..608B} have investigated the lens modelling problem using genetic algorithms and Bayesian techniques, respectively.  An alternate approach to modelling radio images was provided by the \\textsc{LensMEM} algorithm \\citep[][ hereafter WNK94,WKN96]{1994ApJ...426...60W,1996ApJ...465...64W}. This approach used a pixelised source and the novel idea of a mapping matrix (relating pixels in the image and source planes) to generate a source model using an iterative approach. This algorithm is equally suitable for radio and optical data with a point-spread function (PSF) and it forms the basis of the software described in this paper.  In \\S\\ref{sec:how_it_works} the algorithm is reviewed and some significant improvements are described, in particular those proposed by \\citet[][ hereafter SB84]{1984MNRAS.211..111S}. Also discussed is the issue of defining an acceptable goodness-of-fit with non-parametric (i.e. pixelised) sources for both radio and optical data. In \\S\\ref{sec:test_cases} the software is applied to test optical and radio data, showing that for a radio map the lens model parameters can be recovered quite accurately, despite the problems associated with analysing CLEANed images. In addition, it is shown that error estimates for the lens model parameters can be reliably made for optical and radio images using the properties of the $\\chi^2$ surface. In \\S\\ref{sec:examples} \\textsc{Lensview} is applied to the optical arc HST J15433+5352 and the radio ring MG1549+3047. In \\S\\ref{sec:conclusions} some issues relating to the software, and modelling of resolved lenses in general, are discussed and compared to some of the other methods outlined here.  Throughout this paper, we use a flat, $\\Lambda$-dominated cosmology with $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$ and $H_0$=70 km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "\\label{sec:conclusions} Resolved gravitationally lensed images provide the power to discriminate between mass models for the lensing galaxy. However, until recently this potential has been largely untapped. There is now a growing set of algorithms and tools for modelling resolved lensed images. In this section we discuss issues relating to our software, \\textsc{Lensview} and how it compares to other algorithms.  \\emph{Mapping matrix vs non mapping matrix:} The mapping matrix has many convenient properties, not the least of which is that it correctly preserves surface brightness when projecting between the source and image planes. \\citet{2003ApJ...590..673W} use ray-shooting to map source pixels to image pixels. This leads to the question: if a projected image pixel overlaps with several source pixels, which one do you choose? This is most important for image pixels which are in a medium to low magnification region. Such pixels will overlap with several source pixels and their true value has a contribution from all. By ray shooting, only a single pixel is sampled. The effects of this sampling are unknown at present as are the effects of the simple interpolation scheme proposed by \\citet{2004ApJ...611..739T}. In principle, there is no reason why a mapping matrix cannot be incorporated into the semilinear method, but the increased source-to-image pixel relationships of the mapping matrix means the matrix be to inverted by the semilinear method rapidly loses its sparseness. Obviously a mapping matrix is not useful for \\textsc{LensCLEAN} which uses point sources, although the \\textsc{LenTil} algorithm \\citep{2004MNRAS.349....1W} uses something similar to find image positions.  \\emph{Iterative vs non iterative:} The source model for our method builds up through progressive iterations of forward and reverse projection. The semilinear method calculates the source model in a single matrix inversion where the matrix has size $N \\times N$ (with $N$ pixels in the image) and is presumably fairly sparse. The `sparseness' of the matrix must be reduced with an increasingly broad PSF. This will affect running times and possibly be susceptible to accumulated numerical errors. The iterative approach is not significantly affected by the size of the PSF and numerical errors do not accumulate between iterations. The time to project between source and image planes, however, is affected by the fraction of the image plane which covered by the source plane. Nevertheless, the majority of the computational time in our method is taken convolving with the PSF. Typically 100-200 iterations are required.  \\emph{Running times:} Code execution times have not been given for other methods. As mentioned above, \\textsc{Lensview} spends most of its time performing convolutions via the FFT. Hence running times are roughly $O(N \\log N)$ where $N$ is the number of pixels in the image. For the examples in this paper, running on a 2.0GHz Pentium computer, the inner loop is completed in less than one second for J15433 and in approximately seven seconds for MG1549. Hence, large sections of parameter space can be searched in reasonable times with \\textsc{Lensview}.  \\emph{Entropy/regularisation:} The entropy constraint is analogous to regularisation of semilinear inversion although it should be emphasised that the way in which the source is reconstructed in each method is quite different. Any form of regularisation introduces bias into the source, however that is the price one must pay for using real (noisy) images since all image deconvolution methods will require regularisation in some form. The entropy constraint is an excellent form of regularisation because it smooths source pixels which have little or no data constraining them and enforces positivity in the source. Other choices for a source constraint are possible, but we have not explored them here.  \\emph{Images vs visibilities:} Although \\textsc{Lensview} currently works only with images, the extension to use radio visibilities is conceptually straightforward. It is difficult to see how the semilinear method could be extended to use radio visibilities because the matrix inversion would become impractically large- this is exactly why CLEAN and similar algorithms are used to process radio images.  A simple estimate shows that processing the data in image space provides approximately two orders of magnitude increase in the speed of analysis. Fourier transforming between the visibilities and dirty images used in the simulations takes $\\sim 1$ second. Iterative methods must be used with this volume of data, and our models (and those of SB84) usually converge in 100--200 iterations. Assuming no other computational burden, the forward and backward Fourier transforms alone will take several minutes for a single lens model for the inner loop. Conversely, by using the CLEANed image, just the region of interest can be extracted from the larger image for use with \\textsc{Lensview}.  The tests presented here show that the accuracy of the lens model parameters is biased by working in image space, however the systematic errors are small enough that the results are still useful for some applications. \\citet{1996ApJ...464..556E} found similar biases after comparing results of image-space models with visibility-space models. They therefore concluded that using visibilities is preferable to images. \\citet{2004MNRAS.349....1W} subsequently showed that further improvements to LensCLEAN were possible. Since \\citet{1996ApJ...464..556E} did not perform any simulations, it is not known whether their visibility-based results were actually affected by the problems identified by \\citet{2004MNRAS.349....1W}. Hence it is possible that systematic effects were still present.  A major advantage of using visibilities is that a well defined target $\\chi^2$ exists. In some rare cases, e.g. MG1131+0456 \\citep{1988Natur.333..537H}, the entire image generated by the interferometer is contained in a very small area --- just a few square arcseconds. In that case it would be possible to process the \\emph{dirty} image directly using the dirty beam as a PSF. Since there would be no artifacts due to image processing in that case, the noise properties and target $\\chi^2$ of the image would be well defined.   \\emph{Mass distributions:} \\textsc{Lensview} can work with an arbitrary lens mass model with negligible extra computational cost. This appears to be true also for the semilinear method, but is not so for \\textsc{LensCLEAN} which requires solutions to the lens equation for each point source component. Hence \\textsc{LensCLEAN} is substantially slower when using a non-analytic lens model.  The potential applications of \\textsc{Lensview} extend well beyond the simple examples shown in this paper. Some of the most exciting applications are: 1) determining the slope of the mass profile in lens galaxies. Interpretation of $H_0$ measurements rely on this knowledge. \\citet{2005MNRAS.360.1333W} have shown that the optical Einstein ring ER 0047-2808 has a well determined mass profile in the region of the images. 2) properties of the dark haloes of lens galaxies. It is straightforward to model a lens as a stellar (baryonic) component with a dark matter halo. Extended images will allow measurements of the properties of the halo \\citep[e.g.][]{2003ASPC..289..449G}. 3) determination of the lens centre. Again, required for $H_0$ measurements.  We have presented simple elliptical isothermal lens models for the optical lensed arc HST J15433+5352 and the radio Einstein ring MG1549+3047. Our models show that for HST J15433+5352, the data are unable to distinguish between an elliptical model and an elliptical model with external shear. We speculate that this because the source is lying over a cusp, therefore has a very generic image configuration. For MG1549+3047, we show that the previous lens model had an incorrect Einstein radius and therefore the mass enclosed in the ring was overestimated. In addition, we showed that the model requires a small non-singular core, consistent with expectations for a star-dominated mass distribution in the centre of the lens galaxy.  More detailed applications of the software are possible. We look forward to a new era of precision modelling of resolved images."
    },
    "0609/astro-ph0609297_arXiv.txt": {
        "abstract": "Black hole masses are tightly correlated with the stellar velocity dispersions of the bulges which surround them, and slightly less-well correlated with the bulge luminosity. It is common to use these correlations to estimate the expected abundance of massive black holes.  This is usually done by starting from an observed distribution of velocity dispersions or luminosities and then changing variables. This procedure neglects the fact that there is intrinsic scatter in these black hole mass--observable correlations.  Accounting for this scatter results in estimates of black hole abundances which are larger by almost an order of magnitude at masses $>10^{9}M_\\odot$. Including this scatter is particularly important for models which seek to infer quasar lifetimes and duty cycles from the local black hole mass function. However, even when scatter has been accounted for, the $M_\\bullet-\\sigma$ relation predicts fewer massive black holes than does the $M_\\bullet-L$ relation.  This is because the $\\sigma-L$ relation in the black hole samples currently available is inconsistent with that in the SDSS sample from which the distributions of $L$ or $\\sigma$ are based:  the black hole samples have smaller $L$ for a given $\\sigma$, or larger $\\sigma$ for a given $L$. The $\\sigma-L$ relation in the black hole samples is similarly discrepant with that in other samples of nearby early-type galaxies.  This suggests that current black hole samples are biased: if this is a selection rather  than physical effect, then the $M_\\bullet-\\sigma$ and $M_\\bullet-L$  relations currently in the literature are also biased from their true values. ",
        "introduction": "The abundance of supermassive black holes is the subject of considerable current interest (e.g.  Yu \\& Tremaine 2002; Marconi et al. 2004; McLure \\& Dunlop 2004; Shankar et al. 2004; Yu \\& Lu 2004; Ferrarese \\& Ford 2005). Several groups have noted that galaxy formation and supermassive black holes growth should be linked, and many have modeled the joint cosmological evolution of quasars and galaxies (see, e.g., Monaco et al. 2000; Kauffmann \\& Haehnelt 2001; Granato et al. 2001; Cavaliere \\& Vittorini 2002; Cattaneo \\& Bernardi 2003; Haiman et al. 2004; Hopkins et al. 2006; Lapi et al. 2006; Haiman et al. 2006 and references therein). Since the number of black hole detections to date is less than fifty, their abundance is estimated by using secondary indicators. In particular, $M_\\bullet$ is observed to correlate strongly and tightly with the velocity dispersion of the surrounding bulge (e.g. Ferrarese \\& Merritt 2000; Gebhardt et al. 2000; Tremaine et al. 2002). Since detecting bulges is considerably easier than detecting black holes, it has become common to estimate the abundance of black holes by combining the observed distribution of bulge velocity dispersions (e.g. Sheth et al. 2003) with the observed $M_\\bullet-\\sigma$ relation.  A crude estimate follows easily if one is willing to assume that all bulges host black holes, and that the $M_\\bullet-\\sigma$ relation has no intrinsic scatter (e.g. Yu \\& Tremaine 2002; Aller \\& Richstone 2002).    There is some discussion in the literature about whether $L$ or $\\sigma$ is a better predictor of $M_\\bullet$.  There are two parts to this statement which are not always stated explicitly. The first is the assumption that the $M_\\bullet-$observable relation is a single power law; whether this is a better approximation for $L$ than for $\\sigma$ is an open question, although Lauer et al. (2007) argue that the curvature in the $\\sigma-L$ relation for massive early-type galaxies (Bernardi et al. 2007a; Bernardi 2007) suggests that the $M_\\bullet-\\sigma$ relation is unlikely to be a single power-law. In what follows, we will assume the relations in question are indeed single power laws.  The second is the issue of the scatter around the mean relations. It is generally believed that the relation with smaller scatter provides the better estimate of the $M_\\bullet$ distribution. Indeed, Marconi et al. (2004) state that if the scatter around two relations is similar, then both relations should provide equivalent descriptions of the distribution of $M_\\bullet$. One of the goals of the present paper is to show that this is not the whole story.  Provided the intrinsic scatter around the two relations is accurately known, whether or not one relation is tighter than another is irrelevant.  (The only practical difference is that, if the intrinsic scatter is smaller, then observations of fewer objects are required to estimate it reliably.)  Both the $M_\\bullet-\\sigma$ and $M_\\bullet-L$ relations show considerable scatter, not all of which can be accounted-for by measurement errors.  Marconi \\& Hunt (2003) present evidence that the amount by which an object scatters from these relations is correlated with bulge size (half light radius), suggesting that at least some component of the scatter is intrinsic. Gebhardt et al. (2000) suggest that the intrinsic scatter in $M_\\bullet$ at fixed velocity dispersion is of order 0.25~dex, whereas scatter around $M_\\bullet-L_V$ is about 0.35~dex (e.g. Novak et al. 2006). If the intrinsic scatter is indeed this large, then it must be accounted for, especially when estimating the abundances of the most massive black holes ($M_\\bullet \\ge 10^9\\,M_\\odot$).  Section~2 describes a toy model of the effects of scatter which shows that, (i) if intrinsic scatter is ignored, then both the $L$- and $\\sigma$-based predictions will underestimate the true abundance of the most massive black holes; (ii) the observable which correlates most tightly with $M_\\bullet$ will provide the best estimate of the true abundance of the most massive black holes; (iii) if scatter has been correctly accounted for, $\\sigma$- and $L$-based predictors of $M_\\bullet$ abundances should give the same answer. It then shows the $M_\\bullet-\\sigma$, $M_\\bullet-L$ and $\\sigma-L$ correlations, their scatter, and how we use them to estimate black hole abundances.  A direct comparison of the luminosity and velocity dispersion based predictors is provided, both when intrinsic scatter in these relations is accounted for and when it is ignored.  We find that, if scatter is ignored, then the $L$-based method predicts substantially more $10^9 M_\\odot$ objects than does the $\\sigma$-based method.  The toy model suggests that this may be a consequence of ignoring the intrinsic scatter.  However, accounting for this scatter does not eliminate this discrepancy, suggesting that there may be a more serious inconsistency. Section~3 identifies the reason for this discrepancy with the fact that the $\\sigma-L$ relation in the SDSS/Bernardi et al. (2003a) (hereafter SDSS-B07; see Bernardi 2007 for the definition of the SDSS-B07 sample), from which the $L$ and $\\sigma$ distributions are drawn, is rather different from that in the black hole samples, from which the $M_\\bullet-\\sigma$ and $M_\\bullet-L$ relations are derived. A more detailed analysis of the role of selection effects in the $M_\\bullet$ sample is presented in Bernardi et al. (2007b). A final section discusses our findings and summarizes our conclusions. A standard flat $\\Lambda$CDM comological model has been used, with $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$ and $H_0=70$~km~s$^{-1}$~Mpc$^{-1}$. ",
        "conclusions": "It is common to estimate the abundance of supermassive black holes by combining observed correlations between $M_\\bullet$ and bulge luminosity or velocity dispersion, calibrated from relatively small samples, with luminosity or velocity dispersion functions determined from larger samples.  However, the $\\langle M_\\bullet|\\sigma\\rangle$ and $\\langle M_\\bullet|L\\rangle$ relations have intrinsic scatter of about 0.22 and 0.33~dex (Appendix). Accounting for this results in considerably increased estimates of the abundance of black holes with $M_\\bullet \\ge 10^9\\,M_\\odot$, compared to naive estimates which ignore this scatter. Doing so is at least as important as correcting the luminosity function for the fact that the most luminous galaxies have non-standard light profiles (Figure~\\ref{lbulge2bh}). Once this scatter has been accounted for, the $\\sigma$-based estimates of $\\phi(M_\\bullet)$ are in reasonably good agreement with models, such as that of Hopkins et al. (2006), which relate previous QSO and AGN activity to the local black hole mass function. The luminosity-based estimates, on the other hand, are substantially in excess of this model.  These results follow from using a single power-law to parametrize the $M_\\bullet-\\sigma$ and $M_\\bullet-L$ relations.  While this may be too simplistic, this parametrization is not the primary reason why the $L-$ and $\\sigma-$based approaches yield different predictions for black hole abundances. The main cause of the discrepancy is that the $\\sigma-L$ correlation in black hole samples is different from that in the samples from which the luminosity and velocity functions are drawn:  the black hole samples have larger $\\sigma$ for a given $L$ compared to the ENEAR or SDSS-B07 samples or have smaller $L$ for a given $\\sigma$ (Bernardi et al. 2007b). If this is a physical effect, then it compromises the fundamental assumption of black hole demographic studies---that all galaxies host black holes.  If, on the other hand, it is a selection effect, then the $M_\\bullet-\\sigma$ and $M_\\bullet-L$ relations currently in the literature are biased compared to the true relations, making current estimates of black hole abundances unreliable. If black hole masses correlate with bulge luminosity only because of the $M_\\bullet-\\sigma$ and $\\sigma-L$ relations, then the bias in the $\\sigma-L$ relation is not important only if one is using $\\phi(\\sigma)$ to infer black hole abundances:  the $\\phi(L)$-based estimate may be strongly affected.  Identifying the source of the bias is complicated. Residuals from the size-luminosity relation are anti-correlated with residuals from the $\\sigma-L$ relation, as might be expected from the virial theorem (Bernardi et al. 2003b). If the stellar kinematics method of measuring black hole masses favors objects with high surface brightnesses, then one might expect smaller sizes and larger $\\sigma$ at constant $L$:  this would produce a bias in the sense we see. On the other hand, it might be more difficult to measure the influence of the black hole on stellar kinematics if the stellar velocity dispersion is already abnormally large---this would produce a bias in the opposite sense.  Whether either or these effects has played a role in the selection of black samples is an open question. See Bernardi et al. (2007b) for further study along these lines."
    },
    "0609/astro-ph0609774_arXiv.txt": {
        "abstract": "{  We present high-resolution {\\'e}chelle spectroscopy, obtained with the UVES spectrograph on ESO/VLT, of two luminous star clusters in the metal-poor blue compact galaxy ESO\\,338--IG04 at a distance of 37.5 Mpc. Cross-correlating with template stars, we obtain line-of-sight velocity dispersions of 33 and 17 \\kms. By combining with size estimates from Hubble Space Telescope images we infer dynamical masses of $1.3\\times10^7$ \\msun\\ and $4.0\\times10^6$ \\msun for the two clusters, making them among the most massive known. The less massive cluster is the faintest cluster for which a dynamical mass has yet been obtained. In both clusters we detect Balmer absorption lines which we use to estimate their ages.  From the younger ($\\sim 6$ Myr) and more massive cluster, we detect \\heii\\ $\\lambda$4686 emission of intermediate width, indicating the presence of very massive O-stars. Moreover, analysis of the \\oiii\\ $\\lambda$5007 and \\ha\\ emission lines from the region near the younger cluster indicates that it is associated with a bubble expanding at $\\sim40$ \\kms. We also see from the \\nai\\,D absorption lines indications of neutral gas flows towards the younger cluster.  We compare the dynamical masses with those derived from photometry and discuss implications for the stellar initial mass function (IMF) in each cluster. Both clusters are compatible with rather normal IMFs which will favour their long-term survival and evolution into massive bona fide globular clusters. ",
        "introduction": "Super star clusters (SSCs) are important sites for star formation in starburst galaxies (Arp and Sandage \\cite{arp}, Meurer \\etal\\ \\cite{meurer}), and may be the progenitors of globular clusters (GCs). Different observational techniques have been used to investigate the nature and history of these objects. High resolution imaging, notably with the Hubble Space Telescope (HST), has been used to collect photometry and to estimate the sizes of many SSCs (Whitmore \\cite{whitmore_rev}).  Photometry, in combination with spectral evolutionary synthesis models, has also been used to estimate cluster ages and masses ({\\\"O}stlin \\etal\\ \\cite{OBR}).  While many such {\\em photometric} masses have been determined, only a small fraction of extragalactic SSCs and GCs have {\\em dynamical} mass estimates.  Taken together, measurements of the photometric and dynamical masses can allow constraints on the stellar initial mass function (IMF) and the dynamical state in these objects.  The spectrum of a cluster can provide other information about the stellar population.  For example, Gonz{\\'a}lez Delgado \\etal\\ (\\cite{GD}) inferred a population of Wolf-Rayet stars in the massive cluster NGC 1569-A, and Origlia et al.\\ (\\cite{O01}) used ultraviolet spectra to show that the same cluster has a population of O-stars.  High-resolution spectra can be used to measure the velocity dispersion, and, assuming virialisation, to estimate dynamical masses for SSCs and GCs.  Ho and Filippenko (\\cite{HFa}, \\cite{HFb}) were first to use the cross-correlation method described by Simkin (\\cite{Simkin}) and Tonry and Davis (\\cite{TD}) to measure the velocity dispersions of NGC 1569-A (giving a kinematic mass of $3.3\\times10^5$ \\msun) and the SSC in NGC 1705 ($8.2\\times10^4$ \\msun).  The same method has been applied to M82-F by Smith and Gallagher (\\cite{SG}; $1.2\\times10^6$ \\msun), to five clusters in NGC 4214 and NGC 6946 by Larsen \\etal\\ (\\cite{Larsen01}, \\cite{Larsen04}; $0.2-1.8\\times10^6$ \\msun), by Gilbert and Graham (\\cite{GG}) to clusters A1, A2 and B in NGC 1569 ($3.9\\times10^5$ \\msun, $4.4\\times10^5$ \\msun \\ and $2.3\\times10^5$ \\msun, respectively) and by McCrady, Gilbert and Graham (\\cite{MGG}) to two SSCs in M82 ($1.5\\times10^6$ \\msun\\ and $3.5\\times10^5$ \\msun).  Mengel et al.\\ (\\cite{M}) measured the dynamical masses of five SSCs in NGC 4038/4039 ($0.65 - 4.7\\times10^6$ \\msun) by fitting model spectra in the near-IR and optical.  Maraston et al. (2004) presented a dynamical mass estimate of as much as $8\\times10^7$\\msun\\ for the cluster W3 in NGC~7252, which may possibly be a dwarf galaxy.  Bastian \\etal\\ (\\cite{Bastian}) found masses of $1.6\\times10^7$ \\msun\\ for both W30 in NGC 7252 and G114 in NGC 1316, and present a compilation of previously published results (their Table 6).  A possibly related class of objects are the nuclear star clusters found in many bulgeless spiral galaxies (e.g. Walcher \\etal\\ \\cite{walcher}), some of which appear to be very massive but which probably have more extended star formation histories than GCs and SSCs.  ESO\\,338--IG04, also known as Tol\\,1924-416, is a luminous ($M_B=-19$) blue compact galaxy (BCG) --- the closest in a class of galaxies that are rare in the local universe but that become increasingly important towards higher redshifts and the peak in the cosmic star formation rate at $z\\sim 1$ (\\eg\\ Werk \\etal\\ \\cite{Werk}). ESO\\,338--IG04 shows both vigorous star formation and a rich population of more than 100 young SSCs and old globular clusters ({\\\"O}stlin \\etal\\ \\cite{OBR})\\footnote{Following {\\\"O}stlin \\etal\\ (\\cite{OBR}) we adopt a distance of 37.5 Mpc to ESO\\,338--IG04.}. Analysis of the age distribution of these clusters shows evidence that at least one strong starburst occurred in the past, a couple of Gyr ago, and the present starburst has probably been active for about 40 Myr ({\\\"O}stlin \\etal\\ \\cite{OZBR}). In addition ESO\\,338--IG04 hosts a population of cosmologically old (age $\\le 10$~Gyr) globular clusters. What makes ESO\\,338--IG04 particularly interesting as a star cluster formation laboratory is its combination of small reddening, many clusters, and a low metallicity ($12+\\log({\\rm O/H})=8.0$). Compared to other metal-poor star-forming galaxies in the local universe (such as NGC\\,1569, NGC\\,1705, and He\\,2-10) the clusters in ESO\\,338--IG04 are more luminous and numerous.  The present starburst has been triggered by a small merger or interaction with its companion galaxy ({\\\"O}stlin \\etal\\ \\cite{ostlin01}, Cannon \\etal\\ 2004). Hence ESO\\,338--IG04 offers the possibility of directly studying globular cluster formation associated with hierarchical galaxy evolution, at a low metallicity.  In this paper, we report high-dispersion {\\'e}chelle spectroscopy of two of these clusters, labelled Inner-\\#23 and Outer-\\#34 by {\\\"O}stlin et al. (\\cite{OBR}).  Hereafter, we refer to them simply as \\#23 and \\#34. Cluster \\#23, a young blue cluster situated in the centre of the starburst region, is the most luminous in the whole galaxy ($M_V=-15.5$).  Cluster \\#34 is of intermediate age, $\\sim 1$ Gyr, but still very luminous ($M_V=-12.8$), which led {\\\"O}stlin \\etal\\ (\\cite{OBR}) to suspect a mass in excess of $10^7$ \\msun.  Here, we use high-dispersion VLT/UVES spectroscopy to make dynamical mass estimates and new age estimates for both clusters. These two clusters were selected since they represent the most luminous of the young and intermediate age clusters, respectively, and could be observed in a single UVES slit.  The outline of the rest of the paper is as follows. In Sect.\\ 2 we describe the observations and reductions.  In Sect.\\ 3 we show the results obtained; subsection 3.4 in particular describes how we obtained estimates of the velocity dispersions of the clusters through cross-correlation techniques. In Sect.\\ 4 we discuss the inferred dynamical masses and compare with results derived from photometry. Sect.\\ 5 summarises our conclusions. ",
        "conclusions": "We have presented high resolution spectra of two luminous star clusters in the luminous blue compact galaxy ESO\\,338--IG04 (Tololo\\,1924--416).  The spectra have been cross-correlated with template stars observed using the same setup to determine each cluster's line-of-sight velocity dispersion.  Using size estimates from Hubble Space Telescope images, we have used the velocity dispersions to determine the virial masses of the clusters.  Our mass estimates have been compared to masses derived from HST photometry fitted to spectral evolutionary synthesis models. This comparison indicates that both clusters have rather normal IMFs, favouring their survival against internal disruption mechanisms.  One of the clusters (\\#23) is young, with an age of $\\sim6-7$ Myr and mass in excess of $10^7$ \\msun, making it one of the most massive very young clusters known. We find evidence for the simultaneous presence of massive O-stars and red supergiants. The \\oiii\\ $\\lambda$5007 and \\ha\\ emission line profiles from the region surrounding cluster \\#23 indicate the presence of a bubble expanding at a velocity $\\sim40$ \\kms.  In addition, we see signs of neutral gas flows along the line of sight from the cross-correlation analysis of the Na~{\\sc i}\\,D absorption lines.  The IMF of this cluster shows no evidence for any deficiency of low-mass stars, but is consistent with a Salpeter slope over the whole mass range 0.08--120 \\msun. The inferred dynamical mass is larger than expected based on the luminosity of the cluster, which may indicate that the cluster is not yet virialised.  The other cluster (\\#34) is older and has a mass of $4\\times 10^6$ \\msun.  A combined analysis with the photometric data suggests that it has an age of between 0.3 and 1.4 Gyr.  The slope of the IMF at masses less than 1 \\msun\\ must be flatter than the Salpeter value for the virial and photometric masses to be consistent.  Such an IMF is in agreement with that observed for old Galactic globular clusters. This work confirms that this cluster is a bona-fide young globular cluster."
    },
    "0609/astro-ph0609404_arXiv.txt": {
        "abstract": "Thanks to \\int's long exposures of the Galactic Plane, the two brightest Soft Gamma-Ray Repeaters, SGR 1806-20 and SGR 1900+14, have been monitored and studied in detail for the first time at hard-X/soft gamma rays.  This has produced a wealth of new scientific results, which we will review here. Since SGR 1806-20 was particularly active during the last two years, more than 300 short bursts have been observed with \\int~ and their characteristics have been studied with unprecedented sensitivity in the 15-200 keV range. A hardness-intensity anticorrelation within the bursts has been discovered and the overall Number-Intensity distribution of the bursts has been determined. In addition, a particularly active state, during which ~100 bursts were emitted in ~10 minutes, has been observed on October 5 2004, indicating that the source activity was rapidly increasing. This eventually led to the Giant Flare of December 27th 2004, for which a possible soft gamma-ray ($>$80 keV) early afterglow has been detected.  The deep observations allowed us to discover the persistent emission in hard X-rays (20-150 keV) from 1806-20 and 1900+14, the latter being in a quiescent state, and to directly compare the spectral characteristics of all Magnetars (two SGRs and three Anomalous X-ray Pulsars) detected with \\int. ",
        "introduction": "\\label{intro} Soft gamma-ray repeaters (SGRs, for a recent review see \\citet{woodsrew}) are a small group (4--7) of peculiar high-energy sources generally interpreted as ``magnetars'', i.e. strongly magnetised ($B\\sim$10$^{15}$ G), slowly rotating ($P\\sim$ 5-8 s) neutron stars powered by the decay of the magnetic field energy, rather than by rotation \\citep{dt92,pac92,td95}. They were discovered through the detection of recurrent short ($\\sim$0.1 s) bursts of high-energy  radiation in the tens to $\\sim$hundred keV energy range, with peak luminosity up to 10$^{39}$-10$^{42}$ erg s$^{-1}$, above the Eddington limit for neutron stars. The rate of burst emission in SGRs is highly variable. Bursts are generally emitted during sporadic periods of activity, lasting days to months, followed by long ``quiescent'' time intervals (up to years or decades) during which no bursts are emitted. Occasionally SGRs emit ``giant flares'', that last up to a few hundred seconds and have peak luminosity up to  10$^{46}$-10$^{47}$ erg s$^{-1}$. Only three giant flares have been observed to date, each one from a different source (see, e.g., \\cite{mazets} for 0526--66, \\cite{hurley1999} for 1900+14, \\cite{swiftgiant,acsgiant,rhessigiant} for 1806--20).  Persistent (i.e. non-bursting) emission is also observed from SGRs in the soft X--ray range ($<$10 keV), with a typical luminosity of $\\sim$10$^{35}$ erg s$^{-1}$, and, in three cases, periodic pulsations with periods of 5 -- 8 seconds have been detected. Such pulsations proved the neutron star nature of SGRs and allowed the derivation of spin-down at rates of $\\sim$10$^{-10}$ s s$^{-1}$, consistent with dipole radiation losses for magnetic fields of the order of B$\\sim$10$^{14}$-10$^{15}$ G. The X--ray spectra are generally described with absorbed power laws, but in some cases strong evidence has been found for the presence of an additional blackbody-like component with a typical temperature of $\\sim$0.5 keV \\citep{xmm}.  Over the last few years the \\int~ satellite \\citep{integral}, launched in 2002 and operating in the 15 keV-10 MeV energy range, has provided a wealth of new results concerning the two brightest SGRs, 1806--20 and 1900+14. Most aspects concerning the SGRs, short bursts, giant flares, and persistent emission, have been investigated, and new results have been found for each of them. We will review them here. ",
        "conclusions": "Thanks to \\int, and in particular to its imager IBIS, we have been able to study most of the magnetars' phenomenology with unprecedented sensitivity at high energies. One of the most striking results is the discovery, which was particularly unexpected for AXPs, of the persistent hard X-ray emission. This discovery, which can be considered one of the most important \\int~ results at all, represents a new important input for theoreticians who started to include it in the magnetar model (see e.g. \\citet{belo}).  Also, the fact that short bursts evolve with time is a new feature that has to be considered with care within the magnetar model: up to now no clear explanation has been provided for this.  The large number of detected short bursts from \\zerosei~ allowed us a good determination of the shape and slope of their Number-Intensity distribution, showing that a single power law holds over 2.5 orders of magnitude.  In addition, the fact that \\zerosei~ has been particularly active in these last years, also emitting a once-in-a-lifetime event such as the giant flare (and its possible high energy afterglow), has allowed observations of relatively rapid changes of the bursting and persistent emission of a Magnetar and to interpret then with the evolution of a very strong and complicated magnetic field, confirming the magnetic field as the dominant source of energy in Soft Gamma-Ray Repeaters and Anomalous X-ray Pulsars."
    },
    "0609/astro-ph0609318_arXiv.txt": {
        "abstract": "% This paper summarizes the conclusions of an extensive investigation of variable carbon AGB stars in our Galaxy. The zero-point of the period- $M_{bol}$ relation for Galactic C-Miras is found to be close to that in the LMC. The mean age of Galactic C-Miras is $\\sim$ 1.8 Gyr and their initial masses $\\sim 1.8 M_{\\odot}$ with some evidence that age decreases and initial mass increases with increasing pulsation period. Kinematic studies of Galactic C- and O-Miras show that the relative frequency of these two Miras classes depends on the age of the parent population. The lack of carbon stars in the Galactic Bulge seems to be due to a high oxygen abundance. The velocity dispersion of O-Miras allows one to distinguish between two possible models  of Galactic kinematics. ",
        "introduction": "The kinematics of the large amplitude, oxygen-rich AGB variables known as O-Miras have long been known to give valuable clues to the ages and evolution of AGB stars as well as being important for Galactic structure studies (e.g. Feast 1963; Feast \\& Whitelock 2000). Until recently the Galactic carbon-rich Miras (C-Miras) have been much less intensively studied. The present paper summarizes the results of a kinematic and photometric study of Galactic C-Miras leading to a determination of their absolute magnitudes and ages. It then goes on to discuss what C- and O-Miras tell us about Galactic structure and composition. The paper is based on three recently published papers (Whitelock et al. 2006; Menzies, Feast, \\& Whitelock 2006; Feast, Whitelock, \\& Menzies 2006).  It has been known for a long time that an optically selected sample of C-Miras in the LMC shows a well defined $M_{bol}$ - log\\,P relation (Feast et al. 1989). It is now possible to extend this relation to longer periods using data for obscured C-Miras (Whitelock et al. 2003). The relation which extends from about 160 to 1000 days has the form \\begin{equation} M_{bol} = -2.54 \\log\\,P + 1.87 \\end{equation} where the distance modulus of the LMC has been taken to be 18.50mag. The scatter about this relation is only 0.17mag. One aim of the present work was to derive the zero-point of this relation for Galactic C-Miras. ",
        "conclusions": "1. Assuming the $M_{bol}$ - period relation for Galactic C-Miras has the same slope as for those in the LMC, the zero-points in the two systems  are the same within the uncertainties\\\\ 2. Galactic C-Miras have ages of $\\sim 1.8$ Gyr and initial masses of $\\sim 1.8 M_{\\odot}$.\\\\ 3. There is some evidence that the ages of C-Miras decrease with increasing period.\\\\ 4. The ratio of C- to O-Miras in a population depends on the age distribution of that population.\\\\ 5. The lack of carbon stars in the Galactic Bulge is connected to an oxygen overabundance.\\\\ 6. The Galactic O-Miras favour the interpretation of Galactic kinematics of Nordstr\\\"om et al."
    },
    "0609/astro-ph0609632_arXiv.txt": {
        "abstract": "{Sunspot penumbrae harbor highly structured magnetic fields and flows. The moving flux tube model offers an explanation for several observed phenomena, e.g.\\ the Evershed effect and bright penumbral grains. A wealth of information can be extracted from spectropolarimetric observations. In order to deduce the structure of the magnetic field in sunspot penumbrae, detailed forward modeling is necessary. On the one hand, it gives insight into the sensitivity of various spectral lines to different physical scenarios. On the other hand, it is a very useful tool to guide inversion techniques. In this work, we present a generalized 3D geometrical model that embeds an arbitrarily shaped flux tube in a stratified magnetized atmosphere. The new semi-analytical geometric model serves as a frontend for a polarized radiative transfer code. The advantage of this model is that it preserves the discontinuities of the physical parameters across the flux tube boundaries. This is important for the detailed shape of the emerging Stokes Profiles and the resulting net circular polarization (NCP). (a) The inclination of downflows in the outer penumbra must be shallower than approximately $15^\\circ$. (b) Observing the limb-side NCP of sunspots in the \\ion{Fe}{1}~1564.8\\,nm line offers a promising way to identify a reduced magnetic field strength in flow channels. (c) The choice of the background atmosphere can significantly influence the shape of the Stokes profiles, but does not change the global characteristics of the resulting NCP curves for the tested atmospheric models. } ",
        "introduction": "Spectral  line profiles in the penumbra are characterized by line shifts and line asymmetries of their intensity (Stokes I) as well as in polarized light (Stokes Q, U, and V). A suitable measure of the asymmetry of Stokes-$V$ profiles is the area asymmetry or \\emph{net circular polarization}, NCP, of a spectral line, which we define as \\be {\\rm NCP} \\equiv \\int\\limits_{\\delta \\lambda} V(\\lambda)\\, {\\rm d} \\lambda \\, , \\ee where the interval of integration, $\\delta \\lambda$, encompasses the whole line profile. This quantity has the advantage that it is independent of the spatial resolution of the observations since a convolution of a line profile with a Gaussian preserves its area. The first measurements of the net circular polarization in sunspots were reported by \\cite{Illing+al1975}. \\cite{Auer+Heasley1978} showed that these observations could be explained by assuming macroscopic velocity fields. Furthermore, they proved that velocity gradients are a necessary \\emph{and} sufficient condition to produce a NCP.  The advent of modern polarimeters like the Advanced Stokes Polarimeter \\citep[ASP, ][]{Skumanich+al1997ApJS} at the NSO Vacuum Tower Telescope in Sunspot, New Mexico, and the Tenerife Infrared Polarimeter \\citep[TIP, ][]{MartinezPillet+al1999ASP} and the Polarimetric Littrow  Spectrometer \\citep[POLIS, ][]{Beck2005AA_POLIS} at the Vacuum Tower Telescope (VTT) on Tenerife has brought new challenges to observers as well as theoreticians. It is now possible to routinely record maps of the full Stokes vector for a whole sunspot. This facilitates the study of spectral signatures of magnetic flux tubes in the penumbra by analysis of maps of the NCP.  So far, the spectral synthesis has been limited to atmospheric models based on snapshots of the moving tube model simulations by \\cite{Schlichenmaier1998} and to configurations in a two-dimensional plane (``slab models''). Such slab models have been used earlier by \\citet{solanki+montavon1993}, \\citet{Sanchez+al1996} and \\citet{MartinezPillet2000}, and led to the picture of the ``uncombed'' penumbra: Approximately horizontal magnetic flux tubes are embedded in a more inclined background field that is at rest. The sharp gradient or even discontinuity of the velocity of the plasma between the inside of the tube and the static surroundings result in the observed NCP. While two-dimensional models reproduced the observed NCP and its center-to-limb variation, they are limited to the few geometric cases where a flux tube is exactly aligned with the line-of-sight  (LOS). Thus, they do not give any information about the NCP at different positions within the penumbra.  We showed \\citep[][hereafter M02]{ncpletter2002, Mueller+al2002} that a wealth of additional information is contained in NCP maps, which can be retrieved with the help of a three-dimensional model. Discontinuities in the azimuth of the magnetic field (as measured in the observer's coordinate system) lead to a characteristic pattern of the NCP within the penumbra. Specifically, it was shown that the apparent symmetry within the spot is broken due to the following fact: The discontinuities in the magnetic field azimuth are different for two locations that are symmetrically located in the spot with respect to the ``line-of-symmetry'', i.e.\\ the line connecting disk-center and sunspot center (Fig.~\\ref{gfx_los}). These azimuth discontinuities have been confirmed observationally by \\citet{BellotRubio+al2004AA}.  \\begin{figure}[ht] \\begin{center} \\resizebox{0.8\\hsize}{!}{ \\includegraphics{6062f01.eps} } \\end{center} \\caption{\\label{gfx_los}Pictorial representation of the difference in azimuth, $\\Delta \\phi(\\psi)$, between the magnetic field vector inside a flux tube (solid arrows) and that of the background field (dashed arrows). The tubes are nearly horizontal, while the field vector of the background has a steeper inclination, but is also in the plane spanned by the spot axis and the flux tube axis. } \\end{figure}  The impact of this geometric effect on the NCP strongly varies between spectral lines at different wavelengths. A comparison of the NCP of different spectral lines can thus be used to diagnose the magnetic field inclination and orientation of flow channels in the penumbra. As an example, we consider the two iron lines \\ion{Fe}{1}~1564.8\\,nm and \\ion{Fe}{1}~630.25\\,nm. Both lines are Zeeman triplets and have similar Land\\'e factors of g=3 and g=2.5, respectively. It turns out that the infrared line (\\ion{Fe}{1}~1564.8\\,nm) is far more susceptible to the symmetry breaking than the visible line (for details, cf.\\ M02). ",
        "conclusions": "We have presented a generalized geometrical model that embeds an arbitrarily shaped flux tube in a stratified magnetized atmosphere. The new model is a versatile tool to calculate the spectral signature of flux tubes in the penumbra and especially to make predictions about the flow speed and tube inclination from observed maps of the net circular polarization. From the first applications of the new model, we find that \\begin{itemize} \\item[(a)] The inclination of downflows in the outer penumbra must be shallower than approx. 15$^\\circ$. \\item[(b)] Observing the limb-side NCP of sunspots in the \\ion{Fe}{1}~1564.8\\,nm line offers a promising way to identify a reduced magnetic field strength in flow channels. \\item[(c)] The choice of the background atmosphere can significantly influence the shape of the Stokes profiles, but does not change the global characteristics of the resulting NCP curves for the tested atmospheric models. \\end{itemize}  \\noindent In future work, we will compare synthetic NCP maps obtained from this forward-modeling approach with results from Stokes inversion maps of observational data in more detail. This will provide further insight into the structure of penumbral filaments and guide improvements in both modeling and inversion techniques."
    },
    "0609/astro-ph0609154_arXiv.txt": {
        "abstract": "Exploiting a five-year span of data, we present improved timing solutions for the five millisecond pulsars known in the globular cluster NGC\\,6752. They include proper motion determinations for the two outermost pulsars in the cluster, PSR\\,J1910-5959A and PSR\\,J1910-5959C. The values of the proper motions are in agreement with each other within current uncertainties, but do not match (at 4$\\sigma$ and 2$\\sigma$ level respectively) with the value of the proper motion of the entire globular cluster derived in the optical band.  Implications of these results for the cluster membership of the two pulsars are investigated. Prospects for the detection of the Shapiro delay in the binary system J1910-5959A are also discussed. ",
        "introduction": "The globular cluster (GC) NGC\\,6752 is known to host five millisecond pulsars (MSPs) \\citep[][hereafter Paper\\,I]{dlm+01,dpf+02}.  PSR J1910$-$5959B and PSR\\,J1910$-$5959E (hereafter PSR\\,B and E respectively) reside in the central region of the GC and show large negative $\\dot{P}$ values, which are interpreted as an effect of the GC gravitational potential well (Paper\\,I). This in turn implies a large mass-to-light ratio in the central region of NGC\\,6752. \\citet{fps+03} recalculated the center of gravity and studied the luminosity profile of this cluster: combining their HST data with the $\\dot{P}$ value of PSR\\,B and PSR\\,E, they put a firm lower limit on the central mass-to-light ratio of $M/L_{V}\\,\\gapp\\,5.5\\,{\\rm M_\\odot}/{\\rm L_\\odot}$. Also PSR\\,J1910$-$5959D (PSR\\,D) is located close to the GC center. Its $\\dot{P}$ value is positive and of the same order of magnitude of PSR\\,B and E, suggesting that also for PSR\\,D the $\\dot{P}$ value is dominated by the gravitational potential well (Paper\\,I).  PSR\\,J1910$-$5959C\\footnote{Note that, to conform with currently accepted practice, all pulsars associated with the cluster have been given a J2000 name with the same rounded coordinates, corresponding approximately to the cluster centre.} (PSR\\,C) is located at a projected distance $\\theta_{\\perp}\\,=\\,2\\farcm6$ from the GC center (Paper\\,I), which is much larger than the cluster's core radius $r_c\\,=\\,5\\farcs2\\pm 2\\farcs4$ \\citep{fps+03}. The binary pulsar J1910$-$5959A (PSR\\,A) is located at an even larger distance from the GC center ($\\theta_{\\perp}\\,=\\,6\\farcm4$, Paper\\,I), the largest offset known for a GC pulsar.  The positions of PSR\\,A and PSR\\,C are unexpected since mass segregation should have driven the two neutron stars close to the GC center in a time scale ($\\lapp$\\,1\\,Gyr) much shorter than the time since their formation ($\\sim$\\,10\\,Gyr). In particular, \\citet{cpg02,cmp03} explored various scenarios to explain the unusual position of PSR\\,A, invoking a dynamical encounter in the inner region of the GC. The most probable picture is that PSR\\,A was originally in the GC central regions and it has been expelled to the outskirts by the interaction with either a single massive black hole (BH) or a binary [BH\\,+\\,BH] of unequal mass. Timing results in Paper\\,I indicated a low-mass white dwarf as the most probable companion for PSR\\,A. This has been confirmed by \\citet{bvkh03} and \\citet{fpsn03}, who identified with Hubble Space Telescope observations the companion of PSR\\,A with a helium white dwarf star whose mass is $M_{\\it co}\\simeq\\,0.17-0.20\\,{\\rm M_\\odot}$ and whose photometric properties are compatible with its belonging to NGC\\,6752.  The issue of the association of PSR\\,A to NGC\\,6752 has been recently revisited using spectroscopic observations of the optical companion to the pulsar, performed with the ESO-VLT.  \\citet{cfpd06} found full agreement (at $1\\sigma$) between the radial velocity of the center of mass of the binary $\\gamma=-28.1\\pm4.9$ km s$^{-1}$ and the overall cluster radial velocity $v_{\\rm 6752}=-27.9\\pm0.8$ km s$^{-1}$, obtained by \\citet{har96} (catalog revision 2003) averaging various determinations.  This is a strong indication in favour of the association of the pulsar with NGC\\,6752. However, using the same data set, \\citet{bkkv06} compared the systemic velocity of the binary with that of nearby stars which certainly belong to the cluster and concluded that they are only marginally consistent at $2\\sigma$ level.  In this paper we present timing results based on more than five years of regular observations. In particular, with a much longer available data span we have been able to measure proper motions of PSR\\,A and PSR\\,C. The new timing solutions as well as the pulse profiles for all the millisecond pulsars are presented in \\S\\,2, \\S\\,3 reports on the proper motion determinations and implications for the cluster membership of the two pulsars are discussed in \\S\\,4. ",
        "conclusions": "Since the escape velocity from a globular cluster is usually significantly lower than the typical transverse velocity of these stellar systems, it is expected that the proper motion of a cluster pulsar will largely reflect the overall motion of the cluster. For NGC\\,6752, the escape velocity from the central region is $\\sim$\\,30\\,km\\,s$^{-1}$ \\citep{cmp03} and the space velocity is $\\sim$\\,62\\,km\\,s$^{-1}$ with respect to the Solar System barycenter, based on the proper motion measurement by \\citet{dga99} and the distance derived from the distance modulus \\citep{gbc+03}. Observations over a much longer data span may reveal the peculiar (orbital) motion of a pulsar in the cluster's gravitational potential well.  Is it possible that the discrepancy between the proper motions of PSR\\,A and PSR\\,C and the optical proper motion of NGC\\,6752 (\\S\\,3) could be an indication that the two pulsars are not associated to the globular cluster?  In Paper I it was estimated that the probability\\footnote{It is worth nothing that this probability does not account for the similar values of the dispersion measure of PSR\\,A and PSR\\,C. Given the uncertainty in the Galactic electron layer, it is difficult to quantify the probability for this coincidence \\citep{bkkv06}. However, it certainly further decreases the total probability for a chance superposition.} for PSR\\,A to be a Galactic field millisecond pulsar superposed by chance to NGC\\,6752 (at a distance of $6\\farcm4$ off its center) is of order $10^{-5}.$ The compatibility of the measured proper motions of PSR\\,A and PSR\\,C reinforces the unlikeliness for these two MSP to be Galactic field objects by chance superimposed to the globular cluster.  Assuming that both PSRs A and C are members of NGC\\,6752, the discrepancy between pulsar and GC proper motions, measured in the radio and optical band respectively, may result from the different methods used for determining the proper motions in the two spectral bands.  In fact, similar discrepancies have already been noted for the pulsars in 47\\,Tucanae \\citep{fcl+01,fck+03}, in M4 \\citep{tacl99} and, more recently, in M15 \\citep{jcj+06}.  However, the discrepancies in these clusters may not easily be ascribed to a common systematic effect affecting all the optical measurements. The optical proper motion for 47\\,Tucanae was directly measured based on Hipparcos observations. The proper motion determination for M4 \\citep{ch93} was based on the determination of its motion relative to a set of reference field stars whose proper motion relative to the Sun has been in turn obtained by combining their position in the Galaxy (through their parallax) to a dynamic model for the nearby regions of the Galaxy where these reference stars reside.  Finally, the proper motion for NGC\\,6752 \\citep{dga+97} resulted from the comparison of two photographic plates taken 25 years apart, using distant field galaxies as reference objects. In the case of M15, four different optical determinations have been performed \\citep{gclo93,soh+96,obgt97,ch93}, three of which are incompatible with the pulsar proper motions. These three non-matching measures were respectively based on the comparison between photographic plates at different epochs \\citep{gclo93,soh+96} and the use of reference stars from Hipparcos observations \\citep{obgt97}. Only the measurement by \\citet{ch93} is in agreement with the apparent motions of the three pulsars investigated. However it is worth noting that \\citet{ch93} measured the optical proper motion for M15 applying the same method as was used for M4, which in that case led to a discrepant proper motion.  The discrepancy may be alternatively ascribed to very fast peculiar motions of PSR\\,A and PSR\\,C inside the cluster gravitational potential well. Assuming that both the pulsar proper motions and the optical proper motion of the cluster are correct, the relative 2-D velocity vectors of the pulsars with respect to the cluster center are roughly directed towards the cluster inner regions, as is indicated in Figure \\ref{fig:pm-map}. This would mean that [PSR\\,A+WD] cannot be now in the phase of ejection from the cluster and that it is not at its farthest distance from the GC center along its orbit inside the cluster gravitational well.  For $d\\,=\\,4.45\\pm0.15$\\,kpc \\citep{gbc+03} the relative transverse speed of PSR\\,A would be $V_{\\rm rel,A}\\,=\\,51\\pm16$\\,km\\,s$^{-1}$. NGC\\,6752 can provide a gravitational pull strong enough to retain PSR\\,A at its actual location with a peculiar velocity $V_{\\rm rel,A}$ only if the mass enclosed within the pulsar projected position is $M_{\\rm encl}\\,\\geq\\,1.18\\times10^{6}\\,{\\rm M_\\odot}.$ This is in contrast to the total mass value for the cluster obtained with HST observation \\citep{sfs+04}, which is lower by an order of magnitude. Using again the distance modulus in \\citet{gbc+03}, the apparent magnitude given by \\citet{har96} and the colour excess $E_{\\rm B-V}\\,=\\,0.04$ in \\citet{fmf+99}, the resulting overall mass-to-light ratio would be $M_{\\rm encl}/L\\,\\geq\\,8.4\\,{\\rm M_\\odot}/{\\rm L_\\odot}$, which is unreasonably high for a GC, unless we admit an initial mass function much flatter than usually estimated (so producing a very large number of under-luminous stellar remnants) or the presence in the cluster of a significant amount of dark matter. At a more conservative confidence level (4$\\sigma$) for the relative velocity $V_{\\rm rel,A},$ the mass to light ratio would result $M_{\\rm encl}/L\\,\\geq\\,2.5\\,{\\rm M_\\odot}/{\\rm L_\\odot}$, again implying a dynamic mass much higher than the mass derived from optical observations.  A further test of the cluster membership of PSR\\,A and PSR\\,C will be possible in future. It will involve the comparison of the proper motions of PSR\\,A and PSR\\,C with those of the three pulsars close to the cluster core, whose association with NGC\\,6752 is unambigously proved by the very strong gravitational pull affecting the values of their spin period derivative. This task will take some years: our simulations show that, with the present accuracy and collection rate of the TOAs and if the three innermost pulsars display the same proper motion as PSR\\,A, a $3\\sigma$ determination will require a total data span of about $8-10$ years."
    },
    "0609/astro-ph0609681_arXiv.txt": {
        "abstract": "{\\emph{Aims}. We present a systematic theoretical study on 40~polycyclic aromatic hydrocarbons dications (PAHs$^{++}$) containing up to 66~carbon atoms.\\\\ {\\emph{Methods}. We performed our calculations using well established quantum\\textendash chemical techniques in the framework of the density functional theory (DFT) to obtain the electronic ground\\textendash state properties, and of the time\\textendash dependent DFT (TD\\textendash DFT) to evaluate the excited\\textendash state properties.}\\\\ {\\emph{Results}. For each PAH$^{++}$ considered we computed the absolute visible\\textendash UV photo\\textendash absorption cross\\textendash section up to about 30~eV. We also evaluated their vibrational properties and compared them to those of the corresponding neutral and singly\\textendash ionised species. We estimated the adiabatic and vertical second ionisation energy $\\Delta I$ through total energy differences.}\\\\ {\\emph{Conclusions}. The $\\Delta I$ values obtained fall in the energy range 8\\textendash13~eV, confirming that PAHs could reach the doubly\\textendash ionised state in HI regions. The total integrated IR absorption cross\\textendash sections show a marked increase upon ionization, being on the average about two and five times larger for PAHs$^{++}$ than for PAHs$^+$ and PAHs, respectively. The visible\\textendash UV photo\\textendash absorption cross\\textendash sections for the 0, +1 and +2 charge\\textendash states show comparable features but PAHs$^{++}$ are found to absorb slightly less than their parent neutrals and singly ionized species between $\\sim7$ and $\\sim12$~eV. Combining these pieces of information we found that PAHs$^{++}$ should actually be more stable against photodissociation  than PAHs and PAHs$^+$, \\emph{if} dissociation tresholds are nearly unchanged by ionization.}\\\\ ",
        "introduction": "\\label{introduction}  Several spectroscopic features, observed in absorption or in emission in a large number of interstellar environments, are hypothesised to provide strong evidence for the presence of large carbon\\textendash based molecules \\citep{ehr00}. Among them are the diffuse interstellar bands \\citep[DIBs,][]{her95,ful00}, seen in absorption in the visible\\textendash UV; the aromatic infrared bands \\citep[AIBs,][]{pee04,tie05}, dominating the near and mid\\textendash IR spectrum of many interstellar sources; the extended red emission \\citep[ERE,][]{wit04,wit06} and the blue\\textendash luminescence \\citep[BL,][]{vij04,vij05b,vij05a}, wide emission bands peaking in the red and blue parts of the visible spectrum, respectively; the UV\\textendash bump at $\\sim220$~nm and the far\\textendash UV rise in the general interstellar extinction curve \\citep[]{dra03,fit04}.  Based on their spectral properties, their high photostability and their organic nature, polycyclic aromatic hydrocarbons (PAHs) are believed to be abundant in the interstellar medium \\citep[ISM,][]{pug89,all89}. Isolated gas\\textendash phase PAHs absorb efficiently visible\\textendash UV radiation and are therefore likely to contribute to the DIBs spectrum \\citep{leg85,van85,cra85} and to the UV\\textendash bump and far\\textendash UV rise of the extinction curve \\citep{job92a,job92b}. After photon absorption, they undergo a fast radiation-less transition (internal conversion or inter-system crossing), and subsequently relax radiatively by electronic fluorescence and/or phosphorescence (possibly contributing to BL and ERE) and IR emission in their active C\\textendash C and C\\textendash H vibrational modes \\citep[hence producing AIBs,][]{leg84,all85}.  Interstellar PAHs are expected to exist in different charge and hydrogenation states depending on the physical conditions of the host environment \\citep{tie05}. In particular, the possible presence of PAH dications (PAHs$^{++}$) in the ISM, first proposed by \\citet{lea86}, has recently received further interest, following the proposal that PAHs$^{++}$ could be the carriers of ERE \\citep{wit06}. Theoretical works on the IR properties of a few PAHs$^{++}$ \\citep{ell99,bau00,pau02} show that their vibrational patterns are similar to those of their parent singly ionized species (PAHs$^+$). From modelling studies it is predicted that specific values of electron density and radiation field intensity ought to lead to a population of PAHs$^{++}$ in excess of that of PAHs$^+$ \\citep{bak01b,bak01a}. Therefore PAHs$^{++}$ would also contribute to the AIBs.  While the first ionisation energies $I^+$ for more than 200~PAHs containing up to 48 carbon atoms are well documented \\citep{eil81}, only a few experimental determinations of the double ionisation energy $I^{++}$ of some PAHs have been published to date. Using photon impact, electron impact and charge stripping techniques, \\citet{tob94} performed a comprehensive study of $I^+$ and $I^{++}$ for 21 different PAHs up to the size of coronene (C$_{24}$H$_{12}$). More recently, \\citet{sch01} presented a combination between photo\\textendash ionisation, charge stripping, reactivity studies and quantum\\textendash chemical calculations to obtain the single and double ionisation energies of corannulene (C$_{20}$H$_{10}$) and coronene.  From the known values of $I^+$ and $I^{++}$, it was suggested \\citep{lea86,lea87,lea89,lea95b,lea96} that PAHs$^{++}$ could be formed in diffuse HI regions. The proposed formation mechanism is a sequential two\\textendash photons absorption process in which the ionisation of a PAH is followed by a further ionisation of PAH$^+$. The minimum energy needed for the first step is the single ionisation energy $I^+$, while the second step requires a photon energy of at least $\\Delta I$ = $I^{++} - I^{+}$; the maximum amount of internal energy in the newly formed PAHs$^{++}$ is therefore 13.6 eV - $\\Delta I$. Since the total energy of the dication is larger than the sum of the total energies of two singly ionised fragments, PAHs$^{++}$ are metastable \\citep{ros04} and could be dissociated via Coulomb explosion as shown by experiments \\citep{lea89a,lea89b}. However, theoretical calculations on the fragmentation pathways of the benzene dication \\citep{ros04} show that this charge separation mechanism should not be a significant channel in most astrophysical environments. This is consistent with the observed production of significant amounts of PAHs$^{++}$ in soft ionisation experiments \\citep{led99}. Therefore PAHs$^{++}$ could be found in the ISM.    To evaluate the possible contribution of PAHs$^{++}$ to ERE \\citet{wit06} computed the first low\\textendash lying electronic transitions of five representative PAHs$^{++}$ up to the size of ovalene (C$_{32}$H$_{14}$), showing that they fall indeed in the spectral range of the observed ERE. Laboratory measurements of the yield of optical fluorescence by PAHs$^{++}$ are however needed to test this hypothesis. To the best of our knowledge, with the exception of the work by \\citet{wit06}, no other electronic spectra of PAHs$^{++}$ have been published to date. In particular, a detailed study of the electronic excitation properties for a large sample of PAHs$^{++}$ up to the energies excitable in typical interstellar sources, together with their vibrational analyses, has been missing until now.  In the last few years we have started a long\\textendash term project to produce an atlas of synthetic absorption and emission spectra of specific PAHs, to be compared with astronomical observations \\citep{job02,mal03,mul03,mul06b,mul06a,mul06c}. This requires the use of some key molecular parameters of PAHs as a basis to run Monte\\textendash Carlo simulations \\citep{bar83,mul98,job02} of their photophysics in the ISM. We report here these fundamental data for a sample of 40~PAHs$^{++}$ ranging in size from naphthalene and azulene (C$_{10}$H$_{8}$) to circumovalene (C$_{66}$H$_{20}$). Even if PAHs and related species containing from less than 50 to more than 200 carbon atoms are expected to be present in the ISM \\citep{bou99}, we here restricted ourselves to species containing up to a maximum of 66~carbon atoms, since computational costs scale steeply with dimensions. Following our previous work on neutral, cationic and anionic PAHs \\citep{mal04,mal05}, we computed the absolute photo\\textendash absorption cross\\textendash section for the doubly\\textendash ionised species. For these same molecules we also evaluated the IR spectral properties extending the sample of specific PAH$^{++}$ for which they were previously available \\citep{ell99,bau00,pau02}. New calculations, together with our previous results at the same level of theory for the corresponding PAHs and PAHs$^+$, enabled us to estimate, via total energy differences, the adiabatic and vertical $I^+$ and $I^{++}$. Although the computational approach we used for this purpose is not sufficient to predict absolute ionisation energies within chemical accuracy ($\\pm$0.1~eV), it is in reasonable agreement with the available experimental data and can be easily extended to larger PAHs in the near future.  Section~\\ref{computational} contains the technical details for the calculation of the ionisation energies and the vibrational properties (Sect.~\\ref{dft}) and the photo\\textendash absorption spectra (Sect.~\\ref{tddft}). The results obtained are presented in Sect.~\\ref{results} and discussed in Sect.~\\ref{discussion}. Our conclusions are reported in Sect.~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}   We performed a systematic theoretical study on the ionisation energies of PAHs and the IR and visible\\textendash UV absorption properties of their dications. The values we found for the second IEs confirm that PAHs could reach the doubly\\textendash ionised state in H$_\\mathrm{I}$ regions. Combining the IR and visible\\textendash UV absorption properties of PAHs$^{++}$ we predict an increased photostability for these species with respect to their neutral and singly\\textendash ionized parent molecules, unless they have lowest dissociation channels.  The recent paper by \\citet{wit06} has recently renewed the interest in the dications of PAHs, proposing them as plausible candidates of the Extended Red Emission. The fundamental data we present here for a large sample of different PAHs will be useful to model in detail the photophysics of these molecules in interstellar conditions. For instance, we can estimate quantitatively the complete IR emission spectrum of each molecule in the sample as a function of its charge state, for direct comparison with data from the Infrared Space Observatory and Spitzer Space Telescope missions and, in the near future, from the Herschel Space Observatory mission \\citep{mul06a,mul06b,mul06c}. In addition, combined with measurements of the electron recombination rates of singly\\textendash ionized PAHs \\citep{nov05}, and their extension to doubly\\textendash ionized species, which still have to be studied, the present work opens the way for a detailed modelling of the chemistry and photophysics of PAHs in interstellar environments \\citep{bak98,lep01} where they may play a prominent role."
    },
    "0609/astro-ph0609648_arXiv.txt": {
        "abstract": "We present a 20-pc, volume-limited sample of M7--L8 dwarfs created through spectroscopic follow-up of sources selected from the Two Micron All Sky Survey (2MASS) Second Incremental Release Point Source Catalog. In this paper, we present optical spectroscopy of \\newspectra\\ candidate nearby ultracool dwarfs, including \\latenear\\ late-M and L dwarfs likely to be within 20~pc of the Sun and \\latefar\\ more distant late-type dwarfs. We have also identified five ultracool dwarfs with spectral signatures of low-gravity. Combining these data with previous results, we define a sample of \\LFobjects\\ ultracool dwarfs in \\LFsystems\\ systems within 20~pc. These are used to estimate the \\textit{J}- and \\textit{K}-band luminosity functions for dwarfs with optical spectral types between M7 and L8 ($10.5 < M_J < 15$, $9.5< M_{K_S} <13$). We find a space density of $4.9\\times10^{-3}$~pc$^{-3}$ for late-M dwarfs (M7--M9.5) and a lower limit of $3.8\\times10^{-3}$~pc$^{-3}$ for L dwarfs. ",
        "introduction": "Over the last several years, we have been using the Two Micron All Sky Survey \\citep[2MASS,][]{2MASS} to carry out a census of late-type dwarfs in the vicinity of the Sun. In particular, we have undertaken the first comprehensive survey of ultracool M and L dwarfs in the immediate Solar Neighborhood \\citep[hereafter Paper~V]{Cruz03}. \\defcitealias{Cruz03}{Paper~V} This population of cool, low-luminosity dwarfs includes both very low-mass stars and substellar-mass brown dwarfs \\citep{Chabrier00_Review,Burrows01}.  One of the primary aims of our program is the derivation of the luminosity function (the number of objects as a function of absolute magnitude) across the hydrogen-burning limit. The local field luminosity function has been well measured for objects as cool as late-M dwarfs \\citep{PMSU4}, but studies probing to fainter magnitudes have been limited by poor statistics and low sensitivity \\citep{Reid99,K99,K00,NN,Burgasser02_thesis}. Pushing to lower luminosities and larger volumes is crucial to measuring the contribution of brown dwarfs to the local space density.  The initial results from our program to create a statistically well-defined sample of nearby ultracool dwarfs were presented in \\citetalias{Cruz03}. In that paper, we described how we constructed a sample of 630 candidate nearby late-M and L dwarfs from the 2MASS Second Incremental Data Release Point Source Catalog, designated the ``2MU2'' sample. We also presented follow-up spectroscopic observations of a subset of the ultracool candidates. In this paper, we present the remainder of the far-red optical spectroscopic observations and combine these results with previous observations to construct the \\textit{J}- and \\textit{K$_S$}-band luminosity functions. These data include many newly discovered late-type dwarfs, including twelve additions to the core ``20-pc 2MU2'' sample of M7--L8 dwarfs estimated to be within 20~parsecs and used to measure the luminosity function.  This is the ninth paper in the series and follows directly from \\citetalias{Cruz03}. The first three papers and Papers~VII and~VIII \\citep{Paper1, Paper2, Cruz02,Paper7,Paper8} concentrate on proper-motion selected K and M dwarfs. Paper IV \\citep{Paper4} describes the discovery of \\object{2MASS J18353790+3259545}, an M8.5 dwarf within 6~pc of the Sun (which is included in the present analysis). Results from a search for ultracool dwarfs lying close to the Galactic Plane are presented in \\citet[Paper VI]{Paper6}.  Here, we build on \\citetalias{Cruz03} and present the remainder of the far-red optical spectroscopic follow-up of the 2MU2 sample and derive the \\textit{J}- and \\textit{K$_S$}-band luminosity functions based on objects estimated to be within 20~pc. In \\S~\\ref{sec:9sample} we overview the characteristics of the 2MU2 sample including its creation and follow-up status. Spectroscopic observations from seven telescopes over sixteen runs are described in \\S~\\ref{sec:9obs}. The results of the spectroscopy and the methods used to estimate spectral types, absolute magnitudes, and distances are in \\S~\\ref{sec:9results}. Interesting individual objects are presented in \\S~\\ref{sec:individual}. In \\S~\\ref{sec:9LF} we examine the completeness of the 20-pc 2MU2 sample and derive the luminosity function. Finally, we discuss these results in \\S~\\ref{sec:9discussion} and briefly summarize our findings in \\S~\\ref{sec:9summary}. ",
        "conclusions": "\\label{sec:9discussion}  Integrating our results for the luminosity function, we derive a space density of $8.7\\pm0.8 \\times 10^{-3}$~pc$^{-3}$ for M7--L8 dwarfs. Ultracool M dwarfs, spectral types M7 to M9.5, contribute a density of $4.9\\pm0.6 \\times 10^{-3}$~pc$^{-3}$. We find the space density of L dwarfs to be $\\ge3.8\\pm0.6\\times 10^{-3}$~pc$^{-3}$, with L0--L3 dwarfs contributing $1.7\\pm0.4\\times 10^{-3}$~pc$^{-3}$ and late-type L dwarfs responsible for at least $2.2\\pm0.4 \\times 10^{-3}$~pc$^{-3}$. Our results are in good agreement with the space density of $4.5\\times10^{-3}$~pc$^{-3}$ derived for ultracool M dwarfs by \\citet{NN}, but exceed by almost a factor of two their initial estimate of the local L dwarf density. Since we are using a significantly larger sample than \\citeauthor{NN}, our measured densities are likely more accurate.  Putting these results in a wider context, Figure~\\ref{fig:8pcLF} superimposes \\LFj\\ as derived from our ultracool sample against results for the 8-pc sample \\citep{Paper4}. As discussed in \\citet{Paper8}, 95\\% of the 8-parsec stars have reliable trigonometric parallaxes; the \\textit{J}-band photometry is taken either from 2MASS or, for bright stars, drawn from the literature and transformed, if necessary, to the CIT system. The derived space densities are listed in Table~\\ref{tab:8pclf}. While the 8-pc sample includes only a handful of ultracool dwarfs, it is clear that the two datasets are in excellent agreement, with the new data confirming the sharp decline in number densities at $M_J > 10$. In order to get an idea of the space densities in the \\textit{I} band, $M_I\\sim13$ for M7 and $(I-J)$ ranges from 2.4--4 for M7--L8 \\citep{Dahn02}.  The overall morphology of the \\textit{J}-band luminosity function reflects the convolution of the underlying mass function, $\\Psi(M)$, and the $M_J$-mass relation. Qualitatively, \\LFj\\ increases for $0<M_J<7$ since the mass function increases with decreasing mass. At fainter magnitudes, \\LFj\\ turns over not because $\\Psi(M)$ changes drastically, but because the slope of the $M_J$-mass relation changes---while ${{\\delta Mass}\\over {\\delta M_J}} \\sim 0.4 M_\\sun$ mag$^{-1}$ for $M_J < 7$, ${{\\delta Mass}\\over {\\delta M_J}} \\sim 0.07 M_\\sun$ mag$^{-1}$ for $7<M_J<10$ \\citep{Delfosse00}.  Our analysis shows that \\LFj\\ declines sharply with decreasing luminosity beyond $M_J=10$, reaching a minimum at $M_J \\sim13$, roughly corresponding to spectral type L4. Formally, our data indicate that the space densities remain approximately constant at fainter magnitudes; however, since our measurements are lower limits, \\LFj\\ likely increases at $M_J > 14$.  Our field survey has uncovered both very low-mass stars and sub-stellar mass brown dwarfs from the local population. Assuming a typical age of 2--5 Gyrs for a field dwarf, even the most massive brown dwarfs have cooled to L-dwarf temperatures. Moreover, theoretical models \\citep[e.g.,][]{Burrows01} indicate that the hydrogen-burning limit (for solar abundances) corresponds to a temperature of $\\sim$1700~K, or spectral type $\\sim$L4 \\citep{Golimowski04}. Thus, early-type L dwarfs include a mix of stars and brown dwarfs, with brown dwarfs acquiring increasing dominance from L0 to L4 until they are the sole contributors at $M_J\\ga13.5$.  As discussed elsewhere \\citep{Chabrier03, Burgasser04, Allen05}, the morphology of the luminosity function is due to this mix of stars and brown dwarfs in the ultracool regime and is qualitatively in accord with theoretical expectations. The drop in number density from M7 to L4 reflects a further contraction in ${{\\delta Mass}\\over {\\delta M_J}}$. The population of the very lowest-mass stars that appear as L dwarfs span an extremely small range in mass $(0.075<M_\\sun<0.085$) and, as a result, are rare. The brown dwarfs in this effective temperature regime are relatively young and are at the high-mass extreme, near the hydrogen-burning limit.  Brown dwarfs dominate the counts beyond the plateau in \\LFj\\ at $M_J \\sim 13.5$ ($\\sim$L4), and the upturn in number densities reflects the slow down in cooling rates at lower temperatures. For example, a 0.07~$M_\\sun$ brown dwarf takes 2.7~Gyrs to evolve down the L dwarf sequence, but remains a (cooling) T dwarf (T$_{eff}\\approx$1400--600~K) for $\\sim$30~Gyrs, or more than 2 Hubble times; a low-mass, 0.025~$M_\\sun$ brown dwarf spends only 120~Myrs as an L dwarf, but 1.5~Gyrs as a T dwarf \\citep{Burrows01}.  As noted in \\S~\\ref{sec:lf2}, early-type T dwarfs also contribute to the \\LFj\\ at magnitudes fainter than $M_J=14$. Evolutionary models suggest that there may be as many as 80 T0--T5 dwarfs within 20 parsecs \\citep{Burgasser04}, as compared with the 12 late-type L dwarfs contributing to our luminosity function at $M_J>14$. Incorporating those cooler dwarfs in the analysis is likely to lead to a steeply increasing \\LFj.  Clearly, one of the aims in deriving the luminosity function for ultracool dwarfs is setting constraints on the mass function, usually parameterized as a power-law, ${dN \\over dM} \\propto M^{-\\alpha}$, where the Salpeter value is $\\alpha = 2.5$. \\citet{Schultheis06} have recently proposed that $\\alpha >2$ at low masses (implying that brown dwarfs and low-mass stars essentially account for dark matter). This is in contrast to values of $\\alpha \\sim 1$ derived in most other analysis, such as \\citet{PMSU4}\\footnote{Note that \\citeauthor{Schultheis06} are incorrect in suggesting that the nearby star sample employed in this analysis has significant incompleteness.}. A full discussion of this issue is beyond the scope of the present paper. However, we note that the \\citeauthor{Schultheis06} analysis rests on matching models against the \\textit{V}-band luminosity function for nearby stars, and is therefore very weakly constrained for $0.1 M_\\sun < M < 0.15 M_\\sun$ and essentially unconstrained at lower masses. \\citet{Burgasser04} and \\citet{Allen05} have shown that our initial results presented in \\citetalias{Cruz03} (which are in general agreement with the luminosity function derived here) are consistent with a range of values of $\\alpha < 1.5$. While this upper limit is not fully satisfying, it is in disagreement with the steeply re-rising mass function as proposed by \\citeauthor{Schultheis06}.    We have mined the 2MASS Second Incremental Data Release PSC for ultracool dwarfs within 20~parsecs of the Sun. Extensive spectroscopic follow-up has led to the discovery of $\\sim$100 L dwarfs and $\\sim$200~late-M dwarfs---over 50 of these are within 20~pc \\citepalias[and this paper]{Cruz03}, doubling the local census of ultracool dwarfs. In these data we have also uncovered several wide binaries, a handful of young objects, and two slightly metal-poor L dwarfs. Combining our data with previously known nearby late-type dwarfs, we have used \\LFobjects~objects in \\LFsystems~systems to estimate the \\textit{J} and \\textit{K$_S$}-band luminosity functions of ultracool dwarfs in the Solar Neighborhood. This work has provided the first robust estimate of the luminosity function of late-type stars and brown dwarfs. We have measured the density of late-M dwarfs (M7--M9.5) to be $4.9\\pm0.6\\times10^{-3}$~pc$^{-3}$ and the density of L dwarfs to be at least $3.8\\pm0.6\\times10^{-3}$~pc$^{-3}$. Reliable discrimination between different models for the underlying mass function must await observational surveys that probe brown dwarfs at temperatures below $\\sim$600~K, likely near the boundary between T and Y dwarfs."
    },
    "0609/astro-ph0609823_arXiv.txt": {
        "abstract": "In order to constrain the physical processes that regulate and downsize the Active Galactic Nucleus (AGN) population, the predictions of the MOdel for the Rise of GAlaxies aNd Active nuclei ({\\gal}) are compared to luminosity functions (LFs) of AGNs in the optical, soft X-ray and hard X-ray bands, to the local BH--bulge mass relation, and to the observed X-ray number counts and background.  We also give predictions on the accretion rate of AGNs in units of the Eddington rate and on the BH--bulge relation expected at high redshift.  We find that it is possible to reproduce the downsizing of AGNs within the hierarchical $\\Lambda$CDM cosmogony and that the most likely responsible for this downsizing is the stellar kinetic feedback that arises in star-forming bulges as a consequence of the high level of turbulence and leads to a massive removal of cold gas in small elliptical galaxies.  At the same time, to obtain good fits to the number of bright quasars we need to require that quasar-triggered galactic winds self-limit the accretion onto BHs; however, the very high degree of complexity of the physics of these winds, coupled with our poor understanding of it, hampers more robust conclusions.  In all cases, the predicted BH--bulge relation steepens considerably with respect to the observed one at bulge masses $<10^{11}$ \\msun; this problem is related to a known excess in the predicted number of small bulges, common to most similar models, so that the reproduction of the correct number of faint AGNs is done at the cost of underestimating their BH masses.  This highlights an insufficient downsizing of elliptical galaxies, and hints for another feedback mechanism able to act on the compact discs that form and soon merge at high redshift. The results of this paper reinforce the need for direct investigations of the feedback mechanisms in active galaxies, that will be possible with the next generation of astronomical telescopes from sub-mm to X-rays. ",
        "introduction": "\\label{section:introduction}  In the last years evidence has grown of the importance of the interaction between AGNs and their host galaxies.  This interaction is now deemed to be of fundamental relevance for understanding the assembly not only of the supermassive Black Holes (BH) responsible for AGN activity, but also of the galaxies, in particular their spheroidal components (ellipticals or spiral bulges).  Observations of nearby galaxies show the existence of a well-defined correlation between the mass or central velocity dispersion of bulges and the mass of the hosted BHs (Kormendy \\& Richstone, 1995; Magorrian et al., 1998; more recent determinations are found, e.g., in Marconi \\& Hunt, 2003 and H\\\"aring \\& Rix, 2004).  Moreover the observed mass function of these BHs is consistent with that inferred from quasar\\footnote{In this paper we refer to luminous AGNs as quasars, with no reference to their radio emission.}  luminosities, under simple assumptions for the radiative efficiency and accretion rate (Soltan 1982; Cavaliere \\& Padovani 1988; Salucci et al. 1999; Yu \\& Tremaine 2002; Shankar et al. 2004; Marconi et al. 2004; see also Haiman, Ciotti \\& Ostriker 2004). On the other hand, observations at high redshift highlight that quasars and radio-loud AGNs are hosted in elliptical galaxies (see e.g. Dunlop et al. 2003).  An indirect evidence of the BH -- bulge connection relies in the parallel evidence of downsizing of both populations.  The evolution of the LFs of AGNs in the soft and hard X-ray bands reveals that the number density of fainter objects peaks at a lower redshift with respect to brighter ones (see i.e. Ueda et al. 2003; Hasinger, Miyaji \\& Schmidt 2005; Barger et al. 2005; La Franca et al. 2005).  This evidence of downsizing is often referred to as the ``anti-hierarchical'' behavior of BH growth, and is confirmed by the analysis of Merloni (2004), based on the fundamental plane of accreting BHs (Merloni, Heinz \\& di Matteo, 2003), and by Marconi et al. (2004).  Recently, the GOODS collaboration determined the LF of low-luminosity quasars at $z\\sim 4$ (Cristiani et al. 2004; Fontanot et al. 2006a), revealing a dearth of objects with respect to naive extrapolations of Pure Luminosity Evolution models anchored to the SDSS bright quasars (Fan et al. 2003). This evidence, confirmed by the COMBO17 survey at higher luminosities (Wolf et al. 2003) is again in line with the downsizing trend mentioned above.  A similar trend has longly been claimed by some authors to be observed in galaxies (see, e.g., Cowie et al. 1996; Treu et al. 2005; Papovich et al. 2006; Fontana et al. 2006), especially in ellipticals, whose stellar mean ages, metal enrichment and level of alpha-enhancement are known to correlate positively with the stellar mass (see, e.g. Thomas 1999; Renzini 2004).  These evidences point to a formation scenario characterized by a quick burst of star formation rapidly followed by a strong wind which expels the residual Inter-Stellar Medium (ISM); this wind should take place earlier in more massive galaxies (Matteucci 1994) and may be triggered by the quasar shining (Granato et al. 2001).  In this framework it is also possible to explain the high metallicity of quasar hosts (Matteucci \\& Padovani 1993; Hamann \\& Ferland 1999; Romano et al. 2002; D'Odorico et al. 2005).  The influence of AGNs may not be limited to galaxies: jets from radio galaxies are now one of the most promising candidates for quenching cooling flows in galaxy clusters (McNamara et al. 2005; Voit \\& Donahue 2005; Fabian et al. 2006; see also the simulations of Quilis, Bower, \\& Balogh 2001; Dalla Vecchia et al. 2004; Ruszkowski, Bruggen \\& Begelman 2004; Zanni et al. 2005; Sijacki \\& Springel 2006), and thus limit the mass of the most massive ellipticals (Benson et al. 2003; Bower et al. 2006; Croton et al. 2006).  The problem of the demography of the AGN population in the context of galaxy formation in $\\Lambda$CDM models has been addressed by many authors (see, e.g., Kauffmann \\& Haehnelt 2000; Monaco, Salucci \\& Danese 2000; Cavaliere \\& Vittorini 2000, 2002; Cattaneo 2001; Granato et al. 2001, 2004; Mahmood, Devriendt \\& Silk 2004; Bromley, Somerville \\& Fabian 2004; Menci et al. 2003, 2004, 2006; Cattaneo et al. 2005; Vittorini et al. 2005; Bower et al. 2006; Croton et al. 2006; Hopkins et al. 2006; Malbon et al. 2006).  In particular, the downsizing of galaxies and AGNs is not straightforward to reproduce in the hierarchical cosmogonic scenario, where larger dark matter (DM) halos form on average later by the merging of less massive halos.  As far as AGNs are concerned, Monaco, Salucci \\& Danese (2000) and later Granato et al. (2001; 2004) suggested that a delay in the shining of quasars in small bulges, motivated by stellar feedback, could explain this trend.  Unfortunately, a more precise assessment of the success of the models listed above in predicting the downsizing of AGNs is not easy in most cases, as only a few of these papers (like Menci et al. 2004) explicitly compare to data in terms of the number density of AGNs as a function of luminosity.  It is worth noting that in these papers different ideas have been proposed on the origin of the BH -- bulge relation.  Several proposed that the feedback from the AGN is able to self-limit the masses of both spheroids and BHs, (see, e.g., Ciotti \\& Ostriker 1997; Silk \\& Rees 1998; Haehnelt, Natarajan \\& Rees 1998; Fabian 1999; Granato et al. 2004; Murray, Quataert \\& Thompson 2005; Monaco \\& Fontanot 2005). Such self-limitation is supported also by the results of N-body simulations (Di Matteo et al. 2003, 2005; Kazantzidis et al. 2005; Springel, Di Matteo \\& Hernquist 2005).  However, the other authors successfully reproduced the BH -- bulge relation simply assuming a proportionality between star formation (in bulges) and accretion rates, implicitly determined by the mechanism responsible for the almost complete loss of angular momentum of the gas accreting on the BH . The difference between the two approaches (self-limitation or proportionality between star formation and accretion) was recently clarified by Monaco \\& Fontanot (2005).  This paper is the second of a series devoted to describe the {\\gal} code for the formation and evolution of galaxies and AGNs.  The code is described in full detail in Monaco, Fontanot \\& Taffoni (2006; hereafter paper I), while the ability of the model to reproduce the early formation and assembly of massive galaxies is demonstrated in Fontana et al. (2006) and Fontanot et al. (2006b). Here we focus on the self-consistent modeling of accreting BHs and on their feedback on the host galaxies.  By reproducing - or failing to reproduce - the observed properties of the AGN population, in particular the LFs in the optical, soft and hard X-ray bands, and the BH - bulge correlation at $z=0$, we obtain valuable constraints on the physical processes involved.  We also demonstrate our ability to roughly reproduce the observed X-ray counts in the soft (0.5--2 keV) and hard (2--10 keV) bands and background from 0.2 to 300 keV, and give predictions on the BH -- bulge relation at high redshift and on the Eddington ratios of the emitting objects.  The paper is organized as follows. In section~\\ref{section:model} we describe the main properties of the {\\gal} model, the accretion onto BHs and its feedback on the forming galaxy, the runs used in the paper, and the procedure adopted to compute the AGN LFs, number counts and X-ray background.  In section~\\ref{section:results} we compare our results with available data and give some predictions on AGN and BH properties.  The results are discussed in section~\\ref{section:discussion}, while section~\\ref{section:conclusions} gives the conclusions.  Throughout this work we assume the concordance $\\Lambda$CDM cosmological model with $\\Omega_\\Lambda=0.7$, $\\Omega_m=0.3$, $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, and $\\sigma_8=0.9$. ",
        "conclusions": "\\label{section:conclusions}  This paper belongs to a series devoted to presenting {\\gal}, a new model for the joint formation and evolution of galaxies and AGNs, and is focused on the role of feedback in shaping the observed properties of AGNs and their relation with galaxies.  Our analysis confirms that models based on the $\\Lambda$CDM cosmogony are able to roughly reproduce at the same time both the properties of AGNs (presented here) and the properties of galaxy populations (presented in paper I, Fontana et al. 2006 and Fontanot et al. 2006b).  This model has been extended, through a careful modeling of AGN SEDs, to predict many QSO properties, especially in the X-rays.  It reproduces nicely LFs and number counts in the soft and hard X-ray bands, and the measured background from 0.5 to 300 keV, with a possible underestimate of the peak which confirms the claim for a significant though not dominant population of Compton-thick sources. This implies that the bulk of AGN accretion is roughly reproduced.  The agreement between model and data allows us to draw these conclusions.  {\\em Quasar triggered winds are necessary in our model to reproduce the number density of bright quasars.}  The parameter space of quasar winds is discouragingly wide; within a very limited sub-space, characterized by ``only'' six parameters, we find two good solutions, which is a clear sign that our knowledge on the physics of accretion onto BHs and their interaction with galaxies is still too poor to draw firm conclusions.  In any case, the idea that quasar winds are a necessary ingredient of galaxy formation is worth pursuing.  The two wind solutions that we propose are based on the same triggering criterion for the quasar wind, requiring that (i) the accretion rate is high enough to perturb the ISM, (ii) the gas mass to remove is not too high, (iii) the AGN is accreting in a radiatively efficient mode. The first solution is based on a tilt of the relation between bulge star-formation rate and loss rate of angular momentum, compensated by ``drying winds'', where the kinetic energy injected by the central engine causes a complete removal of ISM from the host bulge.  The second solution is based on the ``accreting wind'' mechanism proposed by Monaco \\& Fontanot (2005), where the wind is generated throughout the ISM, so that a fraction of the cold gas is compressed to the center and stimulates further accretion.  A good fit of the AGN LFs is then obtained by assuming, as explained in paper I, that discs with high gas surface density lose angular momentum (and then become bulges) because of the expected change in the feedback regime (Monaco 2004a) and the consequent increase of the velocity dispersion of clouds (as observed by F\\\"orster Schreiber et al. 2006).  {\\em We propose that kinetic feedback in star-forming bulges is a very good candidate for downsizing the AGN population.}  Indeed, the high velocity dispersion of cold gas in star-forming bulges can lead to a massive removal of ISM, leading to a suppression of less luminous AGNs at high redshift, while larger elliptical galaxies can retain their gas more easily.  At lower redshift larger BHs may accrete at lower rates, giving rise to the low-luminosity population dominating the X-ray background.  {\\em We make the following predictions:} (i) the BH-bulge relation is already in place at high redshift, though high BH masses in relatively small bulges may be found at $z>2$ if quasar winds are active; (ii) the Eddington ratios of accreting BHs depend both on redshift and bolometric luminosity, with low values for the faint AGNs at $0<z<1$ that give the bulk of the hard X-ray background (consistent with Ballo et al. 2006); (iii) most BH mass is acquired through accretion more than through mergers, especially if quasar-triggered winds are switched on.  {\\em There is an important point of discrepancy between our model and data, in that the predicted BH -- bulge relation at $M_{\\rm B}<10^{11}$ \\msun, or $M_{\\rm bh}<10^8$ \\msun, is significantly steeper than observed;} in the same mass range the predicted mass function of BHs is too low.  To our understanding, this is connected to a known excess of small bulges with respect to observations (see paper I) and to the excess of small-mass galaxies predicted by most (N-body or semi-analytic) models at $z\\sim1$ (Fontana et al. 2006). We have tested this idea by comparing our model predictions of the Eddington ratios of faint AGNs at $0.4<z<1$ with the data of Ballo et al. (2006); while the agreement is excellent for $M>1.6\\times10^7$ \\msun, smaller BHs accrete at a higher rate than observed, and this compensates for the lack of smaller BHs.  This discrepancy highlights that while kinetic feedback is efficient in downsizing AGNs by quenching the star formation in bulges, stars in small ellipticals may be born in discs, so that another mechanism is needed to produce the required level of downsizing of bulges.  In conclusion, in developing this work we have gained interesting insight into the complex problem of the cosmological rise of AGNs, highlighting two promising astrophysical mechanism (kinetic feedback in star-forming bulges and quasar-triggered galaxy winds) and the need for a third, unknown one to improve the downsizing of elliptical galaxies.  However, our poor understanding of the underlying physics hampers more robust conclusions.  This demonstrates once again that the field of joint galaxy and AGN formation is observationally-driven and we need strategies to single out the physical mechanisms at play. In particular, the main degrees of freedom in the theory are related to the connection between BH accretion and star formation, which can be constrained by estimating accretion rates, star-formation rates and BH masses in active galaxies; to the quasar-triggered winds, which can be constrained by observing warm and cold absorbers or Lyman-$\\alpha$ blobs associated to quasars; and finally to the nature of stellar feedback, which can be constrained by detailed observations of starburst galaxies.  The next generation of telescopes will provide suitable tools to assess these topics."
    },
    "0609/astro-ph0609012_arXiv.txt": {
        "abstract": "Hilbert-Huang Transform (HHT) is a novel data analysis technique for nonlinear and non-stationary data. We present a time-frequency analysis of both simulated light curves and an X-ray burst from the X-ray burster 4U 1702-429 with both the HHT and the Windowed Fast Fourier Transform (WFFT) methods. Our results show that the HHT method has failed in all cases for light curves with Poissonian fluctuations which are typical for all photon counting instruments used in astronomy, whereas the WFFT method can sensitively detect the periodic signals in the presence of Poissonian fluctuations; the only drawback of the WFFT method is that it cannot detect sharp frequency variations accurately. ",
        "introduction": "In X-ray binaries, matter flows onto the neutron star or black hole through an accretion disk and X-rays are emitted by the accretion flow from inner regions that host the strongest gravitational fields in the Universe. The space time characteristics of the X-ray emitting regions lead to the conclusion that the dynamical time scale for the accretion process of X-ray emitting is millisecond. That means, the millisecond X-ray variability is caused by orbital motion, neutron-star spin, disk- and neutron-star oscillations. So the millisecond X-ray variabilities are important for studying strong gravity physics. Kilohertz quasi-periodic oscillations (kHz QPOs) , millisecond pulsations and burst oscillations are millisecond X-ray variabilities. Burst oscillations do not occur in every X-ray burst and are different from burst to burst \\citep{klis00}.  Fourier transformation is a traditional method in the time-frequency analysis of linear, stationary and global time series. In astrophysics, Fourier spectral analysis is also widely used in the time-frequency analysis, although time series in astronomy are sometimes non-stationary. Based on the traditional Fourier spectral analysis, Windowed-Fast-Fourier Transform (WFFT) assumes that the time series to be piecewise stationary. In order to result in good time resolution, the window width of WFFT must be narrow; on the other hand, to achieve a high frequency resolution and high detection sensitivity, a wide window width is required. Therefore, in practice the window width must be chosen carefully depending on the goal of the study.  Both kHz QPOs and burst oscillations were detected in RXTE observations of the atoll source 4U 1702-429 in 1997. There were six type I X-ray burst, with a rise time of $<$ 1 s, and a 10-100 s exponential decay. Time-frequency distributions of the light curve from 4U 1702-429 at time resolution 1/8192 s were presented with WFFT with the unit signal $u(t)$ as the window function, and a window width of 2 s. All power spectra were computed by successively sliding the window at a step of 0.125 s by \\citet{mark99}. They discovered burst oscillations at about 330 Hz in five of the six type I X-ray bursts. In the 1997 July 26 burst a periodic signal was observed whose frequency increased slowly from 328 Hz to 330 Hz.  The advantages of the Hilbert-Huang Transform (HHT) \\citep[see][]{huang98} are that it can be used to analyze nonlinear and non-stationary data without specifying a basis and the data decomposition is made only on physical grounds, compared to Fourier decomposition and wavelet decomposition. Because the light curves obtained with most astronomical instruments involving photon counting devices are characterized by Poissonian fluctuations, it will be very useful if we can analyze such light curves with HHT and obtain detailed physical information such as fast frequency variations. Now, for the light curve of 4U 1702-429, the calculated time-frequency distribution and energy-frequency-time distribution using HHT are quite different from those using WFFT. In this paper we shall investigate the large discrepancies from these two methods by making various numerical tests. The structure of the paper is as follows. In section 2, we introduce the concepts and definitions of {\\it Fourier frequency} and {\\it instantaneous frequency} in the Hilbert Transform. In section 3, we introduce the HHT. In section 4, we make comprehensive comparisons between the HHT and WFFT. Finally we summarize our results and present a discussion in section 5. ",
        "conclusions": "We have reached the following conclusions: (1) Poissonian fluctuations cause large spurious spikes in the calculated Hilbert spectra with the HHT method, sO that in all cases the HHT cannot detect periodic signals reliably as long as the time series is essentially Poissonian; and (2) in comparison, the WFFT method is relatively insensitive to Poissonian fluctuations and can detect periodic signals in all cases with high sensitivity, but the time of frequency jump, should it exists, calculated with WFFT is less accurate than that with HHT .  \\citet{wu04} established the statistical characteristics of white noise through empirical study on Fourier spectra of the IMFs of the numerical white-noise time series with EMD, and provided a method for finding out which IMF components represent signals and which represent noises. However, as we have shown in this paper it is practically impossible to eliminate the spurious spikes in the calculated Hilbert spectra with HHT for Poissonian time series, consequently any empirical estimate of the statistical significance on the frequency component does not help in detecting underlying periodic signals.  It is worth noting, however, that it is possible to combine the HHT and WFFT together to extract as much information as possible from observed time series. For example, the WFFT method may be used to identify all frequency components in the time series and study any slow frequency variations. After the WFFT method identified a frequency jump, as shown in Figure 6, the HHT method can be used to identify the exact time of the sudden jump, as shown in Figure 3 and 4."
    },
    "0609/astro-ph0609538_arXiv.txt": {
        "abstract": "Magnetic precursors of C--shocks accelerate, compress and heat molecular ions, modifying the kinematics and the physical conditions of the ion fluid with respect to the neutral one. Electron densities are also expected to be significantly enhanced in shock precursors. In this Letter, we present observations of strongly polar ion and neutral molecules such as SiO, H$^{13}$CO$^{+}$, HN$^{13}$C and H$^{13}$CN, which reveal the electron density enhancements associated with the precursor of the young L1448--mm outflow. While in the ambient gas the excitation of the ions and neutrals is explained by collisional excitation by H$_2$ with a single density of $\\sim$10$^5$$\\,$cm$^{-3}$, H$^{13}$CO$^{+}$ shows an over excitation in the shock precursor component that requires H$_2$ densities of a factor of $\\geq$10 larger than those derived from the neutral species. This over excitation in H$^{13}$CO$^{+}$ can be explained if we consider an additional excitation by collisions with electrons and an electron density enhancement in the precursor stage by a factor of $\\sim$500, i.e. a fractional ionization of 5$\\times$10$^{-5}$. These results show that multiline observations can be used to study the evolution of the ion and electron fluids at the first stages of the C--shock interaction. ",
        "introduction": "In dense molecular clouds, the fractional ionization of the gas is very low \\citep[$\\leq$10$^{-7}$; Gu\\'elin, Langer \\& Wilson 1982;][]{cas98}. Electron densities are therefore insufficient to excite the first rotational levels of even strongly polar molecules like HCO$^{+}$, HCN or HNC % whose collisional cross sections \\citep[$\\sim$10$^{-6}$--10$^{-5}$$\\,$cm$^3$$\\,$s$^{-1}$; Bhattacharyya, Bhattacharyya \\& Narayan 1981;][]{saha81,fau01} are a factor of $\\geq$100 larger than those of weakly polar species like CO \\citep[$\\sim$10$^{-8}$$\\,$cm$^3$$\\,$s$^{-1}$;][]{saha81}. All attempts to use molecular excitation to measure the effects of collisional excitation by electrons in these regions have been so far unsuccessful \\citep{lan85}.  Modeling of the early stages of C--shocks predicts that the ion and electron densities are enhanced by the magnetic precursor \\citep{dra80}. For molecules like HCO$^{+}$, HCN or HNC, the electron density enhancement by a factor of $\\sim$100 produced by the precursor \\citep[see e.g.][]{flo96}, would make electron collisions competitive with excitation by H$_2$ collisions for the typical densities of dark clouds ($\\sim$10$^5$$\\,$cm$^{-3}$). Since the electron collisional coefficients for high initial \\textit{J} transitions of HCO$^{+}$ are a factor of $\\geq$10 larger than those of neutral molecules like HCN at low temperatures \\citep[$\\sim$10--20$\\,$K;][]{chu74,bha81}, differences in the molecular excitation between ion and neutral species are therefore expected in the precursor stage \\footnote{These excitation differences between ion and neutral molecules reflect a fundamental difference between the corresponding electron--impact cross sections at low energy: for neutrals, cross sections go to zero at threshold while for ions, they are large and finite \\citep[e.g.][]{chu74}.}. Toward the young L1448--mm molecular outflow, the detection of very narrow SiO emission and the enhancement of the ion abundance have been interpreted as signatures of the shock precursor \\citep{jim04}. Multiline observations of strongly polar species like HCO$^{+}$, HNC and HCN toward this outflow, are expected to show differences in excitation between the ambient gas where electron excitation is negligible, and the shock precursor where electron collisions should be important.  In this Letter, we present observations of several rotational transitions of SiO, H$^{13}$CO$^{+}$, HN$^{13}$C and H$^{13}$CN observed toward the ambient and shock precursor components of the L1448--mm outflow. The excitation differences observed between H$^{13}$CO$^{+}$, and SiO, HN$^{13}$C and H$^{13}$CN in the ambient and precursor components, can be explained by the electron density enhancement expected at the first stages of the C--shock evolution.  \\begin{figure} \\epsscale{1.05} \\plotone{f1.eps} \\caption{Low-- and high--\\textit{J} lines of SiO, H$^{13}$CO$^{+}$, HN$^{13}$C and H$^{13}$CN observed toward L1448--mm. Offsets in arcseconds are shown in the upper part of the columns ($\\alpha_{(J2000)}$~=~03$^{h}$25$^{m}$38$_.^s$0, $\\delta_{(J2000)}$~=~30$^{\\circ}$44$'$05$''$). The vertical dotted lines show the ambient gas at 4.7$\\,$km$\\,$s$^{-1}$ and the precursor component at 5.2$\\,$km$\\,$s$^{-1}$ toward L1448--mm (0,0), and at 5.4 and 5.6$\\,$km$\\,$s$^{-1}$ toward L1448--mm (0,--10) and (0,--20), respectively.} \\label{fig1} \\end{figure} ",
        "conclusions": ""
    },
    "0609/astro-ph0609574.txt": {
        "abstract": "{Debris disks are believed to be related to planetesimals left over around stars after planet formation has ceased. The frequency of debris disks around M-dwarfs which account for 70\\% of the stars in the Galaxy is unknown while constrains have already been found for A- to K-type stars. We have  searched for cold debris disks around 32 field M-dwarfs by conducting observations at $\\lambda = 850\\mu$m with the SCUBA   bolometer array camera at the JCMT and at   $\\lambda = 1.2$mm with the MAMBO array at the IRAM 30-m telescopes. This is the first survey of a large sample of M-dwarfs conducted to provide statistical constraints on debris disks around this type of stars. We have detected   a new debris disk around the M0.5 dwarf GJ842.2 at $\\lambda = 850\\mu$m, providing evidence for cold dust at large distance from this star ($\\sim 300$AU). By combining the  results of our survey with the ones of Liu et al. (2004), we estimate for the first time the detection rate of  cold debris disks around field M-dwarfs with ages between 20 and 200~Myr.  This detection rate is $13^{+6}_{-8}\\%$  and is consistent with the detection rate  of  cold debris disks (9 -- 23\\%) around  A- to K-type main sequence stars of the same age. This is an indication that cold disks may be  equally  prevalent across stellar spectral types. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609224_arXiv.txt": {
        "abstract": "We present the analysis of simultaneous multicolour $uvyI_{\\rm C}$ photometry and low-resolution spectroscopy for the rapidly rotating $\\beta$ Cephei star SY Equ. From the photometric time series, we confirm the dominant pulsation frequency, $f_1$ = 6.029 d$^{-1}$, and we find an evidence for two additional modes. In spectroscopy, the highest peak occurs at $f_{\\rm a}$ = 0.197 d$^{-1}$ or its alias 0.803 d$^{-1}$. It can be interpreted either in terms of a binary motion or as the $g$-mode pulsation. In addition, we reveal the pulsation mode with frequency of about 6.029~d$^{-1}$, i.e. the same which dominates photometric variations, and a few new candidates. For the dominant frequency we obtain mode identification from the combined photometric and spectroscopic observations. From non-adiabatic pulsation calculations, we show that the frequency of the dominant mode in SY Equ is consistent with the stellar models of much lower effective temperatures than used in many papers. ",
        "introduction": "SY Equ (HD 203664) is the early-B spectral type star, B0.5\\,IIIn, classified as the $\\beta$ Cephei variable by \\cite{aerts2000} from the Hipparcos photometry. The star is one of the fastest rotators among $\\beta$ Cephei pulsators, having projected rotational velocity of about 200 km~s$^{-1}$ \\cite{little1994}. Moreover, it is a high Galactic latitude star with $b$ = $-$27.5$^{\\circ}$ which locates it more than 1~kpc out of the Galactic plane.  There are several determinations of stellar parameters for SY Equ. \\cite{little1994} derived $\\log T_{\\rm eff}$ = 4.447~dex, $\\log g$ = 3.7~dex from the high-resolution optical spectra and estimated a mass of 14~$M_{\\odot}$ from the evolutionary tracks of \\cite{mame88}. On the other hand, using Geneva photometry, \\cite{aerts2006} determined $\\log T_{\\rm eff}$ = 4.47 $\\pm$ 0.01 dex and $\\log g$ = 3.9 $\\pm$ 0.3~dex. Much lower value of effective temperature, $\\log T_{\\rm eff}$ = 4.401~dex, was obtained by \\cite{keenan1982} from Str\\\"omgren photometry and equivalent width of the H$\\epsilon$ line. We adopt here stellar parameters derived from the International Ultraviolet Explorer (IUE) spectra and visual photometry by \\cite{niemdaszy2005}: $\\log T_{\\rm eff}$ = 4.388 $\\pm$ 0.035 and $\\log g$ = 3.905 $\\pm$ 0.1. These values result in a mass estimate of 11.5 $\\pm$ 1.5~$M_{\\odot}$ if the metallicity $Z$ = 0.02 is assumed. The metallicity in terms of [m/H], obtained from the IUE/INES data, amounts to 0.00 $\\pm$ 0.21~dex, which corresponds to $Z$ = 0.02$^{+0.012}_{-0.007}$. We prefer to rely on this determination because it uses ultraviolet spectral region, where B-type stars emit most of their energy.  In this note we give only a brief outline of our results. The full analysis will be published elsewhere. ",
        "conclusions": "We presented the results of the frequency analysis of both the photometric and spectroscopic time series of SY~Equ. In both datasets, we found the dominant mode at $f_1$ = 6.029~d$^{-1}$, and several new candidate frequencies. Moreover, there is a low-frequency variation at $f_a$ = 0.197~d$^{-1}$ detected in radial velocities, which can be an evidence of the binary motion or the excitation of a gravity mode. For the dominant mode we confirmed the $\\ell$ = 2 identification of \\cite{aerts2006} and identified the azimuthal order, $m$ = $+$2. It means that the dominant frequency is a sectoral prograde mode with $\\ell$ =2. However, one has to be aware that at $v_{\\rm rot} \\approx$ 200 km\\,s$^{-1}$, the rotational mode coupling can take place, if the frequency difference between modes $j$ and $k$ is of the order of angular velocity of rotation, and if the spherical harmonic indices satisfy the relation $l_j = l_k\\pm$ 2, see \\cite{daszy2002}.  Recently,  \\cite{aerts2006} made an attempt of comparisons of theoretical frequencies with the observed value of the dominant mode, but the authors totally ignored the effects of rotation and the instability condition. Our analysis showed that only for lower effective temperatures we enter the instability domain. These temperature values agree very well with the range determined from the ultraviolet IUE spectra. Moreover the rotation effects on frequencies cannot be neglected for such rapid rotator as SY Equ."
    },
    "0609/astro-ph0609230_arXiv.txt": {
        "abstract": "{} {To study the nature of Bulge AGB stars and in particular their circumstellar dust, we have analysed mid-infrared spectra obtained with the ISOCAM CVF spectrometer in three Bulge fields. } {The ISOCAM 5--16.5~$\\mu$m CVF spectra were obtained as part of the ISOGAL infrared survey of the inner Galaxy. A classification of the shape of the 10~$\\mu$m dust feature was made  for each case. The spectra of the individual sources were modelled using a radiative transfer model. Different combinations of amorphous silicates and aluminium-oxide dust were used in the modelling.} {Spectra were obtained for 29 sources of which 26 are likely to be Bulge AGB stars. Our modelling shows that the stars suffer mass loss rates in the range of $10^{-8} - 5 \\times 10^{-7}$ M$_\\odot /$ yr, which is at the low end of the mass-loss rates experienced on the Thermally Pulsing AGB. The luminosities range from 1,700 to 7,700~L$_\\odot$ as expected for a population of AGB stars with M$_{\\rm init}$ of 1.5 - 2~M$_\\odot$. In agreement with the condensation sequence scenario, we find that the dust content is dominated by Al$_2$O$_3$ grains in this sample of low mass-loss rate stars.} {}  \\keywords {stars: AGB -- stars: evolution -- stars: mass loss -- Galaxy: bulge -- dust -- Infrared: stars } ",
        "introduction": "Asymptotic Giant Branch (AGB) stars undergo substantial mass loss ($10^{-8}-10^{-4}$ M$_\\odot$ yr$^{-1}$) which effectively drives their stellar evolution (e.g. Habing 1996).  Radial pulsations lift matter high above the photosphere where it cools down and a complex chemistry - based on either carbon or oxygen preponderance in the photosphere - is initiated, leading to the formation of molecules and dust. Radiation pressure from the central star blows away the dust particles, which in turn collide with the gas causing the massive outflow.  The formation scenario and the nature of dust in the circumstellar environment of oxygen-rich AGB stars is still very much an issue for debate. On the basis of IRAS LRS data, several studies were performed to investigate the dust of large samples of AGB stars of different variability types and mass-loss rates (e.g. Onaka et al. 1989, Stencel et al. 1990, Little-Marenin et al. 1990, Sloan \\& Price, 1995 and 1998, Hron et al. 1997). Another study, using groundbased 10~$\\mu$m spectroscopy, was performed by Speck et al. (2000). In the 10~$\\mu$m region, a sequence of features was found starting from broad ones peaking longward of 11~$\\mu$m and ending with narrow ones peaking at 9.7~$\\mu$m. The interpretation of the broad feature is that it is caused by alumina (amorphous Al$_2$O$_3$) whereas that at 9.7~$\\mu$m is the `classic' silicate feature caused by amorphous silicate dust. Between the two extremes several intermediate types exist and different attempts have been made to classify them,  based on observational characteristics (Little-Marenin et al. 1989, Sloan \\& Price, 1995 and 1998, Speck et al. 2000).  In the theoretical dust condensation sequence (Tielens 1990) it is expected that refractory oxides will form first, close to the star at a relatively high temperature of about 1500 K. If the densities in the cooler regions further out in the circumstellar shell are high enough, gas-solid and solid-solid reactions will further transform these minerals and lead to the formation at lower temperatures of silicate dust. If at some intermediate step in the dust condensation sequence however the densities are too low for a certain reaction to occur, the condensation stops at that point and may not continue to the formation of silicate dust -- this is known as ``freeze-out'' of the dust condensation sequence. Consequently, stars with low mass loss rates will predominantly exhibit an oxide mineralogy -- since the densities are too low for any further step to occur -- while stars with high mass loss rates will form silicate dust in copious amounts. An alternative scenario was suggested by Sloan \\& Price (1998) where the C/O ratio drives the condensation sequence. According to these authors, stars with low C/O ratios produce silicate dust whereas stars with high C/O ratios will produce mainly Al$_2$O$_3$ grains.  Since it is generally believed that both mass loss rates and the C/O ratio of oxygen-rich AGB stars increase as a star evolves on the AGB, studying the evolution of the circumstellar dust as a star evolves on the AGB provides a way to discriminate between the two scenarios. Indeed, in the first scenario (dust condensation sequence and freeze-out) one expects to see oxide dust in the earlier phases of the AGB when the mass loss rate is still low, and an increasing silicate dust content as the mass loss rate increases with the evolution of the star along the AGB. Such a time evolution is exactly opposite to what should happen according to the scenario suggested by Sloan \\& Price (1998). Work done by Heras \\& Hony (2005) on ISO-SWS spectra of a Galactic AGB  star sample and by Dijkstra et al. (2005) on ISO-PHT and ISO-CAM spectra of AGB stars in the Magellanic Clouds support the sequence where the formation of the different dust types is related to the mass-loss rates.  A serious limitation to the previous studies has been that because of the limited sensitivities of groundbased, IRAS and even of the ISO-SWS instruments, only relatively nearby bright objects have been observed spectroscopically. With ISO-PHT and ISO-CAM it was possible to observe stars in the Magellanic Clouds, but only the brightest ones.  The stars in these studies represent a wide range in parameters such as stellar masses, mass-loss rates and variability characteristics.  The evolutionary connection between the different stars and the interpretation of the observational sequences in terms of e.g. dust formation scenarios or stellar evolution has proved difficult to establish.  In this paper we present ISOCAM -- CVF ($5 - 16.5 \\mu$m) spectral measurements on a more homogeneous sample of AGB stars in three ``intermediate'' galactic bulge fields ($|$b$| \\geq 1^{\\circ}$) which were observed as part of the ISOGAL survey (Omont et al. 2003). The ISOGAL survey provided a 5-band point source catalogue including near-infrared (DENIS), 7 and 15~$\\mu$m band photometry.  In total about 16 square degrees towards the inner Galaxy were observed down to a flux level of 10-20~mJy in the mid-infrared, detecting $\\approx 10^5$ sources, mostly AGB stars, red giants and young stars. The survey has resulted in a complete sample of AGB stars in the galactic bulge fields.  Omont et al. (1999) and Ojha et al. (2003, hereafter OOS) demonstrated that the ISOGAL sources in these fields are AGB stars with a range of mass-loss rates between $10^{-8}$ and $10^{-5}$ M$_\\odot$ yr$^{-1}$.  The stars at these latitudes in the direction of the Bulge are believed to belong to the Bulge population which show only a small range in masses. According to Zoccali et al.\\ 2003, the Bulge giants have evolved from a population of stars of $\\approx 1$ M$_\\odot$.  In their analysis of Bulge Mira variables, Groenewegen \\& Blommaert (2005) find slightly higher stellar masses: 1.5 to 2~M$_\\odot$. We do not know the metallicities of our sample of AGB stars. Rich \\& Origlia (2005) find for a sample of M giants in the Baade's Window that their metallicity distribution peaks at slightly subsolar metallicity ($< [{\\rm Fe / H}]> = -0.190 \\pm 0.020$ with a 1$\\sigma$ dispersion of 0.080 $\\pm$ 0.015).  Although, the metallicities of the stars in our sample may be uncertain, we will assume that the main difference between individual objects is their age on the AGB, which is characterized by different luminosities and mass-loss rates.  A further description of the sample is given in Section~4. In the following sections we describe the ISO observations and the data reduction. After the presentation of the spectra in Section~5 we present the results of radiative transfer modelling. These results are discussed in Section~7. ",
        "conclusions": "We have presented the ISOCAM CVF spectra of a sample of AGB stars previously detected in the ISOGAL survey. These are the first mid-IR spectra of Bulge AGB stars with mass-loss rates down to $10^{-8}$ M$_\\odot /$ yr. Modelling of the ISOGAL photometry and the CVF spectra show that our sample consists of AGB stars with luminosities between 1,700 to 7,700~L$_\\odot$, in agreement with the range of luminosities of Thermally Pulsing AGB stars with initial masses of 1.5 to 2~M$_\\odot$. The mass-loss rates of our sample are in the range of $10^{-8} - 5 \\times 10^{-7}$ M$_\\odot /$ yr. We have demonstrated that the availability of the mid-IR spectra, which allow the identification of the circumstellar dust type, plays an important role in determining accurate dust mass-loss rates.  Our sample, that mainly contains low mass-loss rate stars, has a dust content that is dominated by amorphous aluminium oxide dust."
    },
    "0609/astro-ph0609006_arXiv.txt": {
        "abstract": "We derive the electron temperature gradient in the Galactic disk using a sample of \\hii\\ regions that spans Galactocentric distances 0--17 kpc. The electron temperature was calculated using high precision radio recombination line and continuum observations for more than 100 \\hii\\ regions. Nebular Galactocentric distances were calculated in a consistent manner using the radial velocities measured by our radio recombination line survey. The large number of nebulae widely distributed over the Galactic disk together with the uniformity of our data provide a secure estimate of the present electron temperature gradient in the Milky Way. Because metals are the main coolants in the photoionized gas, the electron temperature along the Galactic disk should be directly related to the distribution of heavy elements in the Milky Way. Our best estimate of the electron temperature gradient is derived from a sample of 76 sources for which we have the highest quality data.  The present gradient in electron temperature has a minimum at the Galactic Center and rises at a rate of 287 $\\pm$ 46 \\kkpc.  There are no significant variations in the value of the gradient as a function of Galactocentric radius or azimuth.  The scatter we find in the \\hii\\ region electron temperatures at a given Galactocentric radius is not due to observational error, but rather to intrinsic fluctuations in these temperatures which are almost certainly due to fluctuations in the nebular heavy element abundances. Comparing the \\hii\\ region gradient with the much steeper gradient found for planetary nebulae suggests that the electron temperature gradient evolves with time, becoming flatter as a consequence of the chemical evolution of the Milky Way's disk. ",
        "introduction": "Churchwell \\& Walmsley (1975) pioneered \\hii\\ region radio recombination line (RRL) studies of the relationship between nebular electron temperatures, \\te, and Galactocentric distance, \\rgal, (Churchwell et al. 1978; Wink et al. 1983; Shaver et al. 1983; and others).  Because RRLs are not obscured by interstellar dust, relatively faint \\hii\\ regions at extremely large distances from the Sun could be detected.  They found that there was a Galactic temperature gradient wherein \\te\\ is low in the Galactic Center and increases with \\rgal.  Such a gradient was first observed in other nearby spiral galaxies by Searle (1971), Rubin et al. (1972) and Smith (1975).  This Milky Way electron temperature gradient was confirmed by radio continuum emission (Omar et al. 2002) and by [\\oiii]\\ forbidden line optical observations (Peimbert et al. 1978; Deharveng et al. 2000).  Because heavy elements cool photoionized gas, \\hii\\ region electron temperatures are directly related to the heavy element abundance: low \\te\\ corresponds to higher heavy element abundances because of the greater cooling rate and vice versa.  Consequently, \\te\\ gradients should be inversely related to metal abundance gradients (but see \\S 6).  Because there has been more stellar processing in the inner Galaxy one expects that on average the metallicity decreases as a function of \\rgal.  The existence of a Galactic gradient in \\hii\\ region electron temperature is now firmly established. Nevertheless, there are still uncertainties in the magnitude of the gradient and the possible existence of real variations, both in \\rgal\\ and Galactic azimuth, of nebular \\te.  Extant RRL studies yield \\te\\ gradients that, roughly, vary from 250--440 \\kkpc.  Discrepancies in the results obtained from different studies may be attributed to several factors, including the source sample, the \\te\\ derivation, etc.  Since the exact value of the \\te\\ gradient provides an important constraint on models for Galactic chemical evolution, it must be determined as accurately as possible.  Here we derive \\hii\\ region electron temperatures based on radio recombination line and continuum data for a large sample of nebulae widely distributed across the Galactic disk. The RRL data are of unprecedented sensitivity compared with previous studies.  We examine anew the Galactic temperature gradient and assess the magnitude and origin of the \\te\\ dispersion at a given \\rgal. In \\S\\,\\ref{sec:obs} we describe our \\hii\\ region sample. In \\S\\,\\ref{sec:te} we derive the nebular electron temperatures and discuss non-LTE effects.  Nebular Galactocentric distances, \\rgal, and heliocentric distances, \\dsun, are derived in \\S\\,\\ref{sec:dist}.  In \\S\\,\\ref{sec:tegrad} our new determination of the electron temperature gradient in the Galactic disk is made.  Here we also investigate a possible spatial variation of this gradient.  We discuss the astrophysical implications of our efforts in \\S\\,\\ref{sec:discuss}. ",
        "conclusions": "\\label{sec:discuss}  \\subsection{The Electron Temperature of Galactic H\\thinspace II Regions} \\label{sec:hiite}  The electron temperature of an \\hii\\ region in thermal equilibrium is set by the balance of competing heating and cooling mechanisms.  It is therefore somewhat surprising that there is a Galactic \\te\\ gradient. There are at least four physical properties that could effect \\te: (1) the effective temperature of the ionizing star, $T_{\\rm eff}$, which sets the hardness of the radiation field exciting (and heating) the nebula; (2) the electron density---collisional de-excitation in the high $n_{e}$ \\hii\\ region will inhibit cooling and increase \\te; (3) dust grains which effect the heating and cooling in complex ways; and (4) heavy element abundance which increases cooling and decreases \\te\\ (Garay \\& Rodr\\'{\\i}guez 1983).  Rubin (1985) explored how metallicity, gas density, and the stellar effective temperature effect the average electron temperature in model \\hii\\ regions.  These models predict changes in \\te\\ of 7,000\\kel\\ for a factor of 10 change in metal abundance, 2,900\\kel\\ for a change in density from 100 to \\expo{5}\\percc, and 1,300\\kel\\ for a change in $T_{\\rm eff}$ from 33,000 to 45,000\\kel\\ (B0 to O5 spectral type).  Dust grains are known to play a significant role in the heating and cooling of \\hii\\ regions (e.g., Mathis 1986; Baldwin et al. 1991; Shields \\& Kennicutt 1995).  Photoelectric heating occurs as electrons are ejected from dust grains while the gas is cooled by the collisions of fast particles with grains.  The electron temperature will decrease with distance from the star as the ionizing radiation field is attenuated by dust grains.  But the electron temperature will also increase as coolants are depleted onto dust grains.  Taking these competing factors into account Oliveira \\& Maciel (1986) conclude that dust grains do not significantly contribute to the observed \\te\\ gradient with a maximum variation of 500\\kel. Therefore metallicity is the most sensitive factor that sets the nebular \\te\\ value and so \\hii\\ region metal abundance variations are the best interpretation for the observed \\te\\ gradient.  The electron temperature need not be a constant, however, and the observational methods used to determine \\te\\ are sensitive to different regions of the nebula.  The two main methods used to determine electron temperatures in \\hii\\ regions are (1) recombination line-to-continuum ratios, such as hydrogen RRL and continuum emission; and (2) forbidden line ratios, such as the ratio [\\oiii]$\\lambda$\\,4363 / ($\\lambda$\\,4959 + $\\lambda$\\,5007).  The recombination line method is weighted towards the lower temperature regions with a weak dependence on \\te, while the forbidden line method is weighted towards the higher temperature regions with a strong dependence on \\te\\ (e.g., Peimbert 1967).  Therefore if temperature structure exists in \\hii\\ regions these methods can produce different electron temperatures.  Observations of [\\oiii] and oxygen recombination lines in Orion show such temperature fluctuations (Esteban et al. 1998).  Comparison of hydrogen RRL and [\\oiii] electron temperatures is more complicated since the emission lines arise from different species. The intensities of the RRLs of H relative to the underlying continuum, for example, should give an estimate of the electron temperature in the whole ionized H region (H$^+$ zone), while observations of the [\\oiii] optical emission should give an estimate of the electron temperature in the O$^{++}$ zone.  Photoionization models indicate that the electron temperature should be higher in the outer regions since the ionizing radiation becomes harder with distance from the exciting star and the very efficient coolants like O$^{++}$ are located close to the exciting star (Stasi\\'nska 1980; Garnett 1992). This suggests that $T_e({\\rm O}^{++}) \\leq T_e({\\rm H}^+) \\leq T_e({\\rm O}^+)$ (Stasi\\'nska 1990; Stasi\\'nska \\& Shaerer 1997; Deharveng et al. 2000).  The temperature measured via the RRL, however, is not strictly $T_e({\\rm H}^+)$ and that measured with the optical [\\oiii] forbidden lines is not strictly $T_e({\\rm O}^{++})$. Because the recombination lines are emitted preferentially in low-temperature regions, $T_e ({\\rm radio}) \\leq T_e({\\rm H}^+)$, and because the optical forbidden lines are enhanced in high temperature regions, $T_e$([\\oiii])$ \\ge T_e({\\rm O}^{++}$).  Despite this complexity the electron temperatures determined by these two methods towards the same \\hii\\ region produce values of \\te\\ that are similar to within the uncertainties (Deharveng et al. 2000).  Wink et al. (1983) find no correlation between \\te\\ and either the Lyman continuum flux, $N_{\\rm L}$, used to probe different stellar effective temperatures, or the electron density, $n_{e}$. Theoretically, higher values of $N_{\\rm L}$ and $n_{e}$ should increase \\te\\ (Rubin 1985).  But there are \\hii\\ regions with high $N_{\\rm L}$ and low $n_{e}$ (e.g., the Rosette nebula) that will increase the scatter in any such analysis.  Shaver et al. (1983) found a correlation with \\te\\ by dividing the \\hii\\ regions into two groups of either high or low values of {\\it both} $N_{\\rm L}$ and $n_{e}$. But the correlation seems weak since many of the points at smaller \\rgal\\ where the correlation is best are nebulae that appear to have systematically lower electron temperatures (i.e. the H\\,109\\,$\\alpha$ survey sources).  The ultracompact \\hii\\ region survey of Afflerbach et al. (1996) probes nebula of much higher electron density than other surveys.  Because these observations were made with an interferometer any diffuse gas within their primary beam will be spatially filtered. Their derived electron temperatures are higher by $\\sim$\\,1,000\\kel\\ compared with other surveys of classical \\hii\\ regions.  We explored the effects of $N_{\\rm L}$ and $n_{e}$ on electron temperature in our sample and find no correlation.  We used an approach similar to Shaver et al. (1983) wherein we selected sources either below or above a threshold value of the excitation and electron density ($\\log({\\rm N_{\\rm L}}) = 49.5\\,{\\rm s^{-1}}$ and $n_{e} = 150$\\percc). RRL and continuum surveys of classical \\hii\\ regions with single-dish telescopes probe \\hii\\ region complexes that typically contain more than one ionizing star with different spectral types.  Such regions consist of both low and high density gas (e.g., Balser et al. 1999). But does this account for the observed dispersion of the \\te\\ gradient as has been suggested (e.g., Shaver et al 1983)?  Another possibility is that the dispersion is caused by variations in metallicity.  After all, stellar abundances at a given \\rgal\\ show a dispersion in abundance. Interpretation of the stellar abundances is complicated because samples typically contain stars of varying age which may well have been formed at a different \\rgal\\ than their current location (Edvardsson et al. 1993). Correcting for birth location and age Edvardsson et al. (1993) find a scatter of 0.5 dex in the solar neighborhood.  Friel et al. (2002) find a scatter of about half that in disk star cluster abundances. A scatter of 0.3 dex in heavy element abundance will produce a scatter of 2,000 K in \\te, entirely consistent with our observed scatter.  For \\hii\\ regions, Mehringer et al. (1993) suggest differences in metallicity as the best candidate for the $\\sim 2,000$\\kel\\ change in \\te\\ between components 1 and 2 in the Sgr\\,B complex.  Since our survey spans a large region of the Galaxy, both radially and azimuthally, we were able to explore such effects by comparing \\te\\ at a given \\rgal\\ over a large range of azimuth.  The results were inconclusive.  Sensitive RRL observations in critial areas of the Galaxy could reveal such variations and place important constraints on Galactic chemical evolution models.  \\subsection{The Galactic Center}\\label{sec:gc}  Our nebular sample F fit to the electron temperature gradient excludes the Galactic Center (GC) sources Sgr\\,B2 and G1.13$-$0.1 and includes the outer Galaxy nebulae S209 and S209\\,N.  This sample produces the steepest \\lte\\ gradient, thus providing an extreme constraint on Galactic chemical evolution models.  Early RRL observations of \\hii\\ regions in the GC produced electron temperatures that were higher than those expected from an extrapolation of the electron temperature gradient in the Galactic disk (Mezger et al. 1979; Downes et al. 1980; Wink et al. 1983). Since the chemical evolution of the GC may be different than the disk, it is not surprising that electron temperatures in the GC do not follow the gradient in the disk.  Higher spatial resolution observations reveal a wide range of electron temperatures ($\\sim 4,000-10,000\\,$K) in Galactic Center \\hii\\ regions (Roelfsema et al. 1987; Gaume \\& Claussen 1990; Mehringer et al. 1992,1993; De Pree et al. 1995, 1996; Lang et al. 1997, 2001).  The large range of electron temperatures may be due to the complexity of the GC.  Real variations in metallicity, electron density, and excitation will effect $T_{e}$.  Some of the higher electron temperatures are overestimated when the radio continuum is contaminated with non-thermal emission.  Nevertheless, in some objects very high electron temperatures ($T_{e} \\sim 20,000\\,$K) appear to be real where the gas may not be in equilibrium (Mehringer et al. 1995). Since the zero level for our radio continuum temperature was determined away from the thermal object but not far from the Galactic disk, any smooth non-thermal emission should have been removed.  \\subsection{The Galactic O/H Abundance Gradient}\\label{sec:OtoH}  The electron temperature may be converted to oxygen abundance using a correlation between metal abundance and temperature. Because of the broad interest in abundance gradients, we provide here one such conversion of our \\lte\\ gradient to an O/H gradient, deferring a more extended discussion to a future paper. Using the relation between O/H and $T_e$ of Shaver et al. (1983) yields the result shown in Figure~\\ref{fig:ohgrad}. The O/H gradient is about $-$0.04 dex kpc$^{-1}$ in agreement with results from Afflerbach et al. (1996), Deharveng et al. (2000), Pilyugin (2003) and Esteban et al. (2005). The exact fit to the 76 sample E nebulae in Figure~\\ref{fig:ohgrad} is: \\begin{equation}\\label{eq:ohgradient} {\\rm log\\,(O/H) + 12 = (8.958 \\pm 0.052)~+~(-0.043 \\pm 0.007)\\,\\rgal\\,(kpc)} \\end{equation} with a correlation coefficient of $r=-0.59$.  If our observed \\lte\\ dispersion of $\\sim$\\,1,100 K is caused by real abundance fluctuations, then this would correspond to an [O/H] abundance dispersion of $\\sim$\\,0.16 dex."
    },
    "0609/gr-qc0609041_arXiv.txt": {
        "abstract": "The interaction of a Reissner-Nordstr\\\"om black hole and a charged massive particle is studied in the framework of perturbation theory. The particle backreaction is taken into account, studying the effect of general static perturbations of the hole following the approach of Zerilli. The solutions of the combined Einstein-Maxwell equations for both perturbed gravitational and electromagnetic fields at first order of the perturbation are exactly reconstructed by summing all multipoles, and are given explicit closed form expressions. The existence of a singularity-free solution of the Einstein-Maxwell system requires that the charge to mass ratios of the black hole and of the particle satisfy an equilibrium condition which is in general dependent on the separation between the two bodies. If the black hole is undercritically charged (i.e. its charge to mass ratio is less than one), the particle must be overcritically charged, in the sense that the particle must have a charge to mass ratio greater than one. If the charge to mass ratios of the black hole and of the particle are both equal to one (so that they are both critically charged, or \\lq\\lq extreme''), the equilibrium can exist for any separation distance, and the solution we find coincides with the linearization in the present context of the well known Majumdar-Papapetrou solution for two extreme Reissner-Nordstr\\\"om black holes. In addition to these singularity-free solutions, we also analyze the corresponding solution for the problem of a massive particle at rest near a Schwarzschild black hole, exhibiting a strut singularity on the axis between the two bodies. The relations between our perturbative solutions and the corresponding exact two-body solutions belonging to the Weyl class are also discussed. ",
        "introduction": "The problem of the effect of gravity on the electromagnetic field of a charged particle leading to the consideration of the Einstein-Maxwell equations has been one of the most extensively treated in the literature, resulting in exact solutions (see \\cite{exactsols} and references therein) as well as in a variety of approximation methods \\cite{fermi,whittaker,copson,hannijth,CoW,HR,bicdev,linet,leaute}.  We address in this article the problem of a massive charged particle of mass $m$ and charge $q$ in the field of a Reissner-Nordstr\\\"om geometry describing a static charged black hole, with mass $\\mathcal{M}$ and charge $Q$. We solve this problem by the first order perturbation approach formulated by Zerilli \\cite{Zerilli} using the tensor harmonic expansion of the Einstein-Maxwell system of equations. Note that a similar approach was used by Sibgatullin and Alekseev \\cite{sibgatullin} to study electrovacuum perturbations of a Reissner-Nordstr\\\"om black hole. A manifestly gauge invariant (hamiltonian) formalism was developed by Moncrief \\cite{moncrief1,moncrief2,moncrief3}. The relations between these two different treatments of perturbations have been extensively investigated by Bi{\\v c}\\'ak \\cite{bicak}.  Both the mass and charge of the particle are described by delta functions. The source terms of the Einstein equations contain the energy-momentum tensor associated with the particle's mass, the electromagnetic energy-momentum tensor associated with the background field as well as additional interaction terms, to first order in $m$ and $q$, proportional to the product of the square of the charge of the background geometry and the particle's mass ($\\sim Q^2m$) and to the product of the charges of both the particle and the black hole ($\\sim qQ$). These terms give origin to the so called \\lq\\lq electromagnetically induced gravitational perturbation'' \\cite{jrz2}. On the other hand, the source terms of the Maxwell equations contain the electromagnetic current associated with the particle's charge as well as interaction terms proportional to the product of the black hole's charge and the particle's mass ($\\sim Qm$), giving origin to the \\lq\\lq gravitationally induced electromagnetic perturbation'' \\cite{jrz1}. We give in Section II the basic equations to be integrated in order to solve this problem.  Before proceeding to the solution of the most general problem, we recall a variety of particular approaches followed  in the current literature within the simplified framework of test field approximations \\cite{HR,linet,leaute,bonnor}. In Section III, the test field solutions for a charged particle at rest in the field of a Schwarzschild as well as a Reissner-Nordstr\\\"om black hole are summarized, under the conditions $q/m\\gg1$, $m\\approx0$ and $q\\ll{\\mathcal M}$, $q\\ll Q$. The case of a test charge in the Schwarzschild black hole is reviewed first in Section III A; the corresponding solution is obtained by using the vector harmonic expansion of the electromagnetic field in curved space \\cite{HR}. The conditions above imply the solution of the Maxwell equations only in a fixed Schwarzschild metric, since the perturbation to the background geometry given by the electromagnetic stress-energy tensor is second order in the particle's charge. As a result, no constraint on the position of the test particle follows from the Einstein equations and the Bianchi identities: any position of the particle is allowed, and the solution for the corresponding electromagnetic field has been widely analyzed \\cite{HR,linet}.  In Section III B we recall this same test field approximation applied to the case of a Reissner-Nordstr\\\"om black hole. This treatment is due to Leaute and Linet \\cite{leaute}; in analogy with the Schwarzschild case, they used the vector harmonic expansion of the electromagnetic field holding the background geometry fixed. However, this \\lq\\lq test field approximation'' is not valid in the present context. In fact, in addition to neglecting the effect of the particle mass on the background geometry, this treatment also neglects the electromagnetic induced gravitational perturbation terms linear in the charge of the particle which would contribute to modifying the metric as well. In Section III C we recall how, within this same approximation still neglecting these same feedback terms, Bonnor \\cite{bonnor} studied the condition for the equilibrium involving the black hole and particle parameters $Q, {\\mathcal M}, q, m$ as well as their separation distance $b$. He finds the constraint \\begin{equation} \\label{bonnoreqcond0} m=\\frac{qQb}{{\\mathcal M}b-Q^2}\\left(1-\\frac{2{\\mathcal M}}{b}+\\frac{Q^2}{b^2}\\right)^{1/2}\\ , \\end{equation} If the black hole is \\lq\\lq extreme'' or \\lq\\lq critically charged'' satisfying $Q/\\mathcal{M}=1$, then the particle must also have the same ratio $q/m=1$, and equilibrium exists independent of the separation. In the general non-extreme case $Q/\\mathcal{M}<1$ there is instead only one position of the particle which corresponds to equilibrium, for given values of the charge-to-mass ratios of the bodies. In this case the particle charge-to-mass ratio satisfies $q/m>1$. This condition corresponds to the overcritically charged case for a Reissner-Nordstr\\\"om solution with the same parameters $m$ and $q$ corresponding to a naked singularity.  Already the Leaute and Linet generalization modifies in a significant way the results obtained in the case of a Schwarzschild black hole. Indeed a new phenomenon occurs in this case: as the hole becomes extreme an effect analogous to the Meissner effect for the electric field arises, with the electric field lines of the test charge being forced outside the outer horizon \\cite{elbaproc}. This result by itself justifies the need of addressing this problem in the more general Zerilli approach duly taking into account all the first order perturbations.  The correct way to attack the problem is, thus, to solve the linearized Einstein-Maxwell equations following Zerilli's first order tensor harmonic analysis \\cite{Zerilli}. In Section \\ref{schwpert} we first consider the simplest special case of a massive neutral particle at rest near a Schwarzschild black hole. We first show explicitly that a perturbative solution for this problem free of singularities cannot exist. We then give the explicit form of the perturbation  corresponding to a stable configuration when there is the presence of a \\lq\\lq strut'' between the particle and the black hole, corresponding to a conical singularity. This result is a direct consequence of the classical work by Einstein and Rosen \\cite{einstein}, who considered the Weyl class double Chazy-Curzon solution \\cite{curzon} for two point masses placed on the symmetry axis at a fixed distance. We find it very helpful to use a new gauge condition particularly adapted to this problem which differs from the corresponding Regge-Wheeler gauge in the Schwarzschild case, by adding additional terms to the third Regge-Wheeler gauge condition to take into account the additional terms in the field equations. The resulting perturbed metric we obtain is the linearized form of the exact solution representing two collinear uncharged black holes in a static configuration belonging to the Weyl class \\cite{exactsols}. We review some of their properties and peculiar features of these solutions in the Appendix.  We then turn to the general case of a charged massive particle at rest in a Reissner-Nordstr\\\"om background. This problem needs the simultaneous analysis of both the vectorial perturbation to describe the electromagnetic perturbation and the tensorial perturbation necessary to analyze the perturbation by the particle mass as well as the electromagnetically induced gravitational perturbation and the gravitationally induced electromagnetic perturbation.  The Einstein-Maxwell equations reduce to a system of $6$ coupled ordinary differential equations for $4$ unknown functions determining both the gravitational and electromagnetic perturbed fields. Then, by imposing the compatibility of the system (or equivalently, requiring that the Bianchi identities be fulfilled), we obtain an equilibrium condition for the system which coincides with the condition (\\ref{bonnoreqcond0}) obtained by Bonnor in his simplified approach. This is surprising, since our result has been obtained within a more general framework, and both the gravitational and electromagnetic  fields are different from those used by Bonnor.  We then succeed in the exact reconstruction of both the perturbed gravitational and electromagnetic fields by summing all multipoles. The perturbed metric we derive is spatially conformally flat, free of singularities and thus cannot belong to the Weyl class two-body solutions, which are characterized by the occurrence of a conical singularity on the axis between the bodies. The asymptotic mass measured at large distances by the Schwarzschild-like behaviour of the metric of the whole system consisting of black hole and particle is given by \\begin{equation} M_{\\rm eff}={\\mathcal M}+m+E_{\\rm int}\\ , \\end{equation} where the interaction energy turns out to be \\begin{equation} E_{\\rm int}=-m\\left[1-\\left(1-\\frac{{\\mathcal M}}{b}\\right)\\left(1-\\frac{2{\\mathcal M}}{b}+\\frac{Q^2}{b^2}\\right)^{-1/2}\\right]\\ . \\end{equation}  Recently Perry and Cooperstock \\cite{perry} and Bret\\'on, Manko and S\\'anchez \\cite{manko} have used multi-soliton techniques to study exact electrostatic solutions of the Einstein-Maxwell equations representing the exterior field of two arbitrary charged Reissner-Nordstr\\\"om bodies in equilibrium. They have shown by numerical methods that gravitational-electrostatic balance without intervening tension or strut can occur for non-critically charged spherically symmetric bodies in the case of one black hole with $Q/\\mathcal{M}<1$ and one naked singularity. However, they have not obtained an explicit dependence (in algebraic form) of the balance condition on the separation of the bodies, so we cannot compare our analytical formula with their numerical results. Attention is devoted also to the limiting case in which $Q/\\mathcal{M}=q/m=1$. In this case our solution coincides with the linearized form of the exact solution of Majumdar and Papapetrou \\cite{maj,pap}, where equilibrium exists independent of the separation between the bodies.  Within our approximation for $Q/\\mathcal{M}<1$, we find that a horizon seems to form around the particle very close to it where the gravitational perturbation is not small compared to the background metric and hence the approximation is no longer valid. The study of the full nonlinear problem, namely the corresponding exact two-body solution of the Einstein-Maxwell equations is therefore required. Attempts in this direction are being pursued by Belinski and Alekseev \\cite{belinski} by using soliton techniques. ",
        "conclusions": "The problem of the interaction of a massive charged particle at rest with a Reissner-Nordstr\\\"om black hole has been studied taking into account both electromagnetic and gravitational perturbations on the background fields due to the presence of the particle. Following Zerilli's approach to the perturbations of a charged static black hole, we derived the corresponding solutions of the linearized Einstein-Maxwell equations for both the electromagnetic and gravitational perturbation functions. We were able to exactly reconstruct these functions summing over all multipoles, giving closed form expressions for the components of the perturbed metric as well as the electromagnetic field. A detailed analysis of the properties of this solution will be presented elsewhere."
    },
    "0609/astro-ph0609376_arXiv.txt": {
        "abstract": "We reproduce the expansion velocity--radius ($V_{\\rm{exp}}$--$R_{\\rm{n}}$) relation in planetary nebulae by considering a simple dynamical model, in order to investigate the dynamical evolution and formation of planetary nebulae.  In our model, the planetary nebula is formed and evolving by interaction of a fast wind from the central star with a slow wind from its progenitor, the AGB star.  In particular, taking account of the mass loss history of the AGB star makes us succeed in the reproduction of the observed $V_{\\rm{exp}}$-$R_{\\rm{n}}$ sequence.  As a result, examining the ensemble of the observational and theoretical evolution models of PNe, we find that if the AGB star pulsates and its mass loss rate changes with time (from $\\sim 10^{-6.4}M_{\\odot}$ yr$^{-1}$ to $\\sim 10^{-5}M_{\\odot}$ yr$^{-1}$), the model agrees with the observations.  In terms of observation, we suggest that there are few planetary nebulae with larger expansion velocity and smaller radius because the evolutionary time-scale of such nebulae is so short and the size of nebulae is so compact that it is difficult for us to observe them. ",
        "introduction": "Planetary nebulae (PNe) are the objects apart from the asymptotic giant branch (AGB), evolving to white dwarfs. It is thought that a stellar wind is important for the formation and evolution of PNe. To explain the features of PNe, Kwok et al.  (1978) propose the interacting stellar wind (ISW) theory. After that many authors have studied PNe theoretically and observationally based on the ISW theory.  Observationally, before the ISW theory was presented, a trend of the expansion velocity ($V_{\\rm{exp}}$) of the gas of PNe was found (Kwok et al. 1978), and it has been interpreted in two different ways: one is in terms of a dynamical evolution of PN (Smith 1969; Bohuski and Smith 1974; Sabbadin and Hamzaoglu 1982; Robinson et al. 1982), another is in terms of the evolution of a central star (Renzini 1979).  Based on such studies, Sabbadin et al. (1984; hereinafter SGBO) proposed two models to interpret that trend, the ``two-wind'' model and the ``two-phase'' model.  The former means that the planetary nebula is the region of interaction of a fast wind emitted by the nucleus with a slow wind expelled by its progenitor.  They added the effect of the momentum transfer from radiation field to the model suggested by Kwok et al.(1978) and Kwok(1982).  The latter means that the planetary nebula is a Str$\\ddot{\\rm{o}}$mgren sphere evolving in an expanding nebula which results from a sudden ejection by the AGB star.  They conclude only the ``two-phase'' model agrees with the observational data because the ``two-wind'' model can not reproduce the observed mass ($M_{\\rm{n}}$) and radius ($R_{\\rm{n}}$) relation.  In the two-wind model the resulting mass of nebula is much smaller than that derived from observations (the mass decrement problem), while it matches the $V_{\\rm{exp}}$--$R_ {\\rm{n}}$ relation. The mass decrement problem of the two-wind model is also stressed and examined in Schmidt-Voigt \\& K$\\ddot{\\rm{o}}$ppen (1987a,b).  The ``two-wind'' model had another fault. SGBO didn't take account of the effect of shock sufficiently. As Marten \\& Sch$\\ddot{\\rm{o}}$nberner (1991) stated, the simplification of SGBO corresponds to the fact that the width of the shell is infinitesimally thin. To resolve the mass decrement problem, it is essential to examine the width of the shell (Marten \\& Sch$\\ddot{\\rm{o}}$nberner 1991). Then, we confirm in this paper that it is easy to resolve the mass decrement problem if the width of the shell is determined as the effect of the shock. In this standing point, we try to study the two-wind model to satisfy $M_{\\rm{n}}$--$R_{\\rm{n}}$ and $V_{\\rm{exp}}$--$R_{\\rm{n}}$ relations. We would like to comment that after SGBO, they developed their study on PNe furthermore (e.g. Sabbadin et al. 2005).   We also examine the effect of the time-dependent mass loss rate of the AGB star, then we succeed in reproduction of the observed mass-radius relation simultaneously with the observed $V_{\\rm{exp}}$-$R_{\\rm{n}}$ relation. The importance of the mass-loss to explain the evolution of PNe has been stressed by  Marten \\& Sch$\\ddot{\\rm{o}}$nberner (1991). In this paper, we insist furthermore that our modified ``two-wind'' model reproduces the ensemble of observational data adequately if the radial density profile of AGB matter has a sudden decrement at about 0.1pc from the central star. This decrement is the result of the time-dependent mass-loss history of the AGB star. Thus, we reformulate the model concerning the mass-loss history of the AGB star.  Thus, to explore the evolution and formation of PNe, we revisit the ``two-wind\" model in this paper. In $\\S$2 we describe our model. In $\\S$3 the results of the calculation compared with the observational data are presented.  In $\\S$4, we discuss the results, and summarize our study in $\\S$5. ",
        "conclusions": "In this section, we examine our model compared with the basic model basing on SGBO and discuss the physics of PN evolution.  \\subsection{Difference from a basic model}\\label{sec:difference} As already mentioned, there are two problems in the classical two-wind model by SGBO: (a) the decrement of resulting mass, and (b) the discrepancy in the $R_{\\rm{n}}$-$M_{\\rm{n}}$ relation between observations and the model. Although SGBO's model explains the $2V_{\\rm{exp}}$-$R_{\\rm{n}}$ relation, the authors abandon their idea because of the problems (a) and (b). On the other hand, we take account of the effect of shock and time-depending mass loss rate of the AGB star, then the problems (a) and (b) are solved easily. {}From this point, we convince for the evolution of PN to consider the influence of shock and variation of mass loss rate of the AGB star. That is, the AGB star is pulsating so its mass loss rate varies with time during AGB phase, and during the PN phase the fast wind speed is much larger than the slow wind speed to be a shock: $v \\gg V_{\\rm{agb}}$.  \\subsection{Physical origin of the Vexp-Rn relation} \\begin{figure}[htbp] \\scalebox{0.65}{ \\includegraphics{fig13.eps} \\label{fig:2e-5} } \\caption{2Vexp-Rn:$\\dot{M_{\\rm{agb}}}$=2$\\times10^{-5}M_{\\odot}$ yr$^{-1}$. Horizontal axis is expansion velocity [km s$^{-1}$], and vertical axis is radius [pc]. Plus sign means the data. $\\dot{M_{\\rm{agb}}}$=2$\\times10^{-5}M_{\\odot}$ yr$^{-1}$, $\\alpha$=0.5, $C_{\\rm{s}}$=10 km s$^{-1}$ and other parameter's value are written in the figure.  } \\begin{tabular}{p{7.0cm}} \\\\ \\hline \\end{tabular} \\end{figure} \\begin{figure}[htbp] \\scalebox{0.65}{ \\includegraphics{fig14.eps} \\label{fig:3e-5} } \\caption{2Vexp-Rn:$\\dot{M_{\\rm{agb}}}$=3$\\times10^{-5}M_{\\odot}$ yr$^{-1}$. Same as Fig.13 but in this figure $\\dot{M_{\\rm{agb}}}$=3$\\times10^{-5}M_{\\odot}$ yr$^{-1}$.} \\begin{tabular}{p{7.0cm}} \\\\ \\hline \\end{tabular} \\end{figure} \\begin{figure}[htbp] \\scalebox{0.65}{ \\includegraphics{fig15.eps} \\label{fig:4e-5} } \\caption{2Vexp-Rn:$\\dot{M_{\\rm{agb}}}$=4$\\times10^{-5}M_{\\odot}$ yr$^{-1}$. Same as Fig.13 but in this figure $\\dot{M_{\\rm{agb}}}$=4$\\times10^{-5}M_{\\odot}$ yr$^{-1}$.} \\begin{tabular}{p{7.0cm}} \\\\ \\hline \\end{tabular} \\end{figure} \\begin{figure}[htbp] \\scalebox{0.65}{ \\includegraphics{fig16.eps} \\label{fig:Vagb_log} } \\caption {log Rn-log Mn:$\\dot{M_{\\rm{agb}}}$=2$\\times10^{-5}M_{\\odot}$ yr$^{-1}$. Horizontal axis denotes radius and vertical axis is mass. $\\dot{M_{\\rm{agb}}}$=2$\\times10^{-5}M_{\\odot}$ yr$^{-1}$, $\\alpha$=0.5, $C_{\\rm{s}}$=10 km s$^{-1}$ and other parameter's value are written in the figure. } \\begin{tabular}{p{7.0cm}} \\\\ \\hline \\end{tabular} \\end{figure} \\begin{figure}[htbp] \\scalebox{0.65}{ \\includegraphics{fig17.eps} \\label{fig:3e-5_log} } \\caption{log Rn-log Mn:$\\dot{M_{\\rm{agb}}}$=3$\\times10^{-5}M_{\\odot}$ yr$^{-1}$. Same as Fig.16 but $\\dot{M_{\\rm{agb}}}$=3$\\times10^{-5}M_{\\odot}$ yr$^{-1}$.} \\begin{tabular}{p{7.0cm}} \\\\ \\hline \\end{tabular} \\end{figure} Based on Sect.\\ref{sec:difference}, we discuss physical background of PN evolution in this subsection.  In our model to explain the $2V_{\\rm{exp}}$-$R_{\\rm{n}}$ relation, the greatest point is considering the time-depending mass loss rate of the AGB star.  We also calculate the models which $\\dot{M_{\\rm{agb}}}$ value is a constant. Some results are presented in Fig.13, Fig.14, and Fig.15. The sequences of $\\dot{M_{\\rm{agb}}}$=2$\\times 10^{-5}M_{\\odot}$ yr$^{-1}$ are in Fig.13, those of $\\dot{M_{\\rm{agb}}}$=3$\\times 10^{-5}M_{\\odot}$ yr$^{-1}$ are in Fig.14, and those of $\\dot{M_{\\rm{agb}}}$=4$\\times 10^{-5}M_{\\odot}$ yr$^{-1}$ are in Fig.15.  In these calculations, we set $\\alpha$=0.5, $C_{\\rm{s}}$=10 km s$^{-1}$ and examine the range of $v$=500-10000 km s$^{-1}$ at the steps of 500 km s$^{-1}$. The three cases of $V_{\\rm{agb}}$=10,15,20 km s$^{-1}$ are examined.  In those figures, we plot only the models which satisfy the $R_{\\rm{n}}$-$M_{\\rm{n}}$ relation if the lines turn around the point of $\\rm{log}(R_{\\rm{n}})\\sim$ -1.0 and $\\rm{log}(M_{\\rm{n}})\\sim$ -1.0 in $\\rm{log}(R_{\\rm{n}})$-$\\rm{log}(M_{\\rm{n}})$ plane.  The $R_{\\rm{n}}$-$M_{\\rm{n}}$ relations of the selected models are shown in Fig.16, Fig.17, and Fig.18.  {}From Figs.13, 14, and 15, when $\\dot{M_{\\rm{agb}}}$=const, we can re-product only the region of $2V_{\\rm{exp}} \\leq$ 70 km s$^{-1}$. The models which are able to trace the realm of larger expansion velocity do not match the $R_{\\rm{n}}$-$M_{\\rm{n}}$ data.  Comparing time-depending mass loss models (Fig.6) with steady mass loss models (Figs.13, 14, 15), one can know how important for the evolution of PNe to take account of the variation of AGB mass-loss rate.  In other words, the radial density distribution of ambient matter around PN, $\\rho(r)$, is not in proportion to $r^{-2}$. It needs to fall suddenly at $r\\sim$0.1pc due to the pulsation history during AGB phase. Furthermore, this indicates the difference between the dynamical age and evolutionary age of PNe since the expansion law is affected by the density profile around the proto-PNe. We expect the age-discrepancy problem of PNe (Mellema 1994) may be partially resolved if the origin of the velocity-radius relation is resolved.  By the way, the variety of $v$ contributes greatly to the explanation of data distribution in $2V_{\\rm{exp}}$-$R_{\\rm{n}}$ plane.  In some models, $v$ has a very large value up to 10000 km s$^{-1}$ at $t=0$. However, $v$ may not attain such a large speed in the theory of line-driven wind. We deals with this problem in \\S\\S 4.4 again.  \\subsection{Observational implications} In $2V_{\\rm{exp}}$-$R_{\\rm{n}}$ plane, most of observational data exist in the range of $2V_{\\rm{exp}}$=40-80 km s$^{-1}$ and $R_{\\rm{n}}$=0.05-0.15pc. We present by our calculations that these data are in the realm of $t$=2000-4000 yr(cf. Fig.6 in Sect.2) and it coincides with the fact that PN is an object between the AGB star and white dwarf. Our results suggest the reason why only few data is in the smaller radius region. PNe evolve very quickly across this region, and then it is difficult for us to observe. In the larger radius region, it is also difficult to observe since PNe expand too large and diffuse. By a more highly accurate survey of PNe, in future, we will find that there are a lot of faint PNe and some of them have larger radius among them.  \\subsection{Velocity problem}  As stated in subsection 4.2, we have a time gap problem concerning the fast wind velocity from the central star of PN.  Along the line-driven wind theory, it takes about 8000 yr until $v$ reaches 10000 km s$^{-1}$, while in our calculation it needs to be attained at $t$=0.  However we will solve this problem in two different simple ways: (1) accretion jet and (2) compactness of the central star.  We shall examine the case of (1).  In the line-driven wind it is assumed the PN field is spherical symmetry.  The assumption is not so realistic and many PNe have an asymmetric jet due to its magnetic field.  In this way, the gas of PN is accelerated so quickly that our supposition for $v$ is good.  By this way of thinking, we may say that the reason why few PNe are in the larger expansion velocity and smaller radius region in their early phase ($<$2000 yr) is that the total number of asymmetrically accelerated PNe is small in their early evolution.  The other explanation of the case of (2) is rather simple, although the acceleration mechanism can not be specified. If the central star has very small radius ($\\sim$1000 km), its gravity at the surface of the star is so large. To escape from the central star gravity field, a matter must have a very large velocity.  In fact, a radius of typical white dwarf is in the range of 1000-10000 km.  It is expected that the acceleration mechanism of PNe will be proved by future observations."
    },
    "0609/astro-ph0609695_arXiv.txt": {
        "abstract": "We report the discovery of the two-image gravitationally lensed quasar SDSS J133222.62+034739.9 (SDSS J1332+0347) with an image separation of $\\Delta \\theta=1\\farcs14$. This system consists of a source quasar at $z_s=1.445$ and a lens galaxy at $z_l=0.191$. The agreement of the luminosity, ellipticity and position angle of the lens galaxy with those expected from lens model confirms the lensing hypothesis. ",
        "introduction": "\\label{sec:intro}  Strong gravitational lensing is a powerful probe of the large scale properties of the universe \\citep{kochanek06}. The well-known physics makes it straightforward to use lensing in cosmological estimates of the dark energy \\citep{turner90,fukugita90} and the Hubble constant \\citep{refsdal64}. The gravitational lensing effect depends solely on the mass density of the field, it is also an ideal tool for studies of dark matter in galaxies and clusters of galaxies \\citep[e.g.,][]{oguri04a}. The Cosmic Lens All Sky Survey \\citep[CLASS;][]{myers03,browne03}, which was conducted in the radio band, is the largest current statistical sample of lensed compact sources. However, the CLASS has only 13 lensed radio sources in its well-defined statistical sample. Larger lens samples are critical to using lensed quasars for cosmological and astrophysical tests.  Optical surveys for lensed quasars are both complementary to radio surveys like the CLASS and likely to produce larger lens samples since only one-tenth of quasars are radio-loud \\citep{ivezic02}. Our Sloan Digital Sky Survey Quasar Lens Search \\citep[SQLS;][]{oguri06} is based on the Sloan Digital Sky Survey \\citep[SDSS;][]{york00}, and its catalog of roughly 100,000 spectroscopically-confirmed quasars \\citep[e.g.,][]{schneider05}. So far, the SQLS has succeeded in discovering 14 new lensed quasars \\citep*{inada03a,inada03b,inada03c,inada05,inada06,johnston03,morgan03,pindor04, pindor06,oguri04b,oguri05,burles06}, as well as recovering many previously known lensed quasars \\citep{walsh79,weymann80,surdej87,bade97,oscoz97,schechter98,morgan01}. Indeed, the number of SDSS-discovered lensed quasars is already a significant fraction of all known quasar lenses ($\\sim 80$\\ \\footnote{see CASTLES website, http://cfa-www.harvard.edu/castles/}). The large number of lenses contained in the SQLS and its well-defined selection function means that the lens catalog provided by the SQLS will be very useful for statistical studies of strong lensing \\citep{oguri06}.  In this paper, we report the discovery of a gravitationally lensed quasar, SDSS J133222.62+034739.9 (hereafter SDSS J1332+0347). The quasar was confirmed to be doubly imaged by a low-redshift early-type galaxy, from observations using the Subaru 8.2-meter telescope and the University of Hawaii 2.2-meter (UH88) telescope. This paper is organized as follows. The SDSS data for SDSS J1332+0347, including its selection, is described in \\S \\ref{sec:sdss}. The follow-up images and spectra are shown in \\S \\ref{sec:image} and \\S \\ref{sec:spec}, respectively. In \\S \\ref{sec:model} we model the lens. Finally we summarize our results in \\S \\ref{sec:conc}. Throughout the paper we assume a cosmological model with the matter density $\\Omega_M=0.27$, cosmological constant $\\Omega_\\Lambda=0.73$, and Hubble constant $h=H_0/100\\ {\\rm km\\,sec^{-1}Mpc^{-1}}=0.7$. Magnitudes are measured in AB system. ",
        "conclusions": "\\label{sec:conc} We have presented the extensive follow-up observations of SDSS J1332+0347 using the Subaru telescope and UH88 telescope. We found a bright galaxy (component G) between two stellar components (components A and B) with similar color in both the Subaru/Suprime-Cam $i$-band and UH88/QUIRC $H$-band images. Although the \\ion{C}{3]} emission lines of the two stellar components are distorted by the \\ion{Ca}{2} H and K absorption lines from the bright galaxy at $z_l=0.191$, we can still confirm that the two quasars are at the same redshift, $z_s=1.445$. Subtraction of a scaled spectrum of component B from the component A+G leaves a residual that is consistent with the spectrum of an elliptical galaxy redshifted to $z=0.191$. The observed luminosity, ellipticity and position angle of the bright galaxy are in good agreement with those expected from a standard lens model. All these results lead us to the conclusion that SDSS J1332+0347 is indeed a lensed quasar system, where a source quasar at $z_s=1.445$ is lensed by a bright lens galaxy at $z_l=0.191$ resulting in an image separation of ${\\Delta}{\\theta}=1\\farcs14$. One peculiar characteristic of the lens is that the brighter component A is closer to the lens galaxy than the fainter component B. We have no difficulty reproducing the flux ratio, and note that such {\\it inverted} flux ratios were also observed and easily modeled in the two-image lens HE1104--1805 \\citep{wisotzki93}. This discovery adds another object to the complete sample of the SDSS lensed quasars that can be used to estimate the cosmological model. The brightness of the lens galaxy also makes it a good candidate for dynamical observations; higher-resolution imaging would allow us to determine more accurately the brightness of the components in this lens system, and spectroscopic observations with better signal-to-noise ratios would also allow us to study the interstellar medium in the lens galaxy in detail."
    },
    "0609/astro-ph0609140_arXiv.txt": {
        "abstract": "We develop a semi-analytic model for planet formation during the pre-main sequence contraction phase of a low mass star. During this evolution, the stellar magnetosphere maintains a fixed ratio between the inner disk radius and the stellar radius. As the star contracts at constant effective temperature, the `snow line', which separates regions of rocky planet formation from regions of icy planet formation, moves inward.  This process enables rapid formation of icy protoplanets that collide and merge into super-Earths before the star reaches the main sequence. The masses and orbits of these super-Earths are consistent with super-Earths detected in recent microlensing experiments. ",
        "introduction": "The recent discoveries of super-Earths around low mass stars challenge our understanding of planet formation. From separate microlensing events, \\citet{2006Natur.439..437B} and \\citet{2006ApJ...644L..37G} provide strong evidence that planets with masses $M \\sim$ 5--15\\,$M_{\\oplus}$ are common around M dwarf stars\\footnote{These microlensing observations directly yield the ratio $M_P/M_{\\star}$, where $M_P$ is the mass of the planet and $M_{\\star}$ is the mass of the star. Adopting the most likely parent star -- a K-dwarf or an M-dwarf -- yields the most likely $M_P$.}. With orbital semi-major axes $a \\sim 2.5$--3\\,AU, these planets are probably ice giants roughly similar in structure to Uranus and Neptune in the Solar System.  \\citet{2006ApJ...644L..79B} proposes that these planets form in two stages. After a disk instability produces a gas giant, photoevaporation of the gas giant atmosphere leaves behind an icy core with $M \\sim$ 10--20 $M_{\\oplus}$. This mechanism requires a massive disk to initiate the instability and a nearby O-type star to photoevaporate the gas giant atmosphere. \\citeauthor{2006ApJ...644L..79B} notes that this process should yield (i) super-Earths around M dwarfs formed in rich star clusters and (ii) gas giants around M dwarfs formed in low mass stellar associations.  \\citeauthor{2006Natur.439..437B} suggest that super-Earths favor coagulation models, where collisions of 1--10 km objects eventually produce icy planets with $M \\sim$ 10 $M_{\\oplus}$ at 1--10 AU. Although numerical calculations appear to preclude gas giants at 1--10 AU around M dwarfs \\citep{2004ApJ...612L..73L}, there has been no demonstration that coagulation produces icy planets on reasonable timescales in a disk around an M dwarf.  Here, we develop a semi-analytic coagulation model, and show that contraction of the central star along a pre-main sequence (PMS) Hayashi track sets the initial conditions for planet formation around low mass stars.  Our results indicate that icy protoplanets with $M \\sim$ 0.1--1\\,$M_{\\oplus}$ form in $\\sim$ 0.1--1\\,Myr at 1--4\\,AU. Over 50--500\\,Myr, collisions between protoplanets produce super-Earths with masses similar to those detected in microlensing surveys.  We start with the motivation for our study in \\S \\ref{sec:motivation}, discuss the coagulation model of planet formation and the moving snow line in \\S \\ref{sec:snowline}, develop the disk evolution model in \\S \\ref{sec:evolution}, and apply the model to super-Earth formation in \\S \\ref{sec:superEarths}. We end with a brief summary in \\S \\ref{sec:summary}. ",
        "conclusions": "We have developed an analytic prescription for planet formation by coagulation around low mass stars. The model has two distinctive features that enable formation of super-Earths during the PMS contraction phase.  \\begin{itemize}  \\item We set the inner disk radius at a fixed distance relative to the radius of the central star, $\\xi \\equiv R_{in}/R_{\\star}$. Thus, $R_{in}$ changes as the star contracts to the main-sequence, leading to significant evolution in $\\sigma(a)$.  \\item During PMS contraction, $a_{snow}$ moves inwards by a factor of $\\sim$15--20, producing large enhancements in $\\sigma(a)$ as ices condense out of the nebula. Ice condensation is the key mechanism that allows coagulation to produce super-Earths around low mass stars. This process results in new outcomes for planet formation, including planets with dense, rocky cores and thick, icy surfaces.  \\end{itemize}  We applied this model to super-Earth formation around a $0.25\\,M_{\\odot}$ star. At 1--5\\,AU, isolated oligarchs can grow to masses $\\sim 0.1$--1\\,$M_{\\oplus}$ in $\\sim0.1$--1\\,Myr, short compared to the $\\sim 100$\\,Myr PMS contraction time. These oligarchs merge into super-Earths with masses $\\sim 2$--5\\,$M_{\\oplus}$ as the star contracts to the main sequence. Thus, coagulation can produce planetary systems similar to those detected in recent microlensing events.  Aside from our success in producing icy super-Earths at 1--3\\,AU around low mass stars, the model makes clear predictions for rocky planet formation close to a low mass star.  At $a \\lesssim 1$\\,AU, the isolation masses estimated for rocky planets are $M_{iso} \\lesssim 0.01\\,M_{\\oplus}$. If $\\sim$10 isolated objects merge into a rocky planet, we predict many low mass planets with masses $\\sim$0.1\\,$M_{\\oplus}$ at distances of 0.05--0.5\\,AU around low mass M dwarfs. We plan to explore the consequences of our model in future papers."
    },
    "0609/astro-ph0609189_arXiv.txt": {
        "abstract": "The CDMS experiment aims to directly detect massive, cold dark matter particles originating from the Milky Way halo. Charge  and lattice excitations are detected after a particle scatters in a Ge or Si crystal kept at $\\sim$30\\,mK, allowing to separate nuclear recoils from the dominating electromagnetic background. The operation of 12 detectors in the Soudan mine for 75 live days in 2004 delivered no evidence for a signal, yielding stringent limits on dark matter candidates from supersymmetry and universal extra dimensions. Thirty Ge and Si detectors are presently installed in the Soudan cryostat, and operating at base temperature. The run scheduled to start in 2006 is expected to yield a one order of magnitude increase in dark matter sensitivity. ",
        "introduction": "The Cryogenic Dark Matter Search (CDMS) experiment is designed to search for cold dark matter in form of weakly interacting massive particles (WIMPs). These hypothetical thermal relics of the big bang are expected to have masses and annihilation cross sections at the weak scale if their relic density is around around $\\Omega$=0.2. Examples are the lightest supersymmetric particle (LSP, or neutralino) and the lightest Kaluza-Klein particle (LKP) in theories with large universal extra dimensions.  In CDMS, WIMPs can be detected via their elastic scattering with Ge or Si nuclei. Events with nuclear recoil energies of a few tens of keV and rates below 1~event kg$^{-1}$d$^{-1}$ are expected. The experiment is operated at the Soudan Underground Laboratory in Minnesota (at 2080 m.w.e.), in a dedicated low-background facility. At its core are thirty z-dependent ionization- and phonon-mediated (ZIP) detectors, arranged in 5 towers. Within one tower the ZIPs are stacked 2\\,mm apart, with no intervening material. Such close packing shields the crystals against surface low-energy electrons, and allows to identify multiply scattered events. A particle interaction in a ZIP detector deposits energy into the crystal via lattice vibrations, or phonons, and via charge excitations, or electron-hole pairs. The simultaneous measurement of both phonon and ionization signals allows to accurately measure the recoil energy, and to distinguish between nuclear- and electron recoil events. This distinction becomes possible since nuclear recoils produce fewer charge pairs than electron recoils of the same energy. The ionization yield, defined as the ratio of ionization to full recoil energy, is unity for electron recoils with complete charge collection, and $\\sim$0.3 ($\\sim$0.25) for nuclear recoils in Ge (Si) with recoil energies above 20\\,keV.   \\begin{figure}[h] \\hspace*{-0.5cm}{\\includegraphics[width=.5\\textwidth]{cdms_timing.eps}} \\caption{Ionization yield versus phonon timing parameter for $^{133}$Ba gamma calibration events (dots and crosses) and $^{252}$Cf neutron calibration events (circles). Low-yield $^{133}$Ba events (crosses) have small values of the timing parameter, a typical cut is indicated by the dashed vertical line.} \\label{fig1a} \\end{figure}   ZIP detectors measure the athermal phonos created in a particle interaction using quasiparticle-assisted electrothermal-feedback transition-edge sensors (QETs) photolitographically patterned onto one of the crystal surfaces. The QETs, which are made of 1\\,$\\mu$m wide strips of tungsten connected to eight superconducting aluminum fins are divided into four independent channels, each consisting of 1036 QETs operated in parallel. The narrow tunsten strips form the transition-edge sensors (TESs), which are voltage biased and kept stably within their superconducting to normal transition by electrothermal feedback based on Joule self-heating. Energy deposited in the tungsten electron system raises the temperature of the film, increasing its resistance. The corresponding change in current is detected by a high-bandwidth SQUID array. The phonon pulse rise times are sensitive to the phonon arrival times at each of the four quadrants and allow to localize an event in the x-y plane. In addition, events occurring near the detector's surface display faster rise times than bulk events, allowing for a fiducial volume cut. For the ionization measurement, a drift field of a few V/cm is applied across the crystal using electrodes deposited on the two faces of each detector. The electrodes are segmented radially, in an outer annular guard ring and a central disk-shaped volume, a design optimized to veto events near the edge of the crystals. A detailed description of the CDMS apparatus and shield is given in \\cite{cdms_prd}. ",
        "conclusions": ""
    },
    "0609/hep-th0609094_arXiv.txt": {
        "abstract": "We show that if the visible universe is a membrane embedded in a higher-dimensional space, particles in uniform motion radiate gravitational waves because of spacetime lumpiness. This phenomenon is analogous to the electromagnetic diffraction radiation of a charge moving near to a metallic grating. In the gravitational case, the role of the metallic grating is played by the inhomogeneities of the extra-dimensional space, such as a hidden brane. We derive a general formula for gravitational diffraction radiation and apply it to a higher-dimensional scenario with flat compact extra dimensions. Gravitational diffraction radiation may carry away a significant portion of the particle's initial energy. This allows to set stringent limits on the scale of brane perturbations. Physical effects of gravitational diffraction radiation are briefly discussed. ",
        "introduction": "\\label{intro} Larmor's formula of electromagnetism \\cite{Jackson} states that an electric charge in uniform motion does not radiate. However, there are two ways to have radiation from a charge moving with constant velocity. The first way is to have a particle moving in a medium with velocity exceeding the phase velocity of light in that medium. This gives rise to the well-known Vavilov-Cherenkov radiation \\cite{Jackson,GinzburgFrank,Tamm}. The second way is to consider charge motion in inhomogeneous media. Ginzburg and Frank \\cite{GinzburgFrank} first discussed this effect by investigating a particle in uniform motion which crosses a planar interface between two media with dissimilar refractive index. This kind of radiation is known as transition radiation \\cite{Jackson}. More generally, any motion near finite-size objects also induces radiation, in a process called diffraction radiation \\cite{Bolotovskii}. One of the first experimental verifications on this effect was provided by Smith and Purcell \\cite{SmithPurcell}. The Smith-Purcell experimental set-up consisted of an electron moving close to the surface of a metal diffraction grating at right angles to the rulings. The theory of the Smith-Purcell effect has been discussed by several authors \\cite{FranciaBerg}.  The aim of this paper is to show the existence of gravitational diffraction radiation (GDR) and discuss its physical effects. The model under consideration is a braneworld scenario \\cite{Maartens:2003tw} such as, for example, the Randall-Sundrum I model \\cite{Randall:1999ee}. Braneworld models have attracted a lot of interest in recent years,  revolutionizing our view of how our universe may be described. The central idea of braneworld scenarios is that the visible universe is restricted to a four-dimensional brane inside a higher-dimensional space, called the bulk. The additional dimensions are taken to be compact and other branes may be moving through the bulk. Interactions of the visible brane with the bulk and hidden branes introduce effects not seen in standard physics.  In this set-up, a particle in uniform motion on the visible brane radiates gravitational waves due to the presence of a second (hidden) brane at finite distance, which plays the role of the metal diffraction grating of the Smith-Purcell experiment. GDR on the visible brane is generated by inhomogeneities on the hidden brane due, for example, to bulk-brane interactions and brane fluctuations \\cite{Bando:1999di}. We will show that the amount of GDR depends on the size of extra dimensions and the length scale of brane perturbations. This result is general and independent on the fine details of the model. Without loss of generality, in our computations we will consider a flat five-dimensional spacetime with the extra dimension taking values within the interval $[0,L]$. The distance between the particle and the diffraction grating, $h=L/2$, is the distance between the two branes located at the orbifold fixed points. The model is illustrated pictorially in Fig.\\ \\ref{braneworld1}. The generalization to higher-dimensions is trivial and is discussed at the end of this letter.  \\begin{figure} \\centerline{\\includegraphics[width=0.45\\textwidth]{figure1.eps}\\qquad\\qquad \\includegraphics[width=0.30\\textwidth]{figure2.eps}} \\caption{Pictorial representation of the braneworld scenario considered in the text. The hidden brane at $z=0$ is identified with the brane at $z=L$. The visible brane is located at $z=h=L/2$. The particle moves on the visible brane along $x$ with constant velocity $v$. Only one transverse coordinate, $y$, is shown. The right panel shows the wrapping around the extra dimension $z$. If the two halves of the cylinder are identified, the Randall-Sundrum model I \\cite{Randall:1999ee} is obtained.} \\label{braneworld1} \\end{figure} ",
        "conclusions": "\\label{discussion} In higher-dimensional models of our universe, particles radiate gravitational waves due to the structure of the extra dimensions, specifically the presence of an inhomogeneous hidden brane. We have called this phenomenon ``Gravitational Diffraction Radiation\" because of its similarity with the electromagnetic diffraction radiation which is emitted by a charge moving near a metallic grating. GDR is a general phenomenon of all higher-dimensional braneworld scenarios. It applies to particles propagating both on the brane and in the bulk. As an example of its physical effects, we have discussed the GDR emitted by a particle on the visible brane and derived lower bounds on the scale of brane perturbations.  The above results can be extended in several directions. For instance, this paper considers only inhomogeneities on the hidden brane. The visible brane is also inhomogeneous. Therefore, free motion on the visible brane is not uniform from the five-dimensional perspective. This effect could be comparable to the GDR effect studied here. The derivation of GDR presented above is purely classical. Although the scale of brane inhomogeneities suggests that the system can be treated as classical, the importance of quantum effects should be assessed. It would also be interesting to consider GDR emitted by photons or bulk particles. In the latter case, GDR energy loss would manifest itself as a background of stochastic gravitational waves on the visible brane.  GDR is originated by the lumpiness of the spacetime at small scales. Thus GDR effects may show up in a variety of physical situations besides braneworld models. Emission of radiation by free particles in a cosmic string background was discussed in Ref.\\ \\cite{aliev}. GDR could also play an important role in the very early universe or near the horizon of black holes. Direct signatures are unlikely to be observed in these cases. However, GDR emission could lead to interesting predictions at theoretical level and possibly indirect observable effects."
    },
    "0609/astro-ph0609626_arXiv.txt": {
        "abstract": "{ Using available data for C and M giants with $\\Mbol<-3.6$ in Magellanic Cloud clusters, we derive limits to the lifetimes of the corresponding evolutionary phases, as a function of stellar mass. The C-star phase is found to have a duration between $2$ and $3$ Myr for stars in the mass range from $\\sim1.5$ to $2.8\\,M_\\odot$.  There is also an indication that the peak of C-star lifetime shifts to lower masses (from slightly above to slightly below 2~\\Msun) as we move from LMC to SMC metallicities. The M-giant lifetimes also peak at $\\sim2$~\\Msun\\ in the LMC, with a maximum value of about 4~Myr, whereas in the SMC their lifetimes appear much shorter but, actually, they are poorly constrained by the data. These numbers constitute useful constraints to theoretical models of the TP-AGB phase. We show that several models in the literature underestimate the duration of the C-star phase at LMC metallicities.}  \\authorrunning{L. Girardi \\& P. Marigo} \\titlerunning{TP-AGB lifetimes} ",
        "introduction": "\\label{intro}  Since the work by Frogel et al. (1990, hereafter FMB90), the thermally pulsing asymptotic giant branch (TP-AGB) phase has been recognized to be an important contributor to the integrated near-infrared luminosity of young and intermediate-age stellar populations. Starting from the observed numbers and luminosities of C and M-type giants in Magellanic Cloud star clusters, FMB90 concluded that up to 40 percent of the bolometric cluster luminosity comes from stars with $\\Mbol<-3.6$, which are above the RGB tip and hence belong to the TP-AGB. An interesting plot by FMB90 showed this fraction as a function of SWB cluster type (Searle et al. 1980), which gives approximate estimates of cluster ages. Charlot \\& Bruzual (1991) for the first time used this information to calibrate the amount of TP-AGB stars added into models of the spectrophotometric evolution of galaxies. This was an important piece of information, previously missing, that has contributed to the success of such models in interpreting the observed spectra of galaxies.  Since then several different approaches have been adopted to include the TP-AGB phase into evolutionary population synthesis models. They often use slightly modified versions of the original FMB90 plots to constrain the amount of TP-AGB stars (see Mouhcine \\& Lan\\c con 2002; Maraston 2005, and references therein). Maraston (1998, 2005), for instance, follows an empirical approach based on FMB90 data to add the appropriate ``fuel consumption'' (equivalent to the contribution of these stars to the integrated bolometric light) at LMC metallicities, and then she uses results from uncalibrated TP-AGB evolutionary models from Renzini \\& Voli (1981) to account for the basic dependencies with metallicity. Other authors directly include extended sets of TP-AGB evolutionary models, for a wide enough range of stellar masses and metallicities, into their models. This kind of approach goes from adopting a very simplified description of the TP-AGB (e.g. Bressan et al. 1994) to adding TP-AGB tracks {in which the efficiency of the third dredge-up is} calibrated to reproduce the C-star luminosity functions (CSLF) in the Magellanic Clouds (e.g. Marigo \\& Girardi 2001). {Different schemes are adopted by Bruzual \\& Charlot (2003) and Mouhcine \\& Lan\\c con (2002), who adopt TP-AGB tracks that distinguish between the C- and M-type phases, but which are not directly calibrated on the CSLF: Bruzual \\& Charlot (2003) make a composition of TP-AGB tracks from different sources (namely Vassiliadis \\& Wood 1993, Groenewegen \\& de Jong 1993, and Groenewegen et al. 1995), but do not check the consistency of these composite tracks with the observed properties of resolved AGB populations. Mouhcine \\& Lan\\c con (2002) instead compute their own TP-AGB tracks adopting modified prescriptions but the same values of dredge-up parameters as Groenewegen \\& de Jong (1993); they make several different comparisons with data, avoiding however the explicit check with the CSLFs that was key to Groenewegen \\& de Jong (1993).}  Owing to the many interrelated aspects of TP-AGB evolution and their complex dependence on metallicity and stellar mass, our understanding is that the best approach to be used in population synthesis should be the inclusion of TP-AGB tracks computed in a self-consistent way and {\\em directly} calibrated using a set of observables from the LMC and SMC. Once calibrated, these self-consistent sets of TP-AGB tracks present the potential advantage of better describing the dependence of all TP-AGB properties on the stellar metallicity, and hence they should work better for metallicities significantly different from the LMC and SMC ones. In this context, the main goal of the present paper is to translate FMB90 data into useful quantities -- the TP-AGB lifetimes as a function of stellar mass -- for the direct calibration of TP-AGB tracks of LMC ad SMC metallicities.  Originally, this work has been motivated by the finding by Marigo (2002) that all TP-AGB models computed that far have used a very improper prescription for their low-temperature opacities. She replaced the opacities for solar-scaled chemical mixtures -- still now the standard choice in the literature -- by opacities properly evaluated for the chemical mixtures of evolving TP-AGB envelopes.  One main evolutionary effect of variable molecular opacities is the remarkable reduction of effective temperatures for C stars, which then causes an earlier onset of the superwind regime and a reduction of their lifetimes, when compared to models computed at fixed solar-scaled opacities. This reduction should not affect the previous M-type phase, for which no dramatic change in molecular opacities is expected. These findings also cause serious doubts about the previously-derived behaviours of TP-AGB evolutionary properties -- including lifetimes, termination luminosities, effective temperatures, etc -- with metallicity. It is then evident that new grids of TP-AGB tracks are due, and that the C-star lifetimes, once derived from empirical data, could provide important constraints to them. ",
        "conclusions": "\\label{sec_conclu}  Despite the large error bars, Table~\\ref{tab_binned} and Fig.~\\ref{fig_lifetimes} clearly indicate that C-star lifetimes have values of about $2$ to $3$ Myr, for stars in the mass interval from $\\sim1.5$ to $\\sim2.8\\,M_\\odot$ (Fig.~\\ref{fig_lifetimes}), and for metallicities comprised between the $-0.3$ and $-0.7$~dex implied by LMC and SMC data. There is also an indication that the peak of C-star lifetime shifts to lower masses (from slightly above to slightly below 2~\\Msun) as we move from LMC to SMC metallicities. The M-giant lifetimes also peak at about 2~\\Msun\\ in the LMC, with a maximum value of about 4 Myr. In the SMC the M-giant lifetimes appear much shorter, but actually they are poorly constrained by present data.  These lifetimes correspond to the optically-visible TP-AGB phase. In their study of IR sources, van Loon et al. (2005) find that 1.3--3~\\Msun\\ stars spend of the order of 10--20 percent of their AGB lifetimes as optically-obscured, bright IR objects, with mass losses higher than about $10^{-6}$~\\Msun\\,yr${-1}$. Assuming the typical lifetimes above the RGB tip to be 1~Myr, they derive 1--2$\\times10^5$~yr for the duration of this superwind phase. Moreover, since about $3\\times10^5$~yr is the time required for loosing the envelope masses of the sample stars with the observed (superwind) mass loss rates, van Loon et al. estimate that 30--70 percent of the mass loss of AGB stars occurs as superwind.  If we repeat the same reasoning using our own estimates for the lifetimes in the TP-AGB phase, which are at least twice larger than the value used by van Loon et al. (2005), we find that the superwind phase may account for {\\em all} the mass loss during the TP-AGB phase. Of course, this indication is anyway very uncertain because derived from small numbers of stars.  We have also verified that several models in the literature present C-star lifetimes significantly shorter than the values we find for the LMC, as illustrated in Fig.~\\ref{fig_models}. The comparison is shown only for a subset of TP-AGB models which has already been used, with various procedures, into population synthesis of galaxies.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{lifetime_bw.ps}} \\caption{ The lifetimes of the C-star phase in the LMC as inferred from the cluster data, as a function of turn-off mass (dots with error bars), compared to the predictions of some TP-AGB models in the literature. The comparison is made with models computed for the LMC present metallicity ($Z\\simeq0.008$), whenever available.} \\label{fig_models} \\end{figure}  These models can be roughly classified into two classes: the ones previous to Groenewegen \\& de Jong (1993) are {\\em uncalibrated}, in the sense that they fail in reproducing the C-star luminosities observed in the Magellanic Clouds (as discussed for the first time by Iben 1981). Renzini \\& Voli (1981) TP-AGB models belong to this class; as it can be seen in Fig.~\\ref{fig_models}, they present too low C-star lifetimes as compared to the data, and are systematically shifted towards a much higher range of masses. Notice that we plot just a subset of Renzini \\& Voli's $\\alpha=0$ models (their tables 1a and 1f); their models with $\\alpha>0$, which take into account the effect of hot-bottom burning in more massive stellar envelopes, would present even lower C-star lifetimes for $\\mto\\ga4$~\\Msun. It is clear that these models are highly discrepant with the data.  Groenewegen \\& de Jong (1993) presented the first {\\em calibrated} TP-AGB models, in which the poorly known parameters determining the occurrence and efficiency of third dredge-up events were tuned so as to reproduce the LMC and SMC CSLFs. Indeed, all post-1993 models we plot in Fig.~\\ref{fig_models} have positive C-star lifetimes in the right mass range, with the minimum mass for the present of C-stars being close to 1.5~\\Msun, and a peak lifetime located somewhere between 2 and 3~\\Msun, in rough agreement with the cluster data. This agreement is just expected, since fitting the main features of the CSLF requires imposing efficient dredge-up for masses as low as 1.4~\\Msun, as discussed in detail by Groenewegen \\& de Jong (1993) and Marigo et al. (1999).  However significant discrepancies appear in the C-star lifetimes of calibrated models too. For instance, lifetimes in Groenewegen \\& de Jong (1994) tend to be too short, with maximum values of just 0.6~Myr for models assuming Reimers' (1975) mass loss formula (RE in the plot), and 1.8 Myr for Vassiliadis \\& Wood's (1993) one (VW in the plot). The discrepancies are more evident in the mass range between 1.5 and 2.5~\\Msun, where model lifetimes fall to less than 0.5~Myr, whereas the {ones derived here} are closer to 2~Myr. Notice that the comparison of Fig.~\\ref{fig_models} is not strictly correct because Groenewegen \\& de Jong (1994) tracks do not include overshooting, whereas the LMC cluster data have been age-dated using overshooting models. We estimate that, in order to be put Groenewegen \\& de Jong models in the same scale as the data, their masses should be reduced by a factor of about 20 per cent. This would not solve the differences in lifetimes.  The same discrepancy is shared, to a lower extent, by Mouhcine \\& Lan\\c con's (2002) $Z=0.008$ models, which present the right peak lifetimes of $\\sim2.5$~Myr, but again have too low lifetimes at masses $\\la2$~\\Msun. {Due to the significantly lower metallicities of the old LMC clusters, and to the high metallicity dependence expected for the C-type lifetime, this discrepancy for $M_{\\rm TO}\\la2$~\\Msun\\ would likely be reduced if the comparison was performed with models of smaller metallicity (say $Z=0.004$).}  Notice that Mouhcine \\& Lan\\c con (2002) adopt the same dredge-up parameters as Groenewegen \\& de Jong (1993), although {many} of their model ingredients have changed. Therefore, the Mouhcine \\& Lan\\c con (2002) models are not strictly calibrated on the CSLF, although their C-star luminosities should not be much far from the observed ones.  Marigo (2001) models do not present these discrepancies in their C-star lifetimes at lower masses, and have a peak lifetime of 3.4~Myr, which is well compatible with the data. This could be considered as a set of TP-AGB models which comply with both CSLF and lifetime constraints. However, we know already that these models -- as well as {\\em all} of the previously mentioned ones -- are wrong for a different reason: they do not consider the crucial effect of variable molecular opacities as the chemical composition changes along the TP-AGB evolution (see Marigo 2002). The effect of the variable opacities is illustrated by the models labelled ``Marigo (2002)'' in Fig.~\\ref{fig_lifetimes}; it can be noticed that assuming the same dredge-up parameters as calibrated by Marigo (2001), the change from fixed (solar-scaled) to variable opacities ($\\kappa$-fix and $\\kappa$-var cases, respectively) causes a reduction of $\\sim35$~percent in the C-star lifetimes. These would be still compatible with the cluster data. However, the new models have also lower mean luminosities and hence they fail to reproduce the CSLF. In fact, updating of the input physics of TP-AGB models is not enough, and a re-calibration of dredge-up parameters become necessary in this case.  In conclusion, we have shown that present data for AGB stars in LMC clusters represent useful -- and so far neglected -- constraints to the lifetimes of TP-AGB models. Checking for these constraints should be especially important if the TP-AGB models are to be used into evolutionary population synthesis. In fact, using models with the right luminosities (i.e. calibrated with the CSLF) but with too low lifetimes would lead to underestimate the contribution of TP-AGB stars to the integrated light of single-burst stellar populations.  In a following paper, we will present updated TP-AGB models computed for the variable-opacity case, in which the lifetimes here derived are adopted, together with the observed CSLF in the Magellanic Clouds, in the calibration procedure of the main model parameters."
    },
    "0609/astro-ph0609410_arXiv.txt": {
        "abstract": "{ We investigate the effect of new stellar models, which take rotation into account, computed for very low metallicities (Z = 10$^{-8}$) on the chemical evolution of the earliest phases of the Milky Way. We check the impact of these new stellar yields on a model for the halo of the Milky Way that can reproduce the observed halo metallicity distribution. In this way we try to better constrain the ISM enrichment timescale, which was not done in our previous work (\\cite{Chia06}). The stellar models adopted in this work were computed under the assumption that the ratio of the initial rotation velocity to the critical velocity of stars is roughly constant with metallicity. This naturally leads to faster rotation at lower metallicity, as metal poor stars are more compact than metal rich ones. We find that the new Z = 10$^{-8}$ stellar yields computed for large rotational velocities have a tremendous impact on the interstellar medium nitrogen enrichment for log(O/H)+12 $<$ 7 (or [Fe/H]$<$ $-$3). We show that upon the inclusion of the new stellar calculations in a chemical evolution model for the galactic halo with infall and outflow, both high N/O and C/O ratios are obtained in the very-metal poor metallicity range in agreement with observations. Our results give further support to the idea that stars at very low metallicities could have initial rotational velocities of the order of 600-800\\,km\\,s$^{-1}$. An important contribution to N from AGB stars is still needed in order to explain the observations at intermediate metallicities. One possibility is that AGB stars at very low metallicities also rotate fast. This could be tested in the future, once stellar evolution models for fast rotating AGB stars will be available.}  \\FullConference{International Symposium on Nuclear Astrophysics --- Nuclei in the Cosmos --- IX\\\\ June 25-30 2006\\\\ CERN, Geneva, Switzerland}  \\begin{document} ",
        "introduction": "Recent measurements of nitrogen abundances in metal-poor stars by Spite et al. \\cite{S05} show a large N/O ratio suggesting high levels of production of primary nitrogen in massive stars. Moreover, the N/O abundance ratios in metal-poor stars exhibit a large scatter (roughly 1 dex, much larger than their quoted error bars) although none of the stars measured so far has N/O ratios as low as the ones observed in DLAs.  Last year, Chiappini et al. \\cite{Chia05} studied the implications of this new data set for our understanding of the nitrogen enrichment in the Milky Way (MW). By the time the latter paper was published there was no set of stellar yields able to explain the very metal-poor data of \\cite{S05}. Using the so-called Z=0 population III stellar yields available in the literature (which in some cases predict large N yields for non-rotating massive stars with masses around 20 M$_{\\odot}$ - e.g. \\cite{Chie04}) did not solve the problem either. The main reason is that in population III non rotating models the N production is confined to a very narrow mass range and hence chemical evolution models computed with such prescriptions predict a peak in the [N/Fe] ratios at very low metallicities followed by a strong decrease already before [Fe/H] $\\sim -$3 in disagreement with the observations of \\cite{S05} (see \\cite{Ba06}). This happens even when allowing the Z=0 contribution to be valid up to a threshold (Z$_{tr}$) metallicity above which population III stars stop forming of the order of Z$_{tr}$=10$^{-6}$. This Z$_{tr}$ metallicity is probably already too high as the physics adopted in the Z=0 models would be valid only up to metallicities Z$\\sim$10$^{-10}$, which represents the metallicity below which massive stars first enter the phase of H-burning via the pp chain, followed by the 3$\\alpha$ reaction, which then allows the CNO cycle to proceed, as in Z=0 (population III) stars.  Given the above discussion, \\cite{Chia05} suggested that the most promising way to account for the new data was to assume that stars at low metallicity rotate sufficiently fast to enable massive stars to contribute much larger amounts of nitrogen. It was shown that even when adopting the Meynet \\& Maeder \\cite{M02} stellar yields (these authors had computed stellar yields down to Z=10$^{-5}$) where the nitrogen production is increased in stars of all masses due to rotation, it was not possible to explain the observations of \\cite{S05}. The approach of \\cite{Chia05} was to assume that the stellar yields computed by \\cite{M02} were valid down to Z=10$^{-5}$, but that below this metallicity some mechanism should be able to increase even more the N yields. It was then predicted that massive stars born with metallicities below Z=10$^{-5}$ should produce a factor between 10 and a few times 10$^2$ more nitrogen (depending on the stellar mass) than the values given by \\cite{M02} for Z=10$^{-5}$ and $\\upsilon^{\\rm ini}_{\\rm rot}=300$\\,km\\,s$^{-1}$ in order to reproduce the {\\it mean} locus of the Spite et al. data in a log(N/O) vs. log(O/H)+12 diagram.  The physical motivation for the approach described above would be an increase of the rotational velocity in very metal-poor stars \\cite{M99, M06} and hence an increase in the nitrogen yields. In this framework it is possible to understand the apparent contradictory finding of \\cite{S05} namely, a large scatter in N/O and the almost complete lack of scatter in [$\\alpha$/Fe] ratios of the same very metal-poor halo stars \\cite{Ca04}. As suggested in \\cite{Chia05}, if the nitrogen production in very metal-poor massive stars depends strongly on the rotational velocity of the star, this could explain the scatter in the N/O abundance ratios: the scatter would reflect the distribution of the stellar rotational velocities as a function of metallicity whereas the [$\\alpha$/Fe] ratios would remain unchanged.  Whether the above suggestions were physically plausible remained to be assessed by stellar evolution models computed at lower metallicities, which took rotation and mass loss into account. Novel stellar evolution models have been computed for metallicities Z=10$^{-8}$ \\cite{Hi06} for massive stars. The new calculations show that if the stars at all Z start their life on the zero age main sequence (ZAMS) with on the average a fraction of $\\sim$\\,0.5 the critical velocity, the low Z stars easily reach break-up velocity during MS evolution (see also Meynet et al., this conference). Fast rotation also contributes to produce a more efficient mixing at lower Z, which leads to a large production of N in massive stars. We note that the new models with fast rotation also predict different type of supernova progenitors, final remnants and yields in heavy elements at very low Z (\\cite{M06} and references therein).  In a recent paper, Chiappini et al. \\cite{Chia06} computed chemical evolution models adopting these new stellar prescriptions and show that these new models can explain not only the high N/O in very metal poor stars but also predict an increase in C/O for decreasing metallicity.  However the model computed by \\cite{Chia06} did not include outflows during the halo phase thus leading to an increase of [Fe/H] with time faster than what would be obtained with an outflow model. Moreover, as shown by \\cite{P03}, halo models that can reproduce the metallicity distribution of halo stars should include both inflow and outflow.  Here we compute such a model and compare our new results with the ones we obtained in \\cite{Chia06}. Finally, in \\cite{Chia06} we considered only the so-called {\\it normal} very metal-poor stars, since our goal was to explain the mean ISM enrichment at very low metallicities. Here we place the carbon rich ultra metal-poor stars \\cite{B05} into this context and discuss the implications for the CNO nucleosynthesis. ",
        "conclusions": "In this work, we computed chemical evolution models adopting the very recent calculations of Hirschi \\cite{Hi06} for the evolution of massive stars at very low metallicities under the assumption of an almost constant ratio $\\upsilon^{\\rm ini}_{\\rm rot}/\\upsilon_{\\rm breakup}$ as a function of metallicity (i.e. where the $\\upsilon^{\\rm ini}_{\\rm rot}$ increases towards lower metallicities). In such a framework, massive stars can produce large amounts of nitrogen.  At present, this is the only way to explain the high nitrogen abundances measured recently in {\\it normal} halo stars. This gives further support to the idea that stars rotate faster at very low metallicities. The new stellar evolution models also produce some extra carbon at Z $\\leq$ 10$^{-5}$. As a consequence, an upturn is produced in the C/O at low metallicities. This result is obtained without the need of introducing population III stars (here understood as Z=0 stars with a flatter IMF).  The only alternative to explain the observations showing a high N abundances in very metal poor normal stars\\footnote{Here, \"normal\" stands for the assumption that these stars show a pristine N abundance which reflects the one of  the ISM gas from which these stars formed and did not come from any internal mixing process or binary mass transfer.} would be if AGB stars would contribute to the ISM enrichment before [Fe/H]$\\sim -$4. This would mean that  the timescales for the chemical evolution of the halo would be very different from the ones of \\cite{Chia06}.  Models assuming outflows could delay the ISM enrichment so that AGBs would have time to contribute. Here we show that a model with outflow, which is able to fit the observed halo metallicity distribution and in which the ISM is enriched on longer timescales than the ones computed by \\cite{Chia06}, still gives the same qualitative result. In fact, this model still requires to invoke fast rotation at very low metallicities  to explain the high N/O observed in metal poor halo stars. Future surveys of very metal poor stars (see Beers this conference) will certainly lead to a more complete halo metallicity distribution and will be extremely useful to constrain even better the halo enrichment timescale.  Here we show the results for the mean ISM enrichment history, assuming that at a given metallicity all stars of a given mass rotate at the same velocity. Further computations are in progress, and we envision to compute chemical evolution models in which a distribution of rotational velocities is assumed for each metallicity. These models would predict a scatter in N/O to be compared to the observed one."
    },
    "0609/astro-ph0609760_arXiv.txt": {
        "abstract": "We present a statistical analysis of the photometric properties and spatial distribution of more than 2,800 MgII absorbers with $0.37<z<1$ and rest equivalent width $W_0(\\lambda2796)>0.8$\\AA~detected in SDSS quasar spectra. Using an improved image stacking technique, we measure the cross-correlation between MgII gas and light (in the $g$, $r$, $i$ and $z$-bands) from 10 to 200 kpc and infer the light-weighted impact parameter distribution of MgII absorbers. Such a quantity is well described by a power-law with an index that strongly depends on absorption rest equivalent width $W_0$, ranging from $\\sim-1$ for $W_0\\lesssim 1$\\AA~ to $\\sim -2$ for $W_0\\gtrsim 1.5$\\AA.  At redshift $0.37<z_{abs}\\leq 0.55$, we find the average luminosity enclosed within 100 kpc around MgII absorbers to be $M_g=-20.65\\pm0.11$ mag, which is $\\sim0.5~L_g^\\star$. The global luminosity-weighted colors are typical of present-day intermediate type galaxies.  We then investigate these colors as a function of MgII rest equivalent width and find that they follow the track between spiral and elliptical galaxies in color space: while the light of weaker absorbers originates mostly from red passive galaxies, stronger systems display the colors of blue star-forming galaxies. Based on these observations, we argue that the origin of strong MgII absorber systems might be better explained by models of metal-enriched gas outflows from star-forming/bursting galaxies.  Our analysis does not show any redshift dependence for both impact parameter and rest-frame colors up to $z=1$.  However, we do observe a brightening of the absorbers related light at high redshift ($\\sim 50\\%$ from $z_{abs}\\sim 0.4$ to 1). We argue that MgII absorbers are a phenomenon typical of a given evolutionary phase that more massive galaxies experience earlier than less massive ones, in a \\emph{downsizing} fashion.  This analysis provides us with robust and quantitative constraints of interest for further modeling of the gas distribution around galaxies. As a side product we also show that the stacking technique allows us to detect the light of quasar hosts and their environment. ",
        "introduction": "Quasar absorption lines provide us with a unique tool to probe the gas content in the Universe with a sensitivity that does not depend on redshift.  In order to use this information for investigating galaxy formation and evolution, characterizing the underlying population of galaxies seen in absorption is an important task to achieve.  In 1969, Bahcall \\& Spitzer suggested that metal absorption lines seen in quasar spectra are induced by large halos of gas around galaxies extending up to a hundred kpc.  Since then, the study of quasar absorption lines has become an active field of investigation, but fundamental questions regarding the nature of this gas and its distribution around galaxies still remain.  Studies of absorber-galaxy connections began with the investigation of MgII systems. Such a choice results from a practical constraint: among the dominant ions in HI gas, MgII has the longest-wavelength resonance transition at the doublet $\\lambda\\lambda 2796.35,2803.53$.  Several studies showed that MgII arises in gas spanning more than five decades of neutral hydrogen column density \\citep{bergeron_stasinska_86,steidel_sargent_92,churchill+00}, and at a similar time the first galaxies responsible for some metal absorption features were identified \\citep{bergeron_86,bergeron+87, cristiani_87,bergeron_boisse_91}. \\cite{SDP94} then built a sample of 58 identified MgII systems and found that various types of galaxies with $\\sim L^\\star$ luminosities exhibit MgII absorption up to a radius of $R\\sim 50 h^{-1}$ kpc. For a review of results regarding the connection between MgII absorbers and galaxies see \\citet{churchill+05}. It is also interesting to mention that \\cite{bowen+06} recently reported the detection of MgII absorption around quasars.  While the connection between galaxies and metal absorbers has been clearly established, the origin of the absorbing gas and its properties remains unclear.  Whether these absorption lines trace gas being accreted by a galaxy or outflowing from it is still a matter of debate. Recently the SDSS database has provided us with significantly larger samples of absorbers (by two orders of magnitude), allowing us to accurately measure a number of statistical properties such as the redshift distribution, $dN/dz$, and the distribution of absorber rest equivalent widths, $dN/dW_0$, which is now measured over more than a decade in $W_0$ \\citep{nestor+05,prochter+06}. Unfortunately, physical models or numerical simulations reproducing these observations are still lacking. While efforts on the theoretical side are needed, additional \\emph{statistical} observational constraints are required, especially to characterize the spatial distribution of the gas around galaxies and their corresponding optical properties.  Previous attempts to constrain such properties of MgII absorbers have relied on deep imaging and spectroscopic follow-up of (arbitrarily faint) galaxies in the quasar field in order to identify a potential galaxy responsible for the absorption seen in the quasar spectrum. Such expensive studies have been limited to samples of a few tens of quasars. Moreover recent results have shown discrepancies in the inferred distributions of absorber impact parameters.  Whereas \\cite{SDP94} claimed that $\\sim L^\\star$ galaxies are surrounded by MgII halos extending up to $R\\sim 50 h^{-1}$ kpc, \\cite{churchill+05} recently showed that a more systematic analysis reveals the same type of absorption up to $R\\sim 100 h^{-1}$ kpc in a number of cases. The different results apparently arise from different ways of preselecting the candidate galaxies for the spectroscopic follow-up.  Furthermore, it is interesting to note that the recent detections of MgII around quasars \\citep{bowen+06} reach similar distances. Estimating the distribution of absorber impact parameters \\emph{without the use of any assumption} is therefore needed.  It should be emphasized that the identification of a galaxy responsible for a given absorption feature is in general not a well defined process. Indeed it is always possible that an object fainter than the limiting magnitude, such as a dwarf galaxy, gives rise to the absorbing gas, or that more than one galaxy has a redshift consistent with that of the absorber system which is likely to happen in group environments.  A well-defined quantity that can be obtained from observations and modeled theoretically is the cross-correlation between absorbers and the measured light within a given radius: $\\langle W_0 \\, L\\rangle(\\theta)$.   In \\cite{abs_letter05} we proposed a new approach to measure the systematic photometric properties of large samples of absorbing systems, which combines statistical analysis of both spectroscopic and imaging datasets of the Sloan Digital Sky Survey \\citep[hereafter SDSS][]{york+00}. As MgII lines can only be detected in SDSS spectra at $z\\gtrsim0.4$, the absorbing galaxies are only marginally detected in individual SDSS images at the lowest redshifts, and below the noise level otherwise. However, they do produce systematic light excesses around the background QSOs. Such excesses can be detected with a statistical analysis consisting of stacking a large number of absorbed QSO images. After having demonstrated the feasibility of this technique with EDR data \\citep{abs_letter05}, we now extend our analysis to the SDSS DR4 sample, i.e., a dataset about twenty times larger.  Stacking analysis has the intrinsic limitation of only probing the mean or median value of a quantity for which the underlying distribution is not necessarily known. However, it can lead to detections of low-level signals not detectable otherwise. Moreover, one can also investigate the dependence of an observed quantity $O$ as a function of a set of parameters that describe a given sample. Some properties of the underlying distribution giving rise to $<O>$ can therefore be inferred.  Stacking approaches have provided us with valuable results in a number of areas. For example, \\cite{bartelmann_white_03} have demonstrated that stacking of ROSAT All-Sky Survey X-ray images of high redshift clusters detected in the SDSS can be used to derive their mean X-ray properties. \\cite{zibetti+05} detected low levels of intracluster light by stacking SDSS cluster images, while \\cite{lin_mohr_04} measured their infrared (IR) luminosity by coadding  2MASS data. In the context of galaxies, \\cite{hogg+97} constrained the IR signal from faint galaxies using stacked Keck data. Similarly, \\cite{brandt+01} measured the mean X-ray flux of Lyman Break galaxies, \\cite{zibetti+04} characterized the very low-surface brightness \\emph{diffuse} light in galaxy halos and, more recently, \\cite{white+06} constrained the radio properties of SDSS quasars down to the nanoJansky level.  In spectroscopic studies, stacking techniques have been extensively used to look for weak signals. Composite spectra of SDSS quasars were recently  used to detect weak absorption lines as well as dust reddening effects that are are well below the noise level in individual spectra \\citep{nestor+03,menard05,york+06}.  In order to emphasize the power of the stacking technique using the SDSS, it is instructive to compare the sensitivity that can be reached with respect to that of classical pointed observations. In the present study we make use of SDSS imaging data which is obtained from a 2.5 meter telescope, using a drift-scan technique giving rise to an exposure time of about one minute. The samples of quasars with and without absorbers used for the stacking analysis include approximately 2,500 and 10,000 objects respectively. The careful analysis of the corresponding 0.5 Terabyte of imaging data allows us to reach equivalent exposure times of about 40 and 160 hours respectively, i.e., 4 and 16 nights of observations with the SDSS telescope. Such numbers imply that, as long as systematic effects are small, detecting faint levels of light from high-redshift becomes feasible.   In this paper we present the results of a stacking analysis aimed at characterizing the light associated with MgII absorbers and investigating its properties as a function of absorber rest equivalent width and redshift. As we will show, the sensitivity achieved with this technique even allows us to detect the light from the high-redshift quasar hosts. The paper is organized as follows: in \\S\\ref{data_sec} we present our sample of MgII absorbers; in \\S\\ref{improc_sec} we describe the image processing and the technique adopted to extract photometric information from the stack images; in \\S\\ref{spatial_sec} we analyze the surface brightness profiles of the light in excess around absorbed QSOs and derive the luminosity weighted impact parameter distribution of the MgII systems around galaxies; \\S\\ref{SED_subsec} is devoted to the interpretation of the integral photometry in terms of rest-frame SEDs and luminosity; in \\S\\ref{discussion_sect} we briefly discuss a possible unified scenario for the interpretation of the results; finally, in \\S\\ref{summary} we present a schematic summary of the paper.   ``Standard'' cosmological parameters are assumed throughout the paper, i.e., $\\Omega_{tot}=1.0$, $\\Omega_{\\Lambda}=0.7$, and $h=0.7$. ",
        "conclusions": "\\label{summary}  In this paper we present a statistical analysis that allows us to constrain the mean photometric properties of MgII absorbing galaxies and the impact parameter distribution of the gas.  By stacking a large number of PSF-subtracted images of quasars with absorbers and comparing them with similar but non-absorbed quasars we isolate the excess light associated with MgII absorber systems and investigate its properties. Contrary to a number of previous studies which attempted to identify individual absorbing galaxies, our method measures a more general (and well defined) cross-correlation between gas and light, and does not use any assumption about the extent of the gas and/or a redshift interval within which the galaxy can be found.  Moreover, it is directly applicable to large samples, and can probe large scales. In the present study we use a sample of about 2,800 quasar sightlines with one strong MgII absorber ($W_0(\\lambda2796)>0.8$~\\AA) and a control sample of more than 11,000 reference quasars extracted from SDSS DR4. The large number of systems allows us to study the properties of the absorbing galaxies as a function of redshift and rest equivalent width ($W_0$), thus partly compensating the loss of information inherent to any stacking procedure. Our results are as follows:  \\newcounter{Scount} \\begin{list}{\\arabic{Scount})}{} {\\usecounter{Scount} \\setcounter{Scount}{0} \\item We obtain significant detections of cross-correlations between MgII gas and light out to $\\sim 200$~kpc scales. We show that the surface brightness (SB) distribution of such light is directly proportional to the impact parameter distribution of MgII systems around absorbing galaxies, weighted by galaxy luminosity. \\item Over the range 20-100 kpc, the SB profiles or impact parameter distributions are well described by power laws with an index $\\alpha$ that strongly depends on the rest equivalent width of the absorbers; $\\alpha\\simeq-1$ for low-EW systems ($0.8\\mathrm{\\AA} \\leq W_0 < 1.12\\mathrm{\\AA}$), while $\\alpha\\simeq-2$ for high-EW systems ($W_0 \\geq 1.58\\mathrm{\\AA}$). More accurate fitting formulae are also provided in \\S \\ref{impactpars_sec}.} \\end{list} At low redshift, i.e. $0.37<z_{abs}\\leq0.55$, where the S/N is high enough, we find that \\begin{list}{\\arabic{Scount})}{} {\\usecounter{Scount} \\setcounter{Scount}{2} \\item the mean luminosity enclosed within $10$ kpc $<r<100$ kpc is $M_g=-20.65\\pm0.11$ for the global sample, which is $\\sim0.5~L_g^\\star$ of the field luminosity function in the same redshift regime. \\item \\emph{The spectral energy distribution of stronger systems appears to be significantly bluer}. MgII systems with $W_0>1$\\AA~are, on average, characterized by an SED typical of star-forming/bursting galaxies, while the light of weaker systems is dominated by passive red galaxies (although they might not dominate the number counts).} \\end{list} Moreover, considering the entire redshift range $0.37<z<1.0$, we find:  \\begin{list}{\\arabic{Scount})}{} {\\usecounter{Scount} \\setcounter{Scount}{4} \\item no detectable evolution with redshift of both spatial distribution and \\emph{global} average SED, which is always well fitted by intermediate type spirals. This suggests that MgII systems trace similar galaxies or a similar evolutionary phase of galaxies. \\item The rest-frame optical luminosity increases by $\\sim 50\\%$ from $z\\sim0.4$ to $z\\lesssim1$. This phenomenon is reminiscent of the so-called \\emph{downsizing} which is observed in a number of indicators of galaxy activity. } \\end{list}  We have discussed possible interpretations of these observations including inflowing/outflowing gas and clustering effects.  While detailed modeling will be required to robustly test each model, we point out that a scenario in which metal-enriched gas outflows from star-forming/bursting galaxies is in qualitative agreement with several types of trends seen in the data.  The observational results presented in this paper, namely the cross-correlation between MgII gas and light and the mean colors of absorbing galaxies, provide robust statistical constraints of great interest for further modeling the gas distribution around galaxies. They will help produce a comprehensive and consistent picture of galaxy formation in both absorption and emission.  As a side-product of this study, we have shown that our stacking technique is able to detect the light of QSO hosts and their environment, thus opening new exciting perspectives for the study of the relationships between QSOs and galaxy formation."
    },
    "0609/astro-ph0609283_arXiv.txt": {
        "abstract": "{ We have conducted deep photometric searches for substellar members of the Praesepe (0.5--1~Gyr) and \\so~(3~Myr) star clusters using the Sloan $i'$ and $z'$ broad-band filters, with the 3.5-m and the 5-m Hale telescopes on the Calar Alto and Palomar Observatories, respectively. The total area surveyed was 1177 arcmin$^2$ and 1122 arcmin$^2$ towards the central regions of Praesepe and \\so, respectively. The 5-$\\sigma$ detection limit of our survey is measured at $i'$\\,=\\,24.5 and $z'$\\,=\\,24 mag, which according to state-of-the-art evolutionary models corresponds to masses of 50--55~$M_{\\rm Jup}$ (Praesepe) and 6~$M_{\\rm Jup}$ (\\so), i.e., well within the substellar regime. Besides recovering previously known cluster members reported in the literature, we have identified new photometric candidates in both clusters whose masses expand the full range covered by our study. In \\so, follow-up near-infrared photometry has allowed us to confirm the likely cluster membership of three newly discovered planetary-mass objects. The substellar mass function of \\so, which is complete from the star--brown dwarf borderline down to 7~$M_{\\rm Jup}$, keeps rising smoothly with a slope of $\\alpha$\\,=\\,0.6\\,$^{+0.5}_{-0.1}$ (d$N$/d$M$ $\\sim$ $M^{-\\alpha}$). Very interestingly, one of the faintest Praesepe candidates for which we have also obtained follow-up near-infrared $JHK_s$ photometry nicely fits the expected optical and infrared photometric sequence of the cluster. From its colors, we have estimated its spectral type to be between L4 and L6. If confirmed as a true Praesepe member, it would become the first L-type brown dwarf (50--60~$M_{\\rm Jup}$) identified in an intermediate-age star cluster. Our derivation of the Praesepe mass function, which is based on state-of-the-art evolutionary models, depends strongly on the cluster age. For the youngest possible ages (500--700~Myr), our results suggest that there is a deficit of Praesepe brown dwarfs in the central regions of the cluster, while the similarity between the Praesepe and \\so~mass functions increases qualitatively for models older than 800~Myr. ",
        "introduction": "The knowledge of the substellar mass function of various star-forming regions and very young open clusters ($\\le$\\,10~Myr) can provide key constrains on our understanding of the formation mechanisms giving rise to very low-mass stars and brown dwarfs (masses below 75~$M_{\\rm jup}$, or 0.072~\\Msol, Chabrier \\& Baraffe \\cite{chabrier00}). The study of intermediate-age star clusters (100--1000~Myr) is relevant to track down the dynamical evolution of brown dwarfs and its impact on the shape of the mass function. It is also important to detail the brown dwarfs luminosity and effective temperature changes with time. Furthermore, ultracool L- and T-type brown dwarf members of intermediate-age star clusters would become benchmark objects to understand the very low-mass stellar and substellar populations in the solar neighborhood.  In this paper, we present the results of deep photometric surveys in two open clusters of different ages: \\so~and Praesepe. Our main objectives are: the identification, for the first time, of the Praesepe unambiguous substellar population, and the extension of our previous deep explorations carried out in \\so~to search for planetary-mass candidates below the deuterium burning-mass limit (Zapatero Osorio et al. \\cite{osorio00}). Both clusters have roughly solar metallicity (Caballero \\cite{caballero06}; Hambly et al. \\cite{hambly95a}) and are characterized by very little internal extinction (Lee \\cite{lee68}; Crawford \\& Barnes \\cite{crawford69}). The Praesepe cluster has an age in the interval 0.5--1~Gyr, and is located at a distance of 180~pc (Robichon et al. \\cite{robichon99}), while the youngest cluster, \\so~(1--8~Myr, Oliveira et al. \\cite{oliveira02}; Zapatero Osorio et al. \\cite{osorio02}, with a likely mean age around 3~Myr) lies at 352~pc (Perryman et al. \\cite{perryman97}).  Previous large scale proper motion studies of Praesepe (Hambly et al. \\cite{hambly95b}; Adams et al. \\cite{adams02}) covered most of the lower main sequence of the cluster down to 0.1~$M_{\\odot}$. Deeper surveys by Pinfield et al. (\\cite{pinfield97}, \\cite{pinfield00}), Magazz\\`u et al. (\\cite{magazzu98}), and Chappelle et al. (\\cite{chappelle05}) have extended the Praesepe population down to spectral types M9--L0, which corresponds to a mass of 0.06--0.08~$M_{\\odot}$, i.e., rather close to the substellar borderline. Chappelle et al. (\\cite{chappelle05}) concluded that Praesepe substellar members are likely depleted from the cluster due to dynamical evolution. Our survey does not cover a significant area of Praesepe (thought to be of several tens square degrees), but to the best of our knowledge, is the deepest photometric exploration carried out in the cluster central regions so far.  We present our optical surveys and cluster member candidate selection criteria in Section~2. Follow-up near-infrared photometry (also presented in Section~2) allows us to refine photometric membership status of the candidates. In Section~3 we discuss on the mass of the selected candidates, derive the Praesepe and \\so~mass functions, and compare them to various mass functions available in the literature. We find new planetary-mass candidates (6--13~$M_{\\rm Jup}$) in \\so, and one likely Praesepe brown dwarf (50--60~$M_{\\rm Jup}$) with colors typical of L4--L6 spectral types. ",
        "conclusions": "\\subsection{Cluster membership and mass of the candidates}  Our photometric survey was designed to detect in both $i'$ and $z'$ filters Praesepe brown dwarfs down to about 50 times the mass of the planet Jupiter (1~$M_{\\odot}$\\,=\\,1047~$M_{\\rm Jup}$), or 55~$M_{\\rm Jup}$ (completeness), i.e., well below the substellar mass limit defined at 75~$M_{\\rm Jup}$ (Chabrier \\& Baraffe \\cite{chabrier00}). This prediction is based on the state-of-the-art evolutionary models by Baraffe et al. (\\cite{baraffe98}), an age of 500~Myr, and a distance modulus of $(m-M)$\\,=\\,6.28 mag (180~pc, Robichon et al. \\cite{robichon99}). For the older age of 1~Gyr, the mass limit of our survey is estimated at 60~$M_{\\rm Jup}$, still in the substellar regime. In terms of spectral type, these brown dwarfs would show mid- to late-L spectral appearance with surface temperatures of about 1500\\,K (Vrba et al. \\cite{vrba04}). Similarly, our \\so~survey is sensitive to the same spectral interval, which corresponds to masses of 6--7~$M_{\\rm Jup}$ (limit) and 8~$M_{\\rm Jup}$ (completeness) at the age (3\\,Myr) and distance (352~pc) of the cluster.  \\begin{figure} % \\resizebox{\\hsize}{!}{\\includegraphics{jk_prae.ps}} \\caption{Near-infrared color-magnitude diagram of Praesepe members. Our ``low'' and ``high'' new optical candidates with $JK_s$ photometry are plotted as open triangles and filled circles, respectively, along with their associated error bars. Previously known members and data from the literature are shown with open circles. The field sequence at the distance of the cluster is plotted as a solid line (spectral types are labeled). Masses according to an age of 500~Myr and the models by Chabrier et al. (\\cite{chabrier00}) are indicated to the right side of the diagram.} \\label{prae3} \\end{figure}  \\begin{figure} % \\resizebox{\\hsize}{!}{\\includegraphics{nir_sigori.ps}} \\caption{Optical and near-infrared color-magnitude diagram of \\so~members. Our candidates with $i'z'$ high photometric membership probability are plotted as filled circles, while open triangles stand for low probability ones. Previously identified cluster members are shown with open circles. Overplotted onto the data is the 3-Myr isochrone from Chabrier et al. (\\cite{chabrier00}) and Baraffe et al. (\\cite{baraffe03}). Masses are labeled to the right.} \\label{sigori3} \\end{figure}  \\begin{figure*} % \\caption{Images ($2' \\times 2'$ in extent) of Prae\\,J084111.9$+$192043 taken at four different wavelengths. [Provided in jpeg format in astro-ph.]} \\label{J084111.9} \\end{figure*}  There are various new Praesepe and \\so~candidates discovered in our survey that have been labeled with a ``high'' flag in the optical photometric cluster membership columns of Tables~\\ref{newprae} and~\\ref{newso}. All of them are plotted as filled circles in Fig.~\\ref{iz_cands}, where previously known cluster members from the literature are also shown. Overplotted onto the data are the 500-Myr (Praesepe) and 3-Myr (\\so) isochrones obtained from the NextGen (3000--2000~K, Baraffe et al. \\cite{baraffe98}), dusty (2000--1300~K, Chabrier et al. \\cite{chabrier00}) and COND (1300--800~K, Baraffe et al. \\cite{baraffe03}) evolutionary models. To derive the predicted $i'$ and $z'$ magnitudes from the theoretical $T_{\\rm eff}$'s and luminosities, we have used the relations between the observed Sloan $(i'-z')$ colors of field late-M and L dwarfs and their measured surface temperatures and $J$-band bolometric corrections compiled from the works by Knapp et al. (\\cite{knapp04}), Vrba et al. (\\cite{vrba04}), and Dahn et al. (\\cite{dahn02}). As can be seen from Fig.~\\ref{iz_cands}, the models nicely reproduce the shape and trend of the Praesepe and \\so~optical photometric sequences. Given the age of Praesepe (0.5--1~Gyr), which is close to that of the field, the cluster sequence also matches within the expected photometric errors the sequence delineated by field very low-mass stars and brown dwarfs as it is shown in the Figure. The theoretical 500-Myr isochrone appears slightly overluminous as expected for a younger age. The largest differences in luminosity are observed at colors $(i'-z')$\\,$\\ge$\\,3~mag.  Figures~\\ref{prae3} and~\\ref{sigori3} illustrate the near-infrared color-magnitude diagram of the Praesepe cluster and the optical-near infrared diagram of \\so, respectively. As before, the field sequence moved to the distance of Praesepe is overplotted in Fig.~\\ref{prae3}. It provides a good guideline of the expected photometric trend that true Praesepe members should follow at $J$ and $K_s$ magnitudes. One of the new brightest candidates, Prae\\,J083850.6$+$192317, nicely fits the cluster photometric sequence at both optical and near-infrared wavelengths (see Figs.~\\ref{iz_cands} and~\\ref{prae3}), as also does Prae\\,J084051.2$+$191010, although this candidate was labeled with a ``low'' flag in Table~\\ref{newprae}. These objects are likely Praesepe stars with masses estimated at 0.10--0.15~$M_\\odot$.  Another probable photometric Praesepe candidate is Prae\\,J084039.3$+$192840, which has similar brightness and colors than those of the faintest cluster members identified by Chappelle et al. (\\cite{chappelle05}). Because of its similarity to RPr\\,1 (Magazz\\`u et al. \\cite{magazzu98}) and based on its optical color, we estimate its spectral type to be between M8 and L1, i.e., in the M--L transition. According to the evolutionary models, Prae\\,J084039.3$+$192840 sits on top of the expected star--brown dwarf frontier of the Praesepe cluster.  Very interestingly, we have identified an outstanding Praesepe candidate member that is signficantly fainter and redder than any other candidate: Prae\\,J084111.9$+$192043, which is detected as a very red source in both the Palomar and Calar optical surveys. As can be seen from Fig.~\\ref{iz_cands}, its optical colors are consistent with mid-L spectral type, and models predict a mass of 50--60~$M_{\\rm Jup}$ for it (0.5--1~Gyr). The near-infrared photometry of Table~\\ref{prae_nir_phot} and the compilation of images depicted in Fig.\\ref{J084111.9} confirm its cool nature. Prae\\,J084111.9$+$192043 does follow the expected Praesepe cluster sequence both at optical and near-infrared wavelengths (Figs.~\\ref{iz_cands} and~\\ref{prae3}). We can estimate the number of L4--L6 field dwarfs that may appear in our survey with similar magnitudes to those of Prae\\,J084111.9$+$192043 by taking into account the local number density of these objects in the solar neighbourhood. Cruz et al. (\\cite{cruz03}) found an L0--L4 dwarf space density of 1.9$\\times$10$^{-3}$ objects\\,pc$^{-3}$, while for T dwarfs (which are cooler than the L-type objects), Burgasser (\\cite{burgasser01}) derived a bias-corrected value of 21$\\times$10$^{-3}$ objects\\,pc$^{-3}$. In qualitative terms, there are more T dwarfs than L dwarfs. We will adopt a mean density of 10$\\times$10$^{-3}$ dwarfs with mid-L types per cubic parsec, which is intermediate between the early-L and T classes. For the area explored in the Praesepe survey, we would have expected 0.4 such field objects contaminating our optical search, i.e., less than 1. There is a $\\sim$70\\%~probability that Prae\\,J084111.9$+$192043 is a true Praesepe member. We also note that there is a negligible probability for this object being a background very red galaxy. On the one hand, it remains unresolved in our good-seeing WHT near-infrared images. On the other hand, Fig.~10 of Metcalfe et al. (\\cite{metcalfe06}), which displays the optical near-infrared color-magnitude diagram of galaxies in the interval $H$\\,=\\,16--22.5 mag, shows an empty area around the location defined by the colors and magnitudes of Prae\\,J084111.9$+$192043. Nevertheless, further data are needed to unambiguously confirm its cluster membership, for example, lithium detection in its optical spectrum and proper motion.  Regarding the \\so~cluster, we note that with a few exceptions, there is a one to one relation in the optical and near-infrared photometric cluster membership probability flags given in Table~\\ref{sigori_nir_phot}. By considering the optical and near-infrared photometric data of both known and new \\so~candidates (36 objects out of a total of 43 candidates have photometry at various wavelengths), we estimate that the success rate on the identification of likely cluster members in our survey is about 93\\%~among the ``high'' candidates, and about 38\\%~among the ``low'' ones. Hence, to build reliable cluster luminosity and mass functions we discard the ``low'' optical candidates unless they are proved to follow the near-infrared sequence of \\so.  Of special interest is the population number in the planetary-mass regime of the \\so~cluster. In our survey, we have recovered two previously known cluster candidates whose masses could be below the deuterium burning-mass limit ($z' \\ge 20.5$~mag, $J \\ge 18$~mag). Additionally, we have identified three high-probable and six low-probable new planetary-mass candidates. Near-infrared photometry is available for four of them (two ``high'' and two ``low''). As can be seen from Fig.~\\ref{sigori3}, one low-probable planetary-mass candidate (S\\,Ori\\,J053919.2$-$025910) does not show an infrared color consistent with membership in \\so, and thus, it is not considered further in this paper. On the other hand, S\\,Ori\\,J053844.5$-$025514 and S\\,Ori\\,J053932.4$-$025220, which were given a ``high'' flag based on their optical data (see the right panel of Fig.~\\ref{iz_cands}), also fit the near-infrared cluster sequence. This suggests that they are likely planetary-mass members of the \\so~cluster. In addition, S\\,Ori\\,J054008.5$-$024551 shows infrared colors expected for true cluster members despite the fact it was originally classified as a low-probable optical candidate.  \\begin{figure*} % \\resizebox{\\hsize}{!}{\\includegraphics{dndm_prae.ps} ~ \\includegraphics{dndm_so_v2.ps}} \\caption{Praesepe (500~Myr, left panel) and \\so~(3~Myr, right panel) mass functions from stars to substellar objects. The results from our surveys are depicted with filled dots and histograms. Error bars are Poissonian. In the left panel, we have also plotted the Praesepe stellar mass functions from Adams et al. (\\cite{adams02}, solid line) and Hambly et al. (\\cite{hambly95a}, dotted line). In both panels, and for comparison purposes, the \\so~mass function of Caballero  (\\cite{tesis06}) is shown with a dashed line. It is conveniently scaled to the stellar (Praesepe) and substellar (\\so) regimes. Vertical dotted lines correspond to the hydrogen (0.075~$M_{\\odot}$) and deuterium (0.012~$M_{\\odot}$) burning-mass limits.  } \\label{mf} \\end{figure*}  According to modern evolutionary models, S\\,Ori\\,J053844.5$-$025514, S\\,Ori\\,J053932.4$-$025220, and S\\,Ori\\,J054008.5$-$024551 may have masses between 6 and 10~$M_{\\rm Jup}$ if their cluster membership is finally confirmed. In total, there are 6 \\so~planetary candidates (6--13~$M_{\\rm Jup}$) in the Palomar and Calar surveys that convey our high cluster membership probability criterion and/or fit the near-infrared cluster sequence. Another four lack follow-up near-infrared photometry. Based on the observed colors, their spectral types must range from early-L to mid-L types. From the local space number density of L1--L4 dwarfs (Cruz et al. \\cite{cruz03}), we have calculated that about 2--3 field dwarf interlopers of similar spectral classifications and brightness in the interval $z'$\\,=\\,21.3--21.9 mag may be contaminating our \\so~survey. Hence, the relatively high number of L-type cluster candidates found in our search can be understood in terms of the larger population of \\so~as compared to the field.    \\subsection{Praesepe and \\so~substellar mass functions}  We have obtained the Praesepe and \\so~mass functions by considering the recent evolutionary models of the Lyon group (Baraffe et al. \\cite{baraffe03}, and references therein), cluster ages of 500 and 3~Myr, and distances of 180 and 352~pc, respectively. For each mass, the theoretical surface temperature and luminosity were converted into observables in the manner described in preceeding Sections. Only known and new candidates conveying our photometric criterion of high-probability cluster membership, and candidates that do follow the near-infrared photometric sequences of the clusters were taken into account to derive the mass functions depicted in Fig.~\\ref{mf}. Thus, previously known objects like S\\,Ori\\,42, 44, and~67 are not included in the \\so~mass function calculation. Our data are plotted as filled dots and histograms, and cover the following mass intervals: 0.045--0.13~$M_{\\odot}$ (Praesepe), and 0.007--0.072~$M_{\\odot}$ (\\so). We note that all histogram bins of Fig.~\\ref{mf} lie within the $z'$-band completeness of our search and contain at least 5 cluster member candidates with the only exception of the most massive bins, which remain unfilled due to non-linearity and saturation effects of the CCDs, and the substellar bin of Praesepe.  As discussed by B\\'ejar et al. (\\cite{bejar01}), the number of \\so~members appears to increase toward lower masses with a smooth transition throughout the deuterium burning-mass borderline. Overplotted in Fig.~\\ref{mf} is the cluster stellar and substellar mass function recently obtained by Caballero (\\cite{tesis06}). We remark the very good agreement between Caballero (\\cite{tesis06}) and our derivation for the substellar regime. In the common mathematical linear representation of the mass function, d$N$/d$m$ $\\sim$ $m^{-\\alpha}$, our data can be fit with an exponent $\\alpha$\\,=\\,0.6\\,$^{+0.5}_{-0.1}$, where the uncertainty accounts for the different theoretical mass-luminosity relationships corresponding to an age interval 1--10~Myr.  The substellar census of Praesepe presents a different picture. In the following discussion we assume the age of Praesepe is $\\sim$500\\,Myr. In the left panel of Fig.~\\ref{mf}, we have included the cluster stellar mass functions from 0.1 up to 1~$M_{\\odot}$ obtained by Adams et al. (\\cite{adams02}) and Hambly et al. (\\cite{hambly95a}). The former authors determined that the great majority of the Praesepe stars are mainly contained within a radius of 3.8~deg, while the core radius of the cluster is about 1.1~deg as derived from the King model fitting of the surface density of stars (King \\cite{king62}). Our survey explores regions within the cluster core radius, where we expect the largest number of brown dwarfs to be found, and clearly extends the mass functions of Adams et al. (\\cite{adams02}) and Hambly et al. (\\cite{hambly95a}) toward the substellar domain. For a proper comparison between our Praesepe mass function and the results drawn from the literature, we have normalized the mass functions of Adams et al. (\\cite{adams02}) and Hambly et al. (\\cite{hambly95a}) by taking into account the ratio between the total cluster area and our surveyed area, and the assumption that both brown dwarfs and low-mass stars share a similar spatial distribution in Praesepe. As can be seen from Fig.~\\ref{mf}, Adams et al. (\\cite{adams02}) found that the number of Praesepe stars rises from 1.0 to 0.4~$M_{\\odot}$, where the mass function distinctly becomes flat down to 0.1~$M_{\\odot}$. This contrasts with the continuing rising function below 0.4~$M_{\\odot}$ of Hambly et al. (\\cite{hambly95a}). The two stellar bins derived from our search are consistent with the findings of the latter authors.  In the substellar regime, Chappelle et al. (\\cite{chappelle05}) argued that there is a paucity of brown dwarfs in Praesepe leading to a turnover of the mass function near the substellar boundary. We can test this via the direct comparison of the results of our surveys in Praesepe and \\so. In our search, there are 5 high-probable Praesepe candidates with masses around the substellar limit (70--90~$M_{\\rm Jup}$), while only one high-probable candidate has appeared at much fainter magnitudes and with a likely mass of 52~$M_{\\rm Jup}$ at 500\\,Myr. On the contrary, we estimate that 4--6 Praesepe brown dwarfs should have been found in the mass interval 45--65~$M_{\\rm Jup}$ if the Praesepe and \\so~clusters have similar mass functions. Furthermore, the \\so~photometric sequence of our survey extends well below the deuterium burning mass limit down to about 6~$M_{\\rm Jup}$. The apparent drop of a factor of 5 in the number of Praesepe cluster members between the substellar limit and the brown dwarf regime is indeed indicative of the deficit of massive brown dwarfs in the central regions of Praesepe, confirming the previous conclusions of Chappelle et al. (\\cite{chappelle05}). This result is also in agreement with the findings in the Hyades cluster by Dobbie et al. (\\cite{dobbie02}).  It also points to a marked discrepancy between the \\so~and Praesepe substellar mass functions. As can be seen from Fig.~\\ref{mf}, there is a shallow rising transition between stars and brown dwarfs in \\so, while the substellar bin ($\\sim$55~$M_{\\rm Jup}$) of Praesepe, which is affected by a strong uncertainty since it is based on one single candidate, appears significantly lower than the contiguous, more massive stellar bin. This suggests a sharp change of the Praesepe mass function at the substellar boundary to a rapidly decreasing trend for smaller masses. The population of brown dwarfs as compared to stars seems to be more abundant in \\so~than in the central regions of the Praesepe cluster.  There are several explanations to account for the discrepancy between the \\so~and Praesepe substellar mass functions that we will enumerate below. First, our survey covers a quite small area of the Praesepe cluster, i.e., it might yield results that are non-representative of the entire cluster population. However, we note that other surveys mapping larger areas (e.g., Chappelle et al. \\cite{chappelle05}) reach similar conclusions to ours. Second, Praesepe is a much older star cluster when compared to \\so. Thus, it has likely experienced some mass segregation and dynamical evolution: the smallest cluster members increase their kinetic energy and may have been dispersed from the central regions of Praesepe due to encounters and gravitational interactions with more massive cluster stars. Actually, Adams et al. (\\cite{adams02}) concluded that there has been some evaporation in Praesepe, i.e., the core of the cluster is depleted of low-mass stars. Because of these effects, future searches for the Praesepe brown dwarf population should be conducted in the halo of the cluster.  Third, it might be possible that substellar objects cool down much faster than predicted by current evolutionary models. This would make brown dwarfs fainter than expected at the age of Praesepe, i.e., there is a steeper mass-luminosity relation between the cluster smallest stars in thermonuclear equilibrium and the most massive brown dwarfs. Member candidates like Prae\\,J084111.9$+$192043 would have larger masses. The substellar borderline is accurately determined to happen at M6--M7 spectral types ($\\sim$2800~K) in the Pleiades (Stauffer et al. \\cite{stauffer98}; Mart\\'\\i n et al. \\cite{martin98}), which has an age of 120~Myr. It is expected that the star--brown dwarf frontier lies at cooler types for older ages. Lithium is used as a diagnostic for substellarity (Rebolo et al. \\cite{rebolo92}): the detection of this atomic feature in the spectra of old, ultracool dwarfs implies lithium preservation and a mass below the substellar limit (see Basri \\cite{basri00} for a comprehensive review). As discussed by Kirkpatrick et al. (\\cite{kirk00}), the presence of lithium in the optical spectra of field L0-type dwarfs is $\\sim$0\\%, and slowly rises to 20\\%~by L2, and to 50--70\\%~by L6. Thus, about 50--70\\%~of the field mid-L type dwarfs are bona-fide brown dwarfs for which models predict a mass smaller than about 60~$M_{\\rm Jup}$. This suggests that the substellar borderline is located between spectral types M6 and L6 in both the field and the Praesepe cluster.  Other alternative possibility is the non-universality of the initial mass function. However, to address this issue a wide variety of much younger star clusters have to be studied. On the one hand, the Praesepe stellar mass function between 1 and 0.4~$M_{\\odot}$ resembles that of the \\so~cluster (see Fig.~\\ref{mf}). On the other hand, evidence for mass segregation in the Praesepe cluster has been reported in the literature (e.g., Raboud \\& Mermilliod \\cite{raboud98}), which makes this cluster a non-ideal site to investigate the initial mass function.  Finally, it is possible that the age of the Praesepe cluster is significantly larger than 500\\,Myr. This has an impact on the substellar mass function calculation since theoretical brown dwarf mass-luminosity relationships depend strongly on age. An older age of 800\\,Myr or higher would bring the substellar mass functions of Praesepe and \\so~to a better agreement. We estimate the mass of our coolest Praesepe candidate, Prae\\,J084111.9$+$192043, to be 57~$M_{\\rm Jup}$ at the age of 800\\,Myr, and 63~$M_{\\rm Jup}$ at 1\\,Gyr. However, the mass bin at the substellar limit roughly changes from 500\\,Myr to 1\\,Gyr, suggesting a shallow transition from Praesepe very low-mass stars to massive brown dwarfs, as in \\so. The various determinations of the age of Praesepe available in the literature are generally in the range 500--900\\,Myr, and recent determinations based on the isochrone-fitting procedure yield a likely age of 800\\,Myr (Kharchenko et al. \\cite{kharchenko05}). For this and older ages, the bulk of the cluster brown dwarf population lies at magnitudes fainter than those of Prae\\,J084111.9$+$192043 and will show colors typical of late-L and T-types. Future searches should take this possibility into account."
    },
    "0609/astro-ph0609816_arXiv.txt": {
        "abstract": "We consider essential mechanisms for orbit-shrinkage or ``hardening'' of compact binaries in globular clusters to the point of Roche-lobe contact and X-ray emission phase, focussing on the process of collisional hardening due to encounters between binaries and single stars in the cluster core. The interplay between this kind of hardening and that due to emission of gravitational radiation produces a characteristic scaling of the orbit-shrinkage time with the single-star binary encounter rate $\\gamma$ in the cluster which we introduce, clarify, and explore. We investigate possible effects of this scaling on populations of X-ray binaries in globular clusters within the framework of a simple ``toy'' scheme for describing the evolution of pre-X-ray binaries in globular clusters. We find the expected qualitative trends sufficiently supported by data on X-ray binaries in galactic globular clusters to encourage us toward a more quantitative study. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609027_arXiv.txt": {
        "abstract": "{} {We aim to characterize the $U$-band variability of young brown dwarfs in the Taurus Molecular Cloud and discuss its origin.} {We used the \\xest, where a sample of 11 young \\bona~brown dwarfs (spectral type later than M6) were observed simultaneously in X-rays with \\xmm~and in the $U$-band with the \\xmm~Optical/UV Monitor (OM).} {We obtained upper limits to the $U$-band emission of 10 brown dwarfs ($U$$\\ge$19.6--20.6\\,mag), whereas 2MASS\\,J04141188+2811535 was detected in the $U$-band. Remarkably, the magnitude of this brown dwarf increased regularly from $U$$\\sim$19.5\\,mag at the beginning of the observation, peaked 6\\,h later at $U$$\\sim$18.4\\,mag, and then decreased to $U$$\\sim$18.65\\,mag in the next 2\\,h. The first OM $U$-band measurement is consistent with the quiescent level observed about one year later thanks to ground follow-up observations. This brown dwarf was not detected in X-rays by \\xmm~during the OM observation.} {We discuss the possible sources of $U$-band variability for this young brown dwarf, namely a magnetic flare, non-steady accretion onto the substellar surface, and rotational modulation of a hot spot. We conclude that this event is related to accretion from a circumsubstellar disk, where the mass accretion rate was about a factor of 3 higher than during the quiescent level.} ",
        "introduction": "Young brown dwarfs (BDs) share similar properties with pre-main-sequence low-mass stars, suggesting that they are also undergoing a T~Tauri phase. BDs have accretion disks \\citep[e.g.,][]{jayawardhana03,natta04,mohanty05}, and photometric \\citep{scholz04a,scholz04b}, as well as spectroscopic variability \\citep{scholz05a,scholz06}, pointing to variable mass accretion. Besides, BDs also have an active corona, emitting soft X-rays \\citep[e.g.,][]{neuhaeuser99,preibisch05}.  The gas accreting from the circumstellar disk onto the stellar surface of T~Tauri stars produces a UV/blue excess emission \\citep[see review by][]{bouvier06}. The $U$-band luminosity of T~Tauri stars is directly related to the accretion luminosity \\citep{gullbring98}. However, due to their faintness in the optical, BDs are poorly studied at short wavelengths \\citep[][]{gizis05}. During the \\xest~\\citep[XEST;][]{guedel06b}, the X-ray emission of a sample of 16 \\bona~BDs (spectral type later than M6) was surveyed with \\xmm~\\citep{grosso06}. Eleven of these BDs were located less than $\\sim$$8.5\\arcmin$--\\,$\\sim$12$\\arcmin$ away from the prime target of 9 \\xmm~pointings, which allowed for simultaneous observations with the \\xmm~Optical/UV Monitor \\citep[OM;][]{mason01}, that we report here.  We present the constraints on the $U$-band photometry of these BDs obtained with the OM in Sect.~\\ref{aperture}. In Sect.~\\ref{variability}, we investigate the variability of 2MASS\\,J04141188+2811535 (shortened to 2MASS\\,J0414 henceforth) in the $U$-band and in X-rays. We discuss the origin of the $U$-band event observed from this BD in Sect.~\\ref{discussion}.  \\begin{table*}[!ht] \\caption{The sample of the 11 young BDs, with spectral type later than M6, surveyed with the \\xmm~Optical/UV Monitor (OM) in the XEST \\citep{guedel06b}. Columns~(1) and (2) give the BD names and the 2MASS counterparts, respectively. Column~(4) is the references of the spectral types in Col.~(3), with B02=\\citet{briceno02}, G06=\\citet{guieu06}, L06=\\citet{luhman06}, and M05=\\citet{muzerolle05b}. Column~(5) gives the disk accretion rate when available (\\citealt{mohanty05}; for 2MASS\\,J0414, \\citealt{muzerolle05b}). Column~(6) is the XEST field number, plus the source number if there is an X-ray detection \\citep{grosso06}. Columns~(7) and (8) indicate the duration of one single full-frame/imaging window exposure and the total number of similar exposure obtained on the BD. Column~(9) gives the $U$-band magnitude or the $U$-band limiting magnitude of each window exposure for a 3$\\sigma$ point-source detection (this work). For 2MASS\\,J0422, the limiting magnitude was computed for the broader $UVW1$ filter (200--400\\,nm), which includes the $U$-band.} \\label{table:bd} \\vspace{-0.25cm} \\centering \\begin{tabular}{lclcclccr} \\hline\\hline \\noalign{\\smallskip} \\multicolumn{1}{c}{Name} & 2MASS & \\multicolumn{1}{c}{SpTyp} &   \\multicolumn{1}{c}{Ref.} & $ \\log{\\dot{M}}$ & \\multicolumn{1}{c}{XEST} & \\multicolumn{1}{c}{Window exposure} & $N$ & \\multicolumn{1}{c}{$U$}\\\\ &                &                  &    & M$_\\odot$\\,yr$^{-1}$ & & \\multicolumn{1}{c}{ks} & & \\multicolumn{1}{c}{mag} \\\\ \\multicolumn{1}{c}{(1)} & (2) & \\multicolumn{1}{c}{(3)} & \\multicolumn{1}{c}{(4)} & \\multicolumn{1}{c}{(5)} & \\multicolumn{1}{c}{(6)} & \\multicolumn{1}{c}{(7)} & \\multicolumn{1}{c}{(8)}& \\multicolumn{1}{c}{(9)}\\\\ \\hline 2MASS\\,J0414  & J04141188+2811535   & M6.25 & M05  & -10      & 20     & 2.5 & 10  & 19.5\\\\ 2MASS\\,J0421  & J04215450+2652315   & M8.5 & L06  & \\dotfill & 11     & 5.0 & 5.0 & $\\ge$20.4\\\\ 2MASS\\,J0422  & J04221332+1934392   & M8   & L06  & \\dotfill & 01-062 & 3.5 & 4.0 & $\\ge$19.9\\\\ KPNO-Tau\\,4   & J04272799+2612052   & M9.5  & G06  & -11.1    & 02     & 5.0 & 4.6 & $\\ge$20.4 \\\\ KPNO-Tau\\,5   & J04294568+2630468   & M7.5  & B02  & n.a.     & 15-044 & 1.3 & 3.9 & $\\ge$19.8 \\\\ KPNO-Tau\\,6   & J04300724+2608207   & M9    & G06  & -10.8    & 14     & 5.0 & 2.8 & $\\ge$20.6 \\\\ KPNO-Tau\\,7   & J04305718+2556394   & M8.25 & B02  & -11      & 14     & 5.0 & 2.8 & $\\ge$20.6 \\\\ CFHT-Tau\\,5   & J04325026+2422115   & M7.5  & G06  & n.a.     & 04-003 & 1.5 & 4   & $\\ge$19.8 \\\\ KPNO-Tau\\,9   & J04355143+2249119   & M8.5  & B02  & n.a.     & 08     & 1.6 & 4   & $\\ge$19.6 \\\\ CFHT-BD-Tau\\,2& J04361038+2259560   & M7.5  & B02  & n.a.     & 08     & 1.3 & 8   & $\\ge$19.7 \\\\ CFHT-Tau\\,6   & J04390396+2544264   & M7.25 & G06  & \\dotfill & 05-005 & 5.0 & 3.7 & $\\ge$20.5 \\\\ \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "\\label{discussion}  2MASS\\,J0414 was identified as a (young) BD of the TMC by \\citet{luhman04}. It has a visual extinction of 1.1\\,mag; an M6.25 spectral type, which corresponds to an effective temperature of 2960\\,K; and a bolometric luminosity of 0.015\\,L$_\\odot$ \\citep{luhman04}. From these values the BD radius can be estimated to 0.5\\,R$_\\odot$. Using the evolutionary tracks from \\citet{baraffe98},  the BD mass is 0.07\\,M$_\\odot$. The only published optical spectrum of 2MASS\\,J0414, obtained on December 15, 2003, exhibits a strong excess of continuum at short wavelengths plus a strong H$\\alpha$ emission \\citep[$EW({\\rm H}\\alpha)$=250\\,\\AA;][]{muzerolle05b} indicating accretion from a circumsubstellar disk \\citep{luhman04}, with an estimated accretion rate of $\\sim$$10^{-10}$\\,M$_\\odot$\\,yr$^{-1}$ \\citep{muzerolle05b}. Blueshifted absorption components superimposed on the H$\\alpha$ accretion emission line provide evidence for an outflow \\citep{muzerolle05b}.  2MASS\\,J0414 has $B-V \\sim -0.4$ (Table~\\ref{table:mdm}), which is extremely blue for an object with an M6.25 spectral type. This strong blue excess indicates that this BD is actively accreting. Following the method used in Sect.~\\ref{aperture}, we find that its quiescent level, $U=19.5\\pm0.2$\\,mag, corresponds to $\\dot{M}=1.9\\pm0.4 \\times 10^{-10}$\\,M$_\\odot$\\,yr$^{-1}$, which is consistent with the value of \\citet{muzerolle05b} derived from spectroscopy. Therefore, the high disk accretion rate of 2MASS\\,J0414 is likely the reason why we detected only this BD in the $U$-band.  We consider 3 possible sources of $U$-band variability for this young BD, namely: a magnetic flare, non-steady accretion onto the substellar surface, and rotational modulation by a hot spot. We discuss these 3 possibilities in turn.  \\xmm~observations of coronal flares from active stars have shown an extremely good temporal correlation between UV and X-ray variations \\citep{mitra05}. A similar result was found in the optical $U$-band and X-rays in an \\xmm~observation of Proxima Cen, where several X-ray weak flares were preceded by an optical burst \\citep{guedel04}. We can predict the flare luminosity increase in the X-rays with the UV-X-ray relationship in flares of \\citet{mitra05}. The increase from quiescent to flare peak in the $U$-band is equivalent to a spectral luminosity density of $2.4\\times10^{26}$\\,erg\\,s$^{-1}$\\,$\\AA^{-1}$. The \\citet{mitra05} relationship, based on the $UVW1$ OM filter \\citep{ehle05}, yields $\\sim$$1\\times 10^{29}$\\,erg\\,s$^{-1}$ in the 0.5--8\\,keV energy range, which is a lower limit on the X-ray luminosity increase because the $U$-band filter covers roughly half of the wavelength range of the $UVW1$ filter (200--400\\,nm). Therefore, a UV flare would produce an X-ray flare at least $\\sim$10 times brighter than the upper limit on the X-ray luminosity of 2MASS\\,J0414, hence it would be easily detectable. Yet, we do not detect 2MASS\\,J0414 in X-rays during the $U$-band maximum of luminosity. Moreover, the luminosity increase lasts for at least 6\\,h and exhibits a slow continuous rise. In contrast, the $U$-band flares observed from Proxima~Cen are impulsive, and the duration of the longest flare is shorter than one hour \\citep{guedel04}. We thus conclude that the $U$-band luminosity rise is unlikely to result from a magnetic flare.  Emission line variability occurring on timescales of hours has been reported for young BDs by \\citet{scholz06}. The line variability is interpreted as non-steady accretion onto the substellar surface with the accretion rate varying by one order of magnitude over a few hours. The $U$-band luminosity rise observed for 2MASS\\,J0414 could similarly result from an accretion burst onto the central object.  Assuming a purely radial free-fall motion of the gas from the disk's truncation radius towards the corotating BD, the 6\\,h timescale of the $U$-band rise would correspond to the travel time from a distance of 3.5 substellar radii (i.e., 1.8\\,R$_\\odot$) above the substellar surface\\footnote{The ballistic infall velocity from rest at the disk's truncation radius, $r_{\\rm m}$, is: $V(r)^2=2 G M_\\star \\,(1/r-1/r_{\\rm m})$ \\citep[e.g.,][]{hartmann94}. Therefore, assuming a purely radial infall, the travel time from the disk's truncation radius to the substellar surface is: $t=r_{\\rm m}^{3/2}/\\!\\sqrt{2GM_\\star}\\times(\\sin{2\\theta}-2\\theta+\\pi)/2$, where $\\theta=\\arccos{(\\!\\sqrt{1-R_\\star/r_{\\rm m}})}$.}; the terminal velocity would be about 200\\,km\\,s$^{-1}$. The Keplerian period at the disk's truncation radius is then about 35\\,h, which is in the range of rotational periods observed for BDs \\citep[4--110\\,h][]{scholz05b,zapatero06}. The $U$-band increase of $\\sim$1.1\\,mag would then translate into a mass accretion rate enhanced by a factor of nearly 3 compared to the quiescent level, assuming that the $U$-band luminosity is directly related to the accretion luminosity as observed for the T~Tauri stars \\citep{gullbring98}.  As an alternative to the accretion burst, the 2MASS\\,J0414 light curve may simply reflect the rotational modulation of the object's luminosity by a hot accretion shock located on the substellar surface. Evidence for rotational modulation of the H$\\alpha$ line profile by the accretion flow has previously been reported in another young BD \\citep{scholz05a}. A hot accretion spot would result in the periodic modulation of the light curve on a timescale of the BD rotational period. The 6\\,h rise time of 2MASS\\,J0414's $U$-band light curve suggests a rotational period of the order of 12\\,h, and is thus consistent with rotational modulation.  Assuming a black-body temperature of 8000\\,K for the accretion shock, the hot spot would cover about 0.3\\% of the substellar surface \\citep[e.g.,][]{bouvier93}. This value is similar to the size of accretion shocks on classical T~Tauri stars. The two subsequent $U$-band ground observations, obtained 6.4\\,h apart, show no sign of variability, unlike the OM light curve on this timescale. This indicates that 2MASS\\,J0414 was back to its quiescent level nearly one year after the OM observations. Accretion hot spots have been reported to be variable features in T~Tauri stars, and even in one young BD \\citep[see][ and references therein]{scholz05a}.  We conclude that at least two plausible sources of variability could account for the observed $U$-band light curve of 2MASS\\,J0414: an episodic accretion event onto the central object or the rotational modulation of the BD luminosity by a hot accretion spot. In both cases these observations indicate variable mass accretion rate in this BD.  Our $U$-band survey of 11 BDs (of which 5 are actively accreting and 4 are not accreting) indicates 2MASS\\,J0414 as the brightest source, displaying variability related to accretion from a circumsubstellar disk. 2MASS\\,J0414 is then the best target of this sample of 5 accreting BDs for ground follow-up observations to investigate accretion in substellar objects, for example with simultaneous $U$-band and spectroscopic ground monitoring with a one-hour sampling."
    },
    "0609/astro-ph0609211_arXiv.txt": {
        "abstract": "We report the discovery of 4 candidate intracluster globular clusters (IGCs) in a single deep {\\sl HST ACS\\/} field of the Virgo Cluster. We show that each cluster is roughly spherical, has a magnitude near the peak of the Virgo globular cluster luminosity function, has a radial profile that is best-fit by a King model, and is surrounded by an excess of point sources which have the colors and magnitudes of cluster red giant stars.  Despite the fact that two of our IGC candidates have integrated colors redder than the mean of the M87 globular cluster system, we propose that all of the objects are metal-poor with [M/H] $< -1$.  We show that the tidal radii of our intracluster globulars are all larger than the mean for Milky Way clusters, and suggest that the clusters have undergone less tidal stress than their Galactic counterparts.  Finally, we normalize our globular cluster observations to the luminosity of intracluster stars, and derive a value of $S_N \\sim 6$ for the specific frequency of Virgo intracluster globular clusters.  We use these data to constrain the origins of Virgo's intracluster population, and suggest that globular clusters in our intracluster field have a different origin than globular clusters in the vicinity of M87.  In particular, we argue that dwarf elliptical galaxies may be an important source of intracluster stars. ",
        "introduction": "The constituents of intracluster space can tell us a great deal about the history of galaxies and clusters.  As a cluster forms, tidal interactions between galaxies and with the cluster potential affect the internal structure of galaxies, altering both their morphological and photometric properties \\citep[][and many others]{butcher1978, dressler1980, goto2003, coenda2006}.  At the same, these interactions also liberate material into intergalactic space, thus creating a fossil record of the encounters.  By studying the composition, distribution, and kinematics of these orphaned objects, we can examine the physics of tidal stripping, the distribution of matter in and around galaxies, and the initial conditions and history of cluster formation \\citep[][and many others]{merritt1984, west95, gregg1998, sommer2005}.  Intracluster globular clusters (IGCs) are an especially useful probe of these processes \\citep{west95}.  As a globular cluster evolves, it preserves information about the time of its creation, the chemistry of the gas out of which it formed, and even the gravitational forces to which it has been exposed \\citep[see][and references therein]{ashman98}.  Consequently, a large sample of IGCs can be used to trace the history of galaxy interactions and constrain both the epoch of cluster formation and the system's dynamical history.  Unfortunately, collecting and measuring a large sample of intracluster globular clusters is difficult.  Globular clusters in the halos of galaxies are routinely identified as an excess of point sources above the background, and searches for such objects have been conducted in $\\sim$100 systems out to $\\sim$100~Mpc \\citep[\\eg][]{harris79, kisslerpatig, kundu2002}.  However, as galactocentric distances increase, the surface densities of clusters decrease, so in intracluster space, the identification of globular clusters as point sources is exceedingly difficult.  As a result, there have been only a few, mostly indirect, studies of IGCs \\citep{west03, jordan03, franch03, bassino03}, and their use as cosmological probes has largely been unexploited.  Here we describe the results of a {\\sl Hubble Space Telescope\\/ (HST)} search for intracluster globular clusters (IGCs) in the nearby Virgo Cluster.  At our adopted Virgo distance \\citep[16.2~Mpc, see discussion in][]{VICS1}, globular clusters have half-light radii of $\\gap~0 \\farcs 05$, allowing them to be resolved on images taken with the Advanced Camera for Surveys ({\\sl ACS\\/}).  Moreover, because of the {\\sl ACS'} excellent sensitivity, it is possible to use the instrument to detect individual stars within the clusters and estimate their metallicities via the color of the red giant branch.  In Section~2, we describe our survey, and announce the discovery of four, well-resolved IGC candidates in Virgo.  In Section~3, we discuss the metallicities of these objects, and show that all are metal poor, with photometric properties that differentiate them from the globular clusters of Virgo's central cD galaxy, M87.  In Section~4, we compare the candidate IGCs to Galactic globular clusters and show that their half-light and tidal radii are larger than their Milky Way counterparts.  We attribute these properties to the IGCs' lack of tidal processing, and use the radii to constrain the clusters' origins.  We conclude by estimating the specific frequency of globular clusters in Virgo's intracluster space, and discussing the implications this number has for the origin of IGCs and future IGC surveys. ",
        "conclusions": ""
    },
    "0609/astro-ph0609482_arXiv.txt": {
        "abstract": "Period-luminosity sequences have been shown to exist among the Semi-Regular Variables (SRVs) of the Magellanic Clouds (Wood et al, 1999), the Bulge of the Milky Way galaxy (Glass \\& Schultheis, 2003) and elsewhere. It would clearly be useful to have absolute (trigonometric) calibrations of these relations. This paper investigates whether the sequences can be seen among the M-type giant SRVs of the solar neighbourhood.  Mass loss phenomena among these stars and their dependence on period and spectral sub-type are also discussed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609357_arXiv.txt": {
        "abstract": "The effect of a barred potential (such as the one of the Milky Way) on the galactic orbits of forty-eight globular clusters for which absolute proper motions are known is studied.  The orbital characteristics are compared with those obtained for the case of an axisymmetric galactic potential. Tidal radii are computed and discussed for both the better known axisymmetric case and that including a bar. The destruction rates due to bulge and disk shocking are calculated and compared in both galactic potentials. ",
        "introduction": "A considerable number of absolute proper motions for globular clusters have become available in recent times, largely due to the work of \\citet{D97,D99a,D99b,D00,D01,D03}. With the new proper motions, full space velocities can now be calculated. Then, using a model for the galactic mass distribution, it is possible to study the galactic orbits of a good fraction -almost a third- of the known globular clusters. Vice-versa, through their orbital motions, globular clusters can serve as kinematical probes of the gravitational field of our Galaxy.  Some of the first efforts in orbit computations for globular clusters are due to the work of \\citet{KIH73} and \\citet{KI75} who numerically integrated orbits for M67, NGC 188, Omega Cen, and NGC 2420 using full space motions in a Schmidt-type potential. With a more realistic galactic potential that included a massive halo, \\citet{AM88} and \\citet{A90} computed orbits for a total of 15 clusters -all those with then available absolute proper motions. Using an improved version of their axisymmetric potential, \\citet{AS93} recomputed these orbits and added those of a further six clusters whose absolute proper motions had been determined in the intervening years.  It is only quite recently that vastly improved measures for a significant number of clusters have become available \\citep{D97,D99a,D99b,D00, D01,D03}.  These authors use the Yale proper motion survey of the southern sky and measure motions directly with respect to external galaxies, or in difficult cases, to Hipparcos reference stars. \\citet{D99b} have also used the full space motions (obtained from their proper motions and radial velocities taken from the literature) to compute galactic orbits for 38 clusters in several models for axisymmetric potentials. From the galactic orbits they have derived a number of interesting conclusions.  The importance of using full space motions -instead of only radial velocities- when attempting to derive characteristics of the globular cluster system has been repeatedly emphasized \\citep{W98,D99b,K04}. Indeed, using only partial information -radial velocities- makes it necessary to assume some statistical distribution for two components of the space velocity.  With the availability of full space motions for about a third of the known clusters, this approach seems less and less justified, especially since in some cases its results may be misleading.  However, since evidence for our Galaxy harboring a bar has become increasingly convincing, it appears necessary to study the influence of this bar on the galactic motions of the clusters as obtained from their full space velocities and employing a galactic potential that includes a bar. These orbits can serve the purpose of calculating tidal shock effects on the cluster.  They can also serve as kinematical probes of the inner part of our Galaxy, and so aid in constraining some dynamical parameters of the bar, or at least in establishing whether or not the model potential -especially the bar- is compatible with observed characteristics of the clusters.  In this paper, we analyze the orbits in a galactic potential with a bar of 48 globular clusters with available absolute proper motions. The potential we use is that of \\citet{PMM04}, which is based on the axisymmetric model of \\citet{AS91} and includes a moderately strong bar. We compare the orbits with those obtained in the axisymmetric potential, devoting special attention to orbits that reside entirely within the bar region.  The orbits generally do not show secular changes in the total energy, $E$, or in the z-component of the angular momentum, $h$, both computed in the inertial galactic frame. Indeed, with few exceptions, those quantities are on the average conserved within 10 percent. However in a few cases, sudden changes in $E$ and $h$ do occur. Except for orbits residing entirely within the bar region, the orbits do not differ considerably from those obtained in the axisymmetric case. This similarity also applies to the tidal-shock destruction rates computed in both galactic potentials.  In Section \\ref{gpot} the galactic potential and the observational data to compute the initial conditions of the clusters' orbits are presented.  In Section \\ref{galorb} we show meridional orbits of some clusters, and some effects of the bar are considered.  The effect of the bar on the estimated tidal radii is studied in Section \\ref{tirad}, and compared with the axisymmetric case.  Destruction rates of clusters in both the axisymmetric and non-axisymmetric potentials are studied in Section \\ref{destr}. Section \\ref{concl} presents our conclusions. ",
        "conclusions": "We have obtained orbits for 48 globular clusters in both a barred and an axisymmetric galactic potential. The orbits of outer clusters (those with pericentric distances greater than about 4 kpc) are largely unaffected by the bar. The largest changes were found to occur for 'inner' clusters, i.e., those whose orbits reside mostly within the bar region. The main changes that the bar causes in the orbits are larger vertical and radial excursions. In general the bar causes no net global changes in the energy or the z-component of the angular momentum, computed in an inertial frame. However, there are cases where jumps in these quantities do occur, even causing a temporary reversal of the sense of rotation of the orbit. Tidal radii have been computed with a new expression, and with a numerical evaluation of the relevant quantities along the orbit. No noticeable changes due to the bar were found for most of the clusters. When changes do occur, they generally make the computed tidal radii somewhat larger in the presence of a bar. We find that the destruction rates due to shocks with the galactic bulge and disk are not strongly affected if we consider a barred galactic potential; in some cases (NGC 362, NGC 5897, NGC 6121, and NGC 6584) the total destruction rate due to both effects can even $decrease$ when using the barred potential.  We thus concur with the result of \\citet{LOA92}, that the similarity of destruction rates in axisymmetric and barred potentials does not rule out the existence of a galactic bar. Although we have included adiabatic corrections in the computation of tidal-shock destruction rates, these rates are smaller than those obtained by G\\&O97. Thus, to clarify this difference, the absolute proper motion data should be used with the G\\&O97 approach.    \\appendix"
    },
    "0609/astro-ph0609431_arXiv.txt": {
        "abstract": "The physical structure of a nuclear flame is a basic ingredient of the theory of Type Ia supernovae (SNIa). Assuming an exponential density reduction with several characteristic times we have followed the evolution of a planar nuclear flame in an expanding background from an initial density $6.6\\times 10^7$~\\dens down to $2\\times 10^6$~\\dens. The total amount of synthesized intermediate-mass elements (IME), from silicon to calcium, was monitored during the calculation. We have made use of the computed mass fractions, X$_{IME}$, of these elements to give an estimation of the  total amount of IME synthesized during the deflagration of a massive white dwarf. Using X$_{IME}$ and adopting the usual hypothesis that turbulence decouples the effective burning velocity from the laminar flame speed, so that the relevant flame speed is actually the turbulent speed on the integral length-scale, we have built a simple geometrical approach to model the region where IME are thought to be produced. It turns out that a healthy production of IME involves the combination of not too short expansion times, $\\tau_c\\ge 0.2$~s, and high turbulent intensities. According to our results it could be difficult to produce much more than 0.2 M$_\\sun$~of intermediate-mass elements within the deflagrative paradigma. The calculations also suggest that the mass of IME scales with the mass of Fe-peak elements, making it difficult to conciliate energetic explosions with low ejected nickel masses, as in the well observed  SN1991bg or in SN1998de. Thus a large production of Si-peak elements, especially in combination with a low or a moderate production of iron, could be better addressed by either the delayed detonation route in standard Chandrasekhar-mass models or, perhaps, by the off-center helium detonation in the sub Chandrasekhar-mass scenario. ",
        "introduction": "A distinctive feature of SNIa spectra near maximum light is the presence of intermediate-mass elements from Si to Ca moving at velocities 10000-16000 km.s$^{-1}$. In particular the presence or not of the SiII absorption feature at $6150 \\AA$~ is basic to identify a Type I supernovae as SNIa. According to observations about 0.2-0.4 M$_\\sun$ of IME have to be synthesized during the explosion. An amount of IME larger than 0.2M$_\\sun$~is also needed to explain the observed abundances of these elements in the solar neighborhood \\cite{iwa99}. Therefore, any attempt to build a satisfactory theoretical model of these explosions has to give the right amount of intermediate-mass elements.  In this respect the underproduction of intermediate-mass elements found in current pure deflagration models of Type Ia supernovae is a long standing problem of theoretical modelization of these explosions. On fundamental basis we know that these elements have to be synthesized below densities $\\rho\\simeq 5\\times 10^7$~\\dens in an expanding medium. Nevertheless the resolution achieved by the best multidimensional models carried out to date is far from adequate to solve unambigously this problem. In this work we have taken another route to address this issue. Instead of performing a multidimensional calculation incorporating the unresolved scales as a subgrid we have simulated the evolution of an expanding  nuclear flame in planar aproximation with  very good resolution. We have then used the results to put limits to the amount of IME synthesized in a deflagration. The usefulness of our approximation requires that the synthesis of elements within the mass range $28\\leq$A$\\leq 40$~ takes place, for the most part, during the flamelet regime. In such a regime the flame, although highly wrinkled and distorted, still preserves its identity. Unfortunately the precise value of the minimum density at which the flamelet hypothesis still holds is not known. There are indications that transition from the flamelet to the distributive regime in the combustion diagram could begin around $2\\times 10^7$~\\dens \\citep{nie99, zin05}, and we can take it as a reference.  Keeping that constraint in mind we have proceeded in the following way. First, the detailed structure and evolution of a nuclear flame propagating through an expanding background has been solved for a reasonable set of characteristic expansion times and the nucleosynthetic yields calculated using a simplified $\\alpha$-network of 14 nuclei, $\\S\\ 2$. With that information we provide in $\\S\\ 3$~ a rough calculation of the total amount of IME synthesized in a typical deflagration of a white dwarf as a function of the effective surface of heat exchange between ashes and fuel at the densities of interest, the turbulence intensity and the characteristic expansion time. According to our results a subsonic deflagration is able to produce more than 0.2M$_\\sun$~of IME although only in a limited region of parameter space. This point is discussed in the chapter devoted to the conclusions, $\\S\\ 4$, as well as possibles issues and limitations of our work. ",
        "conclusions": "The amount of intermediate-mass elements ejected in a SNIa explosion is an important diagnosing tool to discriminate between a variety of theoretical models attempting to explain the nature of such explosions. One of the most popular models is the deflagration model, rooted in the subsonic propagation of a nuclear flame through the white dwarf interior. However, with very few exceptions, neither the one-dimensional simulations nor modern three-dimensional calculations of this model have been able to give the right quantity of IME. Hydrodynamical calculations always stubbornly give a too low production. This fact is currently attributed to the low resolution achieved in multidimensional calculations.  In the first part of this paper we have analyzed the synthesis of elements from helium to zinc during the propagation of a nuclear flame in an expanding background. The simulation was carried out in the planar aproximation, thus the resolution was enough to resolve the detailed structure of the flame. In this sense our work represents a continuation of that by Timmes \\& Woosley (1992). As a result we were able to get an idea of the size of the region where the IME were for the most part synthesized: between 0.15 km to 0.75 km depending on the characteristic expansion time. This width is hardly reachable by today multidimensional hydrocodes. It was also possible to give the amount of elements from silicon to calcium cooked behind the flame (Table 1). This information can be used by SNIa modelers to improve the nucleosynthesis output of their hydrodynamical simulations. An interesting result is that IME are massively produced in the range  $\\rho=2\\times 10^7-5\\times 10^7$~\\dens. In fact, according to our calculations 80\\% of the IME were already produced before the density declined below $3\\times 10^7$~\\dens (see Fig. 2). This could have important consequences because the usefulness of our results strongly relies in the basic hypothesis that the flame does not enter the distributive regime above 2-3$\\times 10^7$~\\dens. Finally useful fitting formulae giving the mass of IME synthesized during the expansion as well as to their mass fraction, X$_{IME}$, as a function of the characteristic expansion time $\\tau_c$~were also given (Eqs. 8,9).  In \\S\\ 3 we did use the information given in Table 1 to estimate the quantity of IME ejected during the deflagration of a white dwarf taking advantage of the self-similar properties of the thermonuclear combustion in SNIa. We have taken a geometrical approach to address this problem in which the narrow region where the Si-peak nuclei are typically produced was treated as a tiny shell wrapping around a deformed surface on top of a  volume made of Fe-peak elements. That approach holds as long as the ashes form a well defined and not too dispersed pack limited by an irregular surface. Within our simple model the total amount of elements from silicon to calcium can be obtained from an analytical expression (Eq. 14) which uses four parameters: the ratio between the mass and average density of the NSE-ashes, the fractal dimension of their surface, the velocity of the turbulent combustion at the largest scales, and the characteristic expansion time. In consequence we have tried to delimitate the regions in parameter space able to give the amount of IME compatible with observations (Fig. 4).  As a general conclusion we state the difficulty of pure deflagration models to produce more than 0.2M$_\\sun$~of intermediate-mass elements. In this sense our phenomenological approach may help to understand why the underproduction of Si-peak elements has traditionally been a long standing problem of all sort of hydrodynamical simulations concerning Type Ia supernovae assuming the deflagrative paradigma. Still there remains a significant area in space parameter showing a rich production of IME. In particular the region defined by the unequalities $\\tau_c\\ge 0.2$~s, $D\\ge 2.2$~and $V_8\\ge 1$~seems to be well suited for that purpose. The above lower limit in the fractal dimension of the ahes's surface does not pose a problem as experiments dealing with turbulent combustion indicates a favoured value of D close to 2.3. As an aside, 1D simulations that always assume $D=2.0$ on the resolved length-scales lead unavoidably to an underproduction of IME. Perhaps more demanding is the limitation imposed by the characteristic expansion time. Even though $\\tau_c$~should not be too different from the hydrodynamical time scale, $\\simeq 0.1$~s at the densities of interest, its particular value is model dependent. Hence expansion times with $\\tau_c\\simeq 0.2$~s are probably not unrealistic. Turbulent intensities characterized by $V_8\\ge 1$~are also reachable during the explosion because the rising velocity of the larger floating blobs of incinerated material is of the order, or even greater, than that velocity. Nevertheless, the turbulent cascade generated by these fast moving blobs could fragment the flame in very small pieces, shorter than the flame thickness, invalidating our hypothesis that combustion always takes place in the flamelet regime. Thus, although our main goal was to highlight the favoured regions in parameter space to synthesize a good deal of IME, the presented results also underline the potential importance of the distributed regime for deflagration models, if they were to represent the main channel to normal Type Ia supernovae.  Even if a large mass of IME were attained in deflagration models of SNIa, our study still suggests another potential problem: the production of IME seems to increase monothonically with that of Fe-peak elements. Hence, it would be seemingly difficult for deflagrations to explain peculiar low-luminosity but energetic events like SN1991bg or SN1998de \\cite{str06}."
    },
    "0609/astro-ph0609607_arXiv.txt": {
        "abstract": "{Statistical properties of \\hii\\ region populations in disk galaxies yield important clues to the physics of massive star formation.} {We present a set of \\hii\\ region catalogues and luminosity functions for a sample of 56 spiral galaxies in order to derive the most general form of their luminosity function.} {\\hii\\ region luminosity functions are derived for individual galaxies which, after photometric calibration, are summed to form a total luminosity function comprising 17,797 \\hii\\ regions from 53 galaxies.} {The total luminosity function, above its lower limit of completeness, is clearly best fitted by a double power law with a significantly steeper slope for the high luminosity portion of the function. This change of slope has been reported in the literature for individual galaxies, and occurs at a luminosity of $\\log L = 38.6\\pm0.1$ ($L$ in erg\\,s$^{-1}$) which has been termed the Str\\\"omgren luminosity. A steep fall off in the luminosity function above $\\log L = 40$ is also noted, and is related to an upper limit to the luminosities of underlying massive stellar clusters. Detailed data are presented for the individual sample galaxies.} {The luminosity functions of \\hii\\ regions in spiral galaxies show a two slope power law behaviour, with a significantly steeper slope for the high luminosity branch. This can be modelled by assuming that the high luminosity regions are density bounded, though the scenario is complicated by the inhomogeneity of the ionized interstellar medium. The break, irrespective of its origin, is of potential use as a distance indicator for disc galaxies.} ",
        "introduction": "In their paper giving a major overview of the luminosity functions (LFs) in \\ha\\ emission of the \\hii\\ regions in disk galaxies Kennicutt, Edgar \\& Hodge (1989; hereafter KEH) noted that an important subset of their objects showed LFs with a clear break in slope at a luminosity of $L(\\ha) = 38.7$\\,dex (units of erg\\,s$^{-1}$). At lower luminosities the power law slope is flatter, whereas above the break it is steeper. KEH called those LFs showing a break ``type II\", suggesting that any galaxy will exhibit type II behaviour if it has sufficient high luminosity \\hii\\ regions. A galaxy with sufficient luminous \\hii\\ regions to show a clear break, M51, was observed by Rand (1992), who measured it at $L(\\ha) = 38.6$\\,dex. Rand's work, using CCD data, was especially precise, and revealed, as well as the break, a narrow peak in the LF close to the break luminosity. In the LF of M33, Hodge et al. (1999) found no break, which is not surprising since the function is defined only up to $L(\\ha) = 38.4$\\,dex, with only a dozen regions above $L(\\ha) = 38$\\,dex. Further examples of this type are NGC~6822 and the dwarf galaxy Holmberg~II (Hodge, Strobel \\& Kennicutt 1994), smaller galaxies with small \\hii\\ region numbers at $L(\\ha) > 38$\\,dex.  On the other hand, M101 has some 60 \\hii\\ regions with $L(\\ha) > 38.5$\\,dex and shows a clear break at $L(\\ha) ~ 38.6$\\,dex (Scowen, Dufour \\& Hester 1992).  Our group has published LFs for galaxies selected to have large numbers of high luminosity \\hii\\ regions to test the suggestion that the break luminosity shows low scatter. In three articles (Rozas, Beckman \\& Knapen 1996; Rozas et al. 1999; and Rozas, Zurita \\& Beckman 2000), eight galaxies were measured and all showed type II LFs, most strikingly NGC~7479, which has the largest number of luminous \\hii\\ regions.  In spite of the general result reported in KEH, and in spite of the type II LFs found in individual objects, the question has been raised persistently whether the break is a real feature or an artifact. Two types of doubts can be found in the literature. One claims that although breaks can be found, they are not a universal physical feature, as the break luminosity varies over quite a wide range between objects. One example here is by Thilker, Braun \\& Walterbos (2000) who derived an LF for M51 using an automated method, finding a break at 38.9\\,dex, compared with the value of 38.6 by Rand (1992) who used an interactive, region by region method. The point raised is whether for statistical reasons connected with the chosen binning parameters a clear break luminosity cannot be derived, given the low number of regions per luminosity bin. Others have not found consistent evidence for breaks, and certainly not at a specific luminosity, among these being Gonz\\'alez Delgado \\& P\\'erez (1997) in an \\ha\\ survey of 27 galaxies selected for their nuclear activity. Their sample comprised some 2000 \\hii\\ regions, i.e., $\\approx75$ per galaxy in a luminosity range from $L(\\ha) = 37$\\,dex to 39.5\\,dex. The numbers per bin above 38.5\\,dex are on average quite small for each galaxy which would lead to difficulty in detecting breaks.  On completing a recent imaging survey in \\ha\\ of a set of nearby galaxies (Knapen et al. 2004) we saw that we could make a significant test of the suggestion that a dual slope LF with a well defined break gives a better general description than a single power law slope, and that the original result of KEH, with its clean break at a specific luminosity, is present in a data set with a sufficient statistical base. Our observations have allowed us to isolate and catalogue almost 18,000 \\hii\\ regions in 56 galaxies. Here we give the key results of this study and also offer on-line catalogues for the full set of objects. ",
        "conclusions": "We have analysed continuum-subtracted \\ha\\ images of 57 nearby galaxies, and here present the LFs and catalogues of their populations of \\hii\\ regions. Using the combined data set we have produced a single LF covering the range in \\ha\\ luminosity from $L = 37.5$\\,dex to just over $L = 40$\\,dex. The total number of \\hii\\ regions contributing to this LF is a little under 18,000, which is almost an order of magnitude greater than that of any previously published LF of this type.  The best functional fit of the LF, above the completeness limit, is by a double power law, having a break at $L = 38.6$\\,dex, with the steeper slope to higher luminosity and the shallower slope to lower luminosity. The existence of this break suggests a change in physical regime going from lower to higher \\ha\\ luminosity. If confirmed, and by virtue of its occurrence at high \\hii\\ region luminosities, the break is of potential use as a distance indicator for star-forming disk galaxies.  The number of \\hii\\ regions in our sample is large enough to suggest that the steep fall off in the LF above $L = 40$\\,dex indicates that we are sampling the range where there is a cut-off in the luminosities of the \\hii\\ regions, implying a drop-off in the LF of the underlying massive stellar clusters. This is giving us potentially important information about a physical limit to the masses of stars in young clusters of varying total stellar mass."
    },
    "0609/astro-ph0609577_arXiv.txt": {
        "abstract": "Current experiments collecting high statistics in ultra-high energy cosmic rays (UHECRs) are opening a new window on the universe facing the possibility to perform UHECR astronomy. Here we discuss a large scale structure (LSS) model for the UHECR origin for which we evaluate the expected large scale anisotropy in the UHECR arrival distribution. Employing the IRAS PSCz catalogue as tracer of the LSS, we derive the minimum statistics needed to reject or assess the correlation of the UHECRs with the baryonic distribution in the universe, in particular providing a forecast for the Auger experiment. \\vspace{0pc} ",
        "introduction": "At energies above a few~$\\times 10^{19}\\,$eV, which we will refer to as the ultra-high energy (UHE) regime, protons propagating in the Galaxy retain most of their initial direction. Provided that the extragalactic magnetic field (EGMF) is negligible, UHE protons will therefore allow us to probe into the nature and properties of their cosmic sources. However, due to quite steep CR power spectrum, UHECRs are extremely rare (a few particles km$^{-2}$ century$^{-1}$) and their detection calls for the prolonged use of instruments with huge collecting areas. One further constraint arises from an effect first pointed out by Greisen, Zatsepin and Kuzmin~\\cite{Greisen:1966jv,Zatsepin:1966jv} and since then known as GZK effect: at energies $E\\gtrsim 5\\times 10^{19}$ eV the opacity of the interstellar space to protons drastically increases due to the photo-meson interaction process $p+\\gamma_{\\rm CMB}\\to\\pi^{0(+)}+p(n)$ which takes place on cosmic microwave background (CMB) photons. In other words, unless the sources are located within a sphere with radius of ${\\cal O}$(100) Mpc, the proton flux at $E\\gtrsim 5\\times 10^{19}$ eV should be greatly suppressed. However, due to the very limited statistics available in the UHE regime, the experimental detection of the GZK effect has not yet been firmly established.   A detailed knowledge of the UHECR framework is still missing. Only recently, for instance, magnetic fields were included in simulations of large scale structures (LSS)~\\cite{Dolag:2003ra,ems}. Qualitatively the simulations agree in finding that EGMFs are mainly localized in galaxy clusters and filaments, while voids should contain only primordial fields. However, the conclusions of Refs.~\\cite{Dolag:2003ra} and~\\cite{ems} are quantitatively rather different and it is at present unclear whether deflections in extragalactic magnetic fields will prevent astronomy even with UHE protons or not. Another large source of uncertainty is our ignorance on the chemical composition of UHECRs, mainly due to the need to extrapolate for decades in energy the models of hadronic interactions, though important progress are expected from new high quality data and deconvolution techniques  (see e.g. \\cite{Antoni:2005wq}). Future accelerator measurements of hadronic cross sections in higher energy ranges will also ameliorate the situation, but this will take several years at least.  From now on, therefore, we shall work under the assumptions that UHE astronomy is possible, namely: i) proton primaries; ii) EGMF negligibly small; iii) extragalactic astrophysical sources are responsible for UHECRs acceleration. A possibility favoring these hypotheses is that relatively few, powerful nearby sources are responsible for the UHECRs, and the small scale clustering observed by AGASA~\\cite{Takeda99} may be a hint in this direction. However, the above quoted clustering has not yet been confirmed by other experiments with comparable or larger statistics~\\cite{Abbasi:2004vu,Revenu:2005}, and probably a final answer will come when the Pierre Auger Observatory~\\cite{Auger} will have collected enough data. Independently on the observation of small-scale clustering, one could still look for large scale anisotropies in the data, eventually correlating with some known configuration of astrophysical source candidates. In this context, the most natural scenario to be tested is that UHECRs correlate with the luminous matter in the ``local'' universe. This is particularly expected for candidates like gamma ray bursts (hosted more likely in star formation regions) or colliding galaxies, but it's also a sufficiently generic hypothesis to deserve an interest of its own.  To this aim we use the  IRAS PSCz~\\cite{saunders00a} astronomical catalogue as tracer of Large Scale Structures from which, once that the related selection effects have been taken into account, we derive the related pattern of anisotropies; we then assess the minimum statistics needed to detect the model against the isotropic null hypothesis, in particular providing a forecast for the Auger experiment. Previous attempts to address a similar issue can be found in~\\cite{Waxman:1996hp,Smialkowski:2002by,Singh:2003xr}. Further details of the analysis summarized here can be found in~\\cite{Cuoco:2005yd}. ",
        "conclusions": "UHECRs offer the fascinating possibility to make astronomy with high energy charged particles; in particular, besides small scale clustering, a large scale anisotropy is expected correlating with the local cosmological LSS. However, we showed with a $\\chi^2$ approach that a huge statistics of several hundreds data is required to start testing the model at Auger South.   Untill now, the lack of UHECR statistics has seriously limited the usefulness of such a kind of analysis. However, progresses are expected in forthcoming years. The Southern site of Auger is already taking data and, once completed, the total area covered will be of 3000 km$^2$, thus improving by one order of magnitude present statistics in a couple of years~\\cite{Bertou:2005bx}. The idea to build a Northern Auger site strongly depends on the possibility to perform UHECR astronomy, for which full sky coverage is of primary importance. To this aim, the Japanese-American Telescope Array in the desert of Utah is expected to become operational by 2007~\\cite{Arai:2003dr} offering almost an order of magnitude larger aperture per year than AGASA in the Northern sky."
    },
    "0609/astro-ph0609094_arXiv.txt": {
        "abstract": "Parametric Embedding (PE) has recently been proposed as a general-purpose algorithm for class visualisation. It takes class posteriors produced by a mixture-based clustering algorithm and projects them in 2D for visualisation. However, although this fully modularised combination of objectives (clustering and projection) is attractive for its conceptual simplicity, in the case of high dimensional data, we show that a more optimal combination of these objectives can be achieved by integrating them both into a consistent probabilistic model. In this way, the projection step will fulfil a role of regularisation, guarding against the curse of dimensionality. As a result, the tradeoff between clustering and visualisation turns out to enhance the predictive abilities of the overall model. We present results on both synthetic data and two real-world high-dimensional data sets: observed spectra of early-type galaxies and gene expression arrays. ",
        "introduction": "Clustering and visualisation are two widespread unsupervised data analysis techniques, with applications in numerous fields of science and engineering. Two key strategies can be distinguished in the literature: (1) One is to produce a compression of the data first and then use that to visually detect distinct clusters. A wide range of linear and nonlinear dimensionality reduction techniques developed in machine learning follow this route, including PCA, GTM \\cite{GTM}, etc. (2) The alternative strategy is to cluster the data first and visualise the resulting cluster assignments afterwards. This is more popular in the data mining community \\cite{DM}. A recently proposed method, termed Parametric Embedding (PE) \\cite{PE} proposes to take class posteriors produced by a mixture-based clustering algorithm and project them in 2D for visualisation.  Let us observe however, that for both of these data exploration strategies the two objectives -- clustering and visualisation -- are decoupled and essentially one is entirely subordinated to the other. This is worrying in that inevitably, the errors accumulated in the first stage cannot be corrected in a subsequent stage and may essentially compromise the process of understanding the data.  In this paper we consider the class visualisation problem, as in \\cite{PE} and we identify a setting where a more fruitful coupling between clustering and visualisation can be achieved by integrating them both into a consistent probabilistic model. As we shall see, this is particularly useful in high-dimensional problems, where the number of data dimensions exceeds the number of observations. Such cases are encountered in modern scientific data analysis, e.g. in gene expression analysis, or the analysis of special objects in astronomy.  Our approach is based on a multi-objective formulation in a probabilistic formalism. As we shall see, our model can be decoupled and understood as a sum of two objectives: One term is responsible for clustering and one other for a PE-like projection of the estimated class assignments. These two steps are now interdependent, so that in high dimensional problems the projection step fulfils a regularisation role, guarding against the curse of dimensionality problem.  We use both synthetic data and two real-world high-dimensional data sets in our experiments: Observed spectra (of rare quality and coverage) of early-type galaxies and a benchmark gene expression array data set are used to demonstrate the working of the proposed approach. We find that in both cases we obtain not only a visualisation of the mixture posteriors, but a more predictive mixture model, as measured by out of sample test likelihood, as a result of appropriately combining the objectives of clustering and class projection.  The remainder of the paper is organised as follows: We begin with presenting our probabilistic model in Section 2. The interpretation by which this can be understood as a joint model for mixture based clustering and PE-like projection will become evident in Section 3, where the EM \\cite{Bishop} methodology is used to derive a maximum a posteriori (MAP) estimation algorithm for our model. We then extend this work to take into account additional available knowledge on measurement uncertainties for real-world data analysis. Section 4 presents the application two very different experiments involving galaxy spectra and gene expression. The results are summarised in the concluding section. ",
        "conclusions": "We proposed and investigated a model for class visualisation of explicitly high dimensional data. We have shown this model relates closely to PE \\cite{PE} in that it represents a probabilistic integration of the clustering and visualisation objectives into a single model. We derived empirical Bayesian estimates for our model which make this multi-objective interpretation easy to follow. Although this work may potentially further be enhanced by a fuller Bayesian estimation scheme, the empirical Bayesian methodology has been appropriate for our purposes \\cite{EB_astro} and it allows us to estimate an empirical latent density from a given example set of data and to reason about previously unseen data relative to that. We demonstrated gains in terms of the predictive capabilities of the proposed model over the fully modular and sequential approach to clustering and class visualisation in the case of high dimensional data.  \\subsubsection*{Acknowledgements} This research is funded by PPARC grant PP/C503138/1, `Designer Algorithms for Astronomical Data Mining'. AK also acknowledges partial support from a Wellcome Trust VIP Award (Project 10835).  \\footnotesize"
    },
    "0609/astro-ph0609788_arXiv.txt": {
        "abstract": "We present a detailed study of s-process nucleosynthesis in massive stars of solar-like initial composition and masses 15, 20, 25, and 30 $M_{\\sun}$. We update our previous results of s-process nucleosynthesis during the core He-burning of these stars and then focus on an analysis of the s-process under the physical conditions encountered during the shell-carbon burning. We show that the recent compilation of the $^{22}{\\rm Ne}(\\alpha,\\rm n)^{25}Mg$ rate leads to a remarkable reduction of the efficiency of the s-process during core He-burning. In particular, this rate leads to the lowest overproduction factor of $^{80}$Kr found to date during core He-burning in massive stars. The s-process yields resulting from shell carbon burning turn out to be very sensitive to the structural evolution of the carbon shell. This structure is influenced by the mass fraction of $^{12}$C attained at the end of core helium burning, which in turn is mainly determined by the $^{12}{\\rm C}(\\alpha,\\gamma)^{16}{\\rm O}$ reaction. The still present uncertainty in the rate for this reaction implies that the s-process in massive stars is also subject to this uncertainty. We identify some isotopes like $^{70}$Zn and $^{87}$Rb as the signatures of the s-process during shell carbon burning in massive stars. In determining the relative contribution of our s-only stellar yields to the solar abundances, we find it is important to take into account the neutron exposure of shell carbon burning. When we analyze our yields with a Salpeter Initial Mass Function, we find that massive stars contribute at least 40\\% to s-only nuclei with mass A $\\leq$ 87. For s-only nuclei with mass A $>$ 90, massive stars contribute on average $\\sim$7\\%, except for $^{152}$Gd, $^{187}$Os, and $^{198}$Hg which are $\\sim$14\\%, $\\sim$13\\%, and $\\sim$11\\%, respectively. ",
        "introduction": "\\label{sec:intro}  The s-process nucleosynthesis is the slow neutron-capture process of heavy nuclei in which the neutron-capture rate is slow relative to the beta-decay rate of the heavy nuclei near the line of beta stability \\citep{1957RvMP...29..547B,1957PASP...69..201C}. In this scenario, the seeds to synthesize the heavy nuclei are the iron-group nuclei. It is well known that massive stars of masses above M $\\simeq$ 12$M_{\\sun}$ are the sites of the so-called ``weak component'' of s-process nucleosynthesis. The nuclei produced in this site are restricted to the atomic mass range A $\\simeq$ 65-90 \\citep{1989RPPh...52..945K,1999ARA&A..37..239B}. Many papers have explored this weak component of the s-process, but mainly during the core He-burning phase \\citep{1987ApJ...315..209P, 1989A&A...210..187L, 1990A&A...234..211P, 1991ApJ...371..665R, 1992A&A...258..357B, 2000A&A...354..740R, 2000ApJ...533..998T, 2004ApJ...611..396E}. Few papers have investigated the s-process during the late evolution phases of massive stars, especially during the phase of shell carbon-burning (for short: shell C-burning) in these stars \\citep{1989A&A...210..187L, 1991ApJ...367..228R, 1993ApJ...419..207R}. The study of the s-process during shell C-burning requires special effort, because this burning episode lasts until the end of a massive star's evolution, as indicated by many works dealing with the advanced burning phases of massive stars \\citep{1988PhR...163...13N, 1998ApJ...502..737C, 2000ApJS..129..625L, 2002RvMP...74.1015W, 2004ApJ...611..396E}.  The contribution of the carbon shell to the s-process is important because the nuclei synthesized in this site will be eventually ejected largely unmodified during a supernova explosion of the star due to the lack of a neutron source during the explosive burning. \\citet{1993ApJ...419..207R} argued on the basis of the calculations by \\citet{1988PhR...163...13N} that only stars of mass about 25 $M_{\\sun}$ are able to eject the s-nuclei synthesized in the C-shell. We further quantify this issue in \\S \\ref{sec:expdist}. However, our main goal is to present a detailed study of the s-process during core He-burning and shell C-burning, thereby emphasizing the physical uncertainties influencing the results. We will benefit from our detailed calculations (\\citealp[hereafter EMT04]{2004ApJ...611..396E}), where we have investigated the influence of several important physical quantities on the characteristics of the stellar models during the advanced burning phases of massive stars.  In \\S \\ref{sec:smodels}, we summarize the main features of our previous stellar evolution calculations, which we have carried through the end of central oxygen burning. In \\S \\ref{sec:she}, we present our updated results of the s-process during core He-burning. In \\S \\ref{sec:spcshell}, we present the characteristics of shell C-burning and discuss the results obtained for the s-process during shell C-burning. In \\S \\ref{sec:expdist} we show the location (as function of the interior radius) where nuclei are exposed to neutrons. In this section, we also show the effectiveness of neutron captures in producing heavy elements by calculating the total neutron capture in the stellar models. In \\S \\ref{sec:SolarComparison} we compare the production factor distribution of s-only nuclei from our stellar models with the distribution from the solar abundance. In \\S \\ref{sec:summary} and in \\ref{sec:conclusion} we discuss and summarize the main conclusions of this paper. ",
        "conclusions": "\\label{sec:conlude} Our detailed study of the s-process nucleosynthesis resulting from core He-burning and shell C-burning in massive stars on the basis of the updated nuclear data of some relevant reactions reveal many interesting points which we summarize in the following. \\begin{itemize} \\item The efficiency of s-process nucleosynthesis during core He-burning does not depend on the rate of $^{12}{\\rm C}(\\alpha,\\gamma)^{16}{\\rm O}$ but it is sensitive to the rates of $^{22}{\\rm Ne}(\\alpha,{\\rm n})^{25}{\\rm Mg}$ and $^{16}{\\rm O}({\\rm n},\\gamma)^{17}{\\rm O}$. When we use the updated rates of these two reactions, as described in \\S \\ref{sec:smodels}, we find a significantly reduced efficiency of the s-process during core He-burning (see Table \\ref{tab:hesprocess}).  \\item The s-process in shell C-burning is more complicated and depends on the evolution of the massive star beyond core He-burning. This complexity can be seen from the locations, the number of carbon convective shells, and the thickness of carbon convective shells in our models. These in turn are sensitive to the central carbon mass fraction X$_{12}$ achieved at the end of core He-burning as a result of the reaction $^{12}{\\rm C}(\\alpha,\\gamma)^{16}{\\rm O}$, whose rate is not yet adequately determined. If this rate leads to a relatively lower X$_{12}$, then the s-process occurs later in time, possibly even after central neon burning. Consequently, the neutrons density achieved is high enough (see Fig. \\ref{fig:ndin25n}) to drive the nuclear reaction flow to the Zirconium region. This does not happen when X$_{12}$ is higher, as in our case 25K, since the s-process occurs here earlier in time (before central neon ignition) and at lower temperatures and densities, which result in a smaller neutron density. This explains why the nuclear reaction flow stops essentially in the Strontium region.  \\item Our calculations show that the overabundance of $^{70}$Zn can be used as indicator of the strength of the nuclear reaction flow through the branchings along the s-process path, especially at $^{69}$Zn. We have also found that $^{87}$Rb is strongly produced during shell carbon burning due to the higher rate of neutron-capture of $^{86}$Rb relative to its rate during core helium burning \\citep{1993ApJ...419..207R}.  \\item We measure the s-processing in the core helium, shell helium, and shell carbon burning in massive stars with $N_c = \\int n_c(M_r) \\; dM_r$ and show their relative strengths or importance. We show the s-process contribution from shell carbon burning decreases with increasing mass of the star.  \\item In comparing the yields of s-only nuclei of our stellar models with the solar distribution, we find that it is necessary to include the results of s-processing from shell carbon burning and to mix the yields of all mass range of massive stars to give a reasonable fit to the solar distribution. For s-only nuclei with mass number A$\\leq$87, massive stars contribute at least 40\\% to the solar s-only nuclei. For s-only nuclei with mass number A$>$90, massive stars contribute on average $\\sim$7\\%, except for $^{152}$Gd, $^{187}$Os, and $^{198}$Hg which can be 14\\%, 13\\%, and 11\\%, respectively.  \\end{itemize} \\label{sec:conclusion}"
    },
    "0609/astro-ph0609755_arXiv.txt": {
        "abstract": "Here we consider a scenario in which dark energy is associated with the apparent area of a surface in the early universe. In order to resemble the cosmological constant at late times, this hypothetical reference scale should maintain an approximately constant physical size during an asymptotically de-Sitter expansion. This is found to arise when the particle horizon - anticipated to be significantly greater than the Hubble length - is approaching the antipode of a closed universe. Depending on the constant of proportionality, either the ensuing inflationary period prevents the particle horizon from vanishing, or it may lead to a sequence of ``Big Rips\". ",
        "introduction": "The acceleration of our universe is now a well established phenomenon, and numerous theoretical approaches have attempted to provide an explanation. The cosmological constant is arguably the strongest candidate, although it has a number of well-documented issues, particularly with regard to its size. One of the more attractive methods of attaining such a small number is by relating it to a cosmological length scale.  The  holographic principle speculates that all information stored within some  volume  is  represented on the boundary of that region.  A cosmological variation of the holographic principle was originally proposed by Fischler \\& Susskind \\cite{fischler-1998-}. In this scenario the cosmological horizon acts as a surface which obeys the Bekenstein bound, limiting the entropy of the enclosed volume, such that \\begin{equation} \\label{eq:bek} S \\leq \\frac{A}{4G} . \\end{equation} This concept has been extended by Bousso \\cite{bousso-1999-9907} into a covariant and more general conjecture.  Cohen et al. \\cite{cohen-1999-82} took a slightly different approach, whereby the dark energy density is proportional to the inverse square of some cosmological length. This arises when constraining the energy in some volume to be less than a black hole of the same size.  \\begin{equation} \\label{eq:cohen} L^3 \\rho_\\Lambda \\lesssim L M_P^2 \\end{equation}  \\noindent The most natural choices for $L$ are the particle horizon, and the Hubble length. However Hsu \\cite{hsu-2004-594} has shown that under adiabatic conditions, neither exhibit the equation of state necessary to mimic dark energy. One could enforce $w=-1$, by permitting energy exchange between the dark matter and dark energy, however this approach has not proved particularly successful \\cite{wang-2005-22}.  Li \\cite{li-2004-603} produced a model of holographic dark energy based on Cohen's approach, which has subsequently been explored in further detail \\cite{setare-2006-0609,2005CQGra..22.4895G,2004JCAP...08..013H,2006GReGr..38.1285N} and is compatible with current observational constraints \\cite{zhang-2005-72}. By identifying $L$ as the future event horizon, the required equation of state could emerge without invoking energy exchange. For a saturated inequality, the energy density is given by  \\begin{equation} \\label{eq:holo} \\rho_\\Lambda = 3 d^2 M_p^2 L^{-2} . \\end{equation}  \\noindent The value $d=1$ is usually adopted for the reasons outlined in \\cite{li-2004-603} - in particular, $d \\geq 1$ ensures the entropy doesn't decrease. However, within Li's model there remain two areas of concern. First of all, for self-consistency, there is the requirement that all forms of energy must decay at some time in the future. And secondly, the initial condition of a preset future event horizon appears a little unusual, raising the issue of causality.  Here we take a rather different approach, finding that a more natural form of dark energy may arise when considering the particle horizon. We argue that its associated area may possess greater physical significance than the length scale itself. In Section \\ref{sec:phor} we explore the consequences of attributing an energy density to the area of the particle horizon within a closed universe. Section \\ref{sec:eqstat} establishes the equation of state, and explores different values of the constant of proportionality $d$. In some cases this may lead to a ``Big Rip\" scenario. Section \\ref{sec:obcons} summarises the current and future observational constraints on the model. We speculate on the possible physical interpretations in Section \\ref{sec:physmot}. ",
        "conclusions": "\\label{sec:discholo}  An energy density associated with the size of the particle horizon within a closed universe has been shown to demonstrate a negative equation of state. Precise behaviour depends on the constant of proportionality, $d$. Whilst $d=1$ is the most natural choice, imitating the cosmological constant at late times, current observational constraints slightly prefer $d<1$. This corresponds to a cosmology with a sequence of ``rips\", as the particle horizon repeatedly traverses the closed universe. Such a model has the advantage of addressing the coincidence problem, and has a simple mechanism for lowering the equation of state at late times. We have also established a robust lower bound of $d>0.5$, which must be satisfied in order to ensure $w<1$ in the early universe.  For the case $d=1$ we find cosmic evolution equivalent to the model put forward by Li, which was based on associating the dark energy density with the future event horizon. However by using the particle horizon, this approach is causal, does not require the decay of matter, and has a natural reference scale - the radius of curvature.  As with other holographic models, this approach is rather speculative, but does feature a number of attractive properties. Given the close relationship predicted between $\\Omega_\\Lambda$ and $w$, data from future dark energy surveys such as SNAP should be capable of distinguishing this model from the cosmological constant."
    },
    "0609/physics0609132_arXiv.txt": {
        "abstract": "We built and characterized an optical system that emulates the optical characteristics of an 8m-class telescope like the Very Large Telescope. The system contains rotating glass phase-screens to generate realistic atmosphere-like optical turbulence, as needed for testing multi-conjugate adaptive optics systems. In this paper we present an investigation of the statistical properties of two phase-screens etched on glass-plate surfaces, obtained from Silios Technologies. Those etched screens are highly transmissive (above 85\\%) from 0.45 to 2.5$\\mu$m.  From direct imaging, their Fried parameter r$_0$ values  (0.43$\\pm$0.04\\,mm and 0.81$\\pm$0.03\\,mm, respectively, at 0.633$\\mu$m) agree with the expectation to within 10\\%. This is also confirmed by a comparison of measured and expected Zernike coefficient variances. Overall, we find that those screens are quite reproducible, allowing sub-millimetre r$_0$ values, which were difficult to achieve in the past. We conclude that the telescope emulator and phase-screens form a powerful atmospheric turbulence generator allowing systematic testing of different kinds of AO instrumentation. ",
        "introduction": "The optical image quality and in particular the angular resolution of ground based telescopes is hampered by the refractive index variations of the Earth's atmosphere. These refractive index variations show a temporal and spatial behavior that under certain assumptions is well described by Kolmogorov turbulence \\cite{Kolmogorov41, Tatarskii93, Conan95}.  Adaptive optics (AO) instrumentation, available on almost all 8--10m-class telescopes, can compensate for optical turbulence (seeing) with  certain restrictions. One of these restrictions is the isoplanatic angle, which is a result of using a single reference star for the AO and therefore compensating the integral effect of atmospheric turbulence in one direction only. Multi-conjugate AO systems that compensate the dominant turbulent layers instead, can increase the isoplanatic angle from a few tens of arcseconds \\cite{Weiss02} to up to a few arcminutes at 2.2\\,$\\mu$m \\cite{Berkefeld01}.  With the advent of second generation AO systems like multi-conjugate AO \\cite{Marchetti03}, ground layer AO \\cite{Tokovinin04},  extreme AO \\cite{Koehler04}, multiple field of view AO \\cite{Ragazzoni02}, and multiple object AO \\cite{Neichel05}, it is  becoming more and more important to have a realistic 3D turbulence simulation tool available in the laboratory. Such a tool provides repeatable optical turbulence and therefore well defined atmospheric conditions. It can support the assembly, integration, and verification phase of novel AO instrumentation such that AO performance can be verified before the instrument is attached to the telescope.  In our previous papers \\cite{Butler03,Butler04} we described the design, manufacturing process, and laboratory characterization of ion-exchange phase-screens. While such phase-screen technology allows repeatable optical turbulence generation over a wide spectral range, 0.5--2.5$\\mu$m, it does not allow sub-millimetre r$_0$ values. In this paper we focus on a new type of phase-screen, based on transmissive glass substrates with aberrations imprinted on one surface. The motivation for investigating this new technique is this: For laboratory characterization of astronomical adaptive optics systems designed for D=8--40-m-class telescopes, D/r$_0$ values greater 30 are required. As we can only use scaled-down telescope optics in a laboratory, with typical values of D$_{lab}$=10\\,mm, we need r$_0$ values smaller than D$_{lab}$/30=1/3\\,mm. Glass phase screens with such small r$_0$ values were in the past difficult to fabricate. We identified the surface etching technique as the only tool and therefore best solution so far to this problem.  In the following, we firstly describe in Sec.~\\ref{maps} the design of MAPS, an optical set-up used to test the screens. The design and manufacture of the screens is then outlined in Sec.~\\ref{fabrication}. Our analyses and results are presented in Sec.~\\ref{charac}, and we briefly summarize our findings in Sec.~\\ref{conclus}. ",
        "conclusions": "We have procured and characterized two phase-screens. We find that, within uncertainties, the statistical properties of both screens match those of the computer-generated templates. This suggests that the surface-etching technique allows reproducible phase-screens. In comparison to ion-exchange phase-screens  \\cite{Butler03} we conclude that phase-screens with small r$_0$ values (strongest optical turbulence) are more reliably produced with the etching technology. We want to note that Shack-Hartmann wavefront measurements of phase-screens with a discrete pixel-like fine structure require a minimum Shack-Hartmann lenslet diameter; this diameter has to be at least twice as large as the pixel size on the phase-screen."
    },
    "0609/astro-ph0609080_arXiv.txt": {
        "abstract": "{To quantify the relationships between star formation in cluster galaxies and global cluster properties. Using a subsample of 79 nearby clusters from the RASS-SDSS galaxy cluster catalogue of Popesso et al. (2005a), we perform a regression analysis between the cluster integrated star formation rate ($\\Sigma SFR$) the cluster total stellar mass ($M_{\\star}$), the fractions of star forming ($f_{SF}$) and blue ($f_b$) galaxies and other cluster global properties, namely its richness ($N_{gal}$, i.e. the total number of cluster members within the cluster virial radius, corrected for incompleteness), velocity dispersion ($\\sigma_v$), virial mass ($M_{200}$), and X-ray luminosity ($L_X$).  All cluster global quantities are corrected for projection effects before the analysis. Galaxy $SFR$s and stellar masses are taken from the catalog of Brinchmann et al. (2004), which is based on SDSS spectra. We only consider galaxies with $M_r \\leq -20.25$ in our analysis, and exclude AGNs. We find that both $\\Sigma SFR$ and $M_{\\star}$ are correlated with all the cluster global quantities. A partial correlation analysis show that all the correlations are induced by the fundamental one between $\\Sigma SFR$ and $N_{gal}$, hence there is no evidence that the cluster properties affect the mean $SFR$ or $M_{\\star}$ per galaxy. The relations between $\\Sigma SFR$ and $M_{\\star}$, on one side, and both $N_{gal}$ and $M_{200}$, on the other side, are linear, i.e. we see no evidence that different clusters have different $SFR$ or different $M_{\\star}$ per galaxy and per unit mass. The fraction $f_{SF}$ does not depend on any cluster property considered, while $f_b$ does depend on $L_X$. We note that a significant fraction of star-forming cluster galaxies are red ($\\sim$ 25\\% of the whole cluster galaxy population). We conclude that the global cluster properties are unable to affect the $SF$ properties of cluster galaxies, but the presence of the X-ray luminous intra-cluster medium can affect their colors, perhaps through the ram-pressure stripping mechanism.} ",
        "introduction": "What role does the environment play in the evolution of cluster galaxies? The dependence of the morphological mix from the environmental conditions was qualitatively illustrated in the early study of the Virgo cluster by Hubble and Humason (1931) and has been confirmed in many studies (e.g. Oemler 1974; Dressler 1980; Postman \\& Geller 1984; Dressler et al. 1997). The clear observational evidence is that the high density regions, such as the massive galaxy clusters, are dominated by a quiescent early type galaxy population, while the late type star forming galaxies more likely populate low density regions such as the field. A recently proposed way to study the relation between galaxy population and environmental conditions is the analysis of the ongoing star formation (SF) in galaxies of different environments (see, e.g., Christlein \\& Zabludoff 2005).  The SF rate (SFR) is an important measure of the evolutionary state of a galaxy, and a sensitive indicator of the environmental interactions. Previous studies of cluster galaxy SFRs have sometimes reached conflicting conclusions. The SFRs of cluster galaxies have been found to be reduced (Kennicutt 1983; Bicay and Giovanelli 1987; Kodaira et al. 1990; Moss \\& Whittle 1993; Abraham et al. 1996; Balogh et l. 1998, 2002; Koopmann \\& Kenney 1998; Hashimoto et al. 1998; Gavazzi et al. 2002; Pimbblet et al. 2006), comparable (Kennicutt et al. 1984; Donas et al. 1990; Gavazzi et al. 1991, 1998; Biviano et al. 1997; Moss \\& Whittle 2005), or in some case enhanced (Moss \\& Whittle 1993; Bennet \\& Moss 1998) relative to the SFRs of field galaxies of the same classes.  Several cluster-related environmental processes can affect the SFRs of galaxies.  Some processes mainly affect the gaseous content of a galaxy, such as the ram-pressure stripping (Gunn \\& Gott 1972; Kenney et al. 2004; van Gorkom 2004), re-accretion of the stripped gas (Vollmer et al. 2001), turbulence and viscosity (e.q. Quilis et al. 2001), and starvation/strangulation (Larson et al. 1980). Gravitational processes, which affect both the gaseous and the stellar properties of a galaxy, range from low-velocity tidal interactions and mergers (e.g. Mamon 1996; Barnes \\& Hernquist 1996; Conselice 2006), to high-velocity interactions between galaxies and/or clusters (Moore et al. 1998, 1999; Struck 1999; Mihos 2004). Despite a number of recent studies of nearby and distant clusters, it is not yet clear which of these processes, if any, are dominant.  Clues on the relative importance of the cluster-related environmental processes can be obtained by investigating the evolution of the star-forming properties of cluster galaxies. In this context, the most important evolutionary phenomenon is the Butcher-Oemler (BO hereafter) effect (Butcher \\& Oemler 1978, 1984), i.e. the increasing fraction of blue cluster members with redshift. The BO effect has been confirmed and detailed by many studies since the original works of Butcher \\& Oemler (e.g.  Ellingson et al. 2001; Margoniner et al. 2001; Alexov et al. 2003; De Propris et al. 2003; Rakos and Shombert 2005), although Andreon et al. (2004, 2006) have argued that no cluster-dependent evolution is required to explain the BO effect, which is entirely compatible with the normal color evolution of galaxies in an ageing universe. The BO-effect is purely photometrical. The spectroscopic version of the BO-effect is an excess of emission-line and star-forming galaxies in distant, relative to nearby, clusters, first suggested by Dressler \\& Gunn (1982) and later confirmed by several authors (e.g. Postman et al. 1998, 2001; Dressler et al. 1999; Finn et al. 2004, 2005; Homeier et al. 2005; Poggianti et al. 2006, P06 hereafter)  Most of the analyses so far have concentrated on the comparison of the star-forming properties of individual cluster galaxies with those of field galaxies, and on the variation of the galaxy SFRs on the local density of their environment. However, it is also important to assess the dependence (if any) of the star-forming properties of cluster galaxies on their cluster global properties, such as the mass, velocity dispersion and X-ray luminosity. Should the SFRs of cluster galaxies depend on global properties of their host cluster, results obtained for different individual clusters would not be straightforward to compare, thereby producing apparently discrepant results. Moreover, the relative efficiency of the different evolutionary processes depends on several cluster properties, and investigating the SFRs of cluster galaxies as a function of these properties can help understanding this issue (see, e.g., Pimbblet 2003). Also the evolution of the star forming properties of cluster galaxies must be studied in close connection with the evolution of their host cluster properties. In fact, evolutionary studies of cluster galaxy SFRs may be affected by selection biases if the SFRs depend on global cluster properties, such as their masses. Since in flux-limited surveys more massive clusters are preferentially selected with increasing redshift, a biased estimate of the evolution of the star-forming properties of cluster galaxies may result (see, e.g., Newberry et al. 1998; Andreon \\& Ettori 1999),  Recently, several studies have addressed the dependence of the star-forming properties of cluster galaxies on their host global properties. Several studies have found that the cluster global properties do not affect the star-forming properties of cluster galaxies. In particular, no dependence has been found of either the blue or the late-type galaxy fraction in clusters on cluster velocity dispersions ($\\sigma_v$s) and masses (Goto 2005), nor of the blue fraction with cluster richness, concentration, and degree of subclustering (De Propris et al. 2004). On the other hand, both Margoniner et al. (2001) and Goto et al. (2003) had previously found a dependence of the blue or late-type galaxy fractions on the cluster richness. Goto (2005) has also claimed no dependence on the cluster $\\sigma_v$s and masses of either the total cluster SFR or of the total cluster SFR normalized by the cluster mass, in disagreement with Finn et al. (2005) who have shown that the integrated SFR per cluster mass decreases with increasing cluster mass. Lea \\& Henry (1988), Fairley et al. (2002), and Wake et al. (2005) have all failed to find any dependence of the fraction of blue cluster galaxies with the cluster X-ray luminosity, $L_X$.  Similarly, Balogh et al. (2002) have compared the galaxy SFRs in high-$L_X$ and low-$L_X$ clusters and have found no differences. In the sample of Homeier et al. (2005) there are hints of correlations between the total cluster SFRs and cluster $L_X$s and intra-cluster gas temperatures, $T_X$s, but the trends are not really significant.  Most recently, P06 have found that the fraction of emission-line galaxies (ELGs hereafter) decreases with increasing cluster $\\sigma_v$. The trend is continuous at high-$z$, but is characterized by a break at $\\sigma_v \\sim 500$--600 km~s$^{-1}$ in nearby clusters, where the relation they find is consistent with the results obtained by Biviano et al. (1997).  In this paper we re-address the issue of the dependence of the SFR and the fraction of star forming galaxies on the cluster global properties. At variance with most previous studies, we consider both optical and X-ray cluster global properties, namely the mass, $\\sigma_v$, and $L_X$. While these quantities are correlated (Popesso et al. 2005a, Paper III of this series), it is worthwhile to consider them all, since the star-forming properties of cluster galaxies may show a stronger dependence on one of these properties, thereby pointing to a different physical mechanism affecting their SFRs. E.g., Postman et al. (2005) have recently shown that the fraction of early-type galaxies in distant clusters does depend on $L_X$, but not on $\\sigma_v$, nor on $T_X$.  In our analysis we use a sample of 79 low-redshift clusters taken from the X-ray selected RASS-SDSS galaxy cluster catalog (Popesso et al. 2004, Paper I) and the optically selected Abell cluster sample (Popesso et al. 2006a, Paper V). Besides providing further constraints on the mechanisms of galaxy evolution in clusters, our investigation should be useful for assessing the possible selection effects in the comparison of the star-forming properties of galaxies in nearby vs. distant clusters, as well as in clusters at similar redshifts but with different global properties.  In Sect.~\\ref{s-data} of the paper we describe our dataset.  In Sect.~\\ref{s-sfr} we analyze the relation between the cluster integrated star formation rate and the global properties of the systems. In Sect.~\\ref{s-frac} we apply the same analysis to the fraction of blue cluster galaxies and the fraction of cluster star forming galaxies. Sect.~\\ref{s-conc} contains our conclusions.  Throughout this paper, we use $H_0=70$ km s$^{-1}$ Mpc$^{-1}$ in a flat cosmology with $\\Omega_0=0.3$ and $\\Omega_{\\Lambda}=0.7$ (e.g. Tegmark et al. 2004). ",
        "conclusions": "\\label{s-conc} We have analyzed the relationships between SF in cluster galaxies and global cluster properties, such as cluster $M_{200}$, $\\sigma_v$, $L_X$, and $N_{gal}$.  For our analysis we have used a sample of 79 nearby clusters extracted from the RASS-SDSS galaxy cluster catalogue of Paper III and Paper V. Galaxy $SFR$s and stellar masses are taken from the catalog of Brinchmann et al. (2004), which is based on SDSS spectra. We only consider galaxies with $M_r \\leq -20.25$ in our analysis, and exclude AGNs and composite-spectra galaxies.  All the cluster quantities considered are corrected for incompleteness, when needed, and for projection effects. The de-projection correction is of particular importance in our analysis, since we find that it depends on the cluster mass.  $\\Sigma SFR$ is correlated with all the cluster global quantities mentioned above.  By performing a multiple regression analysis that the main correlation is that between $\\Sigma SFR$ and $N_{gal}$. Since this relation is linear the average SFR of cluster galaxies is the same in different clusters, and is unaffected by either the cluster mass, or its velocity dispersion, or its X-ray luminosity.  We come to essentially the same conclusion when $M_{\\star}$ is considered in lieu of $\\Sigma SFR$. If instead of normalizing $\\Sigma SFR$ with $N_{gal}$ we normalize it with $M_{200}$, we still find $\\Sigma SFR/M_{200}$ does not depend on any cluster global property, except $\\sigma_v$, which we suggest could be evidence of the effect of ram-pressure stripping on the cluster galaxy properties.  Ram-pressure could also be the mechanism able to explain the observed anti-correlation of $f_b$ with $L_X$, since $f_b$ is {\\em not} correlated with either $M_{200}$ or with $\\sigma_v$. On the other hand, the fact that we do not observe any correlation between $L_X$ and $f_{SF}$ is due to the presence of a dominant fraction of {\\em red} star-forming galaxies. They could also be the result of the ram-pressure mechanism, or, in alternative, they could be star-forming galaxies with an anomalous amount of dust.  If global cluster properties affect the star-forming properties of cluster galaxies, their effect is rather marginal, except perhaps on galaxy colors, which seem to be influenced by the presence of the IC diffuse gas.  \\vspace{2cm}  We thank the anonymous referee for useful suggestions which helped us improving the quality of this paper. Funding for the creation and distribution of the SDSS Archive has been provided by the Alfred P.  Sloan Foundation, the Participating Institutions, the National Aeronautics and Space Administration, the National Science Foundation, the U.S.  Department of Energy, the Japanese Monbukagakusho, and the Max Planck Society. The SDSS Web site is http://www.sdss.org/. The SDSS is managed by the Astrophysical Research Consortium (ARC) for the Participating Institutions.  The Participating Institutions are The University of Chicago, Fermilab, the Institute for Advanced Study, the Japan Participation Group, The Johns Hopkins University, Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, University of Pittsburgh, Princeton University, the United States Naval Observatory, and the University of Washington."
    },
    "0609/astro-ph0609563_arXiv.txt": {
        "abstract": "Assuming that there exists a species of heavy sterile neutrinos ($m_{\\nu}>1$ keV) and their decays can serve as a heating source of the hot gas in galaxy clusters, we study how the observational constraints on cooling flow limit the mass of these sterile neutrinos. We predict a relation among the luminosity, total mass and the redshift of a cluster, and we compare this relation with data from 12 clusters to obtain an estimate of the decay rate of the sterile neutrino. ",
        "introduction": "Recent observations indicate that there are not enough cooling flows in galaxy clusters to explain why hot gas in some clusters are so hot \\cite{coolingflow}, \\cite{coolingflow2}. On the other hand, neutrino oscillation experiments (solar, atmospheric and LSND experiments) indicate that 3 ordinary neutrinos cannot explain the three different scales of the mass squared differences between neutrino species $\\Delta m^2$'s, of about $10^{-5}$ eV$^2$, $10^{-2}$ eV$^2$ and 1 eV$^2$ \\cite{oscillation}, suggesting that one or more sterile neutrinos must be added. In this article, we study the possibility that a species of massive sterile neutrinos ($m_{\\nu}>1$ keV) exists and their decays provide a heating source to maintain the temperature of the hot gas. We use this model to predict a relation among the luminosity, total mass and redshift of a cluster, and we obtain an estimate of the decayrate of the sterile neutrinos by using available data from 12 clusters. This idea of the radiative decay of neutrinos as an energy source has been proposed to account for reionization \\cite{reion}.  In standard cosmology, neutrinos with $m_{\\nu}<1$ MeV are classified as hot dark matter (HDM) because they are relativistic at decoupling. As the universe expands, they cool down and become non-relativistic after a redshift of \\begin{equation} z_{rel} \\approx \\frac{m_{\\nu}c^2}{3kT_{\\nu,0}}=2 \\times 10^3 \\left( \\frac{m_{\\nu}}{1~\\rm{eV}} \\right), \\end{equation} where $T_{\\nu,0}$ is the neutrino temperature today. In the non-relativistic regime, the average neutrino speed at neutrino temperature $T_{\\nu}$ is \\cite{Chung} \\begin{equation} <v>=160~{\\rm{km/s}} \\left( \\frac{1~{\\rm{eV}}}{m_{\\nu}} \\right) \\left( \\frac{T_{\\nu}}{1.947~\\rm{K}} \\right). \\end{equation} They do not contribute to structure formation if their average speed is too high. For example, at matter-radiation equality (redshift $\\approx 10^4 h^2$, $h$ being the Hubble parameter), the neutrino speed is close to $c$ if $m_{\\nu} \\ll 10$ eV. By using perturbation theory on neutrino clustering, it was found that massive neutrinos can gravitationally cluster if $m_{\\nu}$ is greater than several eV's \\cite{Chung}. Therefore, neutrinos with mass $m_{\\nu} \\ge$10 eV can be classified as warm dark matter (WDM) rather than HDM. ",
        "conclusions": "It has been a puzzle why some clusters are still so hot even though their cooling times are shorter than the Hubble time. Even though cooling flow models \\cite{xray} can help to explain this phenomenon, significant cooling flow in clusters has not been observed. This is known as the `mass sink problem' \\cite{coolingflow}. Also, the cooling flow model predicts that the temperature of the hot gas should be quite inhomogeneous. However, the observed hot gas temperature profiles are quite homogeneous and gas with temperature less than 1 keV does not exist in the amount predicted. Donahue and Voit summarized the drawbacks of the cooling flow models and suggested that any successful model must explain the mass sink problem and the homogeneous temperature profile with positive core temperature gradients extending to $\\sim 10^2$ kpc \\cite{coolingflow}. Here, we present an alternative solution of the mass sink problem. We have discussed how the radiative decay of sterile neutrinos can heat up the hot gas in clusters. Our model predicts a redshift dependence of $L_x/M$, Eq.~(15), which is consistent with the observed data (see Fig.~1), and we obtain the decay rate of the sterile neutrinos in clusters, which is about $\\Gamma \\sim (4.6 \\pm 1.1) \\times 10^{-17}$ s$^{-1}$. Our picture does not require any cooling flow, and the radiative heating power by neutrino decays results in a quite homogeneous temperature profile. Therefore, our model can fit the observed information and solve the cooling flow problem for a range of $\\Gamma$.  If we can constrain the values of $m_1$ and $\\alpha_1$ \\cite{m1}, we can get a better picture and fix $\\Gamma$. The radius of a $10^{13}$ solar mass $m_{s1}$ neutrino star is of order 10 kpc which is the same scale as a galaxy. Therefore we propose that the galactic dark matter may be composed of $m_{s1}$ and $m_{s2}$ sterile neutrinos. By examining the rotation curve of a galaxy, we can obtain a feasible range of $m_{s1}$. If $m_1$ is below 1 eV, the size of an $m_1$ neutrino star is greater than 1 Mpc, which is the cluster scale. Therefore, the active neutrinos may affect the properties of hot gas and clusters \\cite{manho}. Different neutrinos may correspond to structures in different scales. It would be interesting to study whether our model can solve both the dark matter problem and the cooling flow problem."
    },
    "0609/astro-ph0609049_arXiv.txt": {
        "abstract": "We derive new equations using the {\\it mixed-frame} approach for one- and two-dimensional (axisymmetric) time-dependent radiation transport and the associated couplings with matter.  Our formulation is multi-group and multi-angle and includes anisotropic scattering, frequency(energy)-dependent scattering and absorption, complete velocity dependence to order $v/c$, rotation, and energy redistribution due to inelastic scattering. Hence, the ``2D\" realization is actually ``6 1/2\"-dimensional. The effects of radiation viscosity are automatically incorporated. Moreover, we develop Accelerated-Lambda-Iteration, Krylov subspace (GMRES), Discontinuous-Finite-Element, and Feautrier numerical methods for solving the equations and present the results of one-dimensional numerical tests of the new formalism. The virtues of the mixed-frame approach include simple velocity dependence with no velocity derivatives, straight characteristics, simple physical interpretation, and clear generalization to higher dimensions. Our treatment can be used for both photon and neutrino transport, but we focus on neutrino transport and applications to core-collapse supernova theory in the discussions and examples. ",
        "introduction": "Many phenomena in the Universe must be addressed using the tools of radiation transport and radiation hydrodynamics to achieve a theoretical understanding of their character.  Supernova explosions, gamma-ray bursts, star formation, planet formation, nova explosions, X-ray bursts, Luminous-Blue-Variable (LBV) outbursts, and stellar winds are all time-dependent fluid flow problems in which radiation plays a pivotal, often driving, role. In most circumstances, the radiation is photons, but in studies of the core-collapse supernova mechanism the radiation is neutrinos.  In either case, time-dependent techniques to address radiation transport and the coupling of radiation with matter are of central concern to the theorist whose goal is explaining the transient, dynamical phenomena of the Cosmos.  However, spherically-symmetric algorithms for radiation transport and radiation hydrodynamics, generally necessary to achieve a first-order understanding, are often not sufficient and multi-dimensional approaches are called for. These are not easy, not only to formalate, but to implement. Nevertheless, multi-D radiation transport is emerging as a necessary tool in the theorist's toolbox and computers to evolve the associated equations are becoming available to a wider cohort of researchers.  In this paper, we derive new equations using the {\\it mixed-frame} approach for one- and two-dimensional radiation transport and the associated coupling with matter, significantly extending the pioneering work of Mihalas \\& Klein (1982) to include anisotropic scattering, frequency(energy)-dependent scattering and absorption, complete velocity dependence to order $v/c$, rotation, and energy redistribution due to inelastic scattering.  Moreover, we develop algorithms for solving these equations and present the results of one-dimensional numerical tests.  Hence, we provide both the new formulation and appropriate numerical techniques to solve it. The mixed-frame approach, in which the radiation quantities are defined in the laboratory (Eulerian) frame and the matter and coupling quantities are defined in the comoving frame, has largely been neglected by the radiation transport and atmospheres communities because of their focus on line transfer.  The desire to include narrow spectral lines and to handle Doppler shifts into and out of those lines necessitates many spectral bins and a huge number of angular bins to ensure the lines are resolved on the computational grid.  As a result, most dynamic atmosphere and radiation studies are done using the comoving (Lagrangian) equations of radiation transport.  For instance, the core-collapse supernova community, which has been at the forefront of radiation hydrodynamic developments in astrophysics, has inherited this formulation for their treatment of neutrino transport.  However, many radiation-hydrodynamic problems do not require an exquisite treatment of spectral line transport, but a good treatment of continuum transport.  The core-collapse supernova problem is one such case.  The monochromatic opacities and emissivities of neutrinos are overwhelmingly smooth functions of neutrino energy. Given this, the mixed-frame formulation is ideally suited for supernova theory.  Its virtues vis \\`a vis the comoving-frame approach to the solution of the Boltzmann equation and its related moment equations are numerous: 1) Even in one-dimension, instead of requiring $\\sim$20 velocity-dependent terms (Buras et al. 2006) on the left-hand(streaming) side of the Boltzmann/transport equation, many of which involve spatial velocity derivatives, there are no such terms on the left-hand-side in the mixed frame approach and only {\\it one} grouped linear term (to $O(v/c)$) on the right-hand(source) side;  2) There are no terms with derivatives of the velocity. Therefore, the characteristics of the associated transport equation are all straight lines.  Furthermore, there is no need for the monotonicity in the velocity field required by some implicit solvers; and 3) The mixed-frame method is easily generalized to two and three dimensions, and the associated solvers are straightforward (though more expensive) extensions of those employed in 1D. Note that much has been made of the importance of velocity-dependent terms in the transport equations for the calculation of the neutrino energy deposited in the ``net-gain\" (Wilson 1985) region.  We show that the mixed-frame approach provides the most straightforward perspective from which to understand the physics of this effect and that even its sign is frame-dependent.  In particular, in the mixed-frame formulation the velocity-dependent term augments the net gain during the stalled-shock phase, while in the comoving frame formulation the corresponding terms reduce it.  Hence, any statement concerning the importance of such terms is very frame- and treatment-dependent.  Though the mixed-frame equations depend simply upon velocity, the Lorentz transformations that are the core of this formulation introduce frequency(energy) derivatives of the radiation moments.  These couple adjacent energy groups and might have compromised implicit algorithms that parallelize in group (Livne et al. 2004; Walder et al. 2005; Burrows et al. 2006; Dessart et al. 2006ab; Ott et al. 2006ab).  However, we show that such terms can be handled semi-implicitly, with the logarithmic derivatives of the moments with respect to energy handled explicitly during an otherwise implicit solve. In this way, processors performing updates on only a single group are not coupled during the iteration to adversely affect parallelization and scalability.  Furthermore, a similar semi-implicit tactic works for inelastic scattering terms, since these are sub-dominant in the context of core collapse (Thompson, Burrows, \\& Pinto 2003). Note, however, that parallelization in energy groups is viable for 2D calculations, but not for 3D. An entire 2D hydro grid can now reside on a single processor, but an entire 3D hydro grid requires spatial domain decomposition onto many processors and parallelization in {\\it only} energy groups is not yet viable. Furthermore, given the seven-dimensional nature of 3D radiation transport (3 space + 2 angles + 1 frequency/energy + 1 time), 3D is not yet computationally feasible for astrophysical simulations.  Therefore, we focus in this paper on 1D and 2D mixed-frame formulations.  However, a 2D, mixed-frame, azimuthally-symmetric formalism is still six-dimensional (lacking one spatial dimension and a hemisphere) and is particularly straightforward and powerful. Our approach includes rotation (and, hence, is actually ``6 1/2\"-dimensional). The Eddington factor becomes an ``Eddington tensor\" with five independent entries. The zeroth- and first-moment equations are accurate approximations to the full equations that are closed with Eddington factors that can be calculated at each timestep to provide the full solution, or every $N$ steps to provide an excellent, fast, though approximate, solution.  Since velocity dependence is included in the multi-D context, the effects of radiation viscosity (neutrino or photon) are automatically incorporated into any scheme that includes the hydrodynamics with the radiation force and energy couplings.  Though our treatment can be used for both photon and neutrino transport, it was devised with neutrinos in mind. As a result, much of the discussion and all the tests we perform are in that context. Therefore, a short review of the transport schemes employed to date in supernova modeling is in order.  In the realm of 1D (spherical) neutrino transport, Bowers \\& Wilson (1982), Bruenn (1985),  Wilson (1985), and Mayle, Wilson, \\& Schramm (1987) employed multi-group diffusion codes with flux limiters and did not address angle-dependent transport. Later, various groups achieved a multi-group/multi-angle neutrino Boltzmann capability (Mezzacappa \\& Bruenn 1993; Rampp \\& Janka 2000; Mezzacappa et al. 2001; Thompson, Burrows, \\& Pinto 2003; Liebend\\\"orfer et al. 2001ab, 2004, 2005).  In 2D simulations, LeBlanc \\& Wilson (1970) constructed a gray, flux-limited diffusion code, that nevertheless calculated a two-component vector flux.  Herant et al. (1994) fielded a gray diffusion code with a simple matching to the free-streaming regime.  Burrows, Hayes, \\& Fryxell (1995) developed a gray diffusion code that calculated different solutions for the different $\\theta$ angles (in spherical coordinates) in a 2D hydro simulation, but calculated at each angle a spherical model that employed the matter profiles along that one ray.  They did not calculate lateral transport in the angular direction.  This is the so-called ``ray-by-ray\" approach that is still used by some groups today.  Rampp \\& Janka (2000,2002), Janka et al. (2005ab), and Buras et al. (2006) have updated the ray-by-ray method using sophisticated 1D Boltzmann transport along each ray, but no lateral transport (though they do include in the hydro the effects of lateral radiation pressure and lepton number transport).  The transport is calculated in the comoving frame and mapped to the Eulerian frame of their PPM hydrodynamics. This prescription is viable only if the flow maintains rough sphericity and the core does not move off the center of the grid (as when kicks are imparted to the protoneutron star). However, the ray-by-ray approach is not real 2D transport.  The ZEUS-2D radhydro code of Stone, Mihalas, and Norman (1992), and its update by Hayes and Norman (2003), solve the zeroth- and first-moment equations of the gray transport equation and use a short-characteristics solution of the static transfer equation to obtain the second moment for closure.  These codes are realizations of the variable Eddington factor method and avoid some of the pitfalls of flux limiters, but since they are gray and not multi-group, they are of limited utility for modern supernova simulations.  All the above multi-group codes were formulated in the comoving frame and none of the 2D variants was simultaneously 2D, multi-group, and time-dependent.  A 2D, multi-group, flux-limited capability has recently been developed by Swesty \\& Myra (2005ab, 2006), but they treat the inner core in 1D, not 2D, and to date have not published long-duration core-collapse and post-bounce simulations. Cardall, Lentz, \\& Mezzacappa (2005) and Cardall \\& Mezzacappa (2003) have derived the general form of the neutrino transport equations including higher-order terms in $v/c$ and general relativity, but have not yet developed working implementations. The first bonafide 2D multi-group, multi-angle, time-dependent capability was achieved by Livne et al. (2004) using the implicit, Arbitrary-Lagrangian-Eulerian (ALE) code VULCAN/2D (with remap), for which the radiation field was defined in the laboratory frame, but this code is not fast, does not include the Doppler shifts due to the velocity field (though it does include advection), and does not include energy redistribution.  However, its multi-group, flux-limited diffusion variant is 2D in the entire computational domain and much faster than its multi-angle version. With it, a number of multi-group, fully-2D, radiation/hydrodynamic investigations have been possible, with and without rotation (Walder et al. 2005; Burrows et al. 2006; Ott et al. 2006ab; Dessart et al. 2006ab).  However, as flops become cheaper and computer speeds improve, one will need to do better. This is what motivates the present paper and the development of the new 2D mixed-frame, implicit, multi-group, multi-angle algorithm and its two-moment variants. In the context of hydrodynamics, this code has been christened BETHE\\footnote{{\\bf B}asic (2-Dimensional), {\\bf E}xplicit/Implicit, {\\bf T}ransport and {\\bf H}ydrodynamics {\\bf E}xplosion (Code)}.  In \\S\\ref{sec-form}, we present the transport equation and its general formulation, but quickly in \\S\\ref{mixed} derive, using the appropriate Lorentz transformations, the equations of radiative transfer in the mixed-frame. This section contains our central analytic results.  In \\S\\ref{moment}, we follow with derivations of the associated moment equations.  Then, in \\S\\ref{hydro} we digress into a discussion of the frame transformations of the source terms on the right-hand-sides of the transport equation and of the associated matter energy and momentum equations.  Care in these matters is important to ensure global energy and momentum conservation to $O(v/c)$ and that one employs the correct radiation source terms in the hydro equations. Different realizations of the hydro equations are possible and the form of the radiation energy source term depends upon the form of the hydro energy equation.  For instance, the matter energy equation can be in first-law format (purely Lagrangian with comoving-frame time derivatives) or can include the kinetic energy explicitly (with Eulerian partial time derivatives). For these two formulations, the radiation source terms are different. In \\S\\ref{cylindrical}, we present the mixed-frame formalism in cylindrical (axisymmetric) coordinates and in \\S\\ref{spherical} we present the mixed-frame formalism in spherical coordinates. Then, in \\S\\ref{sec-sph}, we derive the mixed-frame equations in spherical symmetry and explore various solution techniques, including Accelerated-Lambda-Iteration (ALI), a tridiagonal approximate operator, the use of the Krylov subspace algorithm GMRES, the Discontinuous-Finite-Element (DFE) method, and the Feautrier scheme. In \\S\\ref{moment2}, we derive the associated moment equations, provide the matrix representation, and introduce sphericity factors.  In \\S\\ref{sect-imp}, we derive and discuss procedures for the implicit coupling of matter with radiation, including an ALI treatment and the linearization of the energy and composition ($Y_e$) evolution equations.  We follow this in \\S\\ref{sec-num} with a series of numerical tests using the spherical formulation of the mixed-frame transport equations. These include stationary solutions (\\S\\ref{sec-num-st}) and a time-dependent cooling calculation of an idealized protoneutron star with implicit radiation-matter coupling (\\S\\ref{sec-num-imp}).  Section \\ref{conclusion} provides a summary and the Appendix contains a derivation of the mixed-frame treatment of inelastic energy redistribution. ",
        "conclusions": "\\label{conclusion} In this paper, we develop the mixed-frame formalism for radiation transport in both one and two dimensions and provide various solution methodologies and 1D numerical tests.  Velocity-dependence to O($v/c$), anisotropic scattering, energy-dependent cross sections, energy redistribution due to inelastic scattering, and radiation-matter coupling terms are derived and incorporated into the algorithm.  The equations in cylindrical, spherical, and planar coordinates are provided and the effects of bulk and shear radiation viscosity are automatically embedded into the approach. In two dimensions, rotation is consistently included, extending the Eddington factor to an Eddington tensor with five components and introducing an azimuthal component to the radiation flux. The zeroth- and first-moment equations are derived and their roles in the transport solution and the radiation-matter coupling formalism are fully explored.  Various solution philosophies, such as Discontinuous Finite Element (DFE), Feautrier, and short-characteristics, and various convergence approaches, such as the Accelerated-Lambda-Iteration (ALI), Ng acceleration, and GMRES, are also implemented and compared.  Many radiation-hydrodynamic problems do not require an exquisite treatment of spectral line transport, but a good treatment of continuum transport. When such is the case, the mixed-frame approach is clearly superior to comoving-frame formalisms for which the velocity-dependence is complicated. The virtues of the mixed-frame perspective are many and include simple velocity dependence with no velocity derivatives, straight characteristics, simple physical interpretation, and clear generalization to higher dimensions. Moreover, since multi-dimensional radiation-hydrodynamic simulations frequently employ Eulerian grids and hydrodynamics, calculation of the radiation quantities in the laboratory (Eulerian) frame would seem to be natural.  We also stress that one can still use the mixed-frame formalism for treating spectral lines, provided one works with high enough frequency resolution. In the past, this was considered an excessive requirement. However, modern stationary atmosphere models in 1-D already achieve essentially ``full'' frequency resolution, e.g., with frequency spacing of 0.75 fiducial Doppler widths throughout the whole spectrum, that employs some 2 to 3$\\times 10^5$ frequency points (Lanz \\& Hubeny 2003). In this spirit, one can use to advantage the mixed-frame approach as well to for such models and for general mass motions.   The equations and algorithms we have developed can be used for the transport of any radiation, in particular photons and neutrinos. In addition, the two-moment closure approach to approximate transport that springs naturally from our fully angle-dependent formalism is superior to the incorporation of ad hoc flux limiters into the zeroth-moment (energy) radiation equation. Not only is the finite ``speed of light\" effect automatically included in a two-moment closure, but the solution can be made arbitrarily accurate, depending upon how often the correct Eddington tensor is updated.  Radiation transport methods are core tools of the astrophysicist. Time-dependent techniques to address transport and the coupling of radiation with matter are of central concern to the theorist trying to explain the dynamical phenomena of the Universe. Therefore, along with the ongoing advance of the computational arts, the mixed-frame formalism and numerical techniques we have derived and tested in this paper should prove of great value for the detailed future investigation of complex astrophysical problems."
    },
    "0609/astro-ph0609425_arXiv.txt": {
        "abstract": "We investigate matter-enhanced Mikheyev-Smirnov-Wolfenstein (MSW) active-sterile neutrino conversion in the $\\nu_e \\rightleftharpoons \\nu_s$ channel in the collapse of the iron core of a pre-supernova star. For values of sterile neutrino rest mass $m_s$ and vacuum mixing angle $\\theta$ (specifically, $0.5\\,{\\rm keV}< m_s<10\\,{\\rm keV}$ and $\\sin^22\\theta> 5\\times{10}^{-12}$) which include those required for viable sterile neutrino dark matter, our one-zone in-fall phase collapse calculations show a significant reduction in core lepton fraction. This would result in a smaller homologous core and therefore a smaller initial shock energy, disfavoring successful shock re-heating and the prospects for an explosion. However, these calculations also suggest that the MSW resonance energy can exhibit a minimum located between the center and surface of the core. In turn, this suggests a post-core-bounce mechanism to enhance neutrino transport and neutrino luminosities at the core surface and thereby augment shock re-heating: (1) scattering-induced or coherent MSW $\\nu_e\\rightarrow\\nu_s$ conversion occurs deep in the core, at the first MSW resonance, where $\\nu_e$ energies are large ($\\sim 150$ MeV); (2) the high energy $\\nu_s$ stream outward at near light speed; (3) they deposit their energy when they encounter the second MSW resonance $\\nu_s\\rightarrow\\nu_e$ just below the proto-neutron star surface. ",
        "introduction": "The nature of the dark matter and the explosion mechanism for core collapse supernovae are among the most vexing outstanding problems in modern astrophysics. In this paper we point out a curious link between these problems. We show how sterile neutrinos with rest mass and vacuum mixing parameters in a range which makes them viable cold and warm dark matter candidates \\cite{DM1,DM2,XSF,AFP,DH,AF,Kev,BiermannKusenko,Kev2,AbazajianKoushiappas} also can alter energy and lepton number transport physics in collapsing stellar cores, possibly completely altering the standard core collapse/explosion model.  In broad brush, the electron fraction $Y_e$ is lowered, lowering the pressure and, hence, the homologous core mass. This augers against an explosion because the initial bounce shock energy is reduced, leading to the shock stalling closer to the neutron star where shock re-heating by neutrino energy deposition will be less effective. However, we also point out here that alterations in the rate and nature of neutrino energy transport engendered by active-sterile neutrino flavor conversion could go in the direction of helping the explosion by increasing neutrino energy luminosities at a critical epoch when the fate of the supernova shock is determined.  When large magnetic fields are present, active-sterile neutrino flavor transformation involving sterile neutrinos with rest mass $m_s \\sim {\\rm keV}$ can facilitate the generation of asymmetric neutrino emission from the neutron star surface. This, in turn, can generate large \\lq\\lq kicks,\\rq\\rq\\ which may be required to explain the space motions of at least some pulsars\\cite{kicks,kicks2}. Recent work suggest that these kicks can augment the hydrodynamic transport of neutrino energy to the base of the stalled supernova shock, thereby increasing shock re-heating and helping to generate a supernova explosion\\cite{FK}.  The prospect of sterile neutrinos aiding the supernova explosion is a concept that challenges the conventional wisdom that neutrino energy diverted into sterile states is energy unavailable to aid shock re-heating. However, it must be kept in mind that the energy (kinetic plus optical) in the explosion, $\\sim {10}^{51}\\,{\\rm ergs}$, is only $\\sim 10\\%$ of the neutrino sea energy immediately after core bounce, and only a miniscule $\\sim 1\\%$ of the $\\sim {10}^{53}\\,{\\rm ergs}$ in the neutrino sea a few seconds after core bounce ($t_{\\rm pb} > 2\\,{\\rm s}$).  Perhaps energy is not the problem; but the {\\it transport} of neutrino energy well may be the crux problem. Here we show that if neutrinos spend some of their time in sterile states, their transport mean free paths can be re-normalized upward, thereby increasing the neutrino luminosity at the neutrino sphere. A modest increase in neutrino luminosity during the epoch when the stalled bounce shock is being re-heated ($t_{\\rm pb} < 1\\,{\\rm s}$) may be the difference between an explosion and a dud. Though the existence of light sterile neutrinos may be unlikely, the tremendous leverage they would have in altering core lepton numbers and transport properties suggests that we set aside theoretical prejudice for the purpose of exploration.  Active-active channel neutrino flavor mixing, which we know exists, may also affect supernova physics. Indeed, $\\nu_e \\rightarrow \\nu_{\\mu, \\tau}$ flavor transformation in the in-fall phase of stellar collapse was investigated \\cite{FMWS} shortly after the discovery of the Mikheyev-Smirnov-Wolfenstein (MSW) \\cite{MSW,MSW_1} mechanism for in-medium enhancement of neutrino flavor conversion. However, because of the large lepton degeneracy and low entropy in the collapsing core, only un-physically large active neutrino masses can undergo significant matter-enhanced flavor transformation there.  The mass-squared differences, and two of the four vacuum mixing parameters, relevant for the 3-active-neutrino ($\\nu_e, \\nu_\\mu, \\nu_\\tau$) case now are known accurately from experiment \\cite{nuexpt,nuexpt_1}. The mass-squared differences are: the \\lq\\lq solar neutrino scale\\rq\\rq\\ $\\delta m^2 \\approx 8\\times{10}^{-5}\\,{\\rm eV}^2$; and the \\lq\\lq atmospheric scale\\rq\\rq\\ $\\delta m^2\\approx 3\\times{10}^{-3}\\,{\\rm eV}^2$. These are small on the scale of neutrino mass-squared difference necessary to influence core in-fall epoch physics \\cite{FMWS}.  However, the anomalous result of the LSND experiment \\cite{LSND,LSND_1,LSND_2}, if interpreted as arising from vacuum neutrino oscillations, would seem to suggest the existence of a fourth neutrino which, on account of the $Z^0$-width limit, must have interactions considerabley weaker than those of the weak interaction. This is then a so-called \\lq\\lq sterile neutrino,\\rq\\rq\\ perhaps an SU$(2)$ Standard Model singlet. If true, this is an astounding result which calls into question the \\lq\\lq See-Saw\\rq\\rq\\ model for how very massive right-handed neutrino states force neutrino states closely associated with active flavors to be light. The mini-BooNE experiment \\cite{mini} is currently probing the same parameter range as the LSND experiment (and more).  Regardless of the outcome of this experiment, sterile neutrinos with rest masses, and vacuum mixing angles with active species, in the ranges $0.5\\,{\\rm keV} < m_s < 15\\,{\\rm keV}$ and $\\sin^22\\theta \\ge10^{-12}$, respectively, remain a possibility. Currently the only constraints on this speculative sector of particle physics come from X-ray astronomy\\cite{x-ray,x-ray_1}.  The existence of neutrinos with these mass/mixing parameters could subtly or, more likely, drastically alter supernova core and shock re-heating physics as well as the supernova neutrino signal. Lighter sterile neutrinos have been considered in other aspects of the supernova environment\\cite{CFQ,1995hep.ph...11323P,1997PhRvD..56.1704N,1999PhRvC..59.2873M,Fetter:2002xx}. The central point of this paper is that to assess the effects of these sterile neutrinos on supernovae we must include them starting from the outset of core collapse, as many \\lq\\lq down-stream\\rq\\rq\\ aspects of composition, equation of state, and energy/entropy transport which affect the explosion mechanism, nucleosynthesis, and expected neutrino signal could be affected.  In Section II we discuss active-sterile neutrino flavor conversion through MSW resonances in the collapsing core. We examine the prospects for energy and lepton number transport enhancement in Section III. We give a discussion and conclusions in Section IV. Appendix A describes the liquid drop model equation of state and the lepton number re-distribution schemes employed in our one-zone collapse code. ",
        "conclusions": "We have demonstrated that the existence of sterile neutrinos with rest masses in the range of interest for dark matter could have significant implications for the in-fall, collapse epoch of core collapse supernovae. In particular, sterile neutrinos with rest masses in the range  $1\\,{\\rm keV} \\le m_s < 10\\,{\\rm keV}$, and possessing effective vacuum mixing angles characteristic of the channel $\\nu_e\\rightleftharpoons\\nu_s$ which satisfy $\\sin^22\\theta > {10}^{-12}$, could result in enhanced electron capture during in-fall with a concomitant decrease in electron fraction $Y_e$. As discussed in Section II~E, the decrease in $Y_e$ will translate into a decrease in relativistic electron degeneracy pressure in the core. In turn, the consequence of this will be a smaller homologous core and, therefore, a lower initial shock energy and an increased photo-disintegration burden on the shock. Although these alterations certainly auger against obtaining a viable shock and explosion in the framework of current models, it must be kept in mind that there remain significant uncertainties in fundamental nuclear electron capture, equation of state, and transport physics in the post core bounce regime.  \\begin{figure}[htbp] \\includegraphics[width=3.1in]{extrapol_AbazajianFig.eps} \\caption{\\label{fig 9}{\\small Lightly shaded and heavily shaded regions show the ranges of sterile neutrino rest mass $m_s$ (in keV) and effective $\\nu_e\\rightleftharpoons\\nu_s$ vacuum mixing $\\sin^22\\theta$ that give large $Y_e$ reduction and shock re-heating/explosion enhancement, respectively. The line labeled by $\\gamma_{100}>1$ shows the contour of mixing parameters along which the adiabaticity parameter $\\gamma$ at the outer MSW resonance is unity or greater for neutrinos with energy $E_\\nu = 100\\,{\\rm MeV}$. The corresponding line $\\gamma_{10}>1$ shows where $\\gamma$ is unity or greater for neutrinos with energy $E_\\nu=10\\,{\\rm MeV}$. The heavy solid line and the cross-hatched region to the right show the range of parameters ruled out by X-ray observations of the Virgo cluster \\cite{AFP,x-ray}. The dotted line and the region to its right is ruled out by observations of the diffuse X-ray background \\cite{2006MNRAS.370..213B}. The dashed line (labeled LMC) is the claimed X-ray limit from the Large Magellanic Cloud \\cite{Boyarsky}. The dashed and solid horizontal lines are the Lyman-$\\alpha$ forest-derived lower limits on $m_s$ from Ref.~\\cite{Viel} and Ref.~\\cite{Abazajian06}, respectively.}} \\end{figure}  However, we could conservatively conclude that the existence of sterile neutrinos with these properties would result in a supernova paradigm radically different from current models. Additionally, we have outlined a means whereby active-sterile-active neutrino flavor transformation (specifically, $\\nu_e\\rightarrow\\nu_s\\rightarrow\\nu_e$) could alter electron lepton number and energy transport rates. Active neutrinos which spend part of their time propagating as sterile neutrinos will have larger transport mean free paths on average. In short, much of supernova physics \\lq\\lq down stream\\rq\\rq\\ of core bounce could be quite different from standard models.  For example, a higher $\\nu_e$ luminosity as a result of $\\nu_e\\rightarrow\\nu_s\\rightarrow\\nu_e$ would result in a number of changes in the near hydrostatic envelope above the neutrino sphere and below the shock. This would mean a much reduced neutron excess in this region, since the neutron-to-proton ratio is determined locally by a competition between $\\nu_e+n\\rightleftharpoons p+e^-$ and $\\bar\\nu_e+p\\rightleftharpoons n+e^+$. This would have implications for alpha-rich freeze-out, neutrino-heated nucleosynthesis coming from this region. In fact, this effect could be beneficial, increasing $Y_e$ and lowering the yield of neutron-rich, neutron number $N=50$ nuclei. However, the initial neutrino spectra at the neutrino sphere could be altered by the the enhanced deposition of $\\nu_e$'s just below the proto-neutron star surface. This could partially counteract the neutron excess-lowering effect in the envelope by reducing the average $\\nu_e$ energy relative to that for $\\bar\\nu_e$.  In the end, our results call into question the old idea that sterile neutrinos remove energy from the core and, therefore, cannot exist, else there would be no supernova explosions. It must be kept in mind that the supernova explosion energy (kinetic plus optical) is $\\sim {10}^{51}\\,{\\rm erg}$, which is only $\\sim 1\\%$ of the total amount of energy radiated away over some $10\\,{\\rm s}$ as neutrinos of all kinds. It may be possible to \\lq\\lq throw away\\rq\\rq\\ much of this energy and still get an explosion. The timing, rate, and location of energy deposition are the most important determinants of explosion physics and we have demonstrated that some or all of these can be affected by sterile neutrinos with masses/mixings in the range of interest for dark matter.  In Figure~\\ref{fig 9} we summarize the the parameter space of sterile neutrino rest masses and effective vacuum mixing angles (in the $\\nu_e\\rightleftharpoons\\nu_s$ channel) that gives rise to supernova effects. This figure also shows possible constraints stemming from cosmological sterile neutrino radiative decay ($\\nu_s \\rightarrow \\nu_\\alpha + \\gamma$). To obtain significant $\\nu_e\\rightarrow\\nu_s\\rightarrow\\nu_e$ enhancement of $\\nu_e$ luminosity at the neutron star surface (neutrino sphere), high energy $\\nu_s$'s coming from deep in the core must encounter the outer MSW resonance and be converted with high efficiency to $\\nu_e$'s. In the figure, the line labeled by $\\gamma_{100}>1$ shows the contour of mixing parameters along which the adiabaticity parameter $\\gamma$ at the outer MSW resonance is unity for neutrinos with energy $E_\\nu = 100\\,{\\rm MeV}$. Likewise, this figure also shows the corresponding contour, $\\gamma_{10}>1$, for adiabaticity for neutrinos with energy $E_\\nu=10\\,{\\rm MeV}$. The lightly shaded region in this figure, which shows the parameter range for large reduction in $Y_e$, is bounded on the low mixing angle side by the contour $\\gamma_{10}>1$. This is because augmentation of electron capture requires only that $\\nu_e$ neutrinos with energies $E_\\nu \\sim 10\\,{\\rm MeV}$ be efficiently converted  to $\\nu_s$'s.  The existence of relatively small scale structure in the universe, {\\it i.e.,} the Lyman alpha forest, sets an upper limit on neutrino suppression of the matter power spectrum and the dark matter particle collisionless damping scale. In turn, this sets a lower limit on the rest mass of sterile neutrinos if they are to be the bulk of the dark matter. This limit is in dispute. In Figure~\\ref{fig 9} we show both the Abazajian limit \\cite{Abazajian06} and a more conservative limit \\cite{Viel:2005qj}. Likewise, the lack of an observed X-ray line feature in observations of the Large Magellanic Cloud (LMC) has been argued to set an upper limit on sterile neutrino rest mass \\cite{Boyarsky}. We also show current diffuse X-ray background constraints \\cite{2006MNRAS.370..213B}. Perhaps the most stringent constraints come from the lack of observation of a sterile neutrino X-ray decay line from the Virgo cluster \\cite{AFP,x-ray}. This is the heavy solid line and the cross-hatched region to its right.  With all constraints applied, there remains a significant region of unconstrained $m_s\\--\\sin^22\\theta$ parameter space in which there could be large supernova in-fall epoch and shock re-heating effects. Neutrinos with these parameters could also be the dark matter, at least if there is a significant primordial lepton number \\cite{XSF,AFP,AF,DM2} or with particular couplings \\cite{DM2}. These neutrinos fall in the range of parameters that, after core bounce, could engineer large pulsar kicks \\cite{kicks2,kicks}  and modify hydrodynamic neutrino energy transport to the base of the shock \\cite{FK}.  The ultimate conclusion that can be drawn from our work is that stellar collapse is remarkably sensitive to new lepton number-violating physics in the neutrino sector. The reason for this is clear: the energy in core collapse supernovae is gravitational. This gravitational binding energy is stored in degenerate seas of leptons ($\\nu_e$'s and electrons). Any lepton number violating processes in operation can tap into this energy and thereby alter supernova core composition and dynamics. Sterile neutrinos with rest masses $\\sim {\\rm keV}$ which mix with active neutrinos at the level of one part in $\\sim {10}^{11}$ in vacuum are likely to be undetectable in conventional terrestrial experiments but, remarkably, could completely change the current model for core collapse supernovae."
    },
    "0609/astro-ph0609669_arXiv.txt": {
        "abstract": "% We review the theory and observations of star cluster disruption. The three main phases and corresponding typical timescales of cluster disruption are: {\\it I) Infant Mortality} ($\\sim10^7$~yr), {\\it II) Stellar Evolution} ($\\sim10^8$~yr) and {\\it III) Tidal relaxation} ($\\sim10^9$~yr). During all three phases there are additional tidal external perturbations from the host galaxy.  In this review we focus on the physics and observations of Phase I and on population studies of Phases II \\& III and external perturbations (concentrating on cluster-GMC interactions).  Particular attention is given to the successes and short-comings of the Lamers cluster disruption law,  which has recently been shown to stand on a firm physical footing. ",
        "introduction": "The vast majority (perhaps all) of stars are formed in a clustered fashion. However, only a very small percentage of older stars are found in bound clusters.  These two observations highlight the importance of clusters in the star-formation process and the significance of cluster disruption. The process of cluster disruption begins soon after, or concurrent with, cluster formation. \\cite{lada03} found that $\\lesssim10\\%$ of stars formed in embedded clusters end up in bound clusters after $\\sim10^{8}$~yr. \\cite{whitmore03} and \\cite{fall05} have shown that at least 20\\%, but perhaps all, star formation in the merging Antennae galaxies is taking place in clusters, the majority of which are likely to become unbound.  The case is similar in M51, with $>60\\%$ of all young ($<10$~Myr) clusters likely to be destroyed within the first 10s of Myr of their lives \\citep{bastian05}.  On longer timescales, \\cite{oort58} and \\cite{wielen71} noted a clear lack of older ($>$~few Gyr) open clusters in the solar neighbourhood and \\cite{bl03} found a strong absence of older clusters in M51, M33, SMC, and the solar neighbourhood.  The lack of old open clusters in the solar neighbourhood is even more striking when compared with the LMC, which contains a significant number of `blue globular clusters' with ages well in excess of a Gyr (e.g. \\citealt{1966MNRAS.134...59G, degrijs06}).  This difference can be understood either as a difference in the formation history of clusters or as a difference in the disruption timescales.  This later scenario was suggested by \\citet{hodge87}, who directly compared the age distribution of Galactic open clusters and the SMC cluster population. He noted that there are $10-15$ times more clusters with an age of 1 Gyr in the SMC as compared to the solar neighbourhood (when normalising both populations to an age of $10^8$~yr) and concluded that disruption mechanisms must be less efficient in the SMC.  Much theoretical work has gone into the later scenario, with both analytic and numerical models of cluster evolution predicting a strong influence of the galactic tidal field on the dissolution of star clusters (for a recent review see \\citealt{baumgardt06}).  Only recently has there been a large push to understand cluster disruption from an observational standpoint in various external potentials, making explicit comparison with models \\citep{bl03,lamers05a,lamers05b,gieles05a,lamers06}.  We direct the reader to the review by Larsen in these proceedings for a historical look at the observations and theory of cluster disruption.  \\subsection{Phases of cluster disruption} \\label{subsec:phases}  While cluster disruption is a gradual process with several different disruptive agents at work simultaneously, one can distinguish three general phases of cluster mass loss and disruption. As we will see, a large fraction of clusters gets destroyed during the {\\it primary} phase. The main phases and corresponding typical timescales of cluster disruption are: {\\it I) Infant Mortality} ($\\sim10^7$~yr), {\\it II) Stellar Evolution} ($\\sim10^8$~yr) and {\\it III) Tidal relaxation} ($\\sim10^9$~yr). During all three phases there are additional tidal external perturbations from e.g. giant molecular clouds (GMCs), the galactic disc and spiral arms that heat the cluster and speed up the process of disruption. However, these external perturbations operate on longer timescales for cluster populations and so are most important in Phase III. In Fig.~\\ref{fig0} we schematically illustrate the three Phases of disruption and the involved time-scales. Note that the number of disruptive agents decreases in time.  In this review we will focus on the physics and observations of Phase I as well as on recent population studies aimed at understanding Phases II and III on a statistical basis.  For a recent review on the physics of Phases II and III, we refer the reader to \\citet{baumgardt06}.  \\begin{figure} \\begin{center} \\includegraphics[height=8cm]{fig0.ps} \\caption{Schematic overview of the three phases of cluster disruption considered and the responsible physics that drives the disruption. } \\label{fig0}       % \\end{center} \\end{figure}   Before proceeding, it is worthwhile to consider our definition of a cluster. \\cite{schweizer06} defines a cluster to be a {\\it gravitationally bound} stellar association which will survive for 10--20 crossing times.  This definition implies that the stars provide enough gravitational potential to bind the cluster and ignores the role of gas in the early evolution of clusters.  In this review, we will define a cluster as a collection of gas and stars which was {\\it initially gravitationally bound}.  The reason for this definition will become evident in Section~\\ref{infantmortality} ",
        "conclusions": "\\label{sec:discussion}  We showed in \\S\\S~\\ref{sec:populations}\\&\\ref{sec:external} that the simple Lamers disruption law can successfully explain the age and mass distribution of young star clusters populations. Here we will discuss other observations lending support to the Lamers law and some of the standing problems and uncertainties of this scenario which need further attention.  \\subsection{Independent checks on the disruption law}   It is reassuring to see that different datasets of various cluster populations all come to similar conclusions regarding the disruption of clusters.  \\citet{degrijs06} use a variety of studies to look at the cluster population of the LMC.  They also find a lack of old clusters (with respect to what would be expected from a continuous cluster formation rate) and derive $\\gamma=0.56$, again in agreement with other galaxies studied by \\citet{bl03} and \\citet{lamers05a}. Note that a lower value of $\\gamma$ is expected to be observed when the typical \\tdis\\ is of the same order as the oldest clusters in the sample (Fig.~\\ref{fig:mf}), as is the case in LMC.  Outside the local group, the strongly interacting galaxy NGC~6745 has been studied by \\citet{degrijs03} who found evidence for mass dependent disruption, with $\\gamma=0.54$.  The rich cluster system of the intermediate-age merger remnant NGC~1316 shows a clear bimodal colour distribution, with the red component presumably being formed during the merger.  \\citet{goudfrooij04} showed, using deep {\\it HST-ACS} images that if one breaks the red component into `inner' and `outer' regions (with respect to the galactic centre), that the outer region is a continuous power-law while the inner region shows a power-law behavior at the high luminosity end and a flattening at the low luminosity end.  The authors interpret this as evidence for mass-dependent cluster disruption, although no attempt was made to find the characteristic disruption timescale or the value of $\\gamma$.  One standing problem with the Lamers disruption law, also present in other studies on disruption, is whether or not an initial power-law cluster initial mass function (CIMF) can be transformed into a log-normal distribution, which is observed for old globular cluster populations.  The Lamers law can create such a turnover, however the precise value of the turnover mass should be dependent on the ambient density \\citep{lamers05b}, meaning that cluster disruption should be more efficient in the inner regions of a galaxy than in the outer regions.  Thus, without fine tuning the models (e.g. having the same disruption time at all radii due to large radially dependent velocity anisotropies) one would expect a radially dependent turnover peak in the globular cluster mass function, which is not observed.  For a more detailed description of this problem, see the review by Larsen in these proceedings.  Additionally, as noted by \\citet{waters06} the Lamers disruption over-predicts the number of low-mass clusters when applied to old globular cluster populations.    \\subsection{Caveats in the theoretical underpinning} In \\S~\\ref{subsec:tdis_external} we showed that the scaling of \\tdis\\ with \\Mc\\ is a power-law with exponent $\\gamma\\simeq0.6$.  This scaling is similar for two-body evaporation in a tidal field with external perturbations, such as shocks by GMCs and spiral arms, and agrees well with the observations. However, there are still some caveats in the theory explaining this, mostly coming from questions regarding the initial conditions of the simulations.  \\begin{enumerate}  \\item The first caveat stems from the relation between initial mass and radius of the clusters used in the simulations.  If we parameterize this relation as \\rh\\ $\\propto$ \\Mc$^\\lambda$, then \\citet{bm03} use $\\lambda = 1/3$, implying that their clusters fill their tidal radius.  However, observations imply that $\\lambda = 0-0.1$ (with a large scatter) \\citep{larsen04,bastian05}, implying that \\rh\\ is mostly independent of mass. This shallow relation implies that massive clusters are not filling their tidal radius, which would change the dependence of \\tdis\\ with \\Mc\\ \\citep{tanikawa05}.  \\item In the derivation of $\\gamma$ for external shocks (\\S~\\ref{subsec:tdis_external}), only clusters in isolation were considered. How would the presence of a tidal field affect this result?  \\item  How does the relation for \\tdis\\ change if there exists a relation between the concentration parameter and mass of a cluster (i.e. as seen in the Galactic GCs reported by Larsen in these proceedings)?  \\item Could an initial mass-radius relation with $\\lambda=1/3$ be erased during the gas removal phase?  \\item What is the effect of the initial mass/luminosity profile used in the simulations and how does it evolve?  e.g. are clusters born with EFF profiles which are converted into King profiles?  Does the cluster concentration alter its mass loss evolution?  \\item How do the external perturbations and the galactic tidal field cooperate? Can the mass loss due to both effects simply be added linearly?  \\end{enumerate}"
    },
    "0609/astro-ph0609343_arXiv.txt": {
        "abstract": "Structural analysis has been performed for a sample of 15 Southern early-type disk galaxies, mainly S0s, using high resolution $K_s$-band images. The galaxies are mostly barred and many of them show multiple structures including bars and ovals, typical for S0s. The new images are of sufficient quality to reveal new detail on the morphology of the galaxies. For example, we report a hitherto undetected nuclear ring in NGC 1387, a nuclear bar in NGC 1326, and in the residual image also a weak primary bar in NGC 1317. For the galaxies we (1) measure the radial profiles of the orientation parameters derived from the elliptical isophotes, (2) apply Fourier methods for calculating tangential forces, and particularly, (3) apply structural decomposition methods. For galaxies with multiple structures a 2-dimensional method is found to be superior to a 1-dimensional method, but only if in addition to the bulge and the disk, at least one other component is taken into account. {\\it We find strong evidence of pseudo-bulges in S0s}: ten of the galaxies have the shape parameter of the bulge near to $n$ = 2, indicating that the bulges are more disk-like than following the R$^{1/4}$-law. Also, six of the galaxies have either nuclear rings, nuclear bars or nuclear disks. In all non-elliptical galaxies in our sample the $B/T < $ 0.4, as typically found in galaxies having pseudo-bulges. {\\it In two of the galaxies the $B/T$ flux ratio is as small as in typical Sc-type spirals}. This might be the hitherto undiscovered link in the scenario in which spirals are transformed into S0s. {\\it Also, bars in S0s are found to be shorter and less massive, and have smaller bar torques than bars in S0/a-type galaxies}. ",
        "introduction": "Structural components of galaxies like bars, bulges, ovals and bar-related resonance rings are among the key factors for evaluating when present-day galaxies formed and how they evolve over time. Bulges in early-type disk galaxies are generally assumed to be old, formed through hierarchical clustering, but they might also be more recent structures assembled by internal dynamical processes in galaxies. In favorable conditions these processes might lead to transformation of spirals into S0s, through gas stripping in the halo and the disk (Bekki, Warrick $\\&$ Yashuro 2002), or due to strong and sustained star formation in the central regions of galaxies (Kormendy $\\&$ Kennicutt, hereafter KK2004).  \\vskip 0.25cm In these secular evolutionary processes the relative mass of the bulge increases. In case of gas stripping the bulges should have older stellar populations than the disks and their surface brightness profiles are expected to resemble more the $R^{1/4}$ law than the exponential function. However, due to a small amount of gas in S0s, star formation is not expected to be the dominant mechanism for the growth of the bulge, unless exceptionally strong accretion of external gas is taking place in these galaxies. Bulge-like mass concentrations can form also through resonant scattering or through bending instabilities in the central regions of primary bars (Pfenniger 1984; Raha et al. 1991), in which case these so called ``pseudo-bulges'' should have similar stellar populations as the underlying bars. In order to evaluate the possible role of these processes in galaxy evolution, the properties of bars, ovals and bulges need to be systematically studied in the spiral/S0 divide, e.g. through Sa, S0/a, to S0$^+$, S0$^-$ types.   \\vskip 0.25cm Secondary processes in galaxy evolution are generally tied to the growth of a bulge in one way or another, but an alternative approach was also discussed by van den Bergh (1979, 1998), who outlined the DDO system of classification in terms of a nurture theory. In this classification system, S0s form a sequence parallel to spirals with types S0a, S0b, and S0c being the S0 analogues of Sa, Sb, and Sc galaxies. A difficulty with this approach has been that no S0c galaxies have been found so far, that is, clear S0 galaxies showing a bulge to total luminosity ratio similar to an Sc galaxy. However, this might be a bias due to the fact that application of too simple analysis methods can easily hide the real nature of bulges, particularly in S0s which have rich morphological structures.  \\vskip 0.25cm  To our knowledge no previous systematic morphological study of galaxies has been carried out in the spiral/S0 divide in the near-infrared. Also, although 2-dimensional decompositions have become a standard, even in recent studies (de Souza, Gadotti $\\&$ dos Anjos 2004, hereafter SGA2004; Christlein $\\&$ Zabludoff 2004) multiple structures have not been taken into account. The current study is a third paper in a series where morphological properties of a sample of 170 early-type disk galaxies are studied. The description of the sample and the first results are presented in Paper I by Laurikainen, Salo $\\&$ Buta (hereafter LSB2005), and in Paper II by Buta et al. (2006, in press). In paper I we found, for a subsample of 24 S0-S0/a galaxies, that the $B/T$ flux ratios are significantly smaller, and the bulges more disk-like, than generally assumed for their morphological type. In paper II we investigate in detail the Fourier properties of early-type bars and ovals, and identify bars, which are well described in terms of single or double Gaussian components of the relative Fourier amplitudes of density. In this study we analyze the morphological properties of 15 Southern galaxies of the de Vaucouleurs revised Hubble types S0$^-$ to Sa, mainly S0s, with a large range of apparent bar strengths. For these galaxies we use $K_s$-band images to derive the orientation parameters from the elliptical isophotes, apply Fourier methods for calculating non-axisymmetric forces, and perform 2D multicomponent structural decompositions. The multicomponent decompositions are compared with more simple bulge/disk decompositions and with 1D decompositions for the same galaxies.  \\vskip 0.25cm ",
        "conclusions": "A morphological analysis of a sample of 15 early-type disk galaxies, mainly S0s, has been presented in the $K_s$-band. The sample is part of a more complete NIRS0S-sample of 170 S0-S0/a galaxies, selected to have a similar size and selection criteria as in the OSUBGS-sample for spirals. Our main idea is to characterize the properties of bars, ovals and bulges for a large sample of galaxies in the spiral/S0 divide and to use this information to evaluate the origin of S0s and how they are related to spirals. \\vskip 0.25cm  The galaxies were analyzed by applying Fourier methods for calculating the maximum gravitational torques, and by applying a multicomponent 2D method for structural decompositions. A generalized Sersic's function is used for the bulge and exponential function for the disk, whereas the other non-axisymmetric components can be described either by a Sersic's or a Ferrers function. It was also possible to take into account the flattening and the shape parameter of the structures. For comparison 1D decomposition method was also used. We found that if only the bulge and the disk are fitted, 1D and 2D methods give fairly similar results ($<B/T>$ = 0.55-0.61).  \\vskip 0.25cm  Bars in the galaxies in our sample are typically surrounded by lenses, and four of the galaxies have double bar/lens structures. Two of the galaxies classified as S0s in RC3 appeared to be elliptical galaxies (ESO 208-G21, NGC 3706), and three of the non-barred S0s were found to have inner disks (NGC 1411, NGC 3100, NGC 1553). A previously non-detected nuclear ring was discovered in NGC 1387, and a nuclear bar in NGC 1326. We found that bars in S0s are not exactly similar to bars in S0/a-galaxies. They are clearly shorter and less massive and have slightly weaker bar torques than bars in S0/a-galaxies.  \\vskip 0.25cm  One of the most important results of this study is the finding that {\\it 11 non-elliptical galaxies in our sample show strong evidence of having disk-like pseudo-bulges}: these bulges have surface brightness profiles that are more close to an exponential than to a de Vaucouleurs' type profile. Six of the galaxies also have either nuclear rings, nuclear bars or nuclear disks. All galaxies in our sample have $B/T <$ 0.4 as typically found in galaxies having pseudo-bulges. We find $<B/T>$ = 0.25 and $<n>$ = 2, thus confirming the earlier result by LSB2005, showing that bulges in S0s are considerably smaller and have more exponential surface brightness profiles than assumed for their morphological types. A comparison to more simple 2D bulge/disk or 1D-decompositions shows the importance of the multicomponent approach: these two methods give $<B/T>$ = 0.55-0.48, respectively, in both cases the relative mass of the bulge being significantly overestimated. The value $<B/T>$ = 0.55 is close to the value 0.64 obtained by SGA2004 for the early-type disk galaxies in the same wavelength, using a similar 2D bulge/disk approach as used in this study. \\vskip 0.25cm  Quite interestingly, we also find two galaxies, NGC 1411 and NGC 3100, which due to their very small $B/T$ flux ratios might be proverbial analogues of Sc galaxies among classical S0s, as discussed in the DDO galaxy classification scheme by van den Bergh. These galaxies might be manifestations of galaxy evolution in the Hubble sequence, possibly formed by gas stripping mechanisms from late-type spirals."
    },
    "0609/astro-ph0609205_arXiv.txt": {
        "abstract": "We report on an X-ray flare detected on the active binary system II~Pegasi with the Swift telescope.  The event triggered the Burst Alert Telescope in the hard X-ray band on December 16, 2005  at 11:21:52 UT with a 10-200 keV luminosity of 2.2$\\times$10$^{32}$ erg s$^{-1}$ --- a superflare, by comparison with energies of typical stellar flares on active binary systems. The trigger spectrum indicates a hot thermal plasma with T$\\sim$180 $\\times$10$^{6}$K. X-ray spectral analysis from 0.8--200 keV with the X-Ray Telescope and BAT in the next two orbits reveals evidence for a thermal component (T$>$80 $\\times$10$^{6}$K) and Fe K 6.4 keV emission. A tail of emission out to 200 keV can be fit with either an extremely high temperature thermal plasma (T$\\sim$3$\\times$10$^{8}$K) or power-law emission.  Based on analogies with solar flares, we attribute the excess continuum emission to nonthermal thick-target bremsstrahlung emission from a population of accelerated electrons. We estimate the radiated energy from 0.01--200 keV to be $\\sim$6$\\times$10$^{36}$ erg, the total radiated energy over all wavelengths $\\sim$10$^{38}$ erg, the energy in nonthermal electrons above 20 keV $\\sim$3$\\times$10$^{40}$ erg, and conducted energy $<$5$\\times$10$^{43}$ erg. The nonthermal interpretation gives a reasonable value for the total energy in electrons $>$ 20 keV when compared to the upper and lower bounds on the thermal energy content of the flare. This marks the first occasion in which evidence exists for nonthermal hard X-ray emission from a stellar flare. We investigate the emission  mechanism responsible for producing the 6.4 keV feature, and find that collisional ionization from nonthermal electrons appears to be more plausible than the photoionization mechanism usually invoked on the Sun and pre-main sequence stars. ",
        "introduction": "The general model for solar (and stellar) flares involves a release of free energy via magnetic reconnection, which accelerates particles and causes subsequent plasma motions and heating \\citep{dennis1989}. Particle acceleration during solar flares can be diagnosed directly through nonthermal hard X-ray bremsstrahlung emission and radio gyrosynchrotron emission. Until now, evidence for particle acceleration during stellar flares has been either indirect \\citep[through proxies such as optical and UV flare emissions;][]{slh1995,slh2003} or achieved through the detection of nonthermal radio flare emissions \\citep{klein1987}, under the assumption that the observed radiation is gyrosynchrotron emission from accelerated particles. The interpretation of radio observations, while useful for constraining the characteristic magnetic field in the source, is complicated by optical depth effects and spatial inhomogeneities in the emitting source, which are usually not uniquely determined. Previous detections of hard X-ray emission from stellar flares could not demonstrate unequivocally the presence of nonthermal emission; instead, a superhot thermal component could equally well explain the observed 10--50 keV emission \\citep{fs1999,pallavic2001,franciosini2001} without the need for any additional power-law components. The detection and characterization of nonthermal hard X-ray emission during stellar flares would be important, as it would allow a direct investigation of the energy spectrum of accelerated electrons, and a comparison of the total thermal/nonthermal energy budget. Prospects for detecting such emissions have normally been taken to be remote, as during solar flares the nonthermal hard X-ray emission (20 keV-- 1 MeV) is $>$10$^{-5}$ times less intense than soft X-ray (1--10 keV) emission \\citep[see, e.g., the composite solar flare spectrum in Fig.~94 of ][]{aschwanden2002}; this is a consequence of collisional energy losses exceeding bremsstrahlung losses by a factor of 10$^{5}$. This necessitates large flare sizes for hard X-ray detections, which are intrinsically rare, which coupled with a dearth of pointed observations exacerbates the situation.  Iron K fluorescence emission has been seen during solar flares \\citep{parmar1984,zarro1992} and also seen in impulsive flares from some pre-main sequence stars \\citep{coupref}.  The origin on the Sun is attributed to photoexcitation of K shell electrons in photospheric iron by X-ray bremsstrahlung radiation.  The emission mechanism for pre-main sequence stars is the same as for the Sun, but the fluorescence is presumed to arise instead from X-ray irradiation of an accretion disk rather than the stellar photosphere.  There have not been any previous reports of iron K$\\alpha$ fluorescent emission from non-accreting stellar flare sources. Such an identification would provide another link between solar flare physics and the physics of flares on other stars.  II~Peg (HD~224085) is a well-known active binary system which displays evidence for starspots covering a large fraction \\citep[$>$40\\%;][]{marino1999} of the primary's photosphere, and has signatures associated with vigorous magnetic activity across the entire electromagnetic spectrum. At a distance of 42 pc \\citep{hipparcos}, this single-lined spectroscopic binary is composed of a 0.8 M$_{\\odot}$ K2IV primary tidally locked in a 6.7 day orbit with a $\\sim$0.4M$_{\\odot}$ secondary \\citep{berd1998}. Space velocities indicate that this binary is a member of the old disk stellar population \\citep{eggen1978}.  Both photospheric and coronal metal abundances have been determined, leading to an apparent depletion of coronal iron by about a factor of 4 relative to the photospheric value \\citep[][and references therein]{chandraref} \\footnote{This is in contrast with the solar corona, where the coronal iron abundance appears to be {\\it enhanced} by a factor of $\\sim$4 over the solar photospheric abundance \\citep{feldmanlaming2000}.  For more discussion on stellar coronal abundances, see \\citet{screview}.}. Based on the stellar parameters for II~Peg in \\citet{berd1998} and using the bolometric corrections in \\citet{bcref}, we estimate the bolometric luminosity, the power emitted by the two stars over all wavelengths, to be $L_{\\rm bol}\\approx$5.5$\\times$10$^{33}$ erg s$^{-1}$.  II Peg was first detected as a flaring X-ray source by the {\\it Ariel-V} Sky Survey Instrument (1.5-20 keV) in the mid-1970s \\citep{schwartz1981} which detected 2 flares with peak fluxes of $1.1 \\times 10^{-9}$ and $1.5 \\times 10^{-9}$ erg cm$^{-2}$ s$^{-1}$ in the 2-10 keV band, equivalent to peak X-ray luminosities of $1.6 \\times 10^{32}$ and $2.8 \\times 10^{32}$ erg s$^{-1}$. Since then, moderate and large flares from II Peg have been detected by many subsequent X-ray and EUV observatories, including {\\it Ginga} \\citep{doyle1991}, {\\it EXOSAT} \\citep{tagliaferri1991}, {\\it EUVE} \\citep{patterer1993,ostenbrown}, {\\it ASCA} \\citep{mewe1997}, {\\it BeppoSAX} \\citep{covino2000}, and {\\it Chandra} \\citep{chandraref}.  In this paper, we report on a large flare observed on II~Peg by the detectors on the Swift Gamma-Ray Burst Mission \\citep{Swiftref}. Following the discussion in \\citet{schaefer2000} and the characteristics of the flare (radiated energy, luminosity) we term this event a ``superflare''. The event triggered Swift's Burst Alert Telescope (BAT) in the hard X-ray band on December 16, 2005 at 11:21:52 UT\\footnote{http://gcn.gsfc.nasa.gov/gcn/gcn3/4357.gcn3}, and its characteristics in the soft and hard X-ray bands were determined for the extent of time for which hard X-ray emission was detected, $\\approx$ 7000 s. The UV$/$Optical Telescope did turn on during the trigger, but the images were saturated due to II~Peg's brightness (V$\\approx$7.4), and the data were of little use.  In this paper we concentrate on analysis of the X-ray Telescope (XRT) and BAT spectra. ",
        "conclusions": "We have identified nonthermal emission as the most plausible mechanism to explain the hard X-ray emission seen during the rise and decay of a large stellar flare. This flare interestingly also displayed evidence for emission at 6.4 keV whose formation we attribute to the mechanism of collisional ionization. The increased sensitivity of the Swift BAT has enabled a detection at much higher energies than had been possible with previous hard X-ray telescopes. Nonthermal hard X-ray emission has enabled an investigation of the energetics of large stellar flares, without complication from optical depth effects and source inhomogeneities. The characteristics of this flare --- a rare, intense, transient event --- point to the value of triggered observations to study such flares; targeted observations would have had a very low likelihood of observing such an event during the short timescales over which the hard X-ray emission is produced.  Triggered observing modes necessarily miss the preflare emissions, as the flux must pass a threshold value to warrant a telescope slew.  Still, serendipitous science can be obtained, as this Swift-observed flare demonstrates.  Current plans are to increase the trigger sensitivity threshhold of Swift by a factor of 4.5, thus enabling more such opportunities. %  Multi-wavelength observations of solar flares have revealed the dynamical response of the atmosphere to these sudden intense inputs of energy. Future observations of stellar flares like the one discussed here can benefit from the global array of telescopes which have been harnessed for gamma-ray burst studies. It will hopefully be possible to take advantage of optical telescopes to reveal the response of the lower atmosphere and radio telescopes to explore further the action of nonthermal particles. Observations of such events will allow for a comparison of particle acceleration processes in active stars and the Sun.   Support for this work was provided by NASA through Hubble Fellowship grant \\# HF-01189.01 awarded by the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc. for NASA, under contract NAS5-26555.  We are grateful to Randall Smith for his assistance in implementing the thick-target bremsstrahlung code in XSPEC, and in extending the APEC calculations to higher energies. RAO is also grateful for discussions with Joel Allred about hard X-ray observations of solar flares. The authors thank Brian Dennis, the referee, for a close reading of the paper and for suggesting improvements to make the paper appeal to a wider audience."
    },
    "0609/astro-ph0609519_arXiv.txt": {
        "abstract": "{} {The aim of this work is to understand to what extend gravitational interactions between the stars in high-density young stellar clusters, like the Orion Nebula Cluster (ONC), change the angular momentum in their protoplanetary discs.} {Two types of simulations were combined ---  N-body simulations of the dynamics of the stars in the ONC, and angular momentum loss results from simulations of star-disc encounters.} {It is shown that in a star-disc encounter the angular momentum loss is usually larger than the mass loss, so that the disc remnant has a lower specific angular momentum. Assuming an age of 1-2 Myr for the ONC, the disc angular momentum in the higher density region of the Trapezium is reduced by 15-20\\% on average. Encounters therefore play an important part in the angular momentum transport in these central regions but are not the dominant process. More importantly, even in the outer cluster regions the angular momentum loss is on average  3-5\\%. Here it is shown that a 3-5\\% loss in angular momentum might be enough to trigger gravitational instabilities even in low-mass discs - a possible prerequisite for the formation of planetary systems.} {}  \\titlerunning{Cluster as trigger for gravitational instabilities} \\authorrunning{Pfalzner et al.} ",
        "introduction": "Discs are seen as a prerequisite for the formation of planetary systems. There is increasing observational evidence that most, if not all, stars are initially surrounded by discs. For example, Lada et al. (2000) found that 80-85\\% of all stars in the Trapezium are surrounded by discs.   One visually compelling feature in the numerical study of these discs is the development of spiral arms either in self-gravitating high-mass discs or as an effect of encounters \\cite{pfalzner:03}. In a few cases where discs have been observed with a sufficiently high spatial resolution, the existence of such spiral structures have been observationally confirmed \\citep[for example]{clampin:03,corder:05,lin:06}. Spiral arms are always connected to mass and angular momentum transport. Our study of the importance of the encounter-induced angular momentum transport in an cluster environment has two motivations: angular momentum transport (a) in the late stages of accretion in star formation and (b)in spiral arm formation as prerequisite for gravitational instabilities and fragmentation in planet formation scenarios.  The former is connected to the problem that the typical observed angular momentum of the cloud cores is so large \\citep{bodenheimer:araa95} that the star would be spinning  at velocities higher than its break-up velocity if no means for angular momentum loss would be present. One way out is angular momentum transfer to the disc and then from the inner disc regions outwards. Many different mechanisms have been proposed for the latter \\citep[see][and references therein]{larson:mnras02}. Although often discarded as probably not important, encounter-induced angular momentum transport is still often cited as possible mechanism. To clarify this point we study disc angular momentum loss in a cluster environment for the first time in a quantitative, systematic manner.  The second aim of this investigation is to achieve a better understanding of how the cluster environment influences the development of spiral structures in the discs through encounter-triggered angular momentum transport. This might be of special importance as thus such spiral structures are a prerequisite in the theory of planet formation through gravitational instabilities.  The Orion Nebula Cluster (ONC) was chosen as the model cluster for these investigations because it is one of the best-studied regions in our galaxy, and observational constraints significantly reduce the modelling parameters. In addition, its high density suggests that stellar encounters might be relevant for the evolution of circumstellar discs. This is supported by recent investigations by Olczak et al. (2006) indicating that the effect of encounters has - at least for the disc mass loss - previously been underestimated.  The investigation involves three steps --- i) the determination of  the encounter parameters in the ONC dynamics, ii) the angular momentum loss in isolated encounters, and iii) the combination of both results to investigate the angular momentum loss of the discs in the cluster.  For the first step we use the cluster simulations described in Olczak et al. (2006) and star-disc simulations in Pfalzner et al. (2006) to determine a parameter list for the encounters for a model ONC in equilibrium and expansion, i.e. virial ratios of $Q_{vir}$=0.5 and  $Q_{vir}$=1. The results are summarized in Section 2.  This is followed by an investigation of the angular momentum loss in star-disc encounters in Section~3. There have been earlier investigations which looked at angular momentum transport in isolated star-disc encounters \\citep{ostriker:apj94,larwood:97,hall:mnras96,pfalzner:apj04} predominantly in prograde, coplanar encounters for low-mass discs covering only a small parameter range. Under these limitations simple N-body simulations suffice, and hydrodynamic effects and self-gravity within the disc can be largely neglected. Following this approach one has to keep in mind that the results have to be considered as upper limits (see discussion in Olczak et al. (2006) for the disc mass loss). We extended the parameter range covered in Pfalzner et al. (2005a) to larger mass ratios between the interacting stars and penetrating encounters  to cover all situations necessary when modelling the ONC. In this work we neglect photo-evaporation as disc dispersal mechanism \\citep{scally:mnras01}, but this should be included in a future investigation.  In Section~4 the results of Section~2 and 3 are combined to determine the disc angular momentum as a function of time. In Section~5 it is shown how the angular momentum mass loss is influenced by the assumptions made and a discussion is given of how this angular momentum loss can trigger gravitational instabilities. ",
        "conclusions": "\\label{sec:discussion}  Combining simulations that follow the path of the stars in the cluster with angular momentum loss results obtained from simulation of isolated encounters, the encounter-induced angular momentum loss has been investigated for the ONC. The main outcome of these simulations is that for most stars in the ONC, encounters reduce the angular momentum on average only  by 3-5\\% in the remaining disc compared to the initial disc. However, this kind of reduction happens even in the outer parts of the ONC, so that the disc angular momentum is affected by encounters throughout the entire cluster. If one considers only the central dense Trapezium region the SAML increases to the 10-30\\% level. This means that although encounters in the cluster environment do lead to angular momentum loss in the disc, it is not the predominant process for angular momentum transfer.  However, it would be premature to declare encounters as not important for the disc development. On the contrary, encounter-induced angular momentum transport may still play an important part in the disc evolution. As one can see in Fig.~\\ref{fig:interact}, an SAML in the 3-5\\% regime, which we showed to be present in the entire ONC, is sufficient to trigger pronounced spiral arm structures in the disc. As one can see similar structures appear independent of the mass of the perturber given the same angular momentum change.   As mentioned above such spiral structures have been detected in the few cases where discs have been observed with a sufficiently high spatial resolution \\citep[for example][]{clampin:03,corder:05,lin:06}. A fly-by has often been rejected as possible cause of the spiral pattern because no stars have been found in the immediate vicinity of the star-disc system. However, as the results in Section 3 show, angular momentum change can be induced by relatively distant encounters given that the mass ratio between the star and the perturber is large. For example, for $M_2^*/M_1^*$=50 an encounter with a periastron $\\sim$14 times the disc radius could create a spiral arm pattern similar to those shown in  Fig.~\\ref{fig:interact}. Considering the far range of interaction in many cases, a much larger area would have to be considered in a search for possible encounter partners.   What consequences can such spiral patterns in the disc have? The spiral arm pattern above might be the starting point for gravitational instabilities (in above simulations self-gravity of the disc is not considered, so clumping can not occur). So far gravitational instabilities have predominantly been investigated for isolated high-mass discs \\citep[for a summary, see][]{durisen:06}. In this context gravitational instabilities can produce global spiral arms and disc fragmentation into dense clumps and substructures. Possibly these dense clumps may become precursors to gas giant planets \\citep{kuiper:51,cameron:78,boss:98}.  Modern numerical simulations, beginning with \\cite{papaloizou:91}, show that non-axisymmetric disturbances become unstable for $Q=c_s \\kappa / \\pi G \\Sigma <$ 1.5, where $c_s$ is the sound speed, $\\kappa$ the epicyclic frequency, G the gravitational constant and $\\Sigma$ the surface density. When the gravitational instabilities emerge from the linear regime, they may either saturate at nonlinear amplitude or fragment the disc \\citep{durisen:06}. What happens depends on whether a balance can be reached between heating and the loss of disk thermal energy by radiative or convective cooling. Although there is agreement on conditions for fragmentation, it is an open question whether real discs cool fast enough for fragmentation to occur and whether these fragments last long enough to contract into permanent protoplanets \\citep{durisen:06}.   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{avtime.ps}} \\caption{SAML as function of time for an expanding cluster at 0.1pc and 1.0pc.} \\label{fig:angtime_exp2} \\end{figure}  A prerequisite for the gravitational instabilities in isolated discs is a high enough disc mass that self-gravity can carry small perturbations to fragmentation in the non-linear regime. However, observations show that high-mass discs seem to be the exception rather than the rule.   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=0]{encount.eps}} \\caption{Examples of encounters with a (SAML) of 3-5\\%. a) shows an encounter where $M_2^*$= 0.1 $M_{\\sun}$ and $r_p$ = 2.0, b)  an encounter where $M_2^*$= 1 $M_{\\sun}$ and $r_p$ = 3.5 and c) an encounter where $M_2^*$= 5 $M_{\\sun}$ and $r_pd$ = 5.5. } \\label{fig:interact} \\end{figure}  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=0]{encount3.eps}} \\caption{Examples of encounters with a SAML of $\\sim$20\\%. a) shows an encounter where $M_2^*$= 0.1 $M_{\\sun}$ and $r_p$ = 1.5, b)  an encounter where $M_2^*$= 1 $M_{\\sun}$ and $r_p$ = 2.8 and c) an encounter where $M_2^*$= 5 $M_{\\sun}$ and $r_p$ = 4.5. } \\label{fig:interact2} \\end{figure}  Various mechanisms have been suggested for triggering gravitational instabilities, among them formation of massive discs from the collapse of protostars \\citep{laughlin:94,yorke:99,boss:98}, clumpy infall onto the disc \\citep{boss:98}, cooling from stable to unstable state, perturbation by binary companions \\citep{mayer:05} and close encounters with star/disc systems \\citep{boffin:98,lin:98}.  Investigating binary systems \\cite{mayer:05} found that the degree of fragmentation depends on the mass of the disc and the cooling rate. The simulations by Mayer et al.(2005) and Boss (2006) both indicate that apparently, in contrast to isolated discs, in binary systems {\\em lighter discs} are more likely to fragment. This counterintuitive result reflects the fact that intense heating from spiral shocks can suppress the formation of permanent clumps. A balance between compression/shock heating and cooling is required for fragmentation.  The situation in an encounter is similar to that in a binary system, because the secondary star induces a non-axisymmetric perturbation onto the star-disc system, but with the difference that the perturbation is not recurring. At the moment it is unclear whether encounters enhance disc fragmentation or not. Investigations by \\cite{lodato:06} have found that discs with long cooling timescales are actually rendered more stable by encounters. On the other hand, \\cite{shen:06} do get encounter induced fragmentation when they use an isothermal equation of state. In the future it needs to be determined by detailed simulations with the right radiation hydrodynamics which amount of angular momentum change is required  to obtain stable clumps --- is it the stronger interactions (10-30\\% angular momentum change) that we find in the cluster center or rather the weaker (3-5\\% angular momentum change) ones further outside?  There is actually a second means by which spiral structures induced by encounters could support planet formation. In the classical picture of planet \\citep[for a summary see][]{weidenschilling:93} formation sticking of dust grains and further agglomeration leads eventually to planet formation. Rice et al. 2006  showed that dust can accumulated in spiral shocks induced in self-gravitating disc. If this happens in the same fashion in encounter induced spirals, encounters could trigger planet formation via dust agglomeration as well. The higher dust density and possibly the higher velocity dispersion in the spiral arm could assist dust agglomeration. However, the time scale of the lifetime of spiral structures is typically only 1000 to 2000 years which would be very stringent for this planet formation scenario.  From the simulations in this paper we have shown that in systems like the ONC such an encounter situation is a common event throughout the cluster. Therefore, if planets can be formed through gravitational instabilities at all, encounters in a cluster environment are a likely trigger for their occurance."
    },
    "0609/astro-ph0609033_arXiv.txt": {
        "abstract": "Observations of far-infrared (FIR) and submillimeter (SMM) polarized emission are used to study magnetic fields and dust grains in dense regions of the interstellar medium (ISM\\@).  These observations place constraints on models of molecular clouds, star-formation, grain alignment mechanisms, and grain size, shape, and composition. The FIR/SMM polarization is strongly dependent on wavelength.  We have attributed this wavelength dependence to sampling different grain populations at different temperatures.  To date, most observations of polarized emission have been in the densest regions of the ISM\\@.  Extending these observations to regions of the diffuse ISM, and to microwave frequencies, will provide additional tests of grain and alignment models.  An understanding of polarized microwave emission from dust is key to an accurate measurement of the polarization of the cosmic microwave background. The microwave polarization spectrum will put limits on the contributions to polarized emission from spinning dust and vibrating magnetic dust. ",
        "introduction": "The dominant source of Galactic emission at far-infrared (FIR) and submillimeter (SMM) wavelengths ($\\lambda\\sim 50\\,\\micron$ -- 1\\,mm, $\\nu\\sim 300$ -- 6000 GHz) is thermal emission from interstellar dust at temperatures of 10 -- 100 K\\@.  In dense regions of the interstellar medium (ISM) such as molecular clouds, this emission has measurable polarization at almost every point (\\eg\\ \\citealt{archive,dotson06}).  This polarization is due to the alignment of the dust grains with interstellar magnetic fields \\citep{dg51}. As a result of the magnetic alignment mechanism, polarimetry allows detailed studies of magnetic fields (\\eg\\ \\citealt{chussth,mwf01,crutcher}).  Polarimetry can also be used to place constraints on grain properties, cloud environments, and grain alignment efficiency (\\citealt{pspec,whittet01}; \\citealt*{lgm97}).  Observations at multiple wavelengths are especially useful for such studies.  At optical and near-infrared wavelengths, spectropolarimetry puts constraints on grain size, shape, and composition \\citep{aitken96,whittet04,martin}.  In molecular clouds, the FIR/SMM polarization spectrum is observed to fall from $60\\,\\micron$ to $350\\,\\micron$ and rise again to $1300\\,\\micron$.  We attribute this spectrum to correlations between temperature and alignment efficiency \\citep{pspec,mythesis}.  Further tests of grain and alignment models require that measurements of the polarization spectrum be extended to more diffuse regions of the ISM and to longer wavelengths.  Diffuse infrared cirrus clouds present simpler physical environments than dense star-forming molecular clouds.  Long wavelength observations ($\\lambda \\geqslant 850\\,\\micron$) will determine the behavior of the polarization spectrum where thermal emission from dust no longer dominates the total spectral energy distribution.  Experiments such as \\emph{Planck} \\citep{tauber04} and the \\emph{Wilkinson Microwave Anisotropy Probe} (\\emph{WMAP}; \\citealt{hinshaw03}) are designed to make all-sky polarization measurements at these wavelengths.  Dust may also be a major contributor to the observed Galactic emission at microwave frequencies ($\\lambda \\sim 3$ -- 300 mm).  This dust emission will be dominated by electric dipole \\citep{dl98b} or magnetic dipole \\citep{dl99} radiation rather than thermal radiation. As a result, Galactic dust emission may be an important foreground contributor to studies of the cosmic microwave background (CMB\\@). Extension of the polarization spectrum of Galactic clouds from the SMM to the microwave will be key to accurately removing this polarized foreground from the CMB\\@. ",
        "conclusions": "Dust grains aligned with magnetic fields in the interstellar medium are responsible for both polarization by absorption at optical wavelengths and polarization by emission at far-infrared and submillimeter wavelengths.  Observations of the absorption and emission trace different regions of the ISM but both are consistent with the existence of radiative torques which spin grains up to the suprathermal angular velocities needed for magnetic alignment.  In molecular clouds the FIR/SMM polarization is strongly dependent on wavelength.  We have attributed this wavelength dependence to sampling different grain populations at different temperatures.  This wavelength dependence provides an opportunity to study, in different physical environments, the grain properties, the alignment mechanism and its efficiency, and magnetic fields. To date, most observations of polarized emission have been in the densest regions of the ISM\\@.  Extending these observations to regions of the diffuse ISM will provide additional tests of grain and alignment models.  Instruments on board the \\emph{Stratospheric Observatory for Infrared Astronomy (SOFIA)} will have the sensitivity to detect the total (HAWC: \\citealt{hawc}) and polarized (Hale: \\citealt{hale}) flux in diffuse infrared cirrus clouds. Instruments, including \\emph{Planck} and \\emph{WMAP}, designed to separate polarized Galactic foregrounds from the CMB will be able to extend the diffuse cloud polarization spectrum to longer SMM and microwave wavelengths.  The polarization state of the excess microwave emission observed at frequencies of 10 -- 90 GHz is key to an accurate measurement of polarization from the cosmic microwave background.  Total power observations are currently unable to conclusively rule out either of the two leading candidates for this emission: electric dipole emission from small spinning grains and magnetic dipole emission from larger vibrating grains.  Polarization observations made with \\emph{WMAP} and \\emph{Planck} will place limits on the emission contributions from these two sources."
    },
    "0609/astro-ph0609802_arXiv.txt": {
        "abstract": "We describe the statistical properties of narrowband drifting auroral kilometric radiation ('striated' AKR) based on observations from the Cluster wideband receiver during 2002-2005. We show that the observed characteristics, including frequency drift rate and direction, narrow bandwidth, observed intensity, and beaming angular sizes are all consistent with triggering by upward traveling ion solitary structures (`ion holes'). We calculate the expected perturbation of a horseshoe electron distribution function by an ion hole by integrating the resonance condition for a cyclotron maser instability (CMI) using the perturbed velocity distribution. We find that the CMI growth rate can be strongly enhanced as the horseshoe velocity distribution contracts inside the passing ion hole, resulting in a power gain increase greater than 100 dB.  The gain curve is sharply peaked just above the R-mode cut-off frequency, with an effective bandwidth $\\leq$50 Hz, consistent with the observed bandwidth of striated AKR emission. Ion holes are observed \\textit{in situ} in the acceleration region moving upward with spatial scales and speeds consistent with the observed bandwidth and slopes of SAKR bursts. Hence, we suggest that SAKR bursts are a remote sensor of ion holes and can be used to determine the frequency of occurrence, locations in the acceleration region, and lifetimes of these structures. ",
        "introduction": "Auroral kilometric radiation (AKR) bursts exhibit a wide variety of fine structure as seen on frequency-time spectra.  The cyclotron maser instability (CMI) \\citep{w79} is widely assumed to be the basic plasma mechanism responsible for the emission. This mechanism originally assumed a loss-cone electron velocity distribution function, but in situ observations in the acceleration region have shown that a 'horseshoe' or crescent distribution is more accurate \\citep{l90,r93,d98,e00}. The horseshoe distribution, which arises naturally for electron beams in the presence of inhomogeneous magnetic fields \\citep{s05}, provides a robust and efficient free energy source for the CMI mechanism, as shown both by model calculations \\citep[e.g.][]{p84,p99,b00} and in laboratory experiments \\citep[e.g.,][]{s05}. This mechanism has been applied not only to terrestrial AKR emission, but to many astrophysical environments in which the requisite conditions (low density, beamed electrons, inhomogeneous magnetic fields) are thought to be present, e.g. planetary magnetospheres \\citep{z98,f04}, stellar magnetospheres \\citep{l86,k02,w04}, and even relativistic jets in active galaxies \\citep{b05}.  Detailed physical models for the rich variety of observed spectral signatures are still an active area of research. AKR fine structure exhibits many morphologies on time-frequency spectra, including slowing drifting or nearly stationary features lasting tens of seconds \\citep{g79}, drifting features which may be interacting with each other \\citep{ptb01}, and periodically modulated or banded emission \\citep{g82}. Although some AKR bursts appear to be broadband emission extending over 100's of kilohertz \\citep{h01}, the total power is dominated by highly time-variable fine structures, especially during periods of enhanced bursts \\citep{mm02}, suggesting that perhaps all AKR radiation is a superposition of narrowband, short duration fine structures.  \\citet{g79} first suggested that drifting AKR fine structure may be due to localized sources rising and falling along auroral magnetic field lines, with emitted frequency equal to the local electron cyclotron frequency. \\citet{ga81} suggested that the emission is triggered by electrostatic waves which drift along the field at the local ion-acoustic speed. An early detailed model which attempts to explain AKR fine structure is the tuned cavity model of \\citet{c82} in which the source region acts as a waveguide with sharp density boundaries.  The source emits radiation in normal modes analogous to an optical laser, resulting in narrowband emission which drifts as the wave packet propagates into regions of varying width. A related idea is that of \\citet{f95} in which the cavity boundary is oscillating quasi-periodically.  Under these conditions, broadband radio photons will be stochastically accelerated, resulting in quasi-monochromatic discrete tones, whose frequencies slowly drift with changes in the cavity geometry.  A difficulty with these models is that they require special conditions (e.g. sharp density boundaries, oscillating walls) which are not supported by observations \\citep{p02}. \\begin{figure*}[t] \\centerline{\\includegraphics[width=6in] {fig1a.eps}} \\caption{Frequency-time spectra of striated AKR bursts observed on 4 Cluster spacecraft on  August 31, 2002 from 16:14:20 - 16:14:35 UT. The inset shows the projected angular separations of the spacecraft as seen from the source.} \\label{fig-dynspec1} \\end{figure*}   Several authors have reproduced AKR fine structure using electromagnetic particle simulations to model the AKR source region. \\citet{mw91} used a one-dimensional electromagnetic particle-in-cell (PIC) code to model the cyclotron maser instability in a plasma with an inhomogeneous magnetic field. They found that radiation is emitted in individual packets which combine to form drifting features that both rise and fall, consistent with some types of AKR fine structure. Both X- and O-mode drifting fine structures are created, consistent with observations \\citep{b88} but not predicted in feedback models such as \\citet{c82}. A weakness of their model is that predicted drift rates are generally much higher than observed drifts.  \\citet{p02} used a 2-dimensional PIC code with model parameters based on in situ FAST spacecraft measurements of the AKR source region.  The simulated AKR is most strongly amplified longitudinally and consists entirely of short-timescale fine structures. The predicted bandwidth ($\\Delta\\omega/\\omega\\sim10^{-3}$) is somewhat larger than at least some observed AKR fine structure.  More recently, Pottelette and colleagues have investigated the possible connection between AKR fine structure and electron holes \\citep{ptb01,ptbj03} and also tri-polar structures \\citep{pt05}. AKR radiation from small packets  ('elementary radiation sources') associated with the holes are thought to be strongly amplified as the packets slow down and are reflected from the field-aligned electric field.  Since the scale size of electron holes is very small (a few Debye lengths, \\citet{b99}), the observed narrow bandwidth of AKR fine structures is easily accounted for. However, since neither the observed speeds or direction of electron holes (1000 - 2500 km s$^{-1}$ always downward, \\citep{b99}) are representative of fine structure frequency drifts, there is not a straightforward relationship between the observed properties of electron holes and the dynamics of the AKR fine structure.  In this paper, we present evidence that a particular type of AKR fine structure called striped or striated AKR \\citep{m96,mppg00} is triggered by ion solitary structures (ion holes). In section 2 we present new observations of SAKR bursts, including new bandwidth and angular beamwidth measurements.  In section 3 we calculate the change in CMI gain for a density depleted region with a pre-existing horseshoe velocity distribution when perturbed by a passing ion hole. We compare the observed properties of SAKR with those calculated by radiation generated from the CMI instability in an ion hole. In section 4 we discuss some inferred properties of ion holes, such as lifetimes, probability of occurrence as a function of location, and radial dependence of speed, that cannot be inferred from single spacecraft in situ measurements. ",
        "conclusions": "This paper summarizes the observed properties of a distinct form of auroral kilometric radiation fine structure called striated AKR, first described by \\citet{m96}. We present new observational results using the WBD instrument on Cluster which characterize the bandwidth and angular beamsize of individual SAKR sources, as well as derived properties such as intrinsic power and speed along the magnetic field.  Assuming the frequency of SAKR bursts can be identified with the local electron cyclotron frequency, the speed and direction of SAKR sources calculated from their observed drift rates are very similar to those observed for ion solitary structures (ion holes) in the upward current region of the magnetosphere. Hence, we investigated whether SAKR bursts could be the result of enhancement of the cyclotron maser instability by ion holes. Using observed electric field signatures of ion holes in this region, we calculate the perturbation caused by the passage of an ion hole on a 'horseshoe' electron velocity distribution in a dilute ($n_e<<1$ cm$^{-3}$) plasma. The cyclotron maser instability is strongly enhanced inside the ion hole, with power gain exceeding 100 dB in a narrow frequency range just above the x-mode cutoff frequency. These characteristics are in excellent agreement with the observed bandwidth, speed, direction, and flux density of SAKR bursts. Alternative suggestions involving other types of solitary structures to explain AKR fine structure, such as electron holes \\citep{ptb01} or tri-polar structures \\citep{pt05}, are not consistent with SAKR properties, although they may be important in other types of AKR fine structures.  If SAKR bursts are in fact triggered by ion holes, a number of derived properties of ion holes can be deduced which would be difficult or impossible to obtain using {\\it in situ} measurements. These include average lifetimes (a few seconds), evolution of propagation speed (nearly constant over lifetime of hole), and relative numbers versus location (much more common high in the acceleration region than near the base)."
    },
    "0609/astro-ph0609315_arXiv.txt": {
        "abstract": "Regionally averaged relativistic cosmologies have recently been considered as a possible explanation for the apparent late time acceleration of the Universe. This contribution reports on a mean field description of the backreaction in terms of a minimally coupled regionally homogeneous scalar field evolving in a potential, then giving a physical origin to the various phenomenological scalar fields generically called quintessence fields. As an example, the correspondence is then applied to scaling solutions. ",
        "introduction": "Modern cosmology is nowadays settled on observations concerning mainly the distribution of matter and the dynamics of the expansion of the Universe. On the one hand, there are now various cosmological observations supporting a matter distribution that is homogeneous on large scales of order $100$ h$^-1$Mpc. Nevertheless, at late times, the matter distribution is highly structured on smaller scales, with the presence of clusters of galaxies, filaments and voids. Moreover, the statistical isotropy of the Cosmic Microwave Background radiation supports the idea that the Universe is highly isotropic on average, on large scales. Facing these observational issues, one assumes in the standard cosmological framework that the Universe is homogeneous and isotropic on all scales, resulting in a spacetime described by a Friedmann-Robertson-Walker (FRW) metric, the inhomogeneities being perturbations around this homogeneous and isotropic background. Then, all the observables of the Universe on large scales can be deduced from a single degree of freedom: the scale factor of the metric, and the dynamics of the inhomogeneities is well described as long as the density contrast in the matter fields remains small.  On the other hand, many recent observations strongly favor a Universe whose expansion has been accelerating in the recent past and may be accelerating today. In the FRW context this necessarily requires the introduction of exotic sources as for example a cosmological constant or quintessence fields, or a modification of gravity, generally implying the so-called coincidence problem: why is the expansion accelerating approximatelly at the same time when the Universe becomes structured, that is when the density contrast in the matter field is no longer small on a wide range of scales?  Regionally averaged relativistic cosmologies may be able to answer this question by linking the dynamics of the Universe on large scales to its structuration on smaller scales; see interesting discussions of that topic in R\\\"as\\\"anen (2006). It consists in defining cosmologies that are homogeneous on large scales without supposing any local symmetry, thanks to a spatial averaging procedure. It results in equations for a volume scale factor that not only include an averaged matter source term, but also additional terms that can be interpreted as the effects of the coarse-grained inhomogeneities on the large scales dynamics. These additional terms are commonly named backreaction.  In this paper, after introducing the formalism of regionally averaged cosmologies in the first part, we shall propose a correspondence between regionally averaged cosmologies and Friedmannian scalar field cosmologies in the second part, the scalar field being interpreted in this context as a mean field description of the inhomogeneous Universe, that can play the role of a quintessence field. Then, in the third part, as an example of the correspondence, we explicitly reconstruct the mean field theory for the particular class of scaling solutions of the regionally averaged cosmologies, and discuss its properties. This correspondence has been proposed and discussed in Buchert et al. (2006). ",
        "conclusions": "The mean field description of backreaction effects through a scalar field does not only provide a rephrasing of the kinematics of backreaction, but it also justifies the existence of the cosmological effective scalar field, that may be responsible for the cosmic acceleration, on the basis of an underlying fundamental theory, that is Einstein General Relativity: the cosmic quintessence emerges in the process of interpreting the real Universe in a homogeneous context. The study of scaling solutions allowed to understand that the cosmic acceleration is only possible if the mean curvature is strongly coupled to the backreaction and converts into it to maintain it at a high level. Nevertheless, more realistic solutions, with a varying conversion rate must be investigated. Finally, in order to firmly establish that backreaction effects are the source of the acceleration of the expansion, it will be necessary to explicitly compute these effects from generic relativistic models and observations of the large-scale structures of the Universe."
    },
    "0609/astro-ph0609409_arXiv.txt": {
        "abstract": "We present the analysis on the quadrupole phases of the Internal Linear Combination map, ILC(I) and (III) derived by the \\wmap team (1 and 3-year data release). This approach allows us to see the global trend of non-Gaussianity of the quadrupoles for the ILC(III) map through phase correlations with the foregrounds.  Significant phase correlations is found in between the ILC(III) quadrupole and the \\wmap foregrounds phases for K-W band: the phases of the ILC(III) quadrupole $\\xi_{2,1}$, $\\xi_{2,2}$ and those of the foregrounds at K-W bands $\\Phi_{2,1}$, $\\Phi_{2,2}$ display significant symmetry :  $\\xi_{2,1}+\\Phi_{2,1} \\simeq \\xi_{2,2}+\\Phi_{2,2}$, which is a strong indication that the morphology of the ILC(III) quadrupole is mere reflection of that the foreground quadrupole through coupling. To clarify this issue we exploit the symmetry of the CMB power, which is invariant under permutation of the index $m=1\\Leftrightarrow 2$. By simple rotation of the ILC(III) phases with the same angle we reach the phases of foreground quadrupole. We discuss possible sources of phase correlation and come to the conclusion that the phases of the ILC(III) quadrupole reflect most likely systematic effects such as changing of the gain factor for the 3-year data release with respect to the 1-year, rather than manifestation of the primordial non-Gaussianity. ",
        "introduction": "After the release of the Wilkinson Microwave Anisotropy Probe (\\wmap) 1-year results (hereafter WMAP\\,I) Refs.~\\refcite{wmapresults},    % \\refcite{wmapfg},               % \\refcite{wmapsys},              % \\refcite{wmapcl}                % the issue of non-Gaussianity of the CMB has attracted great attention. In the papers Refs.~\\refcite{wmaptacng},      % \\refcite{park},                 % \\refcite{gonzalez},             % \\refcite{eriksen},              % \\refcite{hansen},               % \\refcite{larson},               % \\refcite{land},                 % \\refcite{roukema},              % \\refcite{vectors},              % \\refcite{chenalfven},           % \\refcite{McEwen},               % various kinds of methods have been employed and departure of Gaussianity have been detected. These features can be of primordial origin (Refs.~\\refcite{eriksen},       % \\refcite{jaffe_2005})           % or they could be related to foreground residuals (Refs.~\\refcite{ndv04},         % \\refcite{autox},                % \\refcite{faraday})              % or any sort of systematic effects (\\refcite{hansen}).             %  Recently the \\wmap team released the 3-year data (hereafter WMAP\\,III) Ref.~\\refcite{wmap3ycos}        % producing Internal  Linear Combination map (ILC(III)) available for scientific analysis at the multipole range $\\l\\le 10$. The \\wmap team performed the analysis of Gaussianity for the ILC(III) signal that does not include harmonics $\\l \\le 10$ and showed that the CMB signal is Gaussian Ref.~\\refcite{wmap3ytem}.       % However, as it was cautioned by the \\wmap team, the alignment and planarity at multipoles  $\\l=2, 3$ and 5 mentioned in Refs.~\\refcite{eriksen},        % \\refcite{hansen},               % \\refcite{land},                 % \\refcite{roukema},              % \\refcite{schwarz}               % still present in the map (see recent investigation in Ref.~\\refcite{vectors1}).       %  If breaking of statistical isotropy in the low multipoles is of primordial origin, a very fundamental issue is consequently raised: what cosmological model can provide such a peculiar structure in the CMB angular distribution on the sky~? In Refs.~\\refcite{eriksen},     % \\refcite{jaffe_2006} and        % \\refcite{bridges}               % it has been pointed out that breaking of statistical isotropy in the low multipoles can be explained and in favor of the Bianchi VIIh model. However, Helling, Shupp and Tesilleanu (Ref.~\\refcite{Helling}) % have also pointed out that $\\l=2, 3$ and 5 have anti-alignment in the direction of non-cosmological dipole and argued that accurate technique of the CMB signal subtraction probably needs to be non-linear (with respect to dipole treatment) to prevent any residuals from the dipole in the map.  In addition, methods of foreground separation for the CMB signal very likely leave foreground residuals in the CMB map, a problem that can be illustrated by cross-correlation between the derived signal and the foregrounds. For the WMAP\\,I data this problem was already investigated in Refs.~\\refcite{toh},        % \\refcite{ndv03},            % \\refcite{ndv04},            % \\refcite{autox},            % \\refcite{n2},               % \\refcite{bielewicz}.        %  For WMAP\\,III data, it has been shown in Refs.~\\refcite{chiang_t},   % \\refcite{autox}             % that the low multipoles in ILC(III) clearly display significant cross-correlation with the derived foregrounds, while Ref.~\\refcite{cruz06}       % showed that negative peak at $b=-57^{\\circ}, l=209^{\\circ}$) has no analogue in the foreground map. In Refs.~\\refcite{freeman},    % \\refcite{larson}            % analyzing the \\wmap map-making algorithm (MMA) it has been pointed out that it could produce some residuals from the non-cosmological dipole (related with the motion of our Galaxy in the Local Group) and any non-Gaussian features can ``naturally'' arise as the result of the MMA. Summarizing present status of non-Gaussianity and statistical anisotropy of the CMB for low multipoles $l\\le 10$, we may say that the origin of these peculiarities still is  uncertain.  In this paper we go back to the question raised in Refs.~\\refcite{vectors}      % and \\refcite{blandford},         % if the \\wmap quadrupole is cosmological~? To answer this question, we will implement the analysis of the ILC(I), ILC(III) and de Oliveira-Costa \\& Tegmark (Ref.~\\refcite{doc_teg})  % quadrupole and compare their phases with those of the WMAP\\,III foregrounds\\footnote{{\\tt http://lambda.gsfc.nasa.gov/product/map/m\\_products.cfm}}: combined maps from the synchrotron, free-free, dust emission at K-W band. We will show that combination of the ILC(III) phases $\\xi_{2,m}$ and the foreground phases $\\Phi_{2,m}$ follow the intriguing relation $\\xi_{2,m}+\\Phi_{2,m}\\simeq const $ with accuracy, e.g. within $0.014$ (rad) for the foreground at K band. Since phases of the signal is closely related to morphology of the maps (Ref.~\\refcite{morphology}),                           % detected correlation between the ILC(III) and foreground phases allow us to conclude that it is most likely due to systematic effects (such as estimation of the gain factor and foreground separation) rather then primordial one. ",
        "conclusions": "We have presented the phase analysis of the ILC(I), ILC(III) and de Oliveira-Costa and Tegmark (Ref.~\\refcite{doc_teg}) quadrupoles in comparison with the foreground phases. We have shown that the ILC(III) quadrupole has strong correlation with the foreground phases, mainly with the synchrotron emission per K and KA band foregrounds. We have checked out the possibility that the WMAP third year data release quadrupole can be explained by implementation of the Biachi VIIh model discussed in Refs.~\\refcite{jaffe_2006a} and \\refcite{bridges}. All these models need additional corrections of the phases for $a_{2,1}$ and $a_{2,2}$ components in order to math the phases of the ILC(III). By analysis of the de Oliveira-Costa and Tegmark (Ref.~\\refcite{doc_teg}) quadrupoles  related with a different Galactic mask we can conclude that detected above correlations between the ILC and foregrounds weakly depend on the type of the mask, if the area covered by them is small in comparison with the $4\\pi$. We have shown that transition from the ILC(I) quadrupole to the ILC(III) changes significantly the properties of the CMB-foreground phase correlations through renormalization of the gain factor. We believe that this is the main reason for discussed above phase correlation, clearly demonstrated importance of accurate removal of the systematic effects."
    },
    "0609/astro-ph0609065_arXiv.txt": {
        "abstract": "The metallicity distribution in the Fornax cluster was studied with the XIS instrument onboard the Suzaku satellite. K-shell lines of O and Mg were resolved clearly, and the abundances of O, Mg, Si, S and Fe were measured with good accuracy.  The region within a  $4'$ radius of NGC 1399 shows approximately solar abundances of Fe, Si and S, while the O/Fe and Mg/Fe abundance ratios are about 0.4--0.5 and 0.7 in solar units.  In the outer region spanning radii between $6'$ and $23'$, the Fe and Si abundances drop to 0.4--0.5 solar and show no significant gradient within this region. The abundance ratios, O/Fe and Mg/Fe, are consistent with those in the central region.  We also measured the Fe abundance around NGC 1404 to be approximately solar, and the O, Ne and Mg abundances to be 0.5--0.7 times the Fe level.  The significant relative enhancement of Fe within 130 kpc of NGC 1399 and in NGC 1404 indicates  an origin in SN Ia, in contrast to the species O, Ne, and Mg which reflect the stellar metallicity. The mass-to-light ratios for O and Fe within 130 kpc of NGC 1399 are over an order of magnitude lower than those in rich clusters, reflecting the metal enrichment history of this poor cluster. ",
        "introduction": "The metal abundances in the intracluster medium (ICM) and the hot X-ray emitting interstellar medium (ISM) of early-type galaxies  provide important clues to understand the metal enrichment history and evolution of galaxies. The ICM contains a large amount of metals, which are mainly synthesized by supernova (SN) in early-type galaxies (e.g., Arnaud et al.\\ 1992; Renzini et al.\\ 1993).  The ASCA satellite (Tanaka et al.\\ 1994) first enabled us to measure the distribution of Fe in the ICM\\@ (e.g., Fukazawa et al. 2000; Finoguenov et al. 2000; 2001).  The derived iron-mass-to-light ratio (IMLR) is nearly constant in rich clusters and decreases toward poorer systems (Makishima et al.\\ 2001).  In individual clusters, the IMLR is lower around the center (Makishima et al.\\ 2001).  The Si abundance of the ICM was also measured by ASCA (Fukazawa et al.\\ 1998;2000; Finoguenov et al.\\ 2000; 2001), however, since Fe and Si are both synthesized in SN Ia and SN II, abundance results for these species cannot clearly constrain the nucleosynthesis contribution from different SN types.  In contrast, O and Mg are predominantly synthesized in SNe II. Abundance measurements spanning the range of species from O to Fe are therefore needed to obtain unambiguous information on the formation history of massive stars.  XMM-Newton provided the means to measure O and Mg abundances in some systems, but reliable results have been obtained only for the central regions of very bright clusters or groups of galaxies dominated by cD galaxies (e.g., Xu et al.\\ 2002, Finoguenov et al.\\ 2002, Tamura et al.\\ 2003, Matsushita et al.\\ 2003, 2006; Buote et al (2003)).  In general, the measured abundances of Si and Fe are similar, while the O abundance  is only about half  of the Fe abundance. This indicates that cD galaxies are an important source of Si and Fe which are mainly synthesized in SN Ia. However, our knowledge about the distributions of O and Mg in the ICM is still poor, in contrast to the detailed measurement of the Fe and Si abundances with XMM-Newton (e.g., Tamura et al.\\ 2004).  Regarding early-type galaxies, the metal abundances in the ISM give us important information about the present metal supply into the ICM through SN Ia and stellar mass loss.  In addition, O and Mg abundances should reflect the stellar metallicity and enable us to directly look into the formation history of these galaxies.  In this paper, based on Suzaku (Mitsuda et al.\\ 2006) observations, the abundances of O, Mg, Si, S and Fe of the ICM in the Fornax cluster and those of O, Ne, Mg and Fe of the ISM in NGC 1404 are discussed. The XIS (Koyama et al.\\ 2006) instrument onboard Suzaku has good energy resolution at the O line energy with low background, providing better sensitivity than the EMOS onboard XMM-Newton. The EMOS also has a problem in measuring the Mg abundance in somewhat fainter systems, due to a strong instrumental Al line.  The Mg lines are particularly useful in cluster outskirts, since the strong Galactic O line causes difficulty in the measurement of O emission from nearby clusters.   The Fornax cluster is a nearby poor cluster with an ICM temperature of 1.3-1.5 keV (e.g., Scharf et al.\\ 2005). The X-ray emission shows asymmetric spatial distribution, and the cD galaxy, NGC 1399, is offset from the center (Paolillo et al.\\ 2002; Sharf et al.\\ 2005), which may be related to large scale dynamical evolution such as infall motions of galaxies into the cluster (Dunn \\& Jerjen 2006).  Chandra observations suggests that there may be relative motion between NGC 1399 and the ICM, and the 2nd brightest elliptical galaxy, NGC 1404, is moving supersonically in the ICM (Scharf et al.\\ 2005; Machacek et al.\\ 2005). For the ICM abundances, the Fe and Si abundances within $\\sim 50$ kpc of NGC 1399 were measured with XMM-Newton (Buote et al.\\ 2002). However, due to the high background of XMM-Newton, the O and Mg abundances of the ICM were not determined.  ASCA detected excess hard X-ray emission from several groups of galaxies including the Fornax cluster (Fukazawa et al.\\ 2001; Nakazawa 2001). Dynamical motions in the cluster, driven by merger events, that generate a population of relativistic particles may be the origin of this emission component. Suzaku results on hard emission from Fornax will be presented in a separate article after a detailed study of detector background.   We adopt   solar abundances by Feldman (1992), where the solar O and Fe abundances relative to H are $8.51\\times 10^{-4}$ and  $3.24 \\times10^{-5}$ by number. Recently, the solar photospheric abundances of C, N, O, and Ne decreased by 0.2 dex, considering three-dimensional hydrostatic model atmospheres and non-local thermodynamic equilibrium (Asplund 2005 and references therein). The new solar O and Fe abundances relative to H are 4.90$\\times 10^{-4}$ and $2.95\\times 10^{-5}$, respectively (Lodders 2003). The effect of the new solar abundances will be discussed  in Section 4.    We adopt 19 Mpc (for H$_0$ = 72 km/s/Mpc) for the distance to the Fornax cluster.  The observation of Cepheids with  Hubble Space Telescope yielded the distance of 18.6$\\pm 0.6$ Mpc (Madore et al. 1999). Unless otherwise specified, errors are quoted at 90\\% confidence ",
        "conclusions": "With Suzaku observations of the Fornax cluster, the abundances of O, Mg, Si, S and Fe of the ICM up to 130kpc north and up to 60kpc south of NGC 1399, and O, Ne, Mg and Fe abundances of NGC 1404 were derived accurately. The Fe abundances around NGC 1399 and NGC 1404 are about solar and drop to 0.5 solar in the ICM. The elemental abundance ratios show common values around NGC 1399, NGC 1404, and the ICM: Si and S have similar abundances to Fe and O/Fe, Ne/Fe, Mg/Fe ratios are 0.5--0.7 solar in units of the solar ratio. The O, Ne and Mg abundances of NGC 1404 and the NGC 1399 region are consistent with the stellar metallicity of these galaxies. Most of the Fe around NGC 1399, NGC 1404 and the ICM should have been synthesized by SN Ia. The abundance ratios of O/Fe and Si/Fe in the gas are consistent with a mixture of SN Ia ejecta of the W7 model and metal-poor Galactic stars. Values for the IMLR and OMLR are much smaller than in rich clusters, which indicates that most of the metals may be outside of the region we observed. The metal distribution of the ICM may be used as a tracer of the history of the cluster.  \\bigskip  We would like to thank Prof. R. Kelley for careful internal review of the manuscript. The authors would like to thank ISAS/JAXA and the entire Suzaku team for the opportunity to participate in the development of the Suzaku mission as members of its Science Working Group. This work is supprted by Grant-in-Aid for Science Research of JSPS. JPH acknowledges support from NASA  grant NNG05GP87G."
    },
    "0609/astro-ph0609586_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{intro} A systematic survey of the inner part of the Galaxy performed by the H.E.S.S.\\ Cherenkov telescope system has revealed a number of previously unknown sources of VHE gamma-rays above 100~GeV~\\cite{HESSScanA}~\\cite{HESSScanB}. While in terms of a population approach the sources can be described by common properties like generally rather hard energy spectra (photon index $\\sim$2.3) or a rather narrow distribution in Galactic latitude (rms of $\\sim 0.3^{\\circ}$) the counterpart identification calls for an individual study of these objects. \\begin{table}[t] \\caption{Categories into which the gamma-ray sources will be classified in the following sections.} \\centering \\label{tab:1}       % \\begin{tabular}{c | c c c } \\hline\\noalign{\\smallskip} & Matching position/ & Viable emission & Consistent \\\\ & morphology & mechanism & MWL picture \\\\[3pt] \\tableheadseprule\\noalign{\\smallskip} A & yes & yes & yes \\\\ B & no & yes & yes \\\\ C & yes & yes & no \\\\ D & no & no & no \\\\ \\noalign{\\smallskip}\\hline \\end{tabular} \\end{table} An unambiguous counterpart identification of these (initially) unidentified H.E.S.S. sources requires {\\bf (i) spatial and ideally also morphological coincidence}, {\\bf(ii) a viable gamma-ray emission mechanism} for the object, and {\\bf (iii) a consistent multi-wavelength behaviour} matching the suggested identification and the particle distribution within the source. The H.E.S.S.\\ sources can be classified according to their confidence in identification with known astrophysical objects following the three requirements given above. Table~\\ref{tab:1} summarises the categories. {\\bf Category~A} comprises sources for which the positional and/or morphological match (in case of an extended source) with a counterpart source is excellent and the emission processes can be modelled to provide a consistent picture describing the multi-frequency data. For these sources the association is beyond doubt. For {\\bf Category~B} sources the emission mechanisms can be consistently modelled, however these sources show a less convincing positional and/or morphological match with the potential counterpart. {\\bf Category C} sources on the other hand have a good positional counterpart, they show however a non-consistent multi-wavelength picture, being it because of insufficient data at other wavebands, being it because of a not fully understood emission mechanism. For {\\bf Category D} sources no counterpart candidate exists, these are the classical {\\emph {unidentified sources}}. In the following I will describe examples for sources belonging to each of the 4 categories. The description will focus on Galactic gamma-ray sources, since for extragalactic objects the counterpart identification in the VHE gamma-ray regime has (so far) turned out to be rather unproblematic. ",
        "conclusions": "Table~\\ref{tab:2} summarises the H.E.S.S.\\ gamma-ray sources into the proposed categories as given in table~\\ref{tab:1}. As can be seen from this table, for less than half a firm counterpart can be identified, putting these sources into class~A. The majority of the sources has to be classified as unidentified, since no good counterpart at other wavebands exists at all. It is evident, that the gathering of multi-waveband data will help in identifying these objects. It should however also be noted, that a positional coincidence alone does not help. A consistent modelling of the emission mechanisms at work in generating the gamma-rays must be invoked, before a firm association can be established. For the upcoming GLAST satellite, the situation will be further complicated by source confusion due to the larger point-spread function and also by the stronger diffuse gamma-ray background from decays of neutral pions in the Galactic plane. \\begin{table}[t] \\caption{Summary of the classes for the H.E.S.S.\\ Galactic gamma-ray sources.}  \\centering \\label{tab:2}       % \\begin{tabular}{l | c | c } \\hline\\noalign{\\smallskip} Source & Class & Comment \\\\[3pt] \\tableheadseprule\\noalign{\\smallskip} J1713--397 & A & SNR RX\\,J1713.7--3946\\\\ J0852--463 & A & SNR RX\\,J0852.0--4622\\\\ J0835--455 & A & PWN Vela\\,X \\\\ J1302--638 & A & PWN PSR\\,B1259--63 \\\\ J1420--607 & A & PWN PSR\\,J1420--6048 \\\\ J1514--591 & A & PWN MSH\\,15--5\\emph{2} \\\\ J1747--281 & A & PWN G0.9+0.1 \\\\ J1804--216 & B & PWN or SNR? \\\\ J1825--137 & B & PWN different scale \\\\ J1640--465 & C & SNR? \\\\ J1813--178 & C & Composite SNR? \\\\ J1834--087 & C & SNR? \\\\ J1303--631 & D & \\\\ J1614--518 & D & \\\\ J1632--478 & D & \\\\ J1634--472 & D & \\\\ J1702--410 & D & \\\\ J1708--420 & D & \\\\ J1745--290 & D & Galactic Centre source\\\\ J1745--303 & D & \\\\ J1837--069 & D & \\\\ \\noalign{\\smallskip}\\hline \\end{tabular} \\end{table}  The support of the Namibian authorities and of the University of Namibia in facilitating the construction and operation of H.E.S.S.\\ is gratefully acknowledged, as is the support by the German Ministry for Education and Research (BMBF), the Max Planck Society, the French Ministry for Research, the CNRS-IN2P3 and the Astroparticle Interdisciplinary Programme of the CNRS, the U.K. Particle Physics and Astronomy Research Council (PPARC), the IPNP of the Charles University, the South African Department of Science and Technology and National Research Foundation, and by the University of Namibia. We appreciate the excellent work of the technical support staff in Berlin, Durham, Hamburg, Heidelberg, Palaiseau, Paris, Saclay, and in Namibia in the construction and operation of the equipment."
    },
    "0609/gr-qc0609022_arXiv.txt": {
        "abstract": "When a black hole is in an empty space on which there is no matter field except that of the Hawking radiation (Hawking field), then the black hole evaporates and the entropy of the black hole decreases. The generalised second law guarantees the increase of the total entropy of the whole system which consists of the black hole and the Hawking field. That is, the increase of the entropy of the Hawking field is faster than the decrease of the black hole entropy. In naive sense, one may expect that the entropy increase of the Hawking field is due to the self-interaction among the composite particles of the Hawking field, and that the {\\it self}-relaxation of the Hawking field results in the entropy increase. Then, when one consider a non-self-interacting matter field as the Hawking field, it is obvious that the self-relaxation does not take place, and one may think that the total entropy does not increase. However, using nonequilibrium thermodynamics which has been developed recently, we find for the non-self-interacting Hawking field that the rate of entropy increase of the Hawking field (the entropy emission rate by the black hole) grows faster than the rate of entropy decrease of the black hole along the black hole evaporation in the empty space. The origin of the entropy increase of the Hawking field is the increase of the black hole temperature. Hence an understanding of the generalised second law in the context of the nonequilibrium thermodynamics is suggested; even if the self-relaxation of the Hawking field does not take place, the temperature increase of the black hole during the evaporation process causes the entropy increase of the Hawking field to result in the increase of the total entropy. ",
        "introduction": "\\label{sec-intro}  When a black hole is in an empty space on which there is no matter field except that of the Hawking radiation (Hawking field) \\cite{ref-hr}, then an outgoing energy flow from the black hole into the empty space arises due to the Hawking radiation and the black hole loses its mass energy. This denotes that a black hole evaporates. Concerning the black hole evaporation in the empty space, the generalised second law (GSL) of the black hole thermodynamics is an important and interesting issue \\cite{ref-gsl} \\cite{ref-bht}. The GSL conjectures that the entropy of a black hole $S_g$ is given by the horizon area (divided by 4) and that the total entropy of the whole system, which consists of the black hole and the Hawking field, increases during the black hole evaporation process. Note that, when the GSL holds, the black hole entropy $S_g$ decreases during the evaporation process, but the entropy of the Hawking field increases faster than the decrease of $S_g$.  By the way, the outgoing energy flow from the black hole into the empty space denotes that the Hawking field is in a nonequilibrium state, since no energy flow arises in an equilibrium state. Therefore the black hole evaporation in the empty space should be interpreted as a nonequilibrium process of the whole system. In order to prove the GSL in the context of the black hole evaporation in the empty space, we have to define a nonequilibrium entropy of the Hawking field.  However, because a general definition of the nonequilibrium entropy of matters has not yet been formulated, the existing proofs of the GSL consider an equilibrium between a black hole and a heat bath surrounding it, or assume the existence of a well-defined nonequilibrium entropy of an arbitrary matter \\cite{ref-gsl.2}. Here note that, in the existing proofs, a self-interacting matter field has been considered as the Hawking field. For the black hole evaporation in the empty space, the self-interaction causes the {\\it self}-relaxation of the Hawking field itself during propagating on the empty space to increase the entropy of the Hawking field (if a nonequilibrium entropy were to be defined). Then one may expect naively that the self-relaxation of the Hawking field is the origin of the increase of the total entropy of the whole system. This seems to be a very reasonable picture of the GSL in the context of the black hole evaporation in the empty space. However, I dare to suggest an understanding of the GSL in the context of the nonequilibrium thermodynamics as follows. {\\it The self-interaction of the Hawking field is not necessary, but the increase of the black hole temperature during the evaporation process plays an essential role for the validity of the GSL. That is, due to the temperature increase, the rate of entropy increase of the Hawking field (the entropy emission rate by the black hole) grows faster than the rate of entropy decrease of the black hole, then the total entropy increases without the self-interaction of the Hawking field.} Note that the increase of the black hole temperature along the loss of mass energy is due to the negative heat capacity of the black hole. The negative heat capacity is peculiar to self-gravitating systems \\cite{ref-sgs}. Therefore our understanding of the GSL depends strongly on the feature of the self-gravitation of the black hole.  In order to explain our understanding of the GSL, it is appropriate to consider a non-self-interacting matter field as the Hawking field. For example, photon, graviton, and a massless free Klein-Gordon field ($\\square \\Phi = 0$) can be the non-self-interacting Hawking fields. Recently, in the study on nonequilibrium thermodynamics of a non-self-interacting matter field, it has been reported that a nonequilibrium thermodynamics for a radiation field (photon gas) was formulated in a consistent way \\cite{ref-sst}. With referring to report \\cite{ref-sst}, this paper aims to calculate the time evolution of the total entropy of a black hole and the non-self-interacting Hawking field in a simplified model of the black hole evaporation in the empty space, and to give an understanding of the GSL as explained above. Further, since report \\cite{ref-sst} treats the photons which are collisionless and massless particles, we consider that the Hawking field is not only non-self-interacting but also massless.  This paper is organised as follows. The strategy to reach our understanding of the GSL is explained in section \\ref{sec-strategy}. In section \\ref{sec-sst}, the nonequilibrium entropy obtained in \\cite{ref-sst} is briefly explained. Section \\ref{sec-entropy} is devoted to the application of the nonequilibrium entropy to our strategy. Summary and discussions are in section \\ref{sec-sd}.  We set the fundamental constants unity, $\\hbar = G = c = k_B = 1$. Then the Stefan-Boltzmann constant is $\\sigma = \\pi^2/60$, which is appropriate for a photon gas. When one consider a general non-self-interacting massless field as the Hawking field, replace the Stefan-Boltzmann constant as $\\sigma \\to \\sigma^{\\prime} = N\\, \\pi^2/120 = N\\, \\sigma/2$, where $N = n_b + ( 7/8 )\\, n_f$, and $n_b$ ($n_b = 2$ for photons) is the number of helicities of massless bosonic field and $n_f$ is that of massless fermionic field. ",
        "conclusions": "\\label{sec-sd}  When a black hole evaporates due to the Hawking radiation, the thermodynamic state of the Hawking field is in a nonequilibrium state. Then, referring to a recent development on the steady state thermodynamics for a radiation field, we have analysed explicitly the time evolution of a black hole evaporation process within the framework of the NHR model, and obtained the conclusion that the total entropy $S_{tot}^{(min)}$ increases monotonically for $( \\lambda = )\\, t_{ev}/R_0 \\gg O(1)$, where $t_{ev}$ is the evaporation time and $R_0$ is the initial black hole radius. Therefore, as discussed in section \\ref{sec-strategy}, we can obtain the understanding of the GSL in the context of the steady state thermodynamics as follows. The increase of the black hole temperature due to its negative heat capacity gives the increase of the entropy of the Hawking fields, and the rate of entropy increase of the Hawking fields is faster than the rate of entropy decrease of the black hole. That is, the self-interaction of the Hawking fields is not necessary for the validity of the GSL.  Here we pay attention to the last stage of the evaporation process in the NHR model; the total entropy $S_{tot}$ increases in an accelerating fashion as shown in figure \\ref{pic-total}. According to the relaxation process in the MS model which reaches a total equilibrium state as mentioned in section \\ref{sec-sst} or shown in report \\cite{ref-sst} in detail, one may think the accelerated increase of $S_{tot}^{(min)}$ is strange, because the black hole evaporation process in the framework of the NHR model is also described as the relaxation process of the whole system. However the accelerated increase of $S_{tot}^{(min)}$ never contradict the discussions in section \\ref{sec-sst} and in report \\cite{ref-sst}, because the black hole has the negative heat capacity as shown by equation \\eqref{eq-strategy.capa}. The negative heat capacity is a peculiar property to the self-gravitating systems \\cite{ref-sgs}. While the positive heat capacity of an ordinary system makes the entropy emission rate decrease to zero as the relaxation process proceeds, however the negative heat capacity of the black hole makes the entropy emission rate increase as the black hole evaporation process proceeds as shown by the right panel in figure \\ref{pic-dens}. This increase of the entropy emission rate is the essence of our understanding of the GSL.  Here it is recalled that the negative heat capacity of a self-gravitating gas system (not of a black hole) causes the accelerated increase of its temperature, which is called the {\\it gravothermal catastrophe} \\cite{ref-sgs}. The gravothermal catastrophe is the phenomenon suggested by the ordinary equilibrium thermodynamics and statistical mechanics. Hence the accelerated increase of the total entropy of the black hole and the Hawking fields shown in figure \\ref{pic-total} may be interpreted as a nonequilibrium version of the gravothermal catastrophe.  In this paper, the monotone increasing nature of the total entropy $S_{tot}^{(min)}$ is concluded by the numerical results shown in figure \\ref{pic-total}. Therefore, correctly speaking, our conclusion is not analytically completely supported at present. If the inequality $dS_{tot}^{(min)}/dt > 0$ will be obtained analytically, our understanding of the GSL will have more rigid mathematical reasoning. Further our discussion given in this paper may be understood as a new proof of the GSL, $dS_{tot} > 0$, according to inequality \\eqref{eq-strategy.ineq}.  We turn our discussion to the validity of the assumption $\\lambda = t_{ev}/t_{bh} \\gg O(1)$, which guarantees the validity of the quasi-equilibrium approximation of the black hole. In semi-classical treatment, the black hole temperature is less than the Planck temperature. Then the exotic particles are not expected to be created in the Hawking radiation, and only the ordinary well-known standard particles will be created. Hence we can expect that the number of species of the Hawking fields ($\\sim N$) is less than $108573 = O(10^5)$ (see inequality \\eqref{eq-entropy.range.slow} ). For the case $N \\sim O(10^5)$, the assumption $\\lambda \\gg O(1)$ is valid for a black hole of initial radius $R_0 \\gtrsim O(10)$ in Planck unit. And if we estimate the order of $N$ by the ordinary standard particles (three generations of quarks and leptons and their inner states, gauge particles of four fundamental interactions and their helicities), then we find $N \\sim O(10)$ and the assumption $\\lambda \\gg O(1)$ is valid for a black hole of initial radius $R_0 \\gtrsim O(10^{-1})$ in Planck unit. Hence, when no exotic particle appears in the Hawking radiation, our discussion based on the NHR model seems to be valid for a semi-classical black hole of initial radius $R_0 > 1$ until one Planck time before the evaporation time $t_{ev}$.  Here we notice that, for the case $N \\sim O(10)$, the order relation $\\lambda \\gtrsim O(10^3)$ corresponds to the black hole of initial radius $R_0 \\gtrsim O(10)$. Therefore we find from figure \\ref{pic-total} that, for the black hole of $R_0 \\gtrsim O(10) \\, (\\Rightarrow \\lambda \\gtrsim O(10^3) )$, the normalised total entropy $\\Sigma_{tot}^{(min)}(\\tau)$ evolves in almost the same increasing fashion and reaches almost the same final value $\\simeq 1.33$. Further it should be noted here that, even if we consider the very last stage of the evaporation process where the quasi-equilibrium approximation of the black hole breaks down and the evaporation may proceed to a quantum gravitational process, it is reasonable to expect that the total entropy after the quantum evaporation process finished should be greater than the total entropy before the evaporation process proceeds to the quantum one, since the well-defined entropy has to be a non-decreasing quantity. Hence, because inequality \\eqref{eq-strategy.ineq} indicates that $S_{tot}^{(min)}$ is the lower bound of the total entropy of the whole system, it is suggested that, when the black hole evaporates in the empty space with emitting any self-interacting Hawking field, the final value of the total entropy $S_{tot}$ (not the lower bound $S_{tot}^{(min)}$) should be larger than $1.33 \\times S_g(0)$.  Finally we discuss about the strategy to reach the understanding of the GSL in the context of the steady state thermodynamics. The strategy explained in section \\ref{sec-strategy} has two key assumptions remained to be proven, the additivity of the entropy \\eqref{eq-strategy.total.general} and the existence of a well-defined nonequilibrium entropy $S_m$ of an arbitrary self-interacting Hawking field. However, these two key assumptions seem to be reasonable as are discussed in the following two paragraphs.  On the validity of the additivity \\eqref{eq-strategy.total.general}, we refer the present status of the study on the laboratory system which has self-interactions but not self- and external-gravitational interactions. For example, consider that the system is in a nonequilibrium state which is not so far from an equilibrium state that the heat flux in that system is not extremely strong. Then the so-called extended irreversible thermodynamics \\cite{ref-eit} gives a well-defined nonequilibrium entropy up to the second order in the expansion by the heat flux. Further the nonequilibrium entropy satisfies the additivity as in the ordinary equilibrium thermodynamics. Therefore, we can expect that, when the equilibrium entropy satisfies the additivity even with the gravitational effects, the nonequilibrium entropy also satisfies the additivity. Hence, due to the additivity of the equilibrium entropies of the black hole and the Hawking field \\cite{ref-add}, it seems reasonable to assume the additivity \\eqref{eq-strategy.total.general}.  Next we consider the validity of the existence of a well-defined nonequilibrium entropy $S_m$ of an arbitrary self-interacting Hawking field. As mentioned in the previous paragraph, the nonequilibrium entropy has already been defined well up to the second order in the expansion by the heat flux \\cite{ref-eit}. Further, on the other hand, the evaporation time $t_{ev}$ of a black hole has an extremely long time scale, $\\lambda = t_{ev}/t_{bh} \\gg O(1)$, as discussed above. Therefore it seems appropriate that the black hole evaporation proceeds so slowly that the nonequilibrium state of the self-interacting Hawking field is not so far from an equilibrium state that the expansion of the thermodynamic quantities by the heat flux up to the second order can be a good approximation. Hence it is reasonable to assume the existence of a well-defined nonequilibrium entropy of an arbitrary self-interacting Hawking field.      \\appendix"
    },
    "0609/astro-ph0609703_arXiv.txt": {
        "abstract": "{% The X-ray spectra of luminous Seyfert 1 galaxies often appear to be reflection dominated. In a number of Narrow Line Seyfert 1 (NLS1) galaxies and galactic black holes in the very high state, the variability of the continuum and of the iron line are decoupled, the reflected component being often much less variable than the continuum. These properties have been interpreted as effects of gravitational light bending. In this framework, we present detailed Monte-Carlo simulations of the reflection continuum in the Kerr metric. These calculations confirrm that the spectra and variability behaviour of these sources can be reproduced by the light bending model. As an alternative to the light bending model, we show that similar observational properties are expected from radiation pressure dominated discs subject to violent clumping instabilities and, as a result, have a highly inhomogeneous two-phase structure. In this model, most of the observed spectral and variability features originate from the complex geometrical structure of the inner regions of near-Eddington accretion flows and are therefore a signature of accretion physics rather than general relativity.} ",
        "introduction": "The spectra of  luminous Seyfert galaxies are very well described by photoionized and strongly relativistically blurred reflection models (Fabian et al. 2004, 2005; Crummy et al. 2006; Porquet  2006). In these sources the primary continuum often appears to be strongly suppressed. In a number of Narrow Line Seyfert 1 (NLS1) galaxies and galactic black holes in the very high state, the variabilities of the continuum and of the iron line are decoupled, in apparent contradiction with the predictions of simple disc reflection models (see, e.g.  Miniutti, Fabian \\& Miller 2004; Fabian et al. 2004).  In particular, the monitoring of Seyfert galaxies indicates that the reflection flux  can be weakly variable when the primary emission changes dramatically (Papadakis et al. 2002). Moreover in at least two AGN, MGC 6-30-15 (Miniutti et al 2003) and NGC4051 (Ponti et al. 2006)  and one X-ray binary (XTE J1650-500, Rossi et al. 2005)  the reflection flux is correlated to the primary emission at low fluxes and saturates at higher fluxes. In this paper we discuss two possible models for these puzzling properties.  In the light bending model of Miniutti \\& Fabian (2004) they are interpreted in terms of general relativistic effects, while, in the clumpy disc model of Merloni et al. (2006),  these properties would be the result of the complex structure of near Eddington accretion flows. ",
        "conclusions": ""
    },
    "0609/astro-ph0609535_arXiv.txt": {
        "abstract": "{ We have used NIRSPEC on Keck II to obtain $K$-band spectroscopy of the low mass X-ray binary V616 Mon (= A0620$-$00). V616 Mon is the proto-typical soft x-ray transient containing a black hole primary. As such it is important to constrain the masses of the binary components. The modeling of the infrared observations of ellipsoidal variations in this system lead to a derived mass of 11.0 M$_{\\sun}$ for the black hole. The validity of this derivation has been called into question due to the possiblity that the secondary star's spectral energy distribution is contaminated by accretion disk emission (acting to dilute the variations).  Our new $K$-band spectrum of V616 Mon reveals a late-type K dwarf secondary star, but one that has very weak $^{\\rm 12}$CO absorption features. Comparison of V616 Mon with SS Cyg leads us to estimate that the accretion disk supplies only a small amount of $K$-band flux, and the ellipsoidal variations are not seriously contaminated. If true, the derived orbital inclination of V616 Mon is not greatly altered, and the mass of the black hole remains large. A preliminary stellar atmosphere model for the $K$-band spectrum of V616 Mon reveals that the carbon abundance is approximately 50\\% of the solar value. We conclude that the secondary star in V616 Mon has either suffered serious contamination from the accretion of supernova ejecta that created the black hole primary, or it is the stripped remains of a formerly more massive secondary star, one in which the CNO cycle had been active.} ",
        "introduction": "V616 Mon (= A0620$-$00) is the proto-type for a small family of binaries known as Black Hole X-ray Transients (BHXTs) where the primary is a black hole, and the secondary a relatively low mass star. A sudden increase in the mass transfer rate due to an accretion disk instability has been used to explain the outbursts of these systems (Cannizzo 1993, Menou 2002). After a number of months, the BHXTs return to relative quiescence and are very faint at X-ray and optical wavelengths. During this time, both their broadband spectral energy distributions (SEDs) and moderate resolution spectra are dominated by a late-type secondary star (c.f., Gelino et al. 2001, Gonz$\\acute{a}$lez Hern$\\acute{a}$dez et al. 2004). Because the secondary star is clearly visible, radial velocity observations are possible, allowing the derivation of its orbital and rotational motion. When combined with the orbital period, and an estimate for the orbital inclination, it is possible to put useful limits on the masses of both components in the system. When this is done for V616 Mon, the mass of the black hole primary is 11 M$_{\\sun}$. As such, V616 Mon (P$_{\\rm orb}$ = 7.75 hr) has the largest mass black hole for any X-ray binary whose orbital period is P$_{\\rm orb}$ $\\leq$ 24 hr (Lee et al. 2002). As discussed by Kalogera (1999), it is difficult to produce a binary with a black hole primary that has such an extreme mass ratio.  In the case of V616 Mon, the leading uncertainty in the derivation of the black hole mass is the orbital inclination angle. Measurement of the orbital inclination in any close binary system depends on some assumptions about the nature of the stellar components. But with estimates for the gravity and effective temperature of the stars in such systems, it is possible to use light curve modeling software (e.g., that of Wilson-Divinney 1998 = WD98; see Kallrath et al. 1998) to fit the orbital inclination if it exhibits eclipses, or if one of the components is distorted enough to produce ellipsoidal variations. Since the outbursts of BHXTs are believed to be due to an accretion disk event, the secondary star must be acting as the donor in such systems, and can safely be assumed to fill its Roche lobe. Therefore BHXTs will exhibit ellipsoidal variations if the orbit is inclined to our line of sight, and the secondary star emits a significant fraction of the systemic luminosity. As shown by Shahbaz et al. (1994) and Gelino et al. 2001, V616 Mon exhibits such ellipsoidal variations in the near-IR. Using published parameters for the system derived from earlier radial velocity studies, Gelino et al. concluded that the system has an orbital inclination of 41$^{\\circ}$ $\\pm$ 3.  In that work, Gelino et al. found that the existing $BVRIJHK$ photometry of V616 Mon was consistent with a K4V secondary star, and assumed that the $JHK$ ellipsoidal variations were relatively uncontaminated by emission from the accretion disk and/or hot spot. Hynes et al. (2005) have recently cast doubt on the efficacy of such a model by exploring the possibility that the accretion disk severely contaminates the SED, and might even mimic a K-type secondary star. They justify this exercise by reference to an existing $K$-band spectrum of this source by Shahbaz et al. (1999) that seemed to lack obvious CO absorption features. Below, we present a new, higher resolution and higher S/N infrared spectrum of V616 Mon where we find a late K-type stellar spectrum, one in which the $^{\\rm 12}$CO features are anonymously weak. We then present preliminary stellar atmosphere models for this object, from which we derive estimates for the $^{\\rm 12}$C abundance, and discuss these results in the context of existing BHXT formation scenarios. ",
        "conclusions": "We have obtained a new $K$-band spectrum of V616 Mon using NIRSPEC on Keck II and find that the secondary star is clearly detected, and appears to be a late-type K dwarf with a significant $^{\\rm 12}$C deficit. The weakness of the CO absorption explains the lack of detection of these features in an earlier spectrum, a result which had been used to argue for a significant contamination of the system's $K$-band flux. While we do find that there is some contamination of the infrared spectrum in V616 Mon due to continuing accretion, it does not dramatically change previously estimated values for the orbital inclination, and hence the black hole mass for this system remains large. This source warrants further infrared spectroscopy to conclusively determine whether $^{\\rm 13}$C is truly enhanced, and UV spectroscopy to investigate the N V/C IV ratio that is a strong indicator for the transfer of material that has undergone CNO cycle processing."
    },
    "0609/astro-ph0609645_arXiv.txt": {
        "abstract": "{The inner parts of black-hole accretion discs shine in X-rays which can be monitored and the observed spectra can be used to trace strong gravitational fields in the place of emission and along paths of light rays. This paper summarizes several aspects of how the spectral features are influenced by relativistic effects. We focus our attention onto variable and broad emission lines, origin of which can be attributed to the presence of orbiting patterns -- spots and spiral waves in the disc. We point out that the observed spectrum can determine parameters of the central black hole provided the intrinsic local emissivity is constrained by theoretical models.} ",
        "introduction": "During the last four decades a picture of galactic nuclei has emerged in which massive black holes reside in the centre. In active galactic nuclei (AGN) accretion discs form around the core and provide most of the radiation emerging from that region. According to this scheme, intense X-rays originate just a few gravitational radii from the black hole horizon. Manifestation of strong gravity has been searched in X-ray spectra of AGN with the main aim of finding the firm evidence for the central black hole (BH) and determining its parameters, namely, mass and angular momentum (for a detailed exposition of the subject, see Peterson 1997; Krolik 1999).  Broad spectral features have been expected in X-rays on the basis of the model in which the iron line is formed on the surface of a geometrically thin, optically thick and relatively cold medium after irradiation by a primary source (Fabian et al.\\ 1989). Nowadays, the general relativistic (GR) iron line profiles provide a powerful tool to measure the mass of the black hole in AGN as well as Galactic black-hole candidates. Blandford \\& McKee (1982) and Stella (1990) have developed the reverberation technique that employs a response of the line profile following variations of the illuminating primary source. Along the same line of thought, Matt \\& Perola (1992) proposed to employ variations of the time-integrated line properties -- equivalent width, centroid energy and the line width. The method was developed further by many authors (see Fabian et al.\\ 2000; Reynolds \\& Nowak 2003 for reviews and references). Apart from the mass $\\mbh$, the main parameters of these models are the functional dependency of the intrinsic emissivity over the disc surface, $I(r)$, the specific angular momentum $a$ of the black hole and the inclination $\\theta_{\\mathrm{o}}$ of the accretion disc with respect to observer's line of sight. High resolution in spectral and time domains is crucial to accomplish the analysis.  This paper concentrates on gravitational effects that act on the spectral features by smearing them and moving them among energy bins. In this way gravity exerts the influence on the ultimate form of the observed spectrum. We start with a brief summary of the equations describing intensity (and polarization) propagation in the field of a black hole. Detailed exploration of broad spectral features can be considered as a method complementary to studying the overall X-ray continuum (Shaffee et al.\\ 2006) on one side and narrow spectral features (Turner et al.\\ 2004) on the other extreme.  Relativistic effects are often discussed in terms of geometrical optics with photons travelling through the empty spacetime -- this formulation is suited quite well to various flavours of `hot spots' on the disc surface, but plasma exists also above the disc where it influences the observed signal to a certain degree. Theoretical approaches have been developed that can tackle more complicated situations, such as the case of dispersive media which may be able to address, more accurately, simultaneous observations in mutually remote parts of the electromagnetic spectrum.  In the second part of the paper we discuss flares and spots as a model for X-ray variability: multiple spots are created on the surface of an accretion disc following the intense irradiation. The observed signal is then modulated by relativistic effects. This scheme is testable and it captures many properties of present observations.  Mean spectra of orbiting spots resemble those of axisymmetric rings and we briefly discuss this degeneracy. It is inherent to time integrated data and can be  resolved by increasing sensitivity and time resolution of the observations.  In the present context, the spots represent any kind of localized non-axisymmetric features residing on the disc surface, sharing its orbital motion and modulating the signal. They are thought to result from the illumination of a relatively cold gas of the disc after magnetic reconnection events. Such a possibility was proposed for the X-ray emitting region by Galeev et al.\\ (1979) and developed in many papers (Poutanen \\& Fabian 1999; Merloni \\& Fabian 2001; \\.Zycki 2002; Czerny et al.\\ 2004, and references cited therein).  The flare/spot scenario can be considered as a generalized model of the phenomenological lamp-post geometry where a point-like X-ray source is located at a specified height on the disc axis (Henri \\& Pelletier 1991; Martocchia \\& Matt 1996). Martocchia et al.\\ (2000) and Dov\\v{c}iak et al.\\ (2004b) demonstrated by confidence contour analysis in the Kerr metric that this scheme can be used to determine free parameters {\\em{}if} the resolution of our observations is sufficient and physical mechanisms underlying the line features and the continuum are modeled with sufficient precision, based on a sound understanding of the underlying physics.  New steps towards a detailed physical analysis of the radiation mechanisms forming X-ray spectra have been pursued by several groups (Ballantyne et al.\\ 2001; Nayakshin et al.\\ 2000; Done \\& Nayakshin 2001; R\\'o\\.za\\'nska et al.\\ 2002; Collin et al.\\ 2004; Kallman et al.\\ 2004; Ross \\& Fabian 2005) and adapted specifically for the orbiting spot model (Czerny et al.\\ 2004; Goosmann et al.\\ 2006). A similar line of reasoning stressing that the intrinsic emission of the disc medium has to be accurately modelled has been now adopted also by Brenneman \\& Reynolds (2006). Although the implicit assumptions of the codes and numerical approaches  of different authors are quite divers, useful comparisons were performed  within the overlapping range of the parameter space and can be found in the above-mentioned papers.  More complex geometries of the emitting region should be also explored, such as spiral waves propagating across the disc. This comes in accord with recognition of the importance magnetic torques may have (Krolik et al.\\ 2005), but we can safely conclude that strong gravity of the central BH is very likely the main agent shaping the overall form of the X-ray spectral features from the inner disc. ",
        "conclusions": "The influence of strong gravity on light provides a possibility to study black holes. However, gravity is not the only agent that shapes the observed spectra; emission mechanisms and radiation transfer in the accreting gas need to be understood reliably. Remarkable progress has been achieved on the way to model the physical processes which govern the form of X-ray reflection spectra emerging from different regions of accretion discs. In spite of this advance the current description still contains some elements of the phenomenological approach and the emission from non-axisymmetric patterns cannot be computed entirely `from first principles'. Nonetheless, already at its present formulation the model of a spotted disc helps us to constrain the parameters of accreting black holes. The orbital radius of the peak emission and the disc inclination can be inferred from the variations of the observed line width and the centroid energy. $\\mbh$ can be estimated by comparing the measured orbital period with the value expected for the derived radius. The model of spirals also gives the possibility of tracing the surface structures on accretion discs. Once these are determined the goal of measuring black hole rotation will be feasible with unprecedented accuracy."
    },
    "0609/astro-ph0609473_arXiv.txt": {
        "abstract": "We present accurate simulations of the dynamical bar-mode instability in full General  Relativity focussing on  two aspects which have  not been investigated in detail  in the past. Namely, on  the persistence of the bar deformation once the instability  has reached its saturation and on the  precise  determination of  the  threshold  for  the onset  of  the instability in terms of  the parameter $\\beta={T}/{|W|}$.  We find that generic nonlinear  mode-coupling effects appear  during the development of the instability and these  can severely limit the persistence of the bar deformation  and eventually suppress the  instability. In addition, we observe the dynamics of the instability to be strongly influenced by the  value  $\\beta$ and  on  its  separation  from the  critical  value $\\beta_c$ marking the onset of  the instability.  We discuss the impact these results  have on the  detection of gravitational waves  from this process and  provide evidence that the  classical perturbative analysis of the  bar-mode instability for Newtonian  and incompressible Maclaurin spheroids remains qualitatively valid and accurate also in full General Relativity. ",
        "introduction": "\\label{sec:intro}   It is well known that rotating neutron stars are subject to non-axisymmetric instabilities for non-radial axial modes with azimuthal dependence $e^{i m \\phi}$ (with $m=1,2,\\ldots$) when the instability parameter $\\beta\\equiv T/|W|$ (\\textit{i.e.} the ratio between the rotational kinetic energy $T$ and the gravitational binding energy $W$) exceeds a critical value $\\beta_{c}$.  An exact and perturbative treatment of these instabilities exists only for incompressible self-gravitating fluids in Newtonian gravity (see refs.~\\cite{Chandrasekhar69,ShapiroTeukolsky83}) and this predicts that a dynamical instability should arise for the ``bar-mode'' (\\textit{i.e.} the one with $m$=2) when $\\beta \\geq \\beta_{c} = 0.2738$. On the other hand, the accurate study of the dynamical bar-mode instability, of its nonlinear evolution and of the determination of the threshold for the instability, demand the use of numerical simulations with the solution in three spatial dimensions (3D) of the fully nonlinear hydrodynamical equations coupled to the Einstein field equations.  Despite these requirements, much of the literature on this process has so far been limited to a Newtonian or post-Newtonian (PN) description. While this represents an approximation, these studies have provided important information on several aspects of the instability that could not have been investigated with perturbative techniques. In particular, these numerical studies have shown that $\\beta_{c}$ depends very weakly on the stiffness of the equation of state (EOS) and that, once a bar has developed, the formation of spiral arms is important for the redistribution of the angular momentum (see refs.~\\cite{Tohline85, Tohline86, Brown:2000, HouserCentrellaSmith94, HouserCentrella96, SmithHouserCantrella96, NewCentrellaTohline00, NewShapiro2001, Liu2002}). More recently, instead, it was shown that the threshold for the onset of the dynamical instability can be smaller for stars with a high degree of differential rotation and a weak dependence on the EOS was confirmed in refs.~\\cite{PicketDurisenDavis96, TohlineHachisu90, KarinoEriguchi03, ShibataKarinoEriguchi2002, ShibataKarinoEriguchi2003}. Finally, these Newtonian analyses have also provided the first evidence that an $m$=1-mode dynamical instability may play an important role for smaller values of the critical parameter. This is also referred to as the ``low-$\\beta$'' instability~\\cite{Saijo:2003,Ott05} and it will not be considered here.  Only very recently it has become possible to perform simulations of the dynamical bar instability for old neutron stars in full General Relativity~\\cite{Shibata00b}. These studies have shown that within a fully general-relativistic framework the critical value for the onset of the instability is smaller than the Newtonian one (\\textit{i.e.} $\\beta_{c}\\simeq 0.24-0.25$) and this behaviour was confirmed by PN calculations~\\cite{Saijo2000,Saijo:2001} which also suggested that $\\beta_{c}$ varies with the compactness $M/R$ of the star.  The bar-mode instability may take place in young neutron stars, either as the result of the accretion-induced collapse of a white dwarf~\\cite{Dessart2006} or in the collapse of a massive stellar core. Indeed, recent simulations investigating axisymmetric stellar-core collapse in full General Relativity~\\cite{ShibataSekiguchi2004} have pointed out that for sufficiently differentially rotating progenitors, it is, at least in principle, possible to obtain toroidal protoneutron-star cores with masses between 1.2 and 2.1 $M_\\odot$ which are unstable against bar-mode deformations (it should be mentioned that it is still unclear how likely these high rotation-rates and strongly differential-rotation profiles actually are in nature). Besides pointing out this alternative interesting scenario, the work in ref.~\\cite{ShibataSekiguchi2004} has also suggested that more realistic EOS and neutrino cooling could enhance the process; this scenario has also been considered in the PN approximation in ref.~\\cite{Saijo2005}. An important common feature in all these investigations is that the development of the bar-mode instability was obtained introducing very strong \\textit{ad hoc} $m$=2-mode perturbations.  The recent general-relativistic studies, together with their Newtonian and PN counterparts, have been very helpful in highlighting the main features of the instability. However, several fundamental questions remain unanswered. Most notably: \\textit{i)} What is the role of the initial perturbation?  \\textit{ii)} What is the effect of the symmetry conditions often used in numerical calculations? {\\it iii)} How are the dynamics influenced by the value of the parameter $\\beta$, especially when this is largely overcritical?  Finally and most importantly: \\textit{iv)} How long does a bar survive, once fully developed?  Clearly, the last question has important implications for the possible observational relevance of the gravitational-wave signal emitted through the bar-mode instability as the signal-to-noise ratio (SNR) can increase considerably in the case of a long-lived bar since the SNR grows as the square root of the number of the effective cycles of the signal available for detection. Earlier work on this subject basically suggested that once a bar was formed it would tend to be persistent on the radiation-reaction timescale. Indications and evidences in this direction were presented in a Newtonian framework in ref.~\\cite{Brown:2000} as well as in a PN one in ref.~\\cite{Saijo:2001}. In contrast, fully general-relativistic results, either using suitable symmetry boundary conditions (\\textit{i.e.} the so called $\\pi$-symmetry boundary conditions)~\\cite{Shibata00b} or not~\\cite{ShibataSekiguchi2004}, show a non-persistent bar.  In this work we try to find answers to these important open questions by exploring, in a systematic way, the bar-mode instability for a large number of initial stellar models. In doing this, we intend to go beyond the standard phenomenological discussion of the nonlinear dynamics of the instability often encountered in the literature.  The main results of our analysis can be summarised as follows: \\textit{i)} The initial perturbation (either in the form of an $m$=1 mode or of an $m$=2 mode) can play a role in determining the duration of the bar-mode deformation, but not in determining the growth time of the instability; the only exception to this is represented by models near the threshold; \\textit{ii)} For moderately overcritical models (\\textit{i.e.} with $\\beta \\approxeq \\beta_c$), the use of a $\\pi$-symmetry can radically change the dynamics and extend considerably the persistence of the bar; this ceases to be true for largely overcritical models (\\textit{i.e.} with $\\beta \\gg \\beta_c$), for which even the artificial symmetries are not sufficient to provide a long-lived bar; \\textit{iii)} The persistence of the bar is strongly dependent on the degree of overcriticality and is generically of the order of the dynamical timescale; \\textit{iv)} Generic nonlinear mode-coupling effects (especially between the $m$=1 and the $m$=2 modes) appear during the development of the instability and these can severely limit the persistence of the bar deformation and eventually suppress the {bar deformation}; \\textit{v)} The dynamics of largely overcritical models are fully determined by the excess of rotational energy and the bar deformation is very rapidly suppressed through the conversion of kinetic energy into internal one. In addition, we have also assessed the accuracy of the classical Newtonian stability analysis of Maclaurin spheroids for incompressible self-gravitating fluids~\\cite{Chandrasekhar69}. Overall, and despite having applied it to differentially rotating and relativistic models, we have found it to be surprisingly accurate in determining both the threshold for the instability and the complex eigenfrequencies for the unstable models.  The paper is organized as follows. In Sec.~II we give details on the evolution methods used, while in Sec.~III we discuss the initial models and their properties. In Sec.~IV we introduce the methodology used to analyse the numerical results of the simulations, which are then discussed in Sec.~V in terms of the general dynamics of the instability and of the general properties. In Sec.~VI, we present the features of the instability that are specific to different treatments of the initial conditions, while in Sec.~VII we illustrate two different methods for the determination of $\\beta_c$. Finally, in Sec.~VIII we discuss the impact of our results on the emission of gravitational waves from the unstable models and present in Sec.~IX our conclusions and the prospects of future research.  We have used a space like signature $(-,+,+,+)$, with Greek indices running from 0 to 3, Latin indices from 1 to 3 and the standard convention for the summation over repeated indices. Furthermore, we indicate as ($x$,$y$,$z$) the Cartesian coordinates and we define $r=\\sqrt{x^2+y^2+z^2}$, $\\varpi=\\sqrt{x^2+y^2}$, $\\theta=\\arctan(\\varpi/z)$, $\\phi=\\arctan(y/x)$ for the axial and spherical coordinates. Unless explicitly stated, all the quantities are expressed in the system of adimensional units in which $c=G=M_\\odot=1$. ",
        "conclusions": "\\label{sec:conclusions}  We have presented accurate simulations of the dynamical bar-mode instability in full General Relativity. An important motivation behind this work is the need to go beyond the standard phenomenological discussion of the instability and to find answers to important open questions about its nonlinear dynamics. Among such open problems there are, for instance, the determination of the role of the initial perturbation or of the symmetry conditions, or the influence on the dynamics of the value of the parameter $\\beta$, or, most importantly, the determination of the timescale of the persistence of the bar deformation once this is fully developed. Clearly, this latter question is a very pressing one in gravitational-wave astronomy, as it bears important consequences on the detectability of the whole process.  In order to provide answers to these questions we have explored the onset and development of the instability for a large number of initial stellar models. These have been calculated as stationary equilibrium solutions for axisymmetric and rapidly rotating relativistic stars in polar coordinates. More specifically, the evolved models represent relativistic polytropes with adiabatic index $\\Gamma=2$ and are members of a sequence having a constant amount of differential rotation with $\\hat{A}=1$ and a constant rest mass of $M_0 \\simeq 1.51\\ M_{\\odot}$.  The simulations have been carried out with the general-relativistic hydrodynamics code {\\tt Whisky}, in which the hydrodynamics equations are written as finite differences on a Cartesian grid and solved using HRSC schemes~\\cite{Baiotti03a,Baiotti04a}. The Einstein, equations, on the other hand, have been solved within the conformal traceless formulation implemented within the \\texttt{Cactus} computational toolkit~\\cite{Cactusweb}.  The main results of our analysis can be summarised as follows: \\textit{i)} An initial $m$=1 or $m$=2-mode perturbation can play a role in determining the duration of the bar-mode deformation, but not the growth time of the instability; the only exception to this is represented by models near the threshold; \\textit{ii)} For moderately overcritical models the use of a $\\pi$-symmetry can radically change the dynamics and extend considerably the persistence of the bar; this ceases to be true for largely overcritical models; \\textit{iii)} The persistence of the bar is strongly dependent on the degree of overcriticality and is generically of the order of the dynamical timescale; \\textit{iv)} Generic nonlinear mode-coupling effects (especially between the $m$=1 and the $m$=2 mode) appear during the development of the instability and these can severely limit the persistence of the bar deformation and eventually suppress the {bar deformation}; \\textit{v)} The dynamics of largely overcritical models (\\textit{i.e.} with $\\beta \\gg \\beta_c$) are fully determined by the excess of rotational energy and the bar deformation is very rapidly suppressed through the conversion of kinetic energy into internal one. Interestingly, a similar dynamics for the odd and even modes has been observed also in the case of the low-$\\beta$ instability in recent Newtonian simulations~\\cite{OuTohline2006}. In this case, however, the growth rate of the $m$=1 mode is much smaller and hence the persistence of the bar longer.  Finally, we have considered whether the classical Newtonian stability analysis of Maclaurin spheroids for incompressible self-gravitating fluids is accurate also for differentially rotating and relativistic stars. Overall, we have found the perturbative predictions to be surprisingly accurate in determining the threshold for the instability as well as the complex eigenfrequencies for the unstable models.  While the features of the bar-mode instability discussed here are expected to be rather generic, they have been deduced from the analysis of a small region of the space of parameters.  Work is now in progress to extend the present analysis by considering stellar models with different and larger compactnesses and that are regulated by more realistic EOSs."
    },
    "0609/astro-ph0609190_arXiv.txt": {
        "abstract": "A general approach to analyze the electrodynamics of nuclear matter in bulk is presented using  the relativistic Thomas-Fermi equation generalizing to the case of $N \\simeq (m_{\\rm Planck}/m_n)^3$ nucleons of mass $m_n$ the approach well tested in very heavy nuclei ($Z \\simeq 10^6$). Particular attention is given to implement the condition of charge neutrality globally on the entire configuration, versus the one usually adopted on a microscopic scale. As the limit $N \\simeq (m_{\\rm Planck}/m_n)^3$ is approached the penetration of electrons inside the core increases and a relatively small tail of electrons persists leading to a significant electron density outside the core. Within a region of $10^2$ electron Compton wavelength near the core surface electric fields close to the critical value for pair creation by vacuum polarization effect develop. These results can have important consequences on the understanding of physical process in neutron stars structures as well as on the initial conditions leading to the process of gravitational collapse to a black hole. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609184_arXiv.txt": {
        "abstract": "The ISM of active galaxy centers is exposed to a combination of cosmic ray, FUV and X-ray radiation. We apply PDR models to this ISM with both `normal' and highly elevated ($5\\times 10^{-15}$~s$^{-1}$) cosmic-ray rates and compare the results to those obtained for XDRs. Our existing PDR-XDR code is used to construct models over a $10^3-10^5$ cm$^{-3}$ density range and for 0.16-160 erg s$^{-1}$ cm$^{-2}$ impingent fluxes. We obtain larger high $J$ ($J>10$) CO ratios in PDRs when we use the highly elevated cosmic ray rate, but these are always exceeded by the corresponding XDR ratios. The [CI] 609 $\\mu$m/$^{13}$CO(2-1) line ratio is boosted by a factor of a few in PDRs with $n\\sim 10^3$ cm$^{-3}$ exposed to a high cosmic ray rate. At higher densities ratios become identical irrespective of cosmic ray flux, while XDRs always show elevated [CI] emission per CO column. The HCN/CO and HCN/HCO$^+$ line ratios, combined with high $J$ CO emission lines, are good diagnostics to distinguish between PDRs under either low or high cosmic ray irradiation conditions, and XDRs. Hence, the HIFI instrument on Herschel, which can detect these CO lines, will be crucial in the study of active galaxies. ",
        "introduction": "In centers of late-type galaxies, such as M~82, NGC~253, and Maffei 2, molecular line intensity ratios\\footnote{Intensity ratios of lines a to lines b are related to brightness temperature ratios by the cube of the line frequencies: ($\\nu_{a}/\\nu_{b})^{3}$} are frequently found to be high, e.g. CO(4-3)/CO(1-0)$\\sim 60$ \\citep{Israel1995, White1994, Israel2003} requiring high gas densities $n>10^{5.5}$~cm$^{-3}$, and high temperatures of $T\\geq50$~K. However, FUV photons are easily attenuated by dust and do not penetrate very deep into clouds and the galaxies are observed with beams covering regions typically hundreds of parsecs in size.  At such large spatial scales, excited dense gas is not expected to have very large filling factors. UV radiation seems incapable of maintaining very large gas fractions at temperatures of $T\\sim50-150$~K, and regular PDR models do not explain the observed high ratios.  In addition, \\citet{Israel2002} have observed the [CI] 609~$\\mu$m line in the centers of late-type galaxies, and measured [CI]~609~$\\mu$m/$^{13}$CO(2-1) line intensity ratios in the range 20 -- 60, which are hard to explain by (low-CR) PDR models. The C$^+$/C/CO transition zones in these PDRs contain only a thin layer in which neutral carbon has a high abundance. At the cloud edge, carbon is ionized, and deep in the cloud all carbon is locked up in CO.  Various authors have invoked elevated cosmic ray fluxes caused by greatly enhanced supernova rates in galaxy centers in order to explain large molecular gas masses at high temperatures \\citep{Suchkov1993,Bradford2003}, and to explain high [CI] intensities and column densities relative to $^{12}$CO and $^{13}$CO \\citep{Pineau1992, Schilke1993, Flower1994}. As most cosmic rays are produced in supernovae, their flux is proportional to the star formation rate, which is about $1$~M$_\\odot$~yr$^{-1}$ for the Milky Way.  In circumnuclear starbursts, star formation rates may be two orders of magnitude higher or more.  Such galaxy centers may also contain an embedded accreting black hole producing X-rays.  Like cosmic rays, but unlike UV photons, these X-rays can also penetrate through large column densities ($N_H>10^{24}$~cm$^{-2}$), and can cause the observed high line ratios over areas as large as 500~pc, when the emitted flux is high enough \\citep{Meijerink2005, Meijerink2006}.  In this paper, we investigate whether in galaxy central regions PDRs with and without enhanced cosmic ray fluxes can be distinguished from XDRs on the basis of observable atomic and molecular line ratios. To this end, we calculate line intensities for PDRs with very different cosmic ray rates and compare the results to those obtained for XDRs with the same radiation fields and column densities. ",
        "conclusions": "\\begin{enumerate} \\item CO line intensity ratios increase when cosmic ray ionization rates are enhanced, but they remain smaller than those in XDRs. In particular high-$J$ ($J>10$) CO lines (which will become observable with HIFI in ESA's Herschel space observatory) allow to distinguish between (high-CR) PDRs and XDRs. Using the HIFI time estimator and a beam filling factor of 0.05, we find that a line intensity of $9\\times10^{-6}$~erg~$s^{-1}$~cm$^{-2}$, the largest CO(16-15) intensity produced by our PDR models, will be detectable with HIFI in about 4 hours, while the very bright lines received from highly irradiated XDRs are detectable within minutes.  \\item $[$CI$]$~609~$\\mu$m/$^{13}$CO(2-1) ratios are much higher in high-CR PDRs than in low-CR PDRs at modest densities of $n=10^3$~cm$^{-3}$. At higher densities of $n\\geq10^4$~cm$^{-3}$, this difference vanishes. In XDRs, the ratios are always larger than in PDRs at the same density, independent of incident radiation field.  \\item HCN/CO and HCN/HCO$^+$ line ratios are good diagnostics to distinguish between PDRs and XDRs.  As the ratios obtained for low-CR and high-CR PDRs are different, combination with high-$J$ CO lines is both crucial and profitable in the study of the active galaxy centers.  \\item Although our model results broadly distinguish between low and high CR PDRs and XDRs, there is some degeneracy when constraining the four parameters (density, CR rate, $G_0$ and $F_X$) through individual ratios. Therefore, a combination of various ratios should be considered.  \\end{enumerate}"
    },
    "0609/astro-ph0609467_arXiv.txt": {
        "abstract": "We show that the inclination to the line of sight of bipolar nebulae strongly affects some of their observed properties. We model these objects as having a spherically symmetric Planetary Nebula and a dusty equatorial density enhancement that produces extinction that varies with the viewing angle. Our sample of 29 nebulae taken from the literature shows a clear correlation between the inclination angle and the near-IR and optical photometric properties as well as the apparent luminosity of the objects. As the inclination angle increases --the viewing angle is closer to the equatorial plane-- the objects become redder, their average apparent luminosity decreases, and their average projected expansion velocity becomes smaller.  We compute two-dimensional models of stars embedded in dusty disks of various shapes and compositions and show that the observed data can be reproduced by disk-star combinations with reasonable parameters. To compare with the observational data, we generate sets of model data by randomly varying the star and disk parameters within a physically meaningful range.  We conclude that a only a smooth pole to equator density gradient agrees with the observed phenomena and that thin, equatorially concentrated disks can be discarded. ",
        "introduction": "Planetary Nebulae (PNe) and Symbiotic Nebulae are the visible remains of a heavy mass loss period near the end of the lives of most low and intermediate mass stars (\\cite{G97}). A fraction of PNe is bipolar (BPNe) and this morphology is likely caused by having a binary for central star (\\cite{S04}). BPNe have special properties (\\cite{CS95}) and are important objects for the study of outflows, mass loss, and binarity in PNe Several meetings have been dedicated to asymmetrical PNe (\\cite{HS95}, \\cite{KSR00}, \\cite{MKBS04}), see references herein.  The orientation angle on the plane of the sky of asymmetrical PNe has been studied more than 30 years ago by \\cite{MH75} and more recently by \\cite{P97} who both found an apparent alignment of the long axes of PNe with the plane of the Galaxy. The paper by \\cite{CAM98} --based on a larger sample and using more rigorous statistical methods-- showed that there is no significant alignment and that PNe are essentially randomly distributed on the sky. Assuming that this random distribution also holds for the inclination with respect to the line of sight --where the inclination, i, is taken to be 0$^o$ for a pole-on nebula and 90$^o$ for an object viewed with its main axis parallel to the plane of the sky--we investigate the possible effects of orientation on the observational properties PNe.  \\cite{SHK01} studied the observational orientation effects in six bipolar Proto-PNe and found that objects are redder when viewed edge-on i.e. have the inclination angle i near 90$^o$, as expected from simple geometrical extinction considerations nd assuming some equatorial density enhancement. From these data they concluded that Proto-PNe generally have bipolar shapes and asymmetrical dust disks.  Here we present a more general look at the phenomenon and apply a simple model to a larger sample of objects that we know to be bipolar, to see if we observe differences as a function of inclination angle. We study a range of overall geometries and dust properties of the equatorial tori, different black-body temperature ratios between the dust and central stars, and simulate observational data sets by letting these model parameters vary randomly over a restricted and physically meaningful range. These data sets we then compare directly with our observed results. ",
        "conclusions": "The first and most obvious result is that a sharply changing density distribution cannot reproduce the observed behavior of the luminosity with inclination angle. The only models that agree well with the observations have the smoothly varying structure of the last row of Figure\\,3. We will continue calling this distribution a ``disk''.  The 2-D model copes well with double central stars. By using a hot, compact star with a cooler giant-like companion, we get good agreement with the observations. Since we only have a relatively small sample of random systems the spread in our plots is rather large and a precise and detailed model cannot be selected on this basis. We do, however, confirm that the observed behavior as a function of inclination angle is easily reproducible with simple and physically meaningful and plausible models of bipolar nebulae."
    },
    "0609/astro-ph0609651_arXiv.txt": {
        "abstract": "{ Neutrino transport in spherically symmetric models of stellar core collapse and bounce has achieved a technically complete level, rewarded by the agreement among independent groups that a multi-dimensional treatment of the fluid-instabilities in the post-bounce phase is indispensable to model supernova explosions. While much effort is required to develop a reliable neutrino transport technique in axisymmetry, we explore neutrino physics approximations and parameterizations for an efficient three-dimensional simulation of the fluid-instabilities in the shock-heated matter that accumulates between the accretion shock and the protoneutron star. We demonstrate the reliability of a simple parameterization scheme in the collapse phase and extend our 3D magneto-hydrodynamical collapse simulations to a preliminary postbounce evolution. The growth of magnetic fields is investigated. }  \\FullConference{International Symposium on Nuclear Astrophysics - Nuclei in the Cosmos - IX\\\\ 25-30 June 2006\\\\ CERN}  \\begin{document}  \\begin{figure}[t] \\begin{minipage}{0.45\\textwidth} \\includegraphics[width=\\textwidth]{fig1.eps} \\caption{Electron fraction profiles during core collapse in model G15 \\protect\\cite{Liebendoerfer.Rampp.ea:2005}. Each line shows \\( Y_e \\) as a function of density at a given time. As \\( Y_e(\\rho ) \\) is only a weak function of time and because the accuracy is most important at core-bounce, the parameterization of \\( Y_e \\) is conveniently based on the bounce profile labelled by a circle.} \\label{fig1.eps} \\end{minipage} \\hspace{\\fill} \\begin{minipage}{0.45\\textwidth} \\center{\\includegraphics[width=0.7\\textwidth]{fig2.eps}} \\caption{Strength of the magnetic field at \\( 0 \\) ms, \\( 20 \\) ms, and \\( 40 \\) ms after bounce in two different models. Model 1 (black, cyan, magenta), which starts with a poloidal field and \\( \\Omega=2\\pi \\) rad/s, shows rapid field amplification by winding at the surface of the protoneutron star. This effect is negligible in model 2 (blue, green, red), which started with a toroidal field and \\( \\Omega=0.05 \\) rad/s.} \\label{fig2.eps} \\end{minipage} \\end{figure} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609698_arXiv.txt": {
        "abstract": "{} {We study the relations existing between fluxes emitted at CO(1-0) line, 60 and 100 $\\mu$m wavelengths, B and soft X-ray wavebands for galaxies of all morphological types. The large set of data that we created allows to revisit some of known relations existing between the different tracers of the Interstellar Medium (ISM): the link between the FIR flux and the CO line emission, the relation between X-ray emission in non active galaxies and the blue or FIR luminosity.} {Using catalogues of galaxies and works presented in the literature, we collected fluxes in FIR, 21 cm, CO(0-1) line and soft X-ray for two samples, consisting of normal and interacting galaxies respectively. Joining together these samples, we have data for a total of 2953 galaxies, not all observed in the four above wavebands.} { All the relations found are discussed in the frame of the star formation activity that is the link for most of them. We note that when an active star formation is present, it may link the galaxy fluxes at almost all wavelengths, from X to microwaves. On the contrary, in early-type galaxies where the current star formation rate has faded out the X-FIR fluxes link disappears. This result obtained for early-type galaxies is discussed and explained in detail in the frame of a suitable theoretical model, obtained coupling chemo-dynamical N-body simulations with a dusty spectrophotometric code of population synthesis.} {} ",
        "introduction": "The observations of galaxies at various wavelengths, going from radio to X-ray, allow to study the relationships existing between the various phases of the interstellar gas, and between gas, dust and stars. Some of these relations are already known since many years, such as that between CO and far infrared (FIR) luminosities \\citep{mirabel,sanders,solomon,devereux}. Others, connected with X-ray emission, have been studied more recently \\citep{padovani,david, ranalli}.  \\begin{figure*} \\centering \\resizebox{17cm}{!}{\\includegraphics[angle=0]{SCO_S100.eps}} \\caption{The relation between the flux from the CO(1-0) line and that at 100 $\\mu$m, plotted as a full line, with the representative points of the single galaxies plotted with different symbols. The crosses at the upper left of the figures represent the mean errors for the fluxes in the whole sample. For comparison the relation found by \\citet{bregman} for a sample of early-type galaxies is shown as a dotted line. The early-type galaxies and early spirals, from E to Sb, are plotted in the left panel while the late spirals are plotted in the right panel. } \\label{CO_100} \\end{figure*}  At present, different tracers of the gas are known, such as millimetric lines for the cold molecular gas, the 21 cm line for atomic hydrogen at $\\sim$100 K, IR bands for molecules at thousands of degrees, UV lines and X-ray emission for hotter gas. The dust distribution is traced also by FIR emission at 60 and 100 $\\mu$m if the grains are warm \\citep{bregman} or at 170 $\\mu$m, if they are colder \\citep{popescu}. The diffusion of large archives of observations at the above wavelengths (except for molecular lines) allowed in the last years the compilation of catalogues containing a huge number of galaxies. Using these catalogues and the works presented in the literature, we collected fluxes in FIR, 21 cm, CO(0-1) line and soft X-ray for two wide samples of normal \\citep{normal} and interacting \\citep{interacting} galaxies. Joining together these samples, we have data for a total of 2953 galaxies, not all observed in the four above wavebands. The fluxes measured with the different tracers allow now a study of the link existing between dust, gas and stars based on hundreds of galaxies.  It is known that the fluxes emitted by a galaxy at very different wavelengths may be linked together by means of the star formation mechanism (see \\citet{david,ranalli}). For instance, the formation of massive stars generates the heating of the dust clouds in which they are embedded, by absorption of their UV radiation, and produces a re-emission of this energy in the far infrared. This process links the current star formation rate to the IR emission at 60 and 100 $\\mu$m \\citep{thronson}. The ionizing radiation of stars may produce also the evaporation of the molecular clouds. Inside these clouds, where the particle density is great enough to produce a significant number of collisions between H$_2$ and CO molecules, these latter are excited and produce photons, but in optically thick regions. The warming by the UV stellar light makes these regions less dense, making visible the CO lines at their edge. Because of this mechanism, these lines are considered tracers of the cold molecular hydrogen that does not emit observable lines. The newly formed stars are also responsible of the X-ray emission, produced by very massive stars, by core-collapse SN, and by high mass X-ray binaries. According to the above described mechanisms, we expect that galaxies with active star formation will have a far infrared emission, but also CO and X-ray emissions induced by the more massive stars, linked together by means of different relations.  When the star formation decreases or vanishes, the far infrared emission decreases as well, but it may be fed by the stellar light absorbed and re-emitted in the infrared by dust (cirrus), while low-mass X-ray binaries and Type I SN contribute to the high energy galaxy spectrum. In addition, AGB stars, surrounded by dust, and the cooling flows of the interstellar medium ejected by supernovae may produce additional IR and X emission, between each other.  To study the activity of the galaxies at different wavebands, we collected data on galaxies starting from the original data of fluxes at 60, 100 $\\mu$m, CO(1-0) 2.6 mm and soft X-ray used to compile our catalogues \\citep{norm_cat, inter_cat}. The merging of the two above catalogues produces 1764 known values of far infrared fluxes (1837 have 100 $\\mu$m flux), 391 soft X-ray fluxes and 434 values of the CO(1-0) line luminosity. We extracted from LEDA catalogue \\citep{leda} the values of the distance moduli, blue absolute magnitudes and morphological classification for all of them.  Galaxies with evident sign of interactions or disturbed morphologies according to the catalogues of \\citet{arp,am,vv} are 1038. We shall refer to them as ``perturbed galaxies''. The remaining 1915 galaxies that appear neither morphologically nor dynamically perturbed are called ``normal galaxies''. In our sample, we have 253 galaxies that have spectral classification of the nucleus and 231 of these appear to host an AGN (Seyfert 1, 2 or transition type, Seyfert 3 or Liners) according to the classifications of \\citet{ho} and \\citet{veron}. Most part of the remaining 2722 galaxies lacks of information about nuclear spectrum or have spectra of HII regions (22 starburst spectra). They are not included in any AGN catalogue and for this reason in the following discussion we refer to them as ``non active galaxies\" and to the others as ``active galaxies\". With all these data, we crossed the various tracers to understand and revisit the main relations existing between X, FIR, CO and B luminosities.  \\begin{figure*} \\resizebox{17cm}{!}{\\includegraphics[angle=0]{Fig_X1.eps}} \\caption{{\\it Left:} The X-ray luminosity plotted vs. blue luminosity, in solar units, for late-type, non-active galaxies. The relation corresponding to emission caused by discrete sources, indicated in equation (\\ref{XBlate}) is plotted as a dotted line. {\\it Right:} X-ray luminosity vs. FIR luminosity, in solar units for the late-type, non-active galaxies. The linear relation indicated in equation (\\ref{XFIR}) is plotted as a dashed line.} \\label{X1} \\end{figure*} ",
        "conclusions": "We have been able to describe the relations existing in a galaxy between the various tracers of the ISM and to fix the coefficients of the relations existing between FIR, B and X-ray luminosity, both for early-type and late-type galaxies.  The large set of data we used allowed us to redefine more clearly the relation existing between the CO and the 100 $\\mu$m fluxes. We found that the relation, first obtained by \\citet{bregman} for early type galaxies, is valid also for late type galaxies. In these galaxies, the X-ray flux appears linked also to B and FIR emissions.  The only relation lacking from observations, i.e. the one between L$_X$ and L$_{FIR}$ has been studied by the use of the most recent chemo-dynamical models coupled with dusty evolutionary population synthesis.  The calculated luminosities of the models seem to confirm our hypothesis about a connection between the exhaustion of the star formation and the \"migration\" of the early type galaxies above the linear relation in the L$_X$ vs L$_{FIR}$ plot. In the frame of our assumptions, we may therefore conclude that the prediction of our dusty chemo-dynamical models of galaxy evolution is consistent with the observed lack of a direct relation between $L_{X}$ and $L_{FIR}$ for early type galaxies and is due to the different mechanisms of production of FIR light in galaxies where the active star formation is no longer active. In most of our early-type galaxies the mechanism of IR emission is no more strictly related to the ongoing star formation and to the reprocessing of the radiation in the dense regions where new stars are born. The FIR emission comes therefore most likely from circumstellar dusty shells around AGB stars and from an interstellar diffuse medium due to the outflow of dusty gas from AGB and RGB stars.  Finally, we can summarize that: (i) the SFH of the galaxies seems therefore to have the stronger effect on the position of early-type galaxies in the L$_X$ vs L$_{FIR}$ plot; (ii) other parameters, like the radius of the galaxy and the scale radii of stars and gas, play a secondary role, even if they can significantly contribute to the scatter of the models in the region above the linear relation; (iii) the mass is the main parameter explaining the scatter of the points along the linear relation."
    },
    "0609/astro-ph0609521_arXiv.txt": {
        "abstract": "\\emph{Suzaku} has, for the first time, enabled the hard X-ray variability of the Seyfert 1 galaxy MCG--6-30-15 to be measured. The variability in the 14--45~keV band, which is dominated by a strong reflection hump, is quenched relative to that at a few keV. This directly demonstrates that the whole reflection spectrum is much less variable than the power-law continuum. The broadband spectral variability can be decomposed into two components -- a highly variable power-law and constant reflection -- as previously inferred from other observations in the 2--10~keV band. The strong reflection and high iron abundance give rise to a strong broad iron line, which requires the inner disc radius to be at about 2 gravitational radii. Our results are consistent with the predictions of the light bending model which invokes the very strong gravitational effects expected very close to a rapidly spinning black hole. ",
        "introduction": "The Seyfert 1 galaxy MCG--6-30-15 at $z=0.00775$ has been the centre of much interest since a broad iron line was discovered in its X-ray spectrum with the \\emph{ASCA} satellite (Tanaka et al 1995). Iron line emission is part of the reflection spectrum produced by the hard X-ray power-law continuum in the source irradiating the accretion disc (Guilbert \\& Rees 1988; Lightman \\& White 1988; Ross \\& Fabian 1993) and is broadened and skewed by Doppler, gravitational redshift effects, and light aberration and bending (Fabian et al 1989). The low energy extent of the line can reveal the inner radius of the accretion disc and thus the black hole spin (for reviews see Fabian et al 2000; Reynolds \\& Nowak 2003 and Fabian \\& Miniutti 2006). The iron abundance in MCG--6-30-15 appears to be about 2--3 times the Solar value (Fabian et al 2002), making the iron line particularly strong.  MCG--6-30-15 was observed several more times with \\emph{ASCA} (Iwasawa et al 1996, 1999; Matsumoto et al 2003; Shih et al 2002), by \\emph{BeppoSAX} (Guainazzi et al 1999), \\emph{RXTE} (Lee et al 1999; Vaughan \\& Edelson 2001) and by \\emph{XMM--Newton} (Wilms et al 2001; Fabian et al 2002, 2003; Vaughan \\& Fabian 2004). All of these observations confirmed the presence and general broad shape of the iron line. Evidence that the black hole in MCG--6-30-15 is rapidly spinning has been obtained from the extreme breadth of the line by Iwasawa et al (1996), Dabrowski et al (1997), Wilms et al (2001), Fabian et al (2002) and most recently Reynolds et al (2005), who determine a spin parameter $a=0.989$.  The  \\emph{XMM--Newton} work emphasizes that much of the emission arise from smaller radii, between $2-6$~$r_{\\rm g}$ (where $r_{\\rm g}=GM/c^2$). This may help to explain the otherwise puzzling spectral variations shown by the source. Shih et al (2002), Fabian et al (2002), Fabian \\& Vaughan (2003) and Taylor, Uttley \\& McHardy (2003) found that the spectral variability can be explained by a simple two component model consisting of a highly variable power-law continuum (of almost fixed spectral slope) and a much less variable reflection spectrum, which of course includes the iron line. The reflection (and iron line) strength do not follow the power-law continuum intensity, as expected in a simple reflection picture. The small radius of much of the emission can however explain this behaviour when it is recalled that as well as strong gravitational red-shifting occurring in this region, there is also strong gravitational light bending. This can disconnect variations in the continuum from those of the reflection and has led to the development of the light-bending model (Fabian \\& Vaughan 2003; Miniutti et al 2003, 2004), which is a generalization of earlier work on the strong field regime (Martocchia \\& Matt 1996; Reynolds \\& Begelman 1997).  If the source of the continuum emission (assumed to be an isotropic emitter) changes location close to the black hole, then, even if the continuum has a constant intrinsic luminosity, it appears to the outside observer to change in brightness. This is just due to gravity bending the light rays out of the line of sight by different amounts depending upon the precise location of the source. Much of the radiation is bent down onto the disc, so the observed reflection intensity changes little. The two--component variability pattern and its light bending interpretation have recently found application in many other accreting black hole sources (Fabian et al 2004, 2005; Miniutti, Fabian \\& Miller 2004; Ponti et al 2006).  One striking feature of the model reflection spectrum is a large reflection hump peaking at 20--40~keV. This is where the disc albedo is highest; at lower energies the albedo is reduced by photoelectric absorption and at higher energies it is reduced (at a given energy) by Compton down-scattering of the photons (see e.g. George \\& Fabian 1991). The presence of the Compton hump has been confirmed in the spectrum of MCG--6-30-15 by \\emph{BeppoSAX} and RXTE observations (Guainazzi et al 1999; Lee et al 2000; Fabian et al 2002). What those observations have not done is to determine the variability of the Compton hump and show that it varies in the same way as the rest of the reflection spectrum\\footnote{The possible anticorrelation between the reflection continuum and iron line reported by Lee et al (2000) may in part be due to confusion caused by variation in the fitted power-law index (see discussion of a spurious reflection--photon index correlation found with RXTE data by Vaughan \\& Edelson 2001).}. Here we carry out this important step with the Hard X-ray Detector (HXD) on \\emph{Suzaku}.  \\begin{figure*} \\begin{center} \\includegraphics[width=0.38\\textwidth,height=0.9\\textwidth,angle=-90]{x23lc.ps} \\end{center} \\caption{The light curve from the XIS2/XIS3 detectors in the 0.5--12~keV band. We also show the count rates chosen to select the High Flux (HF) and Low Flux (LF) states. They are chosen so that the HF and LF spectra have approximately the same number of counts ($6.5\\times 10^5$ in the whole 0.5--12~keV band). } \\label{totallc} \\end{figure*}  We concentrate on the \\emph{Suzaku} data above 3~keV since there is a complex warm absorber in MCG--6-30-15 (Otani et al 1996; Lee et al 2001; Turner et al 2003,2004). The difference spectrum between low and high states of the source shows that absorption is minimal above 3~keV (Fabian et al 2002; Turner et al 2004) and High Energy Transmission Grating (HETG) spectra from \\emph{Chandra} show no absorption feature around 6.5~keV which would indicate absorption by species of intermediate ionization which could affect our fits (Young et al 2005). ",
        "conclusions": "The Suzaku observation of MCG--6-30-15 has, for the first time, enabled the hard X-ray variability of the source to be determined. A wide range of techniques including direct spectral fitting, flux-flux plots and rms spectrum all show that the spectral variability on a 45~ks timescale can be decomposed into just two broad components above 3~keV, namely a variable power-law continuum and a harder constant component. This last component has the broad iron line and reflection hump expected from X-ray reflection in the very innermost regions of an accretion disc around a rapidly spinning black hole.  Just the strength and shape of the reflection component predicts a high equivalent width iron line, much stronger than the weak narrow emission component seen in the spectrum. The only way for such a strong iron line to be `hidden' in the spectrum is for it to be smeared into the broad emission structure that we observe.  The behaviour of the power-law and reflection components is consistent with the light bending model. In which case much of the rapid (power-law component) variability seen in the source is due to strong gravitational light bending as the site of the emission region changes location within a few gravitational radii of the black hole. We assume that the emission region is associated with some magnetic structure produced by the strong differential motions in the accretion disc. Note that to produce observed variability the emission region does not need to physically move, just that the main emission site changes position.  The broad iron line is strong in MCG--6-30-15 because both the iron abundance is high, at around 2 times the solar value, and the large degree of light bending around its Kerr black hole has emphasised the reflection spectrum relative to the power-law continuum."
    },
    "0609/astro-ph0609717_arXiv.txt": {
        "abstract": "The timescale of de-leptonization by neutrino loss and associated contraction of a proto-neutron star is short compared to the time to progagate a shock through the helium core of a massive star, and so the de-leptonization phase does not occur in the vacuum of space, but within the supernova ambiance whether or not there has been a successful explosion. Dynamical non--axisymmetric instabilities (NAXI) are predicted for sufficiently strongly differentially rotating proto--neutron stars.  Some modes are unstable for small values of the ratio of rotational kinetic energy to binding energy, ${\\tw} \\gtrsim 0.01$. The NAXI are likely to drive magnetoacoustic waves into the surrounding time--dependent density structure. These waves represent a mechanism of the dissipation of the free energy of differential rotation of the proto--neutron star, and the outward deposition of this energy may play a role in the supernova explosion process. We estimate the power produced by this process and the associated timescale and discuss the possible systematics of the de-leptonization phase in this context. A likely possibility is that the proto-neutron star will spin down through these effects before de--leptonization and produce substantial but not excessive energy input. ",
        "introduction": "Core collapse of massive stars affects a broad range of astrophysics including the production of neutron stars and black holes, the creation and dissemination of heavy elements, and the potential production  of gravitational radiation, but the physics of core collapse supernovae is not yet fully understood. Neutrino processes play a major role and asymmetries are ubiquitous \\citep{wan96, wan01, leo06}.  After its birth as a hot, rapidly spinning proto-neutrons star (PNS) at core bounce, a neutron star that avoids being converted to a black hole will de--leptonize, shrink in radius, and spin up. The evolution in this phase, particularly in the presence of magnetic fields, has not been thoroughly explored, but is critical to establish the rotation and magnetic field distribution of the neutron star as it becomes accessible to observation in the center of the explosion. One of the processes that could enhance the depletion of the angular momentum of the neutron star is the development of non--axisymmetric instabilities (NAXI) with the attendant generation of magnetoacoustic waves. When differential rotation is strong, the value of the ratio of rotational kinetic energy to binding energy, ${\\tw}$, for onset of NAXI can be rather low (see \\S 2).  If the threshold for NAXI is low, these instabilities may be relevant to a significant range of rotating core--collapse situations.  Much of the literature on NAXI in neutron stars ignores the hot PNS conditions just after core collapse, magnetic fields, and the time--dependent density structure that surrounds the new--born neutron star until this matter is swept away in the supernova explosion.  These complications may have a major impact on the rotational and magnetic state of the surviving, observable neutron star. We explore some of the possible effects here.  Section 2 gives a brief summary of the literature of non--axisymmetric instabilities of differentially rotating neutron stars and an estimate of the rate of dissipation of rotation by NAXI. We consider the de--leptonization phase in \\S 3.  Conclusions are given in \\S 4. ",
        "conclusions": "The implication is that if a supernova does not succeed on a time scale of $\\sim 100$ ms, de--leptonization is likely to cause the PNS to contract and spin up and amplify the driving and magnetic dissipation of non--axisymmetric modes that could pump energy into the environment of the neutron star, potentially energizing the explosion if $\\tw \\gtrsim 0.01$ during the evolution.  Even if an explosion has been launched on a time scale of $\\sim 100$ ms, the infall will not have declined to zero in the early part of the contraction process. The ambient density will still give a medium in which to produce and dissipate magnetoacoustic waves generated by non--axisymmetric instabilities.  The contraction phase could thus be important for both the energetics of the explosion and the rotational period of the contracted neutron star, so it is important to understand this phase. We have made some estimates of the energy that could be deposited by magnetoacoustic waves, but caution that this will depend on the actual free energy of differential rotation, the details of the change of structure and hence binding energy during contraction and the mode of the magnetoacoustic flux. If the latter were to emerge up the axis as Poynting flux, it might not contribute to the kinetic energy budget of the expanding ejecta. The contraction could lead to a rapidly rotating neutron star. As we show here, however, a plausible result is a rather rapid spin down prior to substantial contraction followed by contraction at roughly constant $\\tw \\sim 0.01$ corresponding to the threshold for NAXI. That would correspond to a post-contraction rotational period of $\\sim$ 3 ms prior to any final spin down.  The evolution during the de--leptonization phase will depend on the physics and time scales of the dissipation of the free energy of differential rotation. We noted that in our simple model for the time scale associated with the magnetoacoustic flux generated by NAXI, Eqn. \\ref{tau-mhd}, the time scale could be rather short compared to the contraction time.  Eqn. \\ref{tau-mhd} depends on the rotation frequency implicitly through $\\tw \\propto \\Omega^2$ and $\\delta r/r$.  This is in contrast to the dissipation timescale in a model in which the magnetic field serves to provide a magnetic viscosity \\citep{tho05,wil05,ott06}, for which the dissipation power scales as $\\Omega^{3}$ and hence the time scale as $\\Omega^{-1}$.  We note that the time scale associated with this process is still model dependent with \\citet{tho05} and \\citet{ott06} adopting a dissipation power per unit mass of $\\sim \\alpha h_P^2\\Omega(d\\Omega/d ln r)^2 \\sim \\alpha h_P^2\\Omega^3$ and hence a time scale of $\\tau_{diss} \\sim (\\alpha \\Omega)^{-1}(r/h_P)^2$. \\citet{wil05} invoke a threshold magnetic field but effectively adopt a dissipation power per unit mass of $1/2 \\Omega^3 r^2$ and hence a time scale of $\\sim \\Omega^{-1}$. While both of these time scales vary as $\\Omega^{-1}$, that of \\citet{wil05} is shorter by a factor of $\\sim (h_P/r)^2$.  Another important difference between the dissipation by NAXI and magnetoacoustic flux and dissipation by magnetic viscosity is that the magnetic viscosity is, by assumption, dissipated locally as heat, whereas the magnetoacoustic flux may, and probably will, propagate and deposit the dissipated energy non-locally. Deposition near the standing shock, if one remains, is a likely locale for dissipation, since the magnetoacoustic waves are likely to propagate down the density gradient and unlikely to penetrate the entropy jump at the shock. Determining the fraction of the free energy of differential rotation that is dissipated by magnetic or other forms of viscosity as a local heat source and the fraction radiated and deposited non-locally in magnetoacoustic flux will require detailed calculations.  The contraction phase of a PNS within the context of realistic supernova conditions (successful or not) has not been considered in the literature, but in the context of rotating, magnetic, non--axisymmetric structures, this phase may be critical to understand the energetics of the explosion and the post--supernova state of the neutron star.  The contraction phase, however it behaves, will be complete before any neutron star is exposed to external observations."
    },
    "0609/astro-ph0609592_arXiv.txt": {
        "abstract": "We present the first calculation of the kinetic Sunyaev-Zel'dovich (kSZ) effect due to the inhomogeneous reionization of the universe based on detailed large-scale radiative transfer simulations of reionization. The resulting sky power spectra peak at $\\ell=2000-8000$ with maximum values of $[\\ell(\\ell+1)C_\\ell/(2\\pi)]_{\\rm max}\\sim1\\times10^{-12}$. The peak scale is determined by the typical size of the ionized regions and roughly corresponds to the ionized bubble sizes observed in our simulations, $\\sim5-20$~Mpc. The kSZ anisotropy signal from reionization dominates the primary CMB signal above $\\ell=3000$. This predicted kSZ signal at arcminute scales is sufficiently strong to be detectable by upcoming experiments, like the Atacama Cosmology Telescope and South Pole Telescope which are expected to have $\\sim1'$ resolution and $\\sim\\mu$K sensitivity. The extended and patchy nature of the reionization process results in a boost of the peak signal in power by approximately one order of magnitude compared to a uniform reionization scenario, while roughly tripling the signal compared with that based upon the assumption of gradual but spatially uniform reionization. At large scales the patchy kSZ signal depends largely on the ionizing source efficiencies and the large-scale velocity fields: sources which produce photons more efficiently yield correspondingly higher signals. The introduction of sub-grid gas clumping in the radiative transfer simulations produces significantly more power at small scales, and more non-Gaussian features, but has little effect at large scales. The patchy nature of the reionization process roughly doubles the {\\em total} observed kSZ signal for $\\ell\\sim3000-10^4$ compared to non-patchy scenarios with the same total electron-scattering optical depth. ",
        "introduction": "\\label{intro_sect}  The secondary anisotropies of the Cosmic Microwave Background (CMB) are emerging as one of the most powerful tools in cosmology. The small-scale anisotropies are probes of cosmological structures and are thus a valuable tool for studying their formation and properties. The key effect generating small-scale CMB anisotropies is the Sunyaev-Zel'dovich (SZ) effect \\citep{1969Ap&SS...4..301Z}, produced by Compton scattering of the CMB photons on moving free electrons. When this is due to thermal motions the effect is referred to as thermal SZ effect (tSZ), while when it is due to the electrons moving with a net bulk peculiar velocity it is called kinetic SZ effect (kSZ) \\citep{1980MNRAS.190..413S}. The former, coming largely from hot gas in low-redshift galaxy clusters, is the dominant effect, but has a different spectrum than the CMB primary anisotropies. The tSZ spectrum has a characteristic zero at $\\sim217$~MHz, so its contribution can in principle be separated. By contrast, kSZ anisotropies have a spectrum identical with that of the primary anisotropies. Separation of its effects is possible only by virtue of its different spatial structure. We consider the kSZ fluctuations as composed of two basic contributions, one coming from inhomogeneous reionization and the other from the fully-ionized gas after reionization. The calculation of the patchy reionization contribution to the kSZ signal based on the first large-scale radiative transfer simulations of reionization is the main focus of this paper.  This signal provides a signature of the character of reionization that will complement other approaches like observations of the redshifted 21-cm line of hydrogen and surveys of high-z Ly-$\\alpha$ emitters \\citep[e.g.][]{1990MNRAS.247..510S,2000ApJ...528..597T,2002ApJ...572L.123I, 2003MNRAS.341...81I,2003ApJ...596....1C, 2003AJ....125.1006R, 2004ApJ...604L..13S,2004ApJ...608L..77G,2004MNRAS.347..187F, 21cmreionpaper,2005astro.ph.11196M,2006ApJ...646..681S,2006NewAR..50...94B, 2006astro.ph..8032F}. In addition to better understanding of the physics of reionization, such studies could provide better constraints on the fundamental cosmological parameters, in particular on the primordial power spectrum of density fluctuations at smaller scales than are currently available.  The kSZ from a fully-ionized medium has been studied previously by both analytical and numerical means. When linear perturbation theory is used, the effect is associated with the quadratic nonlinearities in the electron density current and is usually referred to as the Ostriker-Vishniac effect \\citep{{1986ApJ...306L..51O},{1987ApJ...322..597V}}. Calculating the full, nonlinear effect is much more difficult to do analytically, although some models have been proposed \\citep[e.g.][]{1998PhRvD..58d3001J,2000ApJ...529...12H, 2002PhRvL..88u1301M,2004MNRAS.347.1224Z}.  This effect involves coupling of very large-scale to quite small-scale density fluctuation modes and thus requires a very large dynamic range in order to be simulated correctly. Current simulations are in rough agreement, but have not quite converged yet \\citep[e.g.][]{2001ApJ...549..681S,2002PhRvL..88u1301M,2004MNRAS.347.1224Z}. All simulations predict significant enhancement of the small-scale anisotropies compared to the analytical theory, due to the nonlinear evolution.  Calculations of the kSZ contribution from patchy reionization are even more challenging. In addition to the above difficulties, such calculations also require detailed modelling of the radiative transfer of ionizing radiation to derive the sizes and distributions of H~II regions in space and time (the reionization geometry) and how these correlate with the velocity and density fields. To date this problem has been studied by only a few recent works. Most of these estimates were done by semi-analytical models, which used various simplified approaches to model the inhomogeneous reionization \\citep{1998ApJ...508..435G, 2003ApJ...598..756S, 2005ApJ...630..657Z,2005ApJ...630..643M}. While such models are useful since they are cheaper to calculate than full radiative transfer simulations, and so allow for fast exploration of the various scenarios, their results must be checked against full detailed simulations in order to ascertain their reliability. The only two existing numerical studies of this effect, \\citet{2001ApJ...551....3G} and \\citet{2005MNRAS.360.1063S} used fairly small computational boxes ($4\\,\\rm h^{-1} Mpc$ and $20\\,\\rm h^{-1} Mpc$, respectively). As a consequence, their results significantly underestimate the kSZ signal, as we discuss below.  There are several upcoming experiments which will search for kSZ effect, in particular the Atacama Cosmology Telescope (ACT)\\footnote{$\\url{http://www.hep.upenn.edu/\\!\\sim\\!angelica/act/act.html}$} and South Pole Telescope (SPT)\\footnote{$\\url{http://spt.uchicago.edu/science/index.html}$}, both of which are expected to be operational by early 2007. These experiments will have $\\sim1'$ resolution and $\\sim\\mu$K sensitivities, which should be sufficient to detect the kSZ signal.  In this paper we present the first calculations of the kSZ effect from patchy reionization based on large-scale radiative transfer simulations. We simulate reionization using a $(100\\,\\rm h^{-1}Mpc)^3$ simulation volume, which is sufficient to capture the relevant large-scale density and velocity perturbations, an important improvement over previous efforts. We use the ray-tracing code C$^2$-Ray \\citep{methodpaper} to follow the radiation from all ionizing sources in that volume identified with the resolved halos, which are of dwarf galaxy size or larger. The halos and underlying density field are provided by a very large N-body simulation with the code PMFAST \\citep{2005NewA...10..393M}. The results of these simulations on the reionization character, geometry and observability of the redshifted 21-cm line of hydrogen were discussed in detail in \\citet[][Paper I]{2006MNRAS.369.1625I} and \\citet[][Paper II]{21cmreionpaper}. A preliminary version of our current results was presented in \\citet{2006astro.ph..7209I}. ",
        "conclusions": "We have derived the kSZ CMB anisotropies due to the inhomogeneous reionization of the universe. This is the first such calculation based on detailed, large-scale radiative transfer simulations of this epoch. As we have shown above, these simulations follow large enough volume to capture the full range of scales relevant to the large-scale reionization geometry. They also include most of the velocity fluctuations power, which is not the case for smaller-box simulations. We have also approximately corrected for the velocity power still missing from the box. The resulting sky power spectra peak at $\\ell=2000-8000$ with maximum values of $[\\ell(\\ell+1)C_\\ell/(2\\pi)]_{\\rm max}\\sim1\\times10^{-12}$. The angular scale of the peak roughly corresponds to the typical ionized bubble sizes observed in our simulations, which is $\\sim5-20$~Mpc, depending on the assumed source efficiencies and the gas clumping at small scales. The kSZ anisotropy signal from reionization dominates the primary CMB signal above $\\ell=3000$. At large scales the patchy kSZ signal depends largely on the source efficiencies and the large-scale velocity fields. It is higher when sources are more efficient at producing ionizing photons, since such sources produce large ionized regions, on average, than less efficient sources. The introduction of sub-grid gas clumping in the radiative transfer simulations produce significantly more power at small scales, but has little effect at large scales. The sub-grid gas clumping also significantly enhances the non-Gaussianity of the kSZ maps, resulting in up to one order of magnitude more of the brightest kSZ regions than a Gaussian would predict. The integrated kSZ signal is strong enough to be detected by upcoming experiments, like ACT and SPT. However, the separation of the patchy reionization signal from the contribution by the fully-ionized gas after reionization seems difficult.  Our current simulations do not include the smallest atomically-cooling ionizing sources, these with total mass $10^8\\,M_\\odot\\lesssim M_{\\rm tot}<2.5\\times10^9\\,M_\\odot$, as these are not resolved in our base N-body simulations. These smaller sources have important effects on the early global reionization history, but are also strongly suppressed due to Jeans-mass filtering in the ionized regions \\citep{selfregulated}. During most of the evolution the reionization process is thus dominated by the larger sources, which are resolved here, and which dictate the large-scale reionization geometry. Consequently, we do not expect that the presence of the low-mass ionizing sources would change our current conclusions significantly."
    },
    "0609/astro-ph0609071_arXiv.txt": {
        "abstract": "For a sample of 43 nearby, late-type galaxies, we have investigated the radial variation of both the current star formation rate and the dust-induced UV light attenuation. To do this we have cross-correlated IRAS images and GALEX observations for each of these galaxies, and compiled observations of the gas (CO and HI) and metal-abundance gradients found in the literature. We find that attenuation correlates with metallicity.  We then use the UV profiles, corrected for attenuation, to study several variants of the Schmidt law and conclude that our results are compatible with a simple law similar to the one of \\citet{kenni98a}, extending smoothly to lower surface densities, but with considerable scatter. We do not detect an abrupt break in the UV light at the threshold radius derived from H$\\alpha$ data (at which the H$\\alpha$ profile shows a break and beyond which only a few HII regions are usually found). We interpret the H$\\alpha$ sudden break not as a change in the star formation regime (as often suggested) but as the vanishingly small number of ionizing stars corresponding to low levels of star formation. ",
        "introduction": "Interstellar dust affects our view of galaxies from the UV to the near-infrared. High-redshift galaxies are commonly studied in the rest-frame UV where dust effects can be  extremely severe, and our estimation of the ``cosmic'' star formation rate (SFR) \\citep[e.g.][]{schimi05} crucially depends on the corrections applied to account for the dust attenuation of starlight.  Radiative transfer models suggest that the ratio of far-infrared to UV radiation is a reliable measure of the UV attenuation \\citep{buat96,witt00,panuzzo03}, depending weakly upon the geometry of stars and dust, the extinction law, or the nature of the underlying stellar population. Since this ratio is not always available, the slope of the UV spectrum reddening with respect to local starbursts \\citep{calzetti94,meurer95,meurer99} has been commonly used as a metric to estimate the amount of attenuation. This allows one to compare the attenuation (as measured by the far infrared/UV ratio) to the slope of the UV spectrum in a diagnostic plot, the commonly-called ``IRX-$\\beta$ relationship''. Several recent works, mostly based on GALEX observations, cast doubt on the IRX-$\\beta$ relation that has been most commonly used during the recent years \\citep{buat05,cortese06,seibert05,gil06}. Independent studies have come to similar conclusions, e.g. \\citet{bell02} for normal spirals and \\citet{bell02b} for individual regions in the LMC.  These works suggests either that there is no relationship (or a very noisy one) or that a relationship exists but differs from the starbursts one, depending on the selection of the sample (infrared vs UV or optically-selected).  The largest angular sized galaxies in the {\\it GALEX Atlas of Nearby Galaxies} \\citep{gil06} allow us to revisit the IRX-$\\beta$ relationship. Indeed, several tens of nearby galaxies observed by GALEX are large enough to have been resolved even by IRAS. The combination of the UV data of GALEX and infrared of IRAS opens up the possibility of looking at the spatially resolved IRX-$\\beta$ relationship across a relatively large sample of galaxies. Even if the number of resolved objects in this study is smaller than the total number of objects used in the global studies, each of our galaxies provides a number of independent points along the radial profile. Regions situated at the same radius within a galaxy are relatively homogeneous (for instance similar metallicities), thus radial profiles are physically relevant.  Another compelling reason to study attenuation\\footnote{ ``Attenuation'' is the loss of flux due to the presence of dust. It is sometimes considered that ``extinction'' should be used only for the dust properties and not the global effect of dust on the galaxy light, also depending on e.g. the geometry of the system. ``Extinction'' is nevertheless often used with the same meaning as ``attenuation''.} in these galaxies is that, as our closest neighbors, they have been extensively surveyed at other wavelengths, and in the published literature there are metallicity gradient and gas density profiles, etc. that allow us to go one step further and try to characterize how the properties of dust might depend on these other physical quantities. We also produce star formation rate radial profiles derived from the UV (and not H$\\alpha$ data), corrected for the dust attenuation for all our galaxies.  These profiles can be used to study how the star formation rate depends on the gas density, {\\it i.e.} the so-called Schmidt law. We will be able to use these star formation rate profiles to finally study the threshold radius beyond which little (if any) star formation is observed. Many studies suggest that the threshold radius corresponds to the point where gas density becomes lower than a critical density for star formation \\citep[e.g.][]{martin01}. It is important to bring as many constraints as possible to bear on the complex problem of star formation and this paper brings its share, by using UV GALEX data. Indeed, while many other studies of the Schmidt law \\citep{kenni98a,boissier03,wongblitz} are based on H$\\alpha$ data, the star formation rate in this paper is derived from the UV, which ``averages'' the star formation on longer time-scales than H$\\alpha$. % Recombination lines as H$\\alpha$ are emitted only in presence of a large ionizing flux, coming mostly from massive stars ($M > 10 M_{\\odot}$) with lifetimes shorter than 2 $\\times 10^7$ years. On the contrary, the UV continuum is emitted by stars with lifetimes up to $10^8$ years \\citep[a review of the various star formation indicators and their associated lifetimes is given in][]{kenni98b}. Because of the scarcity of UV data before GALEX, little work had been done in the UV on the star formation law. Some exceptions are, that of \\citet{buat89b}, \\citet{bersier94} concerning respectively M81 and M51, or \\citet{buat89a} and \\citet{donas87} for studies using sample of nearby galaxies but these papers looked only at the integrated properties of galaxies, and employed rather simple recipes for dust correction. The radial variation of the UV-derived star formation rate, correctly dust-corrected via the far-infrared to UV ratio, is the subject of this paper. This is an important work as we now know that the UV also shows us low levels of star formation \\citep[XUV galaxies, e.g.][]{gildepaz05,thilker05}, in regions with such low dust content that the far-infrared emission is undetected by IRAS, and that the UV is likely to be weakly affected by extinction.    In section \\ref{secdata}, we discuss how the data-sets were assembled and prepared in order to undertake this study of the radial variation of the attenuation, of the star formation, and of their relation to other physical quantities in those same galaxies. In section \\ref{secext}, we present our study of the attenuation derived from these data, and the consequences concerning the resulting star formation law are presented in section \\ref{secstarform}.    \\clearpage \\begin{deluxetable}{l r r r r r r r r r r r } \\tablecaption{\\label{tabrefs}Sample: basic parameters used and references} \\startdata Name      &  E(B-V) &  T  &  D~~   &  PA  &  D25  &  d25 & Res. & \\multicolumn{3}{l}{~~~~~~~~~~References} \\\\ &  (mag)  &     & (Mpc)& (deg)&  (') & (') & ('') & [O/H]     & HI & CO & V(R)   \\\\ \\object{IC 1613}  &   0.025 & 10      &          0.8   &   50  &  16  & 15  & 106    & d      & 24  &    & 24     \\\\ \\object{IC 2574}  &   0.036 &  9      &          4.0   &   50  &  13  &  5  & 110    & e      & 25  &    & 25     \\\\ \\object{M31}      &   0.062 &  3      &          0.8   &   35  & 190  & 60  & 108    & i      & 28  & 28 & 18, 27 \\\\ \\object{M33}      &   0.042 &  6      &          0.8   &   23  &  71  & 42  & 108    & i      & 17  & 19 & 17     \\\\ \\object{M51a}     &   0.035 &  4      &          8.4   &  -17  &  11  &  7  & 117    & i      & 5   & 30 & 40     \\\\ \\object{M64}      &   0.041 &  2      &         17.0   &  -65  &  10  &  5  & 116    &        & 10  & 36 & 10     \\\\  % \\object{M81}      &   0.080 &  2      &          3.6   &  -23  &  27  & 14  & 111    & i      & 5   & 5  & 40     \\\\ \\object{M83}      &   0.066 &  5      &          4.5   &    0  &  13  & 12  & 106    & i      & 6   & 6  & 6      \\\\ \\object{M101}     &   0.009 &  6      &          7.5   &   90  &  29  & 27  &  88    & b      & 23  & 23 & 23     \\\\ \\object{NGC 0055} &   0.013 &  9      &          2.0   &  -72  &  32  &  6  &  96    & i      & 33  &    & 33     \\\\ \\object{NGC 0247} &   0.018 &  7      &          3.1   &   -6  &  21  &  7  & 107    &        & 14  &    & 14     \\\\ \\object{NGC 0253} &   0.019 &  5      &          3.9   &   52  &  28  &  7  & 114    & i      & 34  & 41 & 34     \\\\ \\object{NGC 0300} &   0.013 &  7      &          2.0   &  -69  &  22  & 16  & 106    & i      & 32  &    & 32     \\\\ \\object{NGC 0628} &   0.070 &  5      &         11.2   &   25  &  11  & 10  & 111    &a, h, i & 39  & 36 & 39     \\\\ \\object{NGC 0660} &   0.065 &  1      &         12.4   &  -10  &   8  &  3  & 108    &        &     & 46 & 43     \\\\ \\object{NGC 0891} &   0.065 &  3      &          9.6   &   22  &  14  &  3  & 111    &        & 38  & 37 & 37, 38 \\\\ \\object{NGC 0925} &   0.076 &  7      &          9.3   &  -78  &  11  &  6  & 110    & h, i   &  5  &  5 & 31     \\\\ \\object{NGC 1097} &   0.027 &  3      &         15.3   &  -50  &   9  &  6  & 103    & g      &     & 46 & 29     \\\\ \\object{NGC 1291} &   0.013 &  0      &          9.7   &  -15  &  10  &  8  &  88    &        & 42  &    &        \\\\ \\object{NGC 1512} &   0.011 &  1      &         10.4   &   90  &   9  &  6  &  95    &        &     &    &        \\\\ \\object{NGC 1566} &   0.009 &  4      &         17.5   &   60  &   8  &  7  &  95    & i      &     & 2  &        \\\\ \\object{NGC 2366} &   0.036 & 10      &          3.4   &   25  &   8  &  3  &  95    & f      & 20  &    & 20     \\\\ \\object{NGC 2403} &   0.040 &  6      &          3.2   &  -53  &  22  & 12  & 103    & h, i   & 5   & 5  & 3      \\\\ \\object{NGC 2841} &   0.016 &  3      &         14.1   &  -33  &  08  &  4  & 105    &        & 3,9 & 45 & 3      \\\\ \\object{NGC 2903} &   0.031 &  4      &          8.9   &   17  &  13  &  6  & 106    & h, i   & 5   & 5  & 3      \\\\ \\object{NGC 3198} &   0.012 &  5      &         16.7   &   35  &   9  &  3  & 103    & i      & 4   &    & 8      \\\\ \\object{NGC 3521} &   0.058 &  4      &          9.0   &  -17  &  11  &  5  & 104    & i      &     & 36 & 15     \\\\ \\object{NGC 3621} &   0.080 &  7      &          8.3   &  -21  &  12  &  7  & 105    & i      &     &    &        \\\\ \\object{NGC 3627} &   0.032 &  3      &          9.1   &  -07  &   9  &  4  & 110    &        & 47  & 36 & 47     \\\\ \\object{NGC 4216} &   0.032 &  3      &         17.0   &   19  &   8  &  2  & 108    &        & 16  & 22 &        \\\\ \\object{NGC 4236} &   0.015 &  8      &          4.5   &  -18  &  22  &  7  & 102    &        & 11  &    &        \\\\ \\object{NGC 4395} &   0.017 &  9      &          4.2   &  -33  &  13  & 11  & 104    & h, f   & 44  &    & 44     \\\\ \\object{NGC 4438} &   0.028 &  0      &         17.0   &   27  &   9  &  3  & 106    &        &     &    &        \\\\ \\object{NGC 4559} &   0.018 &  6      &         17.0   &  -30  &  11  &  4  & 104    & i      & 12  &    & 12     \\\\ \\object{NGC 4569} &   0.046 &  2      &         17.0   &   23  &  10  &  4  & 110    &        & 7   & 7  & 7      \\\\ \\object{NGC 4594} &   0.051 &  1      &          9.1   &   90  &   9  &  4  & 108    &        & 1   &    & 1      \\\\ \\object{NGC 4631} &   0.017 &  7      &          9.0   &   86  &  16  &  3  & 106    &        &     &    & 35     \\\\ \\object{NGC 4656} &   0.013 &  9      &          9.0   &   33  &  15  &  3  & 106    &        &     &    &        \\\\ \\object{NGC 4736} &   0.018 &  2      &          5.2   &  -75  &  11  &  9  & 104    & i      & 26  & 36 & 26     \\\\ \\object{NGC 5055} &   0.018 &  4      &          8.2   &  -75  &  13  &  7  & 105    & i      & 9   & 36 & 9      \\\\ \\object{NGC 7331} &   0.091 &  7      &         14.9   &  -09  &  11  &  4  & 104    & i      & 3   & 45 & 3      \\\\ \\object{NGC 7793} &   0.019 &  7      &          2.0   &  -82  &   9  &  6  & 108    & i      & 13  &    & 13     \\\\ \\object{WLM}      &   0.037 & 10      &          1.0   &  -20  &  12  &  4  & 104    & c      &     &    & 21     \\\\  \\enddata \\tablecomments{The columns give (1) Object name, (2) Galactic color excess \\citep{schlegel98} , (3) RC3 morphological type, (4) Distance (Mpc) taken from the GALEX Atlas of nearby galaxies \\citep{gil06}, (5) Position Angle (degree), (6) D25: major axis diameter (minutes), (7) d25: minor axis diameter (minutes), (8) FWHM of the restored 100 microns beam in arcsec (see section \\ref{secrestorebeam} for details), (9) References for O/H gradients, (10), (11), (12) References for HI, CO, and rotation curve data. } \\tablerefs{ {\\bf References for Abundance Gradients [O/H]:} a \\citet{belley92} ;  % b \\citet{kenni03a} ; % c \\citet{lee03aj} ; % d \\citet{lee03aa} ; % e \\citet{masegosa91} ; % f \\citet{roy96} ; % g \\citet{storchi96} ; % h \\citet{vanzee98} ; % i \\citet{zaritsky94}. % {\\bf References for Gas [HI, CO] and Kinematic [V(R)] Data:} 1 \\citet{bajaja84} ; % 2 \\citet{bajaja95} ; % 3 \\citet{begemanthesis} ; 4 \\citet{begeman89} ; % 5 \\citet{boissier03}, and references within ; % 6 \\citet{boissier05}, and references within ; % 7 Boselli et al. (2006), and references within ;  % 8 \\citet{bosmathesis} ; % 9 \\citet{bosma81} ; % 10 \\citet{braun94} ; % 11 \\citet{braun97} ; % 12 \\citet{broeils94} ; % 13 \\citet{carignan7793}; 14 \\citet{carignan247} ; 15 \\citet{casertano91} ; % 16 \\citet{cayatte94} ; % 17 \\citet{corbelli03} ; % 18 \\citet{dame93} ; 19 \\citet{heyer04} ; % 20 \\citet{hunter01} ; % 21 \\citet{jackson04} ; % 22 \\citet{kenney88} ; % 23 \\citet{kenney91} ; % 24 \\citet{lake89} ; % 25 \\citet{martimbeau94} ; % 26 \\citet{mulder93} % 27 \\citet{newton77} ; % 28 \\citet{nieten06} ; % 39 \\citet{ondrechen89} ; % 30 \\citet{paglione01} ; % 31 \\citet{pisano98} ; % 32 \\citet{puche90} ; % 33 \\citet{puche9133} ; % 34 \\citet{puche9134} ; % 35 \\citet{rand94} ;        % 36 \\citet{regan01} ; % 37 \\citet{sakamoto97} ; % 38 \\citet{sancisi79} ; % 39 \\citet{shostak84} ; % 40 \\citet{sofue97} ; % 43 \\citet{sorai00} ; % 44 \\citet{vandriel88} ; % 45 \\citet{vandriel95} ; 46 \\citet{wevers86} ; % 48 \\citet{young82} ; % 50 \\citet{young95} ; % 51 \\citet{zhang93}.}  % \\end{deluxetable}  \\clearpage ",
        "conclusions": "Using the galaxies of large angular size culled from the GALEX Atlas of Nearby Galaxies \\citep{gil06}, we have compiled GALEX UV and IRAS IR images to compute profiles of attenuation and star formation in 43 galaxies. We find typically that UV attenuation decreases, from a few magnitudes of extinction in the centers of galaxies, to low values at larger radii. Combined with the metallicity gradients observed in those same galaxies, this gives rise to a attenuation-metallicity relationship, quantified in section \\ref{secAz}. No clear correlation is found between the attenuation and the gas surface density.  Using the attenuation to correct the FUV profiles for dust, we derive star formation rate profiles with which we test several modern expressions of the classical ``Schmidt law''. Our data overlap with those of \\citet{kenni98a} but extend to significantly lower gas densities suggesting that the UV light in general traces and reveals lower levels of star formation than H$\\alpha$, in coincidence with the recent GALEX discoveries of UV extended (XUV) disks \\citep[e.g.][]{gildepaz05,thilker05}.  We suggest that the much-debated threshold of star formation might have been an observational selection effect caused by the extremely low number of ionizing stars expected to be found beyond the ``threshold'' radius. Indeed, for low levels of star formation, the H$\\alpha$ radiation in one elliptical annuli seems stochastic when the number of massive stars in this annuli becomes close to $\\sim$ 1. The UV radiation coming from less massive and more long-lived stars (thus in larger number) does not present a break in its radial profile at the H$\\alpha$ threshold radius.   While these results are already interesting in their own right, it is obviously time to look to more galaxies, using Spitzer data in combination with GALEX. Spitzer will allow us to include many more galaxies in such studies owing to its better resolution than IRAS. Indeed, a joint analysis of GALEX and Spitzer data is already underway in a SINGS-GALEX collaboration. Note that the SINGS sample will only partially overlap with the one presented here as some of the closest galaxies were not included in SINGS."
    },
    "0609/astro-ph0609247_arXiv.txt": {
        "abstract": "X-ray observations of galaxy clusters potentially provide powerful cosmological probes if systematics due to our incomplete knowledge of the intracluster medium (ICM) physics are understood and controlled. In this paper, we present mock {\\it Chandra} analyses of cosmological cluster simulations and assess X-ray measurements of galaxy cluster properties using a model and procedure essentially identical to that used in real data analysis.  We show that reconstruction of three-dimensional ICM density and temperature profiles is excellent for relaxed clusters, but still reasonably accurate for unrelaxed systems.  The total ICM mass is measured quite accurately ($\\lesssim$6\\%) in all clusters, while the hydrostatic estimate of the gravitationally bound mass is biased low by about 5\\%--20\\% through the virial region, primarily due to additional pressure support provided by subsonic bulk motions in the ICM, ubiquitous in our simulations even in relaxed systems. Gas fraction determinations are therefore biased high; the bias increases toward cluster outskirts and depends sensitively on its dynamical state, but we do not observe significant trends of the bias with cluster mass or redshift. We also find that different average ICM temperatures, such as the X-ray spectroscopic $\\Tspec$ and gas-mass-weighted $\\Tmg$, are related to each other by a constant factor with a relatively small object-to-object scatter and no systematic trend with mass, redshift or the dynamical state of clusters. We briefly discuss direct applications of our results for different cluster-based cosmological tests. ",
        "introduction": "\\label{sec:intro}  X-ray observations of clusters of galaxies can potentially provide plethora of useful cosmological information. The cluster-based cosmological tests include cluster number counts based on the temperature function, baryon fraction and X-ray and SZ Hubble constant measurements.  These cosmological tests provide potentially very powerful constraints on the matter density, $\\Omega_M$, $\\Omega_{\\Lambda}$, $\\sigma_8$, and the equation of state of dark energy, $w$. In all of these cosmological tests, the key observational ingredients are the gas mass ($\\Mgas$), total mass ($\\Mtot$), gas mass fractions ($\\fgas\\! \\equiv\\!  \\Mgas/\\Mtot$), and average cluster temperature ($\\Tave$).  In the era of precision cosmology, it is paramount to achieve accurate measurements of the key cluster parameters, check how they depend on various simplifying assumptions, and control systematics due to our incomplete knowledge of the ICM physics.  There are many important simplifying assumptions used in deriving cluster properties from X-ray data. For example, it is usually assumed that the ICM density is a function of radius only and does not have small-scale substructure. However, substructure in clusters is ubiquitous and this biases the X-ray derived $\\Mgas$ high \\citep{mathiesen_etal99}. Also, the small-scale clumps are often associated with subhalos and thus have lower temperature. Neglecting this can bias the average spectroscopic temperature of the cluster as a whole \\citep{mathiesen_etal01,mazzotta_etal04} which could potentially lead to biases in constraints on cosmological parameters \\citep{rasia_etal05}. Cluster total masses are often estimated assuming the hydrostatic equilibrium for the ICM.  This may underestimate the true gravitational bound mass if non-thermal pressure support is present.  For example, turbulent gas motions can provide about 10\\%--20\\% of the total pressure support even in relaxed clusters and hence bias the hydrostatic estimate \\citep{schuecker_etal04,faltenbacher_etal05,rasia_etal06,lau_etal06}. Spherical symmetry is another commonly used assumption useful in solving for three-dimensional (3D) physical cluster properties from two-dimensional (2D) observed quantities, but could bias results if clusters are aspherical.  X-ray observations with {\\it Chandra} and {\\it XMM-Newton } enable us to study properties of the ICM with unprecedented detail and accuracy and provide important handles on the ICM modeling and associated systematics. Their superb spatial resolution and sensitivity enable accurate X-ray brightness and temperature measurements at a large fraction of the cluster virial radius and also make it simple to detect most of the small-scale X-ray clumps. Despite this recent observational progress, the biases in the determination of the key cluster properties remain relatively uncertain. The main obstacle is that because of their unrivaled statistical accuracy, the X-ray results cannot be contrasted against any other independent observational techniques. In this study, we attempt to check the validity of the X-ray analyses, with a specific focus on the analyses of high-resolution {\\it Chandra} data, by using mock observations of clusters derived from cosmological simulations, for which the true answers are known.  Our simulations properly treat both collisionless dynamics of dark matter and stars and gasdynamics and capture a variety of physical phenomena from the nonlinear collapse and merging of dark matter to shock-heating and radiative cooling of gas, star formation, chemical enrichment of the ICM by supernova and energy feedback. These simulations can therefore be used to test observational biases due to incomplete relaxation of gas, the dynamical state of a cluster, substructure, or non-isothermality. As we show in the forthcoming papers, the current simulations match the observed ICM profiles outside cluster cores \\citep{nagai_etal06c} and the global gas mass fraction \\citep {kravtsov_etal06b} and are therefore sufficiently realistic for a purpose of the current study.  The mock observations of the simulated clusters are generated by exactly reproducing the response properties of the \\emph{Chandra} telescope. The analysis of the mock data is identical to that used for \\emph{real Chandra} observations by \\cite{vikhlinin_etal05c,vikhlinin_etal06}.  The comparison of the true and derived cluster properties provides an assessment of any biases introduced by the X-ray analysis. Some of these biases are potentially redshift-dependent (e.g, those related to higher merger rate or decreased ability to detect small-scale substructures in the X-ray data for high-$z$ systems). To check for any redshift dependence in such biases we use the simulation outputs at $z=0$ and $0.6$. Our results indicate that the X-ray analysis provides accurate reconstruction of the 3D properties of the ICM. The strongest biases we find are those in the hydrostatic mass estimates. They are related to physics explicitly missing from the hydrostatic method (e.g., turbulence), and not to deficiencies of the X-ray analysis. ",
        "conclusions": "\\label{sec:discussion}  We present mock \\emph{Chandra} analyses of cosmological cluster simulations and assess X-ray measurements of galaxy cluster properties.  To test observational X-ray procedures, we construct mock {\\it Chandra} images of the simulated clusters and derive properties of X-ray clusters using a model and procedure essentially identical to those used in real data analysis. The sample includes 16 clusters spanning a broad mass range ($5\\times 10^{13}-2\\times 10^{15}\\,h^{-1}\\,M_{\\odot}$) simulated in the $\\Lambda$CDM cosmology with high spatial resolution and including various physical processes of galaxy formation, such as radiative cooling, star formation and other processes accompanying galaxy formation.  We analyze the simulated clusters at $z=0$ and their most massive progenitors at $z=0.6$.  To test the sensitivity of the X-ray analysis to the cluster dynamical state and substructure, we also distinguish unrelaxed clusters from relaxed systems based on the overall structural morphology of the \\emph{Chandra} images, as usually done to classify observed clusters. Below we discuss and summarize our main findings and conclusions.   \\begin{deluxetable}{cccccc} \\tablecaption{Normalizations of the $M_{500c}-T$ relation\\label{tab:mt}} \\tablehead{ & & & \\multicolumn{3}{c}{$M_5$, $10^{14}\\,h^{-1}\\,M_\\odot$}\\\\ \\cline{4-6}\\\\[-1.7ex] \\colhead{Mass\\tablenotemark{a}} & \\colhead{Redshift} & \\colhead{Sample} & \\colhead{$\\Tmg$} & \\colhead{$\\Tspec$} & \\colhead{$\\Tspec^{[0.5\\r500c]}$} } \\startdata &       & all       & $4.46\\pm0.18$ & $3.78\\pm0.21$ & $3.40\\pm0.17$ \\\\ $M^{\\rm true}_{500c}$&$z=0$  & relaxed   & $4.02\\pm0.10$ & $3.32\\pm0.10$ & $3.00\\pm0.09$ \\\\ &       & unrelaxed & $4.80\\pm0.19$ & $4.14\\pm0.22$ & $3.71\\pm0.18$ \\\\ \\hline &       & all       & $4.65\\pm0.17$ & $4.07\\pm0.17$ & $3.81\\pm0.16$ \\\\ $M^{\\rm true}_{500c}$&$z=0.6$& relaxed   & $4.50\\pm0.12$ & $3.77\\pm0.12$ & $3.51\\pm0.12$ \\\\ &       & unrelaxed & $4.73\\pm0.18$ & $4.22\\pm0.17$ & $3.96\\pm0.16$ \\\\ \\hline \\hline $M^{\\rm est}_{500c}$ &$z=0$  & relaxed   & $3.14\\pm0.03$ & $2.59\\pm0.03$ & $2.35\\pm0.04$ \\\\[1mm] \\hline $M^{\\rm est}_{500c}$  &$z=0.6$& relaxed   & $3.55\\pm0.07$ & $2.97\\pm0.06$ & $2.77\\pm0.06$\\\\[-2mm] \\enddata \\tablenotetext{a}{$M^{\\rm true}_{500c}$ is the true $\\Mtot$ of a cluster enclosed within the true $\\r500c^{\\rm true}$ measured directly in simulations; $M^{\\rm est}_{500c}$ is the estimated hydrostatic mass enclosed within the estimated $\\r500c^{\\rm est}$ from the mock data analysis.} \\end{deluxetable}  \\subsection{Reconstruction of 3D ICM Properties}  Derivation of the cluster 3D properties from projected X-ray data usually relies on simplifying assumptions such as spherical symmetry of a cluster or small-sale uniformity of the ICM. Additional complications arise from a non-ideal response of X-ray telescopes and limited statistical quality of data. This leads to several effects which are hard to reproduce analytically without creating mock X-ray data from realistic cluster simulations. For example, clumps of colder gas associated with massive subhalos can be detected and masked out from a further analysis while smaller subhalos remain undetected and bias projected temperature and X-ray brightness measurements \\citep{mazzotta_etal04,vikhlinin_etal06b}. One of the goals of our work is to include all such effects as completely as possible in generation of the mock X-ray data and thus to test how they bias X-ray measurements and recovery of 3D ICM profiles.  Reassuringly, we find that the ICM models and analysis method presented in \\citet{vikhlinin_etal06} provide a good description of emission measure and projected temperature profiles of the ICM.  For relaxed clusters, these models recover the 3D gas density, temperature, and gas mass profiles with the accuracy of {\\it a few percent} at $0.1<r/\\r500c<2.0$.  The models are flexible enough to describe also non-relaxed systems and provide reasonably accurate gas mass and temperature measurements. A practical implication of these results is that projection effects, cluster substructure, and deviations from spherical symmetry do not strongly affect reconstruction of 3D ICM density, temperature, pressure, and entropy profiles from X-ray data.  Therefore, observational results on ICM profiles can be directly compared to the properties of simulated clusters without a need for detailed mock X-ray analyses. Such a comparison for our sample will be presented in the future paper \\citep{nagai_etal06c}.  The total ICM mass is measured quite accurately in all our simulated clusters. Results are most accurate for relaxed systems: biases and scatter in \\Mgas\\ measurements are accurate to better than $\\sim 3\\%$ percent, independent of a cluster redshift or a radial range within in which the \\Mgas\\ is measured.  For unrelaxed systems, the scatter is larger, as expected \\citep{mathiesen_etal99}, but the bias is still small ($\\lesssim 10\\%$).  Average temperature is an important diagnostic of properties and physical processes of the ICM and a key observable for cosmological applications. Since clusters are not isothermal, the definition of the $\\Tave$ is not unique nor can $\\langle T\\rangle$ be completely accurately derived from data. Mock X-ray analysis is therefore required to study the relation among different $\\langle T\\rangle$ definitions used in observational and theoretical studies.  Our analysis is focused on the comparisons of the true gas-mass-weighted average temperature, $\\TmgSIM$, and the X-ray spectroscopic $\\Tspec$ or gas-mass-weighted $\\Tmg$ temperatures derived from the mock data. We find that these temperatures are not identical, but the difference is mostly a constant factor with a relatively small object-to-object scatter.  For example, $\\Tspec /\\Tmg \\approx 1.14\\pm0.05$ for relaxed clusters, while for unrelaxed objects this ratio is $\\Tspec/\\Tmg \\approx 1.12\\pm0.09$.  All the $\\Tave$ definitions we tested are therefore equivalent, but the difference should be kept in mind when, e.g., observed $M-T$ relations are compared with results of simulations.  A notable exception is the ``emission-weighted'' temperature defined as $\\Tew = \\int T\\, \\rho^2\\,T^{1/2}\\,dV \\big/ \\int \\rho^2\\,T^{1/2}\\,dV$ \\citep{bryan_etal98,frenk_etal99,muanwong_etal01}. Even computed directly from the simulation outputs, this quantity shows a much poorer correlation with $\\TmgSIM$ than any of the observationally derived average temperatures (Fig.\\ref{fig:tbias_z0_true}). Our results thus support the suggestion of \\citet{mazzotta_etal04} that $\\Tew$ should not be used as a measure of the average cluster temperature.   \\subsection{Accuracy of Hydrostatic Mass Estimates} \\label{sec:accur-hydr-total}  One of the major applications of the gas density and temperature profile measurements from the X-ray data is the estimate of a gravitationally bound mass of a cluster assuming that the ICM is in hydrostatic equilibrium in the cluster potential. The X-ray method is by far the most precise way to estimate $\\Mtot$ in individual objects at large radii. This limits our ability to test biases in the X-ray derived $\\Mtot$ estimates through comparisons with independent observational techniques such as weak lensing. The analysis of our mock X-ray data, however, provides important clues.  We find that the hydrostatic mass estimates from our mock data are biased low by about 5\\%--20\\% throughout the virial region (Fig.\\ref{fig:mbias_z0_true}), even in clusters that are identified as relaxed by their X-ray morphology. Such biases seem to be too high to be attributed to inaccuracies in 3D gas density and temperature profile reconstruction. A more plausible explanation is a departure of the ICM from a hydrostatic state. Indeed, an additional pressure component due to residual random bulk motions of the ICM in our simulated clusters completely explains the bias in the hydrostatic mass estimates \\citep[this topic will be discussed in greater detail in][]{lau_etal06}.  Similar results have been also obtained by independent simulations performed with different codes and implementation of physical processes \\citep{rasia_etal06,dolag_etal05b}. The $\\Mtot$ bias derived in the mock X-ray analysis of our relaxed clusters increases towards their outskirts. The average levels are $-8\\%$ at $r=\\R2500c$ and $-13\\%$ at $r=\\r500c$, and quickly become very large beyond $\\r500c$ (Fig.\\ref{fig:CL7csf1_z0}), which qualitatively mimic the radial dependence of turbulent pressure observed in previous simulations \\citep[e.g.,][]{evrard_etal96}. In unrelaxed clusters, not only the departures from hydrostatic equilibrium are strong but also the density and temperature gradients are not measured accurately (Fig.\\ref{fig:CL101csf1_z0}) resulting in $\\Mtot$ biases of 20\\% or more at all radii.  It is unclear to what degree the level of ICM turbulence found in these simulations applies to real clusters because most of the codes model ICM as an ideal inviscid fluid with some amount of numerical viscosity.  Further progress in this area using cosmological simulations will have to rely on better understanding and modeling of viscosity of the weakly magnetized plasma \\citep{ruszkowski_etal04a,ruszkowski_etal04b,dolag_etal05b,sijacki_etal06}. Direct measurements of the ICM turbulence through Doppler broadening of emission lines \\citep{inogamov_etal03} must await a launch of X-ray calorimeters, although some clues could be provided also by observations of X-ray brightness fluctuations \\citep{schuecker_etal04}. Keeping these uncertainties in mind, we can treat the $\\Mtot$ biases quoted above as upper limits.  Cosmic rays and magnetic fields may also provide non-thermal pressure support which is not accounted for in the hydrostatic estimates nor included in our simulations. In addition, temperatures of electron and ion components of the ICM are not necessarily the same, because strong shocks heat primarily ions and the collisional electron-ion equilibration time is long in the cluster outskirts \\citep{markevitch_etal96,fox_etal97}. If $T_e$ (derived from X-ray spectra) is indeed lower that $T_i$ (unobserved), the hydrostatic mass estimates are biased low. We note, however, that at least in one case there is evidence for fast equilibration of the electron and ion temperatures in a post-shock region of the ICM \\citep{markevitch05}.  While direct measurements or theoretical modeling of the non-thermal pressure support may be difficult, it can be constrained by comparisons of the hydrostatic mass estimates with those derived from gravitational lensing for the same objects. An important caveat is that lensing measures the mass in the projected aperture while the mass within the sphere of the same radius is rather uncertain because of projection effects.  \\citet{metzler_etal01} estimate that $M_{500c}$ cannot be determined from weak lensing with better than $\\pm$30\\% uncertainty. Any comparisons of X-ray hydrostatic and weak lensing masses must be done in an average sense, e.g.{} through derived normalizations of the correlation between the $\\Mtot$ and a low-scatter mass proxy such as $Y_X$ \\citep{kravtsov_etal06} or $\\Tspec$.  To summarize, our results indicate that an X-ray analysis of relaxed clusters correctly recovers gas density and \\emph{thermal electron} pressure gradients. The simulations provide an upper limit (5\\%--20\\% depending on radius) for the mass bias due to turbulent motions.  The contribution of other pressure components will have to be constrained by future observations.   \\subsection{Implications for Cluster-Based Cosmological Tests}   \\subsubsection{Tests Based on $\\fgas$}  X-ray measurements of the baryonic mass fraction in massive clusters can be used to constrain $\\Omega_M$ or $h$ \\citep{white_etal93} and also provide a potential standard candle \\citep{sasaki96,pen97,allen_etal04}.  The major observational ingredient in these tests is a measurement of the hot gas mass fraction within a sufficiently large radius. Our analysis provides clues for biases in $\\fgas$ measurements from high-quality \\emph{Chandra} data. We find that the $\\fgas$ determinations are biased high, primarily because the hydrostatic method underestimates the total cluster mass. As was discussed above, the bias in the hydrostatic estimates is related to the physical processes in the ICM (e.g., turbulence, cosmic ray and magnetic pressures etc.) and not to the deficiencies of the X-ray analysis. Further progress in confirming validity of $\\fgas$ measurements will have to rely on independent determinations of cluster total mass (see discussion in \\S~\\ref{sec:accur-hydr-total}). Some sources of the bias (e.g., turbulence) could be minimized if we focus only on relaxed systems. We find that in such systems, both bias and scatter in $\\fgas$ can be controlled to within about 5\\%--8\\% at $\\R2500c$ and 15\\%--20\\% at $\\r500c$, which can be considered the upper limits on the $\\fgas$ measurement bias caused by turbulence (see \\S~\\ref{sec:accur-hydr-total}).  Note that if the ICM turbulence and other processes discussed in \\S~\\ref{sec:accur-hydr-total} do play a role in biasing the hydrostatic mass estimates, the bias can potentially be redshift-dependent since the merger rate increases at higher $z$ \\citep[e.g.][]{gottloeber_etal01}. We do not find any detectable redshift dependence of the bias in the $\\fgas$ measurements in relaxed systems. However, our upper limits, $\\sim\\! 5\\%$, are weaker than the accuracy required, e.g., for using $\\fgas$ as a standard ruler in the dark energy studies \\citep{allen_etal04}. Mock analysis of much larger cluster samples will be required to properly address this issue.  \\subsubsection{Tests Based on the Cluster Mass Function}  Another group of cosmological tests are those based on measurements of the cluster mass function. The relevant issue for our work is how well various mass proxies can be measured. A related question is how well the cluster mass vs.{} proxy relation can be calibrated by the X-ray analysis, which was discussed in \\S~\\ref{sec:accur-hydr-total}.  We did not discuss the easiest X-ray mass proxy, the total luminosity, because it is also the least accurate and least reliably reproduced by numerical simulations. The next, in terms of being easily measured, is the average ICM temperature. The practical averages to use are $\\Tspec$ or $\\TspecIV$\\footnote{Other definitions such as $\\Tmg$ require measurements of the temperature \\emph{profiles} which is feasible only for a few best-observed clusters.}. These definitions are freely interchangeable since the ratio $\\TspecIV/\\Tspec$ is nearly the same in all clusters (Fig.\\ref{fig:tbias_z0_true}, Table~\\ref{tab:tbias1}).  One disadvantage of the spectral temperature as a mass proxy is that its measurement is somewhat sensitive to the cluster dynamical state and small-scale substructure. This sensitivity arises because $\\Tspec$ is rather sensitive to the presence of lower-temperature components. This effect contributes to the systematic difference in the normalizations of the $M-T$ relation between relaxed and non-relaxed clusters and also slight deviations from the self-similar evolution (Table~\\ref{tab:mt}).  Some implications of this effect for the cosmological measurements are discussed by \\citet{rasia_etal05}. The prospects for calibrating the $\\Tspec$ measurement biases by cosmological numerical simulations are mixed. The effects related to the cluster dynamics, such as merger activity and shocks etc.{}, are treated sufficiently accurately. However, the biases caused by the presence of undetected cold clumps will require a more realistic modeling of the ICM cooling and associated feedback.  More accurate X-ray proxies for cluster total mass is the mass of the hot gas, $\\Mgas$, and the ``X-ray $Y$-parameter'', $Y_X=\\Mgas\\times\\Tspec$ \\citep{kravtsov_etal06}. The properties of $Y_X$ (this is the best X-ray mass proxy known) are extensively discussed elsewhere \\citep{kravtsov_etal06}, so we concentrate on $\\Mgas$ here. The $\\Mgas$ is measured very reliably. The average bias is $<5\\%$ with only a few strong outliers (Fig.\\ref{fig:mbias_z0_true}) and there is no detectable redshift dependence. A similar accuracy should in principle be achieved for the $\\Mtot$ estimates since to the first order, $\\Mtot\\propto\\Mgas$. Note that the biases in $\\Mgas$ measurements are mostly due to large-scale deviations of the cluster body from spherical symmetry. This effect should be reliably treated in the current generation cosmological simulations, and so the accuracy of the $\\Mgas$-based total mass estimates can be improved yet further. All we need is a larger sample of simulated clusters. A major obstacle, however, is the uncertainties in the mass fraction of cold baryons (stars and molecular gas). This fraction can be substantial, as indicated by our simulations \\citep{kravtsov_etal05} and significant deviations of the observed $\\fgas$ in clusters from the cosmic mean \\citep[e.g.,][]{ettori03,vikhlinin_etal06}. Further progress in this area requires significant improvements in the reliability of the stellar mass measurements and treatment of the star formation in the simulations. We will further discuss these issues in the forthcoming papers.  \\smallskip  To briefly summarize our results, we find that the analysis of deep X-ray observations of galaxy clusters can reliably recover its intended targets: the distribution of density and temperature of the hot ICM. The accuracy of estimating the total cluster mass from these parameters is limited by additional physical processes in the ICM, such as turbulence, magnetic and cosmic ray pressure, and possible departures from the electron-ion equilibrium. Our work provides a realistic estimate of the effects of turbulence, which are found to cause underestimation of $\\Mtot$ by 5\\%--20\\% in clusters that can be visually classified as relaxed (this can be considered as an upper limit for the bias if the ICM viscosity is non-negligible). Constraining the role of other effects will probably have to rely on accurate independent mass measurements (e.g., by weak lensing) in a large, complete cluster sample."
    },
    "0609/astro-ph0609137_arXiv.txt": {
        "abstract": "The consortium for Galactic studies with the Arecibo L-band Feed Array (ALFA) is conducting a neutral hydrogen (HI) survey of the whole Arecibo sky (declination range from $-1$\\degree~to 38\\degree), with high angular (3.5$'$) and velocity resolution (0.2 \\kms). The precursor observations with ALFA of a region in the Galactic anti-center reveal numerous isolated, small (a few pc in size), and cold ($T_{\\rm k}<400$ K) HI clouds at low negative velocities, distinctly separated from the HI disk emission (`low-velocity clouds', LVCs). These clouds are most likely located in the transition region between the Galactic disk and halo (at scale heights of 60--900 pc), yet they have properties of typical cold neutral clouds. LVCs are colder and, most likely, smaller and less massive than Lockman's clouds in the disk/halo interface region of the inner Galaxy. Our observations demonstrate that the cloudy structure of the interface region is most likely a general phenomenon, not restricted to the inner Galaxy. LVCs have sizes and radial velocities in agreement with the expectations for clouds formed in low-temperature fountain flows, although we measure a factor of ten higher HI column densities. Alternatively, LVCs could represent the final stages of the infalling intergalactic material in the on-going construction of the Galaxy.  In the same dataset at higher negative velocities, we have discovered a `companion' HI cloud located 50$'$ southwest of CHVC186+19-114. CHVC186+19-114 is a typical compact high velocity cloud (HVC) with a well-defined core/envelope structure. The companion cloud has a diameter of only $7'\\times9'$, and is one of the smallest HVCs known, most likely stripped from the main cloud through the interactions with the halo medium. ",
        "introduction": "The Galactic halo, a hot, ubiquitous multi-phase gas, with a low density, and temperature around $10^{6}$ K, that surrounds our Galaxy is crucial to Galactic activity. As well as acting as a thermal insulator, the halo constrains models of the interstellar medium (ISM) in the Galactic disk through its impact on the  dynamics of the hot intercloud component. From the theoretical point of view, the Galactic fountain model \\citep{Shapiro76} was the first to address interactions between the disk and the halo. In this model, the hot gas flows from the disk into the halo, starts to cool, and falls back onto the disk. One of the later models, so called chimney model of the ISM \\citep{Norman89}, suggested that the disk and the halo are connected through chimneys -- large conduits resulting from superbubbles bursting out of the disk and forming collimated structures that can reach heights of $\\sim1$ kpc. From the observational point of view, numerous ISM structures (shells, superbubbles, worms, chimneys, plumes) were found in the vertical gas distribution of our own Galaxy and several nearby galaxies, and have been postulated to play an important role in the transfer of mass and energy between galactic disks and their halos. However, the details of how these structures replenish the hot halo gas are still not fully understood \\citep{Dove00}.   In understanding the ways in which the Galactic disk and halo are related the transition region between the Galactic thin disk and the halo, also known as the Galactic thick disk, is of special interest. This region extends to heights up to 1--1.5 kpc above the plane. The disk/halo interface region was discovered several decades ago \\citep{Shane67,Lindblad67} and was considered to represent a smooth envelope of neutral hydrogen (HI) surrounding the Galactic spiral structure. Further studies \\citep{Albert83,Danly89,Lockman86} found occasional isolated halo clouds at lower heights and suggested that this component could be pervasive, yet patchy, and most likely dynamically connected with the disk.   Galactic fountain and chimney models, as well as recent numerical simulations \\citep{Avillez00}, predicted the existence of small-scale, cloudy structure in the interface region, with clouds ranging in size from a few parsecs to  tens or hundreds of parsecs. Recent observations with the Green Bank telescope (Lockman 2002) have actually revealed dramatic small-scale structure of the interface region in the inner Galaxy, composed of numerous discrete HI clouds. While this confirms the expectations from theoretical models and simulations, the full extent and the origin of this cloudy component is still not understood.   Another set of attractive questions concerning the  disk/halo interface reagion comes from the idea that the formation of galaxies is an ongoing process \\citep{Oort66} and that the accretion of intergalactic material (IGM) serves as the main source of new star formation fuel \\citep{Maller04,Murali02}. However, the details of the transformation from the smooth and hot IGM on kpc scales, down to clumpy and cold gas on star-forming scales, still need to be worked out. For example, what is the internal structure of IGM clouds at different stages of their infall? As searches for accretion remnants around other galaxies are difficult due to sensitivity and area limitations (e.g., Fraternali et al. 2002), our own Galaxy is still one of the best places to start seeking answers to these questions.   HI clouds in the Galactic halo at heights $>1$ kpc have been studied observationally for several decades and are often  classified into several distinct groups based on how much their velocities deviate from Galactic rotation (intermediate and high velocity clouds, HVCs), and their relative size (compact and extended clouds). Velocities reach approximate $+/- 450$ \\kms~in the LSR reference frame and their sizes vary from over 100 deg$^{2}$ to compact HVCs with sizes of 4--30 arcmin \\citep{Putman02,Bruns04a}. In addition, almost all HVCs, especially the more compact ones, clearly show head/tail or core/envelope structures, with heads or cores having a narrow linewidth while tails and envelopes having a significantly broader velocity structure, possibly showing signs of interaction with the ambient medium (e.g. Br\\\"{u}ns \\& Mebold 2004). It is still not clear whether this range of observational properties, especially the low-end of the cloud size and velocity distributions, is caused by observational limitations. What determines the size and morphology of the observed halo clouds?  Are these remnants of very different physical processes, or do different `types' of clouds represent different stages of the same phenomenon --- accretion onto the Galaxy, for example?  The questions above are just a subset of issues that still await to be addressed. The main purpose of this paper is to summarize the ongoing observational efforts to understand the properties of gas clouds from the halo to the plane by studying the disk/halo interface of our Galaxy. This work is inspired and enabled by the new multi-beam system, very recently installed at the Arecibo telescope\\footnote{The Arecibo Observatory is part of the National Astronomy and Ionosphere Center, operated by Cornell University under a cooperative agreement with the National Science Foundation.}, which allows sensitive, large-scale HI surveys with an angular resolution of 3.5$'$. Previous large-scale surveys of the link between the disk and the halo have been limited by resolution (Leiden-Dwingeloo Survey, 36$'$; \\cite{Hartmann97}), or an inability to trace features from the halo into Galactic emission (Putman et al. 2002). Recent high-resolution Galactic Plane surveys \\citep{Taylor03,McClure05} have been limited to tracing the lower velocity and latitude HI emission and thus also do not provide adequate information on the disk/halo interface.  The structure of this paper is organized as follows. We start in Section~\\ref{s:galfa_background} by briefly introducing the consortium for Galactic studies with the Arecibo L-band Feed Array (GALFA) and emphasizing the importance of the Arecibo telescope for large-scale Galactic surveys. Section~\\ref{s:obs} describes the observing strategy adopted for Galactic HI spectral-line surveys with ALFA. In particular, the fast scanning mode utilizing the basket-weave technique, and a newly developed method for obtaining reference HI spectra by multiple frequency switching, are applicable to radio telescopes in general. In Section~\\ref{s:background} we summarize major previous observational studies of HI in the Galactic anti-center. We then present our results in Sections~\\ref{s:results} and \\ref{s:hvc}, and provide extensive discussion of observational and theoretical issues in Sections~\\ref{s:discussion-obs} and \\ref{s:discussion-theory}. We summarize our results and point out future studies in Section~\\ref{s:summary}. ",
        "conclusions": "\\label{s:discussion-obs} \\begin{table*} \\footnotesize \\caption{Derived (median) properties of HI clouds and filaments at low and high negative velocities. } \\centering \\label{table2} \\begin{tabular}{lcccc} \\noalign{\\smallskip} \\hline \\hline \\noalign{\\smallskip} & T$_{\\rm k,max}$ & $N({\\rm HI})_{\\rm peak}$ & $n \\times D$ & $nT_{\\rm k,max} \\times D$\\\\ &  (K)    &($10^{19}$ cm$^{-2}$)     & (cm$^{-3} \\times {\\rm kpc}$) & ($10^{3}$ K cm$^{-3} \\times {\\rm kpc}$)  \\\\ \\hline Low-velocity Clouds & 200--1300 & 1.0--3.0 & 1.3-5.5 & 0.3-4 \\\\ Low-velocity Filaments& 9000--20000 & 4--12 & 0.7-1.5 & 10-20 \\\\ \\hline CHVC186+19-114 Core & 350--1000  &2--11 & 0.9 & 1    \\\\ CHVC186+19-114 Envelope  &5600  &2--13 & 0.8 & 4  \\\\ CHVC186+19-114 Companion&5000  &5 & 3.2 & 16  \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{table*}   \\subsection{Distances and HI masses of low-velocity clouds}  A major unknown parameter of LVCs is their distance. As we have shown in Section~\\ref{s:properties_lvc}, 80\\% of LVCs in our sample have $T_{\\rm k}<400$ K and FWHM of only 3--4 \\kms. The peak $N({\\rm HI})$ for all clouds is (1--3)$\\times10^{19}$ cm$^{-2}$. These linewidths and column densities are very similar to those of ordinary CNM clouds at $|b|>10$\\degree~studied through HI absorption measurements \\citep{Heiles03b}, suggesting that the anti-center LVCs are likely to be ``true'' CNM clouds. As mentioned earlier, LVCs often appear related to larger, warmer filamentary structures, with a typical FWHM of 20-30 \\kms.  In the Galaxy, the CNM and WNM can co-exist only over a certain, well-defined range of thermal pressures, $P_{\\rm min}$ to $P_{\\rm max}$ [this was first emphasized by \\cite{Field69} and most recently re-examined  by \\cite{Wolfire03}]. This co-existence of cold LVCs with warm filamentary structures can be used to estimate cloud distances. As the ISM is highly turbulent, often turbulent pressure is more important than the thermal pressure. However, as nicely explained in \\cite{Wolfire03}, on the scales of typical CNM clouds thermal pressure dominates, and CNM clouds are expected to be embedded in a spatially isobaric medium with their surface layers having the same thermal pressure as the ambient warm medium. \\cite{Wolfire03} calculated $P_{\\rm min}$ and $P_{\\rm max}$ in the Galactic midplane for several Galactocentric radii ($R_{\\rm g}$). Both quantities decrease with $R_{\\rm g}$. For example, at $R_{\\rm g}=8.5$ kpc, $P_{\\rm min}=1960$ K cm$^{-3}$ and $P_{\\rm max}=4810$ K cm$^{-3}$, while at $R_{\\rm g}=18$ kpc, $P_{\\rm min}=272$ K cm$^{-3}$ and $P_{\\rm max}=1220$ K cm$^{-3}$.    We can now calculate for our LVCs their volume density ($n$) and thermal pressure ($P$) using: $n  \\propto N({\\rm HI})_{\\rm peak} \\times D^{-1}$ and $P \\propto nT_{\\rm k,max}$, where $T_{\\rm k,max}=21.86 {\\Delta v}^2$. Both $n$ and $P$ depend on the cloud distance $D$, also these equations assume that clouds are spherically symmetric (please note that sheets will have higher pressure, while cigars oriented towards the observer will have lower pressure). Estimates for both $n \\times D$ and $P \\times D$ are given in Table 2, together with the estimated maximum kinetic temperature and the peak HI column density. We can now compare measured values for $P \\times D$, for each cloud, with the thermal pressure allowed for the CNM and WNM co-existence to estimate $D$. If we assume the typical thermal pressure in the Solar neighborhood of 3000 K cm$^{-3}$, then clouds in our sample have the distance range 0.1--1.3 kpc, with the median distance being 0.3 kpc. The lowest thermal pressure that allows the CNM and the WNM to co-exist is $\\sim270$ K cm$^{-3}$ \\citep{Wolfire03}, this places an upper limit on the distance to our clouds of $\\sim3.6$ kpc (this value is the median for the whole cloud sample). The highest thermal pressure that allows the CNM and the WNM to co-exist is $\\sim5000$ K cm$^{-3}$ (for $R_{\\rm g}\\ga8.5$ kpc as our clouds are in the direction of the Galactic anti-center), this places a lower limit on the distance to our clouds of $\\sim0.2$ kpc. If the thermal pressure is, however, $>5000$ K cm$^{-3}$ then cloud distances could be $<0.2$ pc and our clouds may be very closeby. Such high thermal pressures are not expected to be common in the Galaxy for $R_{\\rm g}>8.5$ kpc from the theoretical point of view. If present they would indicate the existence of very short-lived, transient clouds. The observed morphology of LVCs, in the form of discrete, isolated clouds with sharp edges and appearing over many velocity channels, does not suggest their transient character. However, future direct distance measurements are essential to confirm cloud distances. We also caution that our distances were estimated under the assumption of the {\\it local} pressure equilibrium at cloud surfaces. Several recent observations and numerical simulations have challenged the global pressure equilibrium condition for the interstellar gas, however still mildly favoring the local pressure equilibrium state (for example see Mac Low et al. 2005).   Table 2 also gives $P \\times D$ for low-velocity filaments, which appear higher than for the clouds and are most likely due to a non-spherical, sheet-like morphology. This is supported by the visual appearance of filaments which have a far more elongated morphology than the clouds.  If we assume for now that the LVCs have distances in the range 0.2--3 kpc, then their height from the plane is in the range $60$--900 pc. As the typically assumed scale height (half-thickness) of the Galactic thin disk is $z_{\\rm thin}=100$ pc \\citep{Belfort84,Kulkarni88,Avillez00}, most of the LVCs appear located in the transition region between the thin disk and the hot Galactic halo. With a typical angular radius of 9$'$, LVCs' linear size is 0.5--8 pc, while their HI mass ranges from $3 \\times 10^{-2}$ to 7 M$_{\\odot}$. If the clouds are embedded in the warm medium with a temperature of 6000 K, then the mass-loss rate due to evaporation (calculated using eq. [47] in \\cite{McKee77}) is $1.5\\times10^{24}$ gr yr$^{-1}$, and hence the cloud evaporation timescale is $\\ga100$ Myr.   Several previous observational studies have found cold clouds at significant heights above the disk in our own Galaxy and in several other spiral galaxies. \\cite{Crawford02} obtained NaI and CaII observations towards the opening of the Local Interstellar Chimney and found cold neutral clouds with low negative velocities, located at $0.3<|z|<2.5$ kpc. They suggested that these clouds most likely belong to a scattered population of infalling diffuse clouds, consistent with a Galactic Fountain model. \\cite{Howk05} discussed evidence for cold, CNM-like, gas in the disk/halo transition region of several spiral galaxies based on optical imaging of extra-planar dust. They suggested that this gas was most likely expelled from the thin disk through a quiescent process. \\cite{Richter03} found evidence for H$_{2}$ in dense, mostly neutral phase, and closely linked to CNM clouds in IVCs at even higher $z$ heights. These results support our distance estimate, however the formation mechanism(s) for cold clouds in the disk/halo transition regions remain under debate (see Section 9.1).    \\subsection{Is there a link between our low-velocity clouds and Lockman's clouds?} \\label{s:low-galactic}  {\\rotate \\begin{table*} \\footnotesize \\caption{Median properties of `different' types of HI clouds in the Galactic disk/halo interface region. } \\centering \\label{table1_02} \\begin{tabular}{ccccccccccc} \\noalign{\\smallskip} \\hline \\hline \\noalign{\\smallskip} Reference&Number of& $l$     & $b$        & V$_{\\rm dev}$ & Size & FWHM &$N({\\rm HI})_{\\rm peak}$ & $D$  & $|z|$ & $M_{\\rm HI}$ \\\\ & clouds      &(\\degree)& (\\degree)  & (\\kms)        & (pc)& (\\kms) &($\\times10^{19}$ cm$^{-2}$) &(kpc)& (pc)& (M$_{\\odot}$)\\\\ \\hline Lockman (2002)&38& 28--29  &$-4$ to $-12$& 15 to 40\\tablenotemark{a}     &24 & 12&2  &7.5  &950& $\\la50$  \\\\ Stil et al. (2005)&17&18--67 &$-1.3$ to $1.3$& 25 to 60     & $\\sim10$\\tablenotemark{b}&$6$ &20 & 8&80& 60 \\\\ This work &12&183--188 & 12 to 23   &$-25$ to $-40$     & 0.5--8&4& 2  &0.2--3&60--900& $0.03-7$     \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular}  \\tablenotetext{a}{One half of clouds in this study has V$_{\\rm dev}\\le15$ \\kms. Median V$_{\\rm dev}=13$ \\kms.} \\tablenotetext{b}{Most of the clouds in this study appear significantly elongated.} \\end{table*} }   One of the most important results of this study is that GALFA observations clearly demonstrate the cloudy structure of the Galactic disk/halo interface region in the outer Galaxy. The cloudy disk/halo interface, emphasized recently by \\cite{Lockman02} in the inner Galaxy, is obviously not restricted just to the inner Galaxy and could be quite extended radially. We compare here properties of the disk/halo interface clouds in the inner and outer Galaxy.   Recent HI observations with the Green Bank telescope showed that the halo in the inner Galaxy is not smooth but populated with discrete HI clouds \\citep{Lockman02}. These clouds, often referred to as ``Lockman's clouds'', are located $\\sim900$ pc below the plane. However, they appear to kinematically follow the HI distribution in the disk. We summarize typical properties of this cloud population in Table 3. Frequently, clouds are organized into larger structures and connected by diffuse envelopes and/or filaments which can have sizes up to $\\sim1$ kpc. In addition, high resolution observations \\citep{Lockman04} show that most Lockman's clouds have significant sub-structure, suggestive of the presence of a two-phase medium.  \\cite{Stil05} discovered an additional set of 17 isolated HI clouds in the inner Galaxy, with velocities up to 60 \\kms~greater than the maximum velocity allowed by Galactic rotation (see Table 3). Despite their forbidden velocities, these clouds follow the main disk emission and therefore cannot be classified as high-velocity gas. Located within the Galactic thin disk, at $|z|\\la80$ pc, these fast-moving clouds demonstrate that interstellar clouds can have very large random velocities along the line-of-sight. Stil et al. (2005) suggested that the fast-moving clouds, at $|z|\\la80$ pc, and Lockman's clouds, at $|z|\\sim950$ pc, may belong to the same population of clouds that is widely spread throughout the Galaxy. Both works focused exclusively on clouds at tangent points in the inner Galaxy and therefore have well-determined cloud distances.    Table 3 lists properties of LVCs in the outer Galaxy to facilitate their comparison with HI clouds in the inner Galaxy. This table gives: the number of clouds in each study, the range of Galactic longitude and latitude, cloud deviation velocities (as defined in Section 5.4), median linear sizes, FWHM,  peak HI column densities,  typical distances, heights above the plane, and estimated HI masses.  Several potentially important differences are noticeable.  (1) It is striking that our LVCs, which lie outside the Solar circle, are on average smaller, colder, and less massive than clouds in the inner Galaxy. For our selected sample, most clouds have FWHM of $3-4$ \\kms, a factor of three smaller than that of Lockman's clouds. Angular sizes of LVCs are in the range 6$'$--12$'$ (with a few exceptional examples having a size of only 4$'$), while Lockman's clouds have a median size of $\\sim11'$.  Linear sizes and HI masses of LVCs are of course uncertain as they depend on cloud distances. However distances of 6--8 kpc, which would bring HI masses close to those of Lockman's clouds, would imply the thermal pressure of 100 K cm$^{-3}$ which is below the minimum pressure at which the CNM can exist in thermal equilibrium \\citep{Wolfire03}. Cloud angular sizes in our study though could be partially a selection effect, as we have focussed on well-isolated compact clouds for which Lockman's angular resolution is less adequate.   (2) $|V_{\\rm dev}|$ of LVCs is on average higher than that of Lockman's clouds. The median deviation velocity for Lockman's sample is $13$ \\kms, while LVCs have $\\langle |V_{\\rm dev}| \\rangle=30$ \\kms. If real, this difference could be important, suggesting that interface clouds at larger galactocentric radii ($>8$ kpc) have larger deviation velocities than clouds at smaller radii ($\\la5$ kpc). Again, the difference in $|V_{\\rm dev}|$ is at least partially due to our selection criterion --- to avoid confusion with the disk emission we have selected clouds with $|V_{\\rm dev}|\\ga25$ \\kms, but clouds with lower $|V_{\\rm dev}|$ also exist in the data set. Larger cloud samples from the GALFA surveys will be able to clarify whether clouds with larger $V_{\\rm dev}$ are more common in the outer Galaxy.  (3) It is also important to notice that LVCs have negative $V_{\\rm dev}$, meaning they are moving toward us at velocities too fast to be explained by the Galactic rotation. Lockman's clouds, however, have systematically positive $V_{\\rm dev}$, meaning they are moving away from us with velocities too fast to be explained by the Galactic rotation.   In conclusion, if our clouds would have distances of 6--8 kpc then their properties would be similar to those of Lockman's clouds, but this would imply the thermal pressure below what is traditionally allowed for the CNM. If, however, our clouds are at distances $\\la2-3$ kpc, then they are colder, smaller, and less massive than Lockman's clouds. Taking into account various selection effects and different resolution provided by the Arecibo and Green Bank telescopes, clouds in the inner and outer Galaxy may belong to the same population with a range of properties ($z$, sizes, HI masses). However, distinctly different formation mechanisms can not be excluded either at this stage. Halo clouds with sizes similar to those of our LVCs were reported recently by \\cite{Kalberla05} and Dedes et al. (in preparation). Using the Effelsberg radio telescope these authors found about 20 HI clouds at $R_{\\rm g}=13-17$ kpc and $z=3-5$ kpc. However, when viewed at high resolution with the Very Large Array, these clouds resolve into smaller cores with a FWHM of 3--7 \\kms, a peak HI column density of a few $\\times10^{19}$ cm$^{-2}$, and diameters of a few pc.     \\label{s:discussion-theory}  \\subsection{Galactic Fountain and ballistic clouds}   One of the most common models for the interaction between the Galactic disk and the halo is based on Galactic fountain flows \\citep{Shapiro76,Houck90}. In this model, hot gas from the disk is pushed up by the stellar activity in the form of buoyant outflows, it travels through the halo, cools down, and as a result of thermal instabilities rains back onto the disk in the form of cold clouds. In their best model for low-temperature fountains (model 3; for a halo temperature of $3 \\times 10^{5}$ K and a volume density $10^{-3}$ cm$^{-3}$) \\cite{Houck90} suggested that about 70 cold halo clouds should be present above a disk area of 400 pc$^{2}$, having a typical size of $\\sim3$ pc, a temperature of 500 K, and a column density of $2.5 \\times 10^{18}$ cm$^{-2}$. A peak in cloud distribution is expected at the location of cloud formation, $z\\sim1$--2 kpc, however if drag is important then the density of clouds at $z<0.5$ kpc could be significant too. The expected radial velocities for fountain clouds are in the range from $-100 \\sin b$ to $+50 \\sin b$ \\kms. 3-D simulations of a Galactic fountain by \\cite{Avillez00} also show the formation of numerous small cloudlets, with sizes down to their resolution limit of 1.25 pc, resulting from a breakup of collimated structures and also larger clouds being swept up by blast waves.  Typical sizes of LVCs are surprisingly similar to what is expected for fountain clouds. In addition, at $b=17$\\degree, the expected velocity range for fountain clouds is $-30$ to 15 \\kms~(note that clouds at the positive velocities would be very difficult to detect due to confusion with Galactic HI), again very much in agreement with what we find for LVCs. The observed $N({\\rm HI})$ of LVCs is about ten times higher that what is expected for fountain clouds, although \\cite{Houck90} note that $N({\\rm HI})$ could be higher if clouds are concentrated at lower $z$ heights.  Another important evidence in favor of the fountain-style clouds is hinted by cloud deviation velocities. \\cite{Collins02} investigated a ballistic infall of fountain clouds and showed that a vertical velocity gradient (with scale height $z$) is expected for halo clouds; this is a consequence of a decrease in the gravitational acceleration vector with height from the disk. Hence, at any galactocentric radius, the LSR velocity of clouds in the Halo, at $z>1$ kpc, will be different from the LSR velocity of clouds in the disk. However, the steepness of this vertical velocity gradient is a function of the galactocentric radius at which clouds originated before being kicked upward from the disk. For example, for the input circular velocity of 230 \\kms~and at a galactocentric radius of $R_{\\rm g}=8$ kpc, the expected $\\Delta V_{\\rm lsr}$ is $\\la30$ \\kms~when $z$ changes from $z=0$ to $z\\sim2$ kpc (assuming an initial cloud kick velocity of 100 \\kms). Note that viscous interactions between Halo clouds, and/or more complex interactions between the Halo medium and clouds can decrease $\\Delta V_{\\rm lsr}$. However, for very small galactocentric radii, $\\Delta V_{\\rm lsr} \\approx 0$ \\kms, for any $z$. The expected LSR velocity difference ($\\Delta V_{\\rm lsr}$) for clouds at higher $z$ relative to clouds in the disk, is therefore something we can directly compare with the cloud deviation velocity, $V_{\\rm dev}$.  The anti-center LVCs have  $\\langle |V_{\\rm dev}| \\rangle=30$ \\kms, in agreement with what is expected for fountain clouds at heights $z=$1--2 kpc. The expected circulation timescale (time for clouds to return back to the disk) for the fountain clouds is about 50 Myr \\citep{Houck90,Collins02}. As the cloud evaporation timescale ($\\ga100$ Myr) is significantly longer than the expected circulation timescale, clouds can reach the disk safely before being evaporated. In Section~\\ref{s:low-galactic} we hinted at a potentially interesting difference in $|V_{\\rm dev}|$ for clouds in the inner and outer Galaxy --- most halo clouds in the inner Galaxy appear to have smaller $|V_{\\rm dev}|$ than clouds in the outer Galaxy. This potential trend in $|V_{\\rm dev}|$ is an important tool for investigating cloud origin --- if halo clouds are a result of a Galactic fountain we would expect $|V_{\\rm dev}|$ to increase with galactocentric radius. Future work with larger cloud samples will investigate statistical significance of this trend.  In conclusion, LVCs have typical linear sizes and radial velocities in good agreement with expectations from the Galactic fountain model. This would suggest that Galactic fountain flows, and the circulation of fountain clouds, are (at least) as common in the outer Galaxy as in the inner Galaxy. This is a surprising result when one takes into consideration that the Galactic star formation rate (SFR) strongly depends on galactocentric radius, based on various observations as well as predictions from chemical and spectrophotometric models for the evolution of spiral galaxies (for references see \\cite{Boissier99} and \\cite{Boissier03}). The star formation rate peaks around $R_{\\rm g}\\sim4$ kpc and then steadily decreases towards larger radii, at  $R_{\\rm g}\\sim12$ kpc for example, SFR$\\sim0.3\\times{\\rm SFR}_{\\odot}$. However, the decrease in SFR at large radii is alleviated partially by the decrease of the gravitational potential which enables vertical fountain flows to propagate easier.   \\subsection{Connection between filaments and clouds}  Most of LVCs appear related to large, diffuse HI filamentary features. Diffuse filaments and clouds both appear only over a restricted range of LSR velocities, $-20$ to $-30$ \\kms. This is suggestive of their possible physical association.  An association between larger filaments and small, cold clouds is expected as a result of shell or chimney fragmentation. \\cite{Norman89} expect that clouds formed from chimney walls will have velocities $|v| \\sim$10--50 \\kms. \\cite{McClure-Griffiths06} found several cold HI clouds associated with caps of a large expanding shell and suggested that shell fragmentation is responsible for the cloud formation. Similarly, \\cite{Koo06} proposed that small, fast-moving clouds could be fragments of old supernova remnant shells.   Visual appearance of filaments and clouds (Figure 4) strikingly resembles recent numerical simulations by \\cite{Audit05} of dynamically triggered condensation of the WNM into small CNM clouds. In these simulations, a collision between incoming turbulent WNM streams creates a thermally unstable region of higher density and pressure but lower temperature, which starts to fragment into cold structures. This thermally unstable WNM has a filamentary morphology, and is seeded with numerous small (size of $\\sim0.1$ pc) and cold clouds ($n\\sim50$ cm$^{-3}$, $T\\sim80$ K). The whole process of fragmentation is promoted and controlled by turbulence, and cloud morphology and properties, as well as the fraction of the condensed cold gas, are governed by turbulent properties. Based on this apparent morphological similarity it may be reasonable to expect that external (dynamical) forcing (through, for example, collisions, or interactions with the dense halo gas) of some types of vertical WNM flows (e.g. chimneys or walls of expanding shells) can trigger condensation of warm HI into cold clouds.    \\subsection{Infalling multi-phase gas?}   The idea that galaxies acquire their fresh star formation fuel through the infall of clouds has attracted a lot of attention, especially recently with advances in numerical simulations of galaxy formation \\citep{Maller04}. As the hot gas within the cooling radius ($\\sim150$ kpc for the Milky Way size galaxy) is a subject to thermal instabilities, it fragments into smaller, warm ($T\\sim10^{4}$ K),  pressure-supported clouds.  The warm clouds fall into the Galaxy, after undergoing cloud-cloud collisions and ram pressure stripping, and are tidally broken into smaller clouds before impacting the Galaxy. The cloud formation and infall are expected to be balanced at the present time, and models predict abundant evidence for this ongoing process in the Galactic halo, in the form of numerous condensed clouds.    Kaufmann et al. (2006) investigate the formation of a galactic disk, through N-body and SPH simulations, and show that cloud accretion is not spherical but happens along the angular momentum axis. This results in outer disks preferentially being populated with small, cold accretion remnants, while the central regions below and above the plane channeling hot gas. In this scenario, cloudy structure of the disk/halo interface region is expected in the inner and outer galactic disks but due to a very different phenomenon --- while outer parts are driven by the accretion process, the inner parts could be a result of a fountain process.   However, what exactly happens to clouds during the final stages of their infall is still mainly unexplored, both theoretically and observationally. How do large and warm clouds transform into `cloud forms' that can easily be integrated and digested  by the disk, to provide a smooth build up of fresh star formation fuel? \\cite{Maller04} suggest that clouds get shredded into smaller but still warm ($\\sim10^{4}$ K) condensations, with further cooling being stopped by the extragalactic ionizing background. Clouds in outer disks in simulations by Kaufmann et al. (2006) are also warm, 100--600 pc in size, however apparently they can cool and condense further to reach thermal equilibrium at a temperature $<10^{4}$ K.    CHVCs and our LVCs may be able to fit into this framework as representing later stages of the cloud infall process. CHVC186+19-114 shown in Figures 8--11, typical of several hundreds of compact HVCs, could represent a cloud that has condensed in the halo and has begun to infall, but has not yet reached the disk. The low-velocity filamentary/cloudy structure, shown in Figures 2--7, could represent the remaining structure of the accreted cloud as it is being almost integrated with the Galactic disk. This scenario seems plausible when one considers the similarities between these two types of objects (see Table 1): CHVC186+19-114's `core' structure has velocity linewidths and HI column densities comparable to those found in LVCs, while typical HI linewidths and column densities of the filamentary structures at about $-25$ \\kms~are very similar to those of CHVC186+19-114's envelope. If the LVCs trace the final stages of the infall process then they suggest cloud sizes and temperatures significantly smaller from what is currently found in numerical simulations. However, details of heating/cooling and self-shielding effects during the cloud infall still await to be addressed. Observationally, constraining metallicities of halo clouds would be the crucial test of the cloud infall hypothesis.       The consortium for Galactic studies with ALFA is conducting an HI survey of the whole Arecibo sky (declination range from $-1$\\degree~to 38\\degree) with high angular (3.5$'$) and velocity resolution (0.2 \\kms). At the time this article is being written about 20\\% of the survey has already been completed. In this paper we focused only on one of the main science drivers for GALFA --- the Galactic disk/halo interface region. With only about 200 square degrees imaged in the Galactic anti-center region, primarily for the purpose of hardware and software testing, several interesting and new phenomena have emerged.  We have found numerous small (most likely just a few pc in size) and cold ($T_{\\rm k}<400$ K) HI clouds at low negative velocities distinctly separated from the HI disk emission (`low-velocity clouds' or LVCs). Over a restricted range of velocities and exactly the same range as for LVCs, we have also found larger sheet-like filamentary structures. Most likely, the LVCs and filaments are physically related. While the distances to LVCs will have to be better constrained in the future, preliminary estimates, based on the thermal pressure requirements  for the CNM and the WNM coexistence, range from 0.2 to 3 kpc. This implies that the clouds'  height from the plane is in the range  60 to 900 pc, suggesting that most LVCs are located in the transition region between the Galactic disk and halo, yet they have properties of typical CNM clouds. LVCs are colder and, most likely, smaller and less massive than Lockman's clouds in the disk/halo interface region in the inner Galaxy. In addition, LVCs may have larger absolute deviation velocities than Lockman's clouds. Nevertheless, the existence of a large number of LVCs in the outer Galaxy (at a galactocentric radius $>8$ kpc) demonstrates that the cloudy and frothy character of the interface region is more likely to be a general phenomenon, not restricted exclusively to the inner Galaxy.  In the same dataset but at more extreme negative velocities we have studied CHVC186+19-114, a typical CHVC with a well-defined core/envelope structure. As we have imaged a large area around the CHVC, we discovered a companion HI cloud located 50$'$ southwest of CHVC186+19-114, most likely physically related to CHVC186+19-114. Velocity profiles suggest that the companion cloud was stripped off the main cloud's envelope due to the interactions with the halo medium. This is the first time that a companion is detected around an isolated CHVC and may represent evidence for the breakup of CHVCs into smaller condensations. Based on its size, the companion cloud could be classified as a member of the population of UCHVCs, discovered only very recently.   We have discussed briefly three possible scenarios for the formation and maintenance of the clumpy Galactic disk/halo interface region. The first scenario is provided by Galactic fountains, buoyant hot flows from the disk which cool and rain back in the form of cold clouds. Observational properties (cloud size and radial velocity) of LVCs are in a good agreement with predictions of low-temperature fountains models. The second scenario is motivated by numerical simulations by \\cite{Audit05} of dynamically triggered condensation of the WNM into small CNM clouds. Vertical WNM flows propagating upward from the Galactic plane, in the form of expanding shells or chimneys, could be triggered into formation of CNM clouds at larger $z$ heights. This comparison is primarily based on morphological resemblance at this stage as detailed models do not exist. Finally, the third scenario investigates whether the low-velocity clouds/filaments could represent the final stages of the infalling IGM. This has been encouraged by the recent models and simulations \\citep{Maller04,Kaufmann06} which emphasize the multi-phase and multi-scale character of the galaxy formation process. While simulations show encouraging qualitative resemblance at this stage the detailed physical processes during the final stages of cloud infall are largely unexplored. We hope that our current, and especially future high-resolution datasets, will stimulate and guide further theoretical and numerical work in this area.  Future GALFA datasets will clearly find more disk/halo interface clouds and new HVCs, especially in the regime of the smallest HVCs which is largely unexplored. With the larger data sets we will be able to confirm whether the cloudy structure exists at other Galactic longitudes, and also the statistical significance of the potential differences between interface clouds in the inner and outer Galaxy. Also, future GALFA datasets will be important for establishing spatial distribution and properties of the smallest HVCs."
    },
    "0609/astro-ph0609301_arXiv.txt": {
        "abstract": "Early-type galaxy velocity dispersions and luminosities are correlated.  The correlation estimated in local samples ($\\le 100$~Mpc) differs from that measured more recently in the SDSS. This is true even when systematics in the SDSS photometric and spectroscopic parameters have been accounted-for. We show that this is also true for the ENEAR sample if galaxy luminosities are estimated using distances which have been corrected for peculiar motions.  We then show that, because the estimate of the `true' distance is derived from a correlation with velocity dispersion, in this case the $D_n-\\sigma$ relation, using it in the $\\sigma-L$ relation leads to an artificially tight relation with a biased slope. Making no correction for peculiar velocities results in a $\\sigma-L$ relation which is very similar to that of the SDSS, although with larger scatter.  We also measure the $\\sigma-L$ correlation in a mock ENEAR catalog, in which the underlying galaxy sample has the same $\\sigma-L$ correlation as seen in the SDSS.  The mock catalog produces the same $D_n-\\sigma$ relation as the data, the same biased slope when $D_n-\\sigma$ distances are used to estimate luminosities, and good agreement with the input $\\sigma-L$ relation when redshift is used as the distance indicator. This provides further evidence that the true $\\sigma-L$ relation of ENEAR galaxies is indeed very similar to that of SDSS early-types. Our results suggest that local $\\sigma-L$ relations which are based on Fundamental Plane distances should also be re-evaluated. Our findings also have important implications for black hole demographics; the best direct estimates of the masses of supermassive black holes come from local galaxies, so estimates of the black hole mass function are more safely made by working with the $M_\\bullet-\\sigma$ correlation than with $M_\\bullet-L$. ",
        "introduction": "The luminosities and velocity dispersions of early-type galaxies are strongly correlated:  their logarithms follow an approximately linear relation (e.g. Faber \\& Jackson 1976). There is a long and complicated history of what the slope of this relation is, mainly due to the fact that, if there is intrinsic scatter around this correlation, then there are at least three interesting slopes:  that obtained from fitting the mean luminosity at each velocity dispersion, $\\langle\\log L|\\log\\sigma\\rangle = a_{L|\\sigma}\\log\\sigma + b_{L|\\sigma}$, the mean velocity dispersion as a function of luminosity, $\\langle\\log\\sigma|\\log L\\rangle = a_{\\sigma|L}\\log L + b_{\\sigma|L}$, and the slope of the principal axis of the joint distribution of $L$ and $\\sigma$ (e.g. Lynden-Bell et al. 1988; Saglia et al. 2001). The first of these allows one to use $\\sigma$ to predict $L$, whereas the second must be used if one wishes to predict $\\sigma$ from $L$.  Until recently, authors were not careful to distinguish between these cases.  Because $L$ depends on the distance to the source, whereas $\\sigma$ does not, it is straightforward to estimate $\\langle\\log\\sigma|\\log L\\rangle$ from flux limited samples; naive estimates of the other two correlations are compromised by selection effects. Most studies of this correlation, based on local samples, agree that the typical velocity dispersion at fixed luminosity, $\\langle\\log\\sigma|\\log L\\rangle$, scales as $L^{1/4}$: $a_{\\sigma|L}\\approx 1/4$. The common unfortunate abuse of jargon is to say that luminosity scales as the fourth power of velocity dispersion; in fact, $a_{L|\\sigma}$ is considerably shallower. In what follows, we will be almost exclusively concerned with how well $L$ can be used as a predictor for $\\sigma$: $\\langle\\log\\sigma|\\log L\\rangle$.  Prior to the year 2000, studies of the $\\sigma-L$ correlation were restricted to local samples (from within about 50 Mpc) containing $\\sim 100$ objects.  The SDSS provided a sample of early-type galaxies that was larger by about two orders of magnitude (Bernardi et al. 2003a):  $\\sim 9000$ objects drawn from a volume which extended to considerably larger distances (median redshift $z\\sim 0.1$).  However, the $\\sigma-L$ relation in this sample, \\begin{equation} \\Bigl\\langle\\log_{10}\\sigma|M_r\\Bigr\\rangle_{\\rm SDSS-B03} = 2.203 - 0.102\\,(M_r+21), \\label{sdss03} \\end{equation} with an intrinsic scatter of about 0.07~dex is inconsistent with that found in local samples. For example, at $\\log_{10}(\\sigma/{\\rm km~s}^{-1}) \\ge 2.4$, the relation which Forbes \\& Ponman (1999) found best fits the Pruniel \\& Simien (1996) sample of 236 local early-type galaxies differs from the SDSS fit by more than 0.05~dex. Expressed in terms of absolute magnitudes, the fits differ by more than 0.5~mags at $M_r\\le -23$.  This is substantially larger than expected given the measurement errors.  What causes this difference?  Because the SDSS luminosities and velocity dispersions are obtained from an automated pipeline, (i.e., light profiles and spectra were not inspected individually), it may be that, say, the SDSS measurements differ systematically from those estimated for local objects in the literature.  Indeed, it is known that the SDSS photometric reductions underestimate the luminosities of objects in crowded fields (Mandelbaum et al. 2005; Bernardi et al. 2006a; Hyde et al. 2006; Lauer et al. 2006). See the Appendix for a more detailed discussion. The Appendix also shows that the SDSS also slightly overestimates the velocity dispersions at small $\\sigma$. In what follows, we will refer to the sample in which these systematics have been accounted-for as the SDSS-B06 sample (see the Appendix). In the SDSS-B06 sample, \\begin{equation} \\Bigl\\langle \\log \\sigma|M_r\\Bigr\\rangle_{\\rm SDSS-B06} = 2.190 - 0.100\\,(M_{r}+21) \\label{sdssVL} \\end{equation} with intrinsic scatter of 0.07~dex; it happens that this is not very different from the relation obtained by Bernardi et al. (2003b).  A common problem from which all local samples suffer is that, while the apparent magnitude of an object can be measured quite accurately,  the absolute magnitude is more difficult because it depends on the distance to the galaxy.  The true distance is difficult to measure because the redshift, which is usually well-determined, is a combination of the distance to the galaxy and the component of its peculiar velocity which is directed along the line of sight to the observer: $cz = Hd + v_{\\rm los}$. Typical velocities are expected to be of order a few hundred km~s$^{-1}$, so the redshift is a reliable distance indicator only beyond about $100h^{-1}$Mpc.  Most of the objects in the SDSS lie well beyond this distance, whereas all local samples are shallower. Thus, a legitimate concern is whether uncertainties in estimating the true distance are driving the difference between the SDSS measurement and those which are based on more local samples.  To address such concerns, we have studied the $\\sigma-L$ relation in the definitive sample of nearby early-type galaxies---that assembled in the ENEAR database (da Costa et al. 2000; Bernardi et al. 2002a; Alonso et al. 2003; Wegner et al. 2003). ENEAR contains about 1000 objects out to 7000~km~s$^{-1}$, for which measured redshifts and estimated distances are available. Section~2 shows how the $\\langle\\log\\sigma|\\log L\\rangle$ relation changes if one uses the redshift rather than the estimated distance when computing $L$. It highlights the fact that, if $\\sigma$ played a role in determining the distance, e.g. if the distance comes from a Fundamental Plane or $D_n-\\sigma$ analysis, then the $\\sigma-L$ relation may be biased. In Section~3 we discuss our results. ",
        "conclusions": "We showed that the ENEAR $\\sigma-L_d$ relation, where $L_d$ indicates that the luminosity is based on a distance estimate which has been corrected for peculiar motions, is almost certainly biased.  This is because the distance indicator used to estimate $L_d$ depends on $\\sigma$. We also showed that if the redshift is used as the distance indicator (i.e., no correction for peculiar motions is made) then the resulting $\\sigma-L_z$ relation, while noisier, is less likely to be biased. The ENEAR $\\sigma-L_z$ relation is in excellent agreement with that measured in the SDSS, where neglecting peculiar velocities is an excellent approximation. Thus, it appears that the discrepancy between the SDSS $\\sigma-L$ relation and that in local samples is due to the use of a distance indicator which correlates with $\\sigma$.   Our results are of particular interest for the problem of estimating black hole abundances from the observed distribution of luminosities or velocity dispersions (e.g. Yu \\& Tremaine 2002; Shankar et al. 2004; Tundo et al. 2006).  The first method requires $\\langle M_\\bullet|L\\rangle$ whereas the second requires $\\langle M_\\bullet|\\sigma\\rangle$.  The two approaches will only give the same estimate of $\\phi(M_\\bullet)$ if the $\\langle \\sigma|L\\rangle$ relation of the black hole sample is the same as that of the sample from which the luminosity function $\\phi(L)$ and velocity dispersion function $\\phi(\\sigma)$ are drawn. Current black hole samples are relatively local, so their $\\sigma-L$ correlations are very different from that of the SDSS. Bernardi et al. (2006b) suggests that current black hole samples are biased towards objects with abnormally large velocity dispersions for their luminosities. If this is a selection rather  than physical effect, then the $M_\\bullet-\\sigma$ and $M_\\bullet-L$ relations currently in the literature are also biased from their true values. Bernardi et al. find that the bias in the $\\langle M_\\bullet|\\sigma\\rangle$ relation is likely to be small, whereas the $\\langle M_\\bullet|L\\rangle$ relation is biased towards predicting more massive black holes for a given luminosity. Therefore, the estimate based on the velocity function (e.g. that of Sheth et al. 2003) is to be preferred.  The $\\langle\\sigma|L\\rangle$ relations show evidence for a flattening at large $L$, a fact we do not use here, but which is relevant to studies of BCGs (e.g. Bernardi et al. 2006a) and may be relevant to studies of black hole demographics."
    },
    "0609/astro-ph0609670_arXiv.txt": {
        "abstract": "{ The mechanisms governing the star formation rate in spiral galaxies are not yet clear. The nearby, almost face-on, and interacting galaxy M51 offers an excellent opportunity to study at high spatial resolutions the local star formation laws.} {In this first paper, we investigate the correlation of H$_2$, \\HI, and total gas surface densities with the star forming activity, derived from the radio continuum (RC), along radial averages out to radii of 12\\,kpc.} {We have created a complete map of M51 in $^{12}$CO 2--1 at a resolution of 450\\,kpc using HERA at the IRAM-30m telescope. These data are combined with maps of \\HI\\ and the radio-continuum at 20\\,cm wavelength. The latter is used to estimate the star formation rate (SFR), thus allowing to study the star formation efficiency and the local Schmidt law $\\Sigma_{\\rm SFR}\\propto\\Sigma_{\\rm gas}^n$. The velocity dispersion from CO is used to study the critical surface density and the gravitational stability of the disk.} { { The total mass of molecular material derived from the integrated $^{12}$CO 2--1 intensities is $2\\,10^9$\\,\\msun. The $3\\sigma$ detection limit corresponds to a mass of $1.7\\,10^5$\\,\\msun. The global star formation rate is 2.56\\,\\msun\\,yr$^{-1}$ and the global gas depletion time is $0.8$\\,Gyr. \\HI\\ and RC emission are found to peak on the concave, downstream side of the outer south-western CO arm, outside the corotation radius. The total gas surface density $\\Sigma_{\\rm gas}$ drops by a factor of $\\sim20$ from 70\\,\\msun\\,pc$^{-2}$ at the center to 3\\,\\msun\\,pc$^{-2}$ in the outskirts at radii of 12\\,kpc.  The fraction of atomic gas gradually increases with radius. The ratio of \\HI\\ over H$_2$ surface densities, $\\Sigma_{\\rm HI}/\\Sigma_{\\rm H2}$, increases from $\\sim0.1$ near the center to $\\sim20$ in the outskirts without following a simple power-law. $\\Sigma_{\\rm HI}$ starts to exceed $\\Sigma_{\\rm H2}$ at a radius of $\\sim4$\\,kpc. The star formation rate per unit area drops from $\\sim400$\\,\\msun\\,pc$^{-2}$\\,Gyr$^{-1}$ in the starburst center to $\\sim2$\\,\\msun\\,pc$^{-2}$\\,Gyr$^{-1}$ in the outskirts.  The gas depletion time varies between 0.1\\,Gyr in the center and 1\\,Gyr in the outskirts, and is shorter than in other non-interacting normal galaxies. Neither the \\HI\\ surface densities nor the H$_2$ surface densities show a simple power-law dependence on the star formation rate per unit area. In contrast, $\\Sigma_{\\rm gas}$ and $\\Sigma_{\\rm SFR}$ are well characterized by a local Schmidt law with a power-law index of $n=1.4\\pm0.6$.  The index equals the global Schmidt law derived from disk-averaged values of $\\Sigma_{\\rm gas}$ and $\\Sigma_{\\rm SFR}$ of large samples of normal and starburst galaxies. The critical gas velocity dispersions needed to stabilize the gas against gravitational collapse in the differentially rotating disk of M51 using the Toomre criterion, vary with radius between 1.7 and 6.8\\,kms$^{-1}$. Observed radially averaged dispersions derived from the CO data vary between 28\\,kms$^{-1}$ in the center and $\\sim$8\\,kms$^{-1}$ at radii of 7 to 9\\,kpc. They exceed the critical dispersions by factors $Q_{\\rm gas}$ of 1 to 5. We speculate that the gravitational potential of stars leads to a critically stable disk. }} {} ",
        "introduction": "%   \\begin{figure}[h] \\centering \\includegraphics[angle=-90,width=8cm]{5579fig1.eps} \\caption{Map of $^{12}$CO 2--1 integrated intensities [Kkms$^{-1}$] showing M51, i.e. NGC\\,5194 and its companion galaxy NGC\\,5195 in the north-east.  The image has a resolution of $11''$ and is constructed from a masked moment calculation \\citep{adler1992} to minimize the noise contribution from emision-free channels when integrating over the full kinematic extent of M51, 350\\,kms$^{-1}$ $<$ v$_{\\rm lsr}$ $< 600$\\,kms$^{-1}$ .  Color coded intensities are on a logarithmic scale. The cross marks the 0/0 center position at $\\alpha=13\\fh29\\fm52\\fs7$, $\\delta=47\\degr11\\arcmin43\\arcsec$ (eq.2000). The noise level is $1\\sigma=0.65$\\,Kkms$^{-1}$. Contours range in units of 1$\\sigma$ from 1, 3, 6, 16 to 56 in steps of 10. The peak intensity is 47.3\\,Kkms$^{-1}$. All intensities are on the $T_{\\rm A}^*$ scale. } \\label{fig_mapco21} \\end{figure}     \\begin{figure}[h] \\centering \\includegraphics[angle=-90,width=8cm]{5579fig2.eps} \\caption{VLA map of integrated \\HI\\ intensities [Jy/beam] at $13''$ resolution \\citep{rots1990} in colors.  Contours show integrated $^{12}$CO 2--1 intensities (cf.\\,Fig.\\,\\ref{fig_mapco21}). \\label{fig_maphi} } \\end{figure}  \\begin{figure}[h] \\centering \\includegraphics[angle=-90,width=8cm]{5579fig3.eps} \\caption{Map of radio continuum intensities at 20\\,cm (VLA C$+$D, $15''$ resolution) in units of Jy/beam \\citep{patrikeev2006}. Contours show integrated $^{12}$CO 2--1 intensities (cf.\\,Fig.\\,\\ref{fig_mapco21}). \\label{fig_map20cm} } \\end{figure}   M51 is an interacting, grand-design spiral galaxy at a distance of only 8.4\\,Mpc seen nearly face-on (Table\\,\\ref{tab_m51_properties}). It is rich in molecular gas, most of which is found in the two very prominent spiral arms which are presumably caused by the tidal interaction { \\citep{tully1974c,howard1990}. } Kinematic studies reveal unusually large streaming motions implying a strong density wave and the presence of galactic shocks \\citep{aalto1999}. The \\CII\\ map of \\citet{nikola2001} shows two secondary lobes of emission to the northeast and southwest near the corotation radius of the density wave pattern, presumably due to cloud-cloud collisions, stimulating star formation. \\citet{calzetti2005} combine { Spitzer infrared data of the Infrared Array Camera (IRAC) and the Multi Imaging Photometer for Spitzer (MIPS) with UV data from the Galaxy Evolution Explorer (GALEX). These data are used} to discuss the various tracers of star formation in M51. Enhanced star formation activity at the two lobes seen in \\CII\\ is probably triggered by the interaction \\citep{nikola2001}.  This is also indicated by disk simulations of \\citet{toomre_toomre1972} which show tidal tails emerging from the disk at the positions of enhanced star formation.  VLA observations of the \\HI\\ line show a prominent tail outside the main disk extending to the south and east, and covering more than $25'$, i.e. 100\\,kpc, in projected distance \\citep{rots1990} \\citep[cf.][]{tilanus_allen1991}.  \\citet{meijerink2005} find that a major fraction of cold dust emission at 850$\\,\\mu$m stems from an extended exponential disk with a scale height of 5.45\\,kpc, possibly also tracing total gas column densities. The kinematics of H$\\alpha$ emission of M51 were recently studied by the SINGS team \\citep{daigle2006}, improving on the work of \\citet{tilanus1991}.  Its proximity and inclination make M51 an ideal target to study the efficiency at which it forms stars from the gas under the influence of the interaction, and to compare with the star formation laws commonly found in spiral galaxies. Spiral galaxies often follow two empirical laws: { as star formation is fueled by the interstellar gas reservoir, the star formation rate (SFR) is proportional to a power of the total gas surface density: $\\Sigma_{\\rm SFR}\\propto\\Sigma_{\\rm gas}^n$ .  This appears to hold globally when $\\Sigma_{\\rm gas}$ and $\\Sigma_{\\rm SFR}$ are averaged over the entire star-forming region of a galaxy \\citep[][ and others]{schmidt1959,kennicutt1998} with $n=1.4$.  This may be understood assuming that $\\Sigma_{\\rm SFR}$ is proportional to the gas surface density over the free-fall time $\\tau_{\\rm ff}$ and that the free-fall time and the gas surface density are simple functions of the average gas density $\\rho$: $\\tau_{\\rm ff}\\propto{\\rho}^{-1/2}$ and $\\Sigma_{\\rm gas}\\propto\\rho$. This simple model leads directly to the Schmidt law with an index of 1.5 \\citep[e.g.][]{elmegreen1994}.  More recent studies show that the global Schmidt law also appears to hold locally in individual spiral galaxies when studying radially averaged profiles of $\\Sigma_{\\rm gas}$ and $\\Sigma_{\\rm SFR}$ \\citep{wong_blitz2002,boissier2003}. The second law is that the total gas surface density is of the order of a critical gas surface density given by the \\citet{toomre1964} criterion for gravitational stability \\citep{martin_kennicutt2001,li2006}. }  The relative importance of the different components of the interstellar medium was addressed by \\citet{nikola2001} when trying to explain the mapped \\CII\\ emission. In a recent study comprising all major cooling lines of the molecular gas, i.e. those of \\CII, \\OI, CO, and \\CI, at selected spiral arm positions and the center of M51, \\citet{kramer2005} conclude that the bulk of the emission stems from clumpy photon dominated regions. Only about 15--30\\% of the \\CII\\ emission stems from dense \\HII\\ regions.  The \\CI\\ and CO emission of the center region have also been studied by \\citet{israel2006,israel_baas2002,gerin_phillips2000}. \\citet{schinnerer2004} mapped two distinct regions in the spiral arms of M51 in HCN, HCO$^+$ and other tracers of the chemistry.  { \\citet{rydbeck2004} presented a large CO 1--0 map of M51 obtained at the Onsala-20m telescope.  } The inner region of M51 were previously mapped in CO 1--0 with the FCRAO-14m, IRAM-30m, and NRO-45m telescopes by \\citet{scoville_young1983,lord_young1990,gb1993no1,nakai1994,kuno1995,kuno1997}. These single-dish observations reveal the large-scale emission for which interferometric observations are not sensitive. OVRO was used to create CO 1--0 maps by \\citet{rand_kulkarni1990} and \\citet{aalto1999}.  \\citet{helfer2003} used BIMA in combination with NRAO-12m single-dish data to map the inner part of M51 in CO 1--0.   Here, we present a complete IRAM-30m CO 2--1 map of M51, improving on the previous data of \\citet{gb1993no1,gb1993no2}.  \\begin{center} \\begin{table}[h*] \\caption[]{\\label{tab_m51_properties} {\\small Basic properties of M51. References: $^a$ RC3 catalogue \\citet{devaucouleurs1991}, % $^b$ \\citet{feldmeier1997}, note that \\citet{takats2006} recently reported a distance of only $7.1\\pm1.2$\\,Mpc, $^c$ \\citet{tully1974b}. }} \\begin{tabular}{lrrrrr} \\hline \\hline & M51 \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} RA(2000)                     & 13:29:52.7 \\\\ DEC(2000)                    & 47:11:43   \\\\ Type                         & SA(s)bc pec$^{a}$ \\\\      % Distance [Mpc]               & $8.4^{b}$ \\\\             % $11''$ correspond to         & 448\\,pc \\\\                  % Heliocentric velocity [kms$^{-1}$]  & 463$^{a}$ \\\\       % Position Angle [deg]         & $170^{c}$ \\\\                % Inclination [deg]            & $20^{c}$ \\\\               % \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{table} \\end{center} ",
        "conclusions": "% \\label{sec-discussion}  {  \\subsection{Toomre stability}  Several authors have argued that galactic disks are self-regulated through gravitational instabilities to have their Toomre $Q$ parameter of the order of 1 \\citep[see e.g.][]{combes2001}. As soon as gas dissipation leads to $\\Sigma_{\\rm gas}>\\Sigma_{\\rm crit}$ at a given radius, the disk becomes gravitationally unstable, leading to density waves transferring angular momentum, causing dissipation and heating of the gas and thus an increase of the velocity dispersion, stabilizing the disk again.  Models indeed indicate that the disk settles at the border of instability \\citep{lin_pringle1987}.  However, in general the velocity dispersion and surface density of stars may play an important role in the stability analysis. \\citet{bottema1993} discuss the dispersion of stellar disks and the self-regulating mechanism that keeps $Q$ near 1 over the entire disk of spiral galaxies. To take this into account, a total $Q$ parameter has been used by e.g. \\citet{combes2001}: $Q_{\\rm tot}^{-1} = Q_{\\rm gas}^{-1} + Q_{\\rm star}^{-1}$. Neglecting stars is valid only when the gas surface density divided by the gas dispersion dominates all other $\\Sigma/\\sigma$ terms.  Observationally, $Q_{\\rm gas}$ has often been found to be larger than 1 by factors of a few.  \\citet{kennicutt1989} derived radial averages of the Toomre-parameter of a sample of 15 galaxies, assuming a constant velocity dispersion of 6\\,kms$^{-1}$, and found that star formation is restricted to regions where $Q_{\\rm gas}<1.5$.  A similar conclusion was drawn from a study of 32 nearby spiral galaxies by \\citet{martin_kennicutt2001} with few exceptions.\\footnote{Both, \\citet{kennicutt1989} and \\citet{martin_kennicutt2001} included M51 in their sample, using FCRAO $50''$ CO 1--0 data combined with H$\\alpha$ data. However, the Toomre parameter and its variation with radius in M51 are not explicitly discussed. }  Similarly, \\citet{wong_blitz2002} also assume a constant velocity dispersion and find $\\Sigma_{\\rm gas}\\approx\\Sigma_{\\rm crit}$ over a wider range of radii for some of their galaxies while others show deviations of factors of $\\sim2$.  While in the inner 4\\,kpc of M51, the gas dispersion is large and $Q_{\\rm gas}$ lies between 2 and 5, the stars may contribute here, and make the disk marginally stable. Streaming motions are a consequence of spiral waves and gravitational instabilities \\citep{combes2001}. M51 shows strong streaming motions of $60-150$\\,kms$^{-1}$ in the plane of the galaxy in the inner spiral arms as can be seen in the interferometric CO maps of \\citet{aalto1999}. The broadening of the velocity dispersions due to streaming motions is naturally taken into account here.  However, the observed dispersion is also broadened by the very steep, unresolved rotation curve in the center of M51. A more careful stability analysis would need to correct the observed dispersions for rotation gradients, take into account the influence of stars, and study individual regions of M51 at high spatial resolution.   \\subsection{The local Schmidt law and star formation efficiency}  Recent smoothed particle hydrodynamics (SPH) simulations of gravitational instability of isolated disk galaxies comprising a dark matter halo, a disk of stars, and isothermal gas by \\citet{li2005,li2006} find a slope of 1.31 of the local Schmidt law (Table\\,\\ref{tab_schmidt}), close to the slope of the global Schmidt law, which is also reproduced, and close to the slope of 1.35 found in M51. On the other hand, the slope of 1.3 is at the low-end of the local Schmidt law slopes observed so far.  We find a normalization factor $A$ of the local Schmidt law in M51 of 1.12\\msun pc$^{-2}$ Gyr$^{-1}$ (Table\\,\\ref{tab_schmidt}), similar to $A$ of gas-rich models in \\citet{li2006}. This is a factor 5 larger than the average value of the \\citet{li2006} models. It is also a factor 5 larger than the normalization found in the large sample of \\citet{kennicutt1998}.  Interestingly, the normal galaxies of the \\citet{kennicutt1998} sample show on average a factor $\\sim3$ longer gas consumption times than found in M51. In Section\\,\\ref{sec-sfr-schmidt} we also noticed that the Milky Way \\citep{misiriotis2006} and the 6 CO-bright galaxies studied by \\citet{wong_blitz2002} exhibit consumption times which are larger than the consumption time found in M51 by about one order of magnitude. Apparently, the star formation efficiency of M51 is higher than in many other normal galaxies.  The reason may be the interaction of M51 with its neighbouring galaxy NGC\\,5195, as has been proposed before as mentioned in the Introduction \\citep[e.g.][]{nikola2001,howard1990,toomre_toomre1972}. This is supported by the recent simulations of \\citet{li2004} which show that interaction and merging can lead to instable disks and strong starbursts.    }"
    },
    "0609/astro-ph0609446_arXiv.txt": {
        "abstract": "% We are examining mass loss from globular clusters giant stars, focussing on metallicity dependance. We present three sets of observations: TIMMI-2 mid-IR spectra of 47 Tuc, UVES high-resolution optical spectra of several clusters, and an infrared atlas of $\\omega$ Cen using the Spitzer Space Telescope. ",
        "introduction": "Mass loss from giants is of great importance to stellar evolution --- low-mass stars lose up to 0.2 M$_{\\odot}$ on the giant branch. Metal-poor globulars giants can probe metallicity dependances of this mass loss, exposing the Universe's chemical enrichment history, and revealing how low-metallicity stars can produce large quantities of dust (\\emph{e.g.} \\citeauthor{BWvL+06}~\\citeyear{BWvL+06}). ",
        "conclusions": ""
    },
    "0609/astro-ph0609736_arXiv.txt": {
        "abstract": "We report improved transformation equations between the $u'g'r'i'z'$ and $UBVR_CI_C$ photometric systems.  Although the details of the transformations depend on luminosity class, we find a typical rms scatter on the order of 0.001 magnitude if the sample is limited to main sequence stars.  Furthermore, we find an accurate transformation requires complex, multi-color dependencies for the bluer bandpasses. Results for giant stars will be reported in a subsequent paper. ",
        "introduction": "Stellar photometry has a long and successful history in the study of stellar properties.  However, high precision photometry is necessary in order to enable accurate measurements of extinction, metallicity, and tests of stellar evolutionary models.  Several filter systems have been developed over the last few decades with strengths and weaknesses associated with each (see \\citep{Straizys99} for a review).  Recently, the Sloan Digital Sky Survey (hereafter SDSS; \\citep{York00}) has greatly enhanced the available photometric data for Galactic stars and thereby promises to significantly add to our understanding of the Galaxy's stellar populations and formation history.  Although the $u'g'r'i'z'$ filter system was defined by \\citet[][F96]{Fukugita96} for the study of large scale structure and quasars in the northern galactic cap, these data have also provided a wealth of photometric data for Galactic stars \\citep[e.g.][]{Newberg99}.  Given the great potential of the SDSS data, we have begun a photometric study of star clusters in the $u'g'r'i'z'$ filter system \\citep{Rider04,Moore05,Smith04}. However, most of the existing photometric data and stellar evolutionary models are limited to the Johnson- Cousins $UBVR_CI_C$ filter system.  Thus accurate transformation equations are necessary in order to evaluate the utility of the $u'g'r'i'z'$ filter system.  As part of the development of the SDSS photometric system, F96 derived preliminary transformation equations between the $UBVR_CI_C$ and $u'g'r'i'z'$ filters.  These transformations were derived from the expected bandpasses, the predicted response of the SDSS detectors, and the spectral energy distribution of a metal poor, F subdwarf, specifically BD+17$^{\\circ}$4708.  This star was chosen as the primary standard due to the availability of high precision, absolute spectrophotometry (see F96).  The initial transformation equations of F96 were primarly intended to enable the simulation of the colors of various celestial objects.  As a result, F96 assumed that the color transformation equations (e.g. $u' - g'$ vs. \\ub) would be linear between the two photometric systems.  More accurate empirical transformations were derived by \\citet{Smith02} from $u'g'r'i'z'$  photometry of a subset of the $UBVR_CI_C$ standard stars from \\citet{Landolt92}.  \\citet{Smith02} also assumed a linear transformation and made no distinction between luminosity classes among their standard stars. Finally, more complex transformations for $UBV$ to $u'g'r'$ filters were presented by \\citet{Karaali05}.  These equations include two Johnson colors for the $u' - g'$ and $g' - r'$ relations, but do not make a distinction between luminosity classes. In this paper we re-examine \\citet{Smith02} data with the goal of deriving more accurate transformation equations between these two photometric systems.  In \\S2 we describe our analysis procedure and \\S3 contains our results. In \\S4 we present our conclusions and describe our ongoing efforts to expand the analysis to other luminosity classes. ",
        "conclusions": "Comparing the residuals and the actual plots for each color transformation demonstrates that our new transformation equations give better results for MS stars than \\citet{Karaali05} and \\citet{Smith02}. Additional work remain to fully develop the complete set of transformation equations for all luminosity classes.  First, more data is needed for extremely red MS stars ($B - V$ $>$ 1.3).  Also, an investigation of the effects of metallicity on the transformations for each class of star must be undertaken.  This means that other isochrone models of $u'g'r'i'z'$ interest, with existing $UBVR_CI_C$ data (such as Y$^2$; \\citep{Yi01}), must be identified and transformed.  This is a complicated and time consuming prospect, but with the advent of high-precision photometry, it must be completed in order to deliver the accuracies needed when describing stellar systems.  We plan to pursue this endeavor in future work. A beginning step is the development of a full metallicity and luminosity photometric study in the $u'g'r'i'z'$ filter system.  This study has all ready begun using the 0.6m telescope at Red Buttes Observatory in Laramie, Wyoming to support the extension of the Sloan Digital Sky Survey."
    },
    "0609/hep-ph0609086_arXiv.txt": {
        "abstract": "We propose a hybrid inflation model with a complex waterfall field which contains an interaction term that breaks the $U(1)$ global symmetry associated to the waterfall field charge.  We show that the asymmetric evolution of the real and imaginary parts of the complex field during the phase transition at the end of inflation translates into a charge asymmetry. The latter strongly depends on the vev of the waterfall field, which is well constrained by diverse cosmological observations. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609785_arXiv.txt": {
        "abstract": "{ The nature of the Hyades stream, or Hyades moving group, is a long-standing question of Galactic Astronomy. While it has become widely recognized that the Hercules stream, an unbound group of stars lagging behind galactic rotation and moving outward in the galactic disk, is associated with the outer Lindblad resonance of the rotating galactic bar, there is still some debate about the nature of the more prominent low-velocity stream sharing the kinematics of the Hyades open cluster. Is this stream caused by additional non-axisymmetric perturbations of the galactic potential, such as transient or quasi-stationary spiral waves, or by the on-going evaporation of the Hyades cluster? A simple observational test has been designed to answer that question, i.e. to determine whether the Hyades stream is primarily composed of coeval stars originating from the Hyades cluster, or of field stars. Using the Geneva-Copenhagen survey of F and G dwarfs, we compare the mass distribution and metallicity of the stream to those of field disk stars. If the Hyades stream is composed of stars trapped at resonance, its mass distribution should obey the present-day mass function (PDMF) of the disk, and its metallicity should reflect its origin in the inner regions of the Galaxy. On the other hand, if it is an evaporated cluster, we expect a different mass distribution, depending on the inital mass function (IMF) of the cluster, and on the proportion of evaporated stars as a function of mass. We find that extreme conditions have to be adopted for the selective evaporation and IMF of the cluster to make the observed mass ditribution of the stream only roughly consistent (at a one-sigma level) with the coeval evaporated cluster scenario. The observed mass distribution is in much better agreement with the PDMF of the field. We also note that the peculiar metallicity of the stream is inconsistent with that of a field population from the solar neighbourhood trapped in the primordial cluster during its formation process and subsequently evaporated. These observations thus favour a resonant origin for the Hyades stream, as suggested in Famaey et al.\\ (2005). ",
        "introduction": "It has been known for a very long time that a spatially unbound group of stars in the solar neighbourhood is sharing the same kinematics as the Hyades open cluster (Hertzsprung 1909, Str\\\"omberg 1922, Eggen 1958, Perryman et al. 1998). Assuming that it was a vestige of an initially more massive cluster which partly evaporated with time, Eggen christened this kinematically cold group the Hyades supercluster. More generally, it is called the Hyades stream, or Hyades moving group.  During the last fifteen years, Eggen's hypothesis that kinematic groups of this type are in fact cluster remnants has been largely debated, because they may also be generated by a number of global dynamical mechanisms. A rotating bar at the centre of the Milky Way may e.g. cause the velocity distribution in the vicinity of the outer Lindblad resonance to become bimodal, due to the coexistence of orbits elongated along and perpendicular to the bar's major axis. Today, this mechanism is thought to actually account for the Hercules stream (Dehnen 2000), a group of stars lagging behind the galactic rotation and moving outward in the disk. Moreover, the perturbation that the triaxial bar induces on a flat axisymmetric disk does produce some chaos: Fux (2001) showed that when the bar is taken into account, the chaotic regions, decoupled from the regular regions, are more heavily crowded in the region of the Hercules stream in velocity space. On the other hand, other streams, including the Hyades, could still be linked with the spirality of the Galaxy, because any perturbation of the axisymmetric potential is likely to buffet the stars.  Eggen's scenario and the dynamical perturbations are, of course, not mutually incompatible phenomena. Clusters are known to evaporate over time, and there must be some intermediate state when a group of stars no longer spatially identifiable still share similar velocities. Indeed, disk stars (most of which move on quasi-circular epicyclic orbits) which formed at the same place and time, and which stayed in the same region of the Galaxy after a few galactic rotations must necessarily have the same period of revolution around the Galactic center, and thus the same guiding-center which in turn implies the same tangential velocity (Woolley 1961). On the other hand, our Galaxy is known to have large spiral perturbations that are likely to have a kinematic effect on the velocity of stars. In the past, the importance of stirring by spiral structure has been underestimated because it was thought that spirals heated the disk strongly, and because the amount of heating is observationally constrained. However, Sellwood \\& Binney (2002) showed that the dominant effect of spirals is to stirr without heating.  Thus, the Hyades stream may well be an outward-moving stream of stars on horseshoe orbits that cross the corotation of the spiral pattern. The kinematics of K and M giant stars in the solar neighbourhood seem to comfort this scenario, and suggest that the Hyades stream is composed of metal-rich stars with a wide range of ages (Famaey et al. 2005, see also Chereul \\& Grenon 2001). However, metallicities were available for too few giants, and individual age estimates were not precise enough to reach a decisive conclusion. This motivates the present analysis, where we examine whether properties of F and G dwarfs from the Geneva-Copenhagen survey (Nordstr\\\"om et al. 2004) are in better agreement with this dynamical scenario, with Eggen's scenario, or with a mix of both. ",
        "conclusions": "Using the Geneva-Copenhagen catalogue of F and G dwarfs (Nordstr\\\"om et al. 2004), we have analyzed the mass function of the Hyades stream, an overdensity of stars in the $UV$-plane kinematically associated with the Hyades cluster.  In order to be compatible at a one-sigma level with the predicted mass function for stars evaporated from the primordial Hyades cluster, one needs a rather extreme scenario, with a very significant preferential evaporation from the cluster of 1~$M_\\odot$ stars over 1.5~$M_\\odot$ stars, and with a very steep IMF, not flatter than a power-law exponent of -3.0 (see Sects. 4.1 and 4.2). The initial total mass and virial radius of the cluster should also be very high ($M>1400$ $M_\\odot$, $r>7.5$ pc).  On the other hand, the observed mass function (see Sect. 4.2, Eq.~\\ref{Eq:value}) is perfectly compatible with the hypothesis that the stream is uniquely composed of field-like stars obeying the PDMF. However, because the Hyades cluster is known to share the kinematics of the stream and to evaporate over time, the most likely scenario is that the stream is indeed mainly composed of field-like stars (about 85\\% of the stream for low-mass stars), but also partly of coeval stars evaporated from the primordial Hyades cluster (about 15\\% of the stream for low-mass stars).  The peculiar metallicity of the stream at all masses (see Sect. 4.3) is not compatible with a scenario where the field-like stars would have been trapped in the primordial Hyades cluster during its formation process, and subsequently evaporated. Those stars were thus most probably trapped at resonance by a spiral perturbation. Indeed, a series of strong transient spirals with their mean corotation at the solar galactocentric radius are known to produce small-scale structure in the local velocity distribution (De Simone et al. 2004). Stars on horseshoe orbits that cross the corotation could wander over 2-3 kpc in much less than 1 Gyr (Sellwood \\& Binney 2002; L\\'epine et al. 2003), while staying on quasi-circular orbits, not elongated enough to betray their place of birth. This is consistent with a galactocentric origin near $R=6$ kpc for the Hyades stream, that could explain its metallicity excess at low masses. On the other hand, the Hyades stream could also correspond to nearly closed orbits trapped at the $4:1$ inner Lindblad resonance of a two-armed spiral density wave (Quillen \\& Minchev 2005). In any case, the prominence of the Hyades dynamical stream raises the question of the amplitude of the spiral perturbation needed to produce such a stream. The most detailed models of gas flows in the Galaxy (Bissantz et al. 2003) indicate that the amplitude of the spiral structure in the mass density is larger by a factor $1.5$ than its amplitude in the near-infrared luminosity density, while such a large amplitude is also needed to produce the large non-axisymmetric motion of the star-forming region W3OH (Xu et al. 2006). This requires the baryonic disk to be massive even near the Sun, and is very constraining for the dark matter distribution in the Galaxy (see e.g. Famaey \\& Binney 2005). On the other hand, the probable resonant origin of the Hyades stream cautions against naive backward integration of individual orbits in an axisymmetric potential when describing the past evolution of such streams, and it bares the importance of evaluating the impact of radial migrations on the chemical evolution of the Galaxy (see also Haywood 2006)."
    },
    "0609/gr-qc0609017_arXiv.txt": {
        "abstract": "{The dynamics of density and metric perturbations is investigated for the previously developed model where the decay of the vacuum energy into matter (or vice versa) is due to the renormalization group (RG) running of the cosmological constant (CC) term. The evolution of the CC depends on the single parameter $\\nu$, which characterizes the running of the CC produced by the quantum effects of matter fields of the unknown high energy theory below the Planck scale. The sign of $\\nu$ indicates whether bosons or fermions dominate in the running.  The spectrum of perturbations is computed assuming an adiabatic regime and an isotropic stress tensor. Moreover, the perturbations of the CC term are generated from the simplest covariant form suggested by the RG model under consideration. The corresponding numerical analysis shows that for $\\nu>0$ there is a depletion of the matter power spectrum at low scales (large wave numbers) as compared to the standard $\\Lambda$CDM model, whereas for $\\nu<0$ there is an excess of power at low scales. We find that the LSS data rule out the range $\\left|\\nu\\right|> 10^{-4}$ while the values $\\left|\\nu\\right|\\leq 10^{-6}$ look perfectly acceptable. For $\\nu<0$ the excess of power at low scales grows rapidly and the bound is more severe. From the particle physics viewpoint, the values \\ $|\\nu|\\simeq 10^{-6}$ correspond to the ``desert'' in the mass spectrum above the GUT scale $M_X\\sim 10^{16}\\,GeV$. Our results are consistent with those obtained in other dynamical models admitting an interaction between dark matter and dark energy. We find that the matter power spectrum analysis is a highly efficient method to discover a possible scale dependence of the vacuum energy.} ",
        "introduction": "\\label{sect:Intro}  The analysis of cosmic perturbations\\,\\cite{peebles,LL,weinberg} represents one of the main tests for the candidate cosmological models and especially for the models of a time-dependent Dark Energy (DE). The recent data on the accelerated expansion of the Universe from the high redshift Type Ia supernovae \\cite{Supernovae}, combined with the measurements of the CMB anisotropies \\cite{CMBR,WMAP, WMAP3Y} and LSS\\,\\cite{cole,tegmark} indicate that the DE is responsible for most of the energy balance in the present Universe. The natural candidate to be the DE is the cosmological constant, mainly if we take into account that from the quantum field theory (QFT) point of view the zero value of the CC would be extremely unnatural\\,\\cite{weinRMP,nova,cosm}. The solution of the standard CC problems, such as the origin of the fine-tuning between vacuum and induced counterparts and the problem of coincidence, is not known. Historically the first attempts to solve some of these problems went along the lines of dynamical adjustments mechanisms\\,\\cite{Dolgov}. Nowadays the discussion is approached from many different perspectives and has generated an extensive literature, see e.g. \\,\\cite{CCRev,Copeland06} and references therein. One of the most interesting aspects of these discussions is the possibility to have a slowly varying DE density of the Universe\\,\\cite{RatraPeebles}. This option cannot be ruled out by the analysis of the present data\\,\\cite{Alam,Jassal}. In case that the new generation of cosmic observational experiments will indeed detect a variable DE, the natural question to ask is: would this fact really mean the manifestation of a qualitatively new physical reality such as quintessence \\cite{quint}, Chaplygin gas \\cite{chap}, extra dimensions \\cite{extra} or of low-energy quantum gravity \\cite{PVT,pol,antmot,reuter}?  Let us emphasize here that one cannot honestly address this question without also exploring the possibility of a time-variable CC whose running is due to known physical effects, that is, due to the quantum effects of matter fields on a curved classical background\\,\\cite{book}.  The most natural manner to investigate the possible running of the CC produced by the quantum effects of matter fields is the renormalization group (RG) in curved space-time \\cite{lam,nova,cosm}. Unfortunately, the existing calculational methods are not sufficient for performing a complete theoretical investigation of this problem \\cite{apco}. At the same time, the phenomenological proposal \\cite{nova,cosm,RGCC,IRGA03} shows that a positive result can not be ruled out. The hypothesis of a standard quadratic form of decoupling for the quantum effects of matter fields at low energies \\cite{nova,babic} leads to consistent cosmological models of a running CC with  potentially observable consequences \\cite{CCfit,AHEP03,Gruni,SS1,SolSte}. Among the many possibilities one can distinguish essentially two different models of running CC in QFT in curved space-time \\footnote{ Generalizations of this kind of semiclassical RG cosmological models (which may e.g. help to alleviate the cosmic coincidence problem) are also possible, see \\cite{LXCDM1}}. The first one \\cite{RGCC,IRGA03,CCfit} admits that there is energy exchange between matter and vacuum sectors of the theory\\,\\footnote{See \\,\\cite{BarrowClifton} for a general discussion of energy exchange between cosmic fluids.}. This leads to the vacuum energy density function $\\rho_\\La=\\rho_\\La(H)$ associated to the cosmological term, $\\Lambda$, while the Newton constant $G$ remains invariable\\,\\footnote{Let us notice that here we follow \\cite{nova} and identify the cosmic energy scale with the Hubble parameter $H$, despite other choices are also possible \\cite{nova, cosm, reuter,babic,baier}.}. The second model does not admit a matter-vacuum energy exchange, but has both parameters scale-dependent $G=G(H)$ and $\\rho_\\La=\\rho_\\La(H)$\\,\\cite{Gruni}. In the present work we shall concentrate on the analysis of the perturbations for the first model only, and postpone the investigation of the second model for another publication. Let us remark that the analysis of cosmic perturbations in the first model has already been addressed in \\cite{OphPel} using an indirect procedure based on bounding the amplification of the density matter spectrum in the recombination era caused by vacuum decay into CDM. This amplification is necessary to compensate for the dilution of the $\\delta\\rho_m/\\rho_m$ matter spectrum at low redshifts resulting from the enhanced matter-radiation density $\\rho_m$ -- associated to the decay of the vacuum into CDM. A qualitatively similar analysis has been performed in \\cite{WangMeng}, and also in \\cite{AlcLim} using the so-called CMB shift parameter\\,\\cite{WMAP}. It is worth noticing that the upper bounds for the running obtained in all these references disagree between themselves, in some cases at the level of several orders of magnitude. This severe divergence illustrates the importance and necessity of performing a direct calculation of the perturbations, which we undertake here thoroughly. For example, the estimate used in \\cite{OphPel} was derived from the phenomenological study presented in Ref.\\,\\cite{Peacock}, where the matter power spectrum of density fluctuations measured by the 2dF Galaxy Redshift Survey (2dFGRS) was compared with the corresponding matter power spectrum derived from the measurements of the CMB anisotropies. The comparison served to place a bound at the level of $10\\%$ on the maximum possible difference between the two derivations of the matter power spectrum\\,\\cite{Peacock}. By requiring that the amplification of the density fluctuations of the RG model at the recombination era do not surpass this bound the authors of \\cite{OphPel} were able to obtain a corresponding bound on the the RG parameter of\\,\\cite{RGCC}, specifically $\\nu<10^{-3}$. Of course this bound is only approximate as it does not come from a direct calculation of the density perturbations in the RG model. Notwithstanding, as we shall see in what follows, there is a qualitative agreement between these results and the ones obtained in the direct calculation that we present here. In both cases the RG models with Planck-scale mass particles are not favored and one is led to conclude that a significant energy gap (or ``desert'') below the Planck scale, $M_P$, is necessary to avoid contradictions between the calculated spectrum of the density perturbations and the observational data. The favored mass scale turns out to be three orders of magnitude below $M_P$, and can be identified with the typical GUT scale $M_X\\sim 10^{16}\\,GeV$. In this paper we will provide a rigorous quantitative derivation of this result. Let us also remark that the analysis of perturbations presented here is not covered by the many works in the literature dealing with perturbations in a Universe with dynamical dark energy -- see e.g. \\,\\cite{Copeland06} and references therein. In fact, the nature of the DE under consideration is not directly related to the properties of dynamical scalar field(s), but to the QFT running of the cosmological constant. Therefore, our approach is qualitatively new and has not been dealt with in the literature before. It constitutes, to our knowledge, the first computation of density perturbations in the presence of a running cosmological term. This is important in order to assess the impact on structure formation from a possible scale dependence of the vacuum energy. \\\\ \\indent The paper is organized as follows. In the next section we summarize the necessary information about the variable CC model under discussion. The reader may consult the papers \\cite{nova,RGCC,CCfit,AHEP03} for further technical details and extensive discussions. In section 3 we derive the equations for the perturbations. Section 4 is devoted to the numerical solution of these equations and to their comparison with the 2dFGRS\\,\\cite{cole}. Finally, in the last section we draw our conclusions. ",
        "conclusions": "\\label{sect:Conclusion}  The possibility that the cosmological constant $\\Lambda$ can be a running parameter within QFT in curved space-time was investigated extensively in \\,\\cite{nova}. The phenomenological implications of these renormalization group (RG) cosmologies on supernovae observations have been amply addressed in different works preceding this one \\,\\cite{RGCC,IRGA03,babic,CCfit,AHEP03,Gruni,SS1,SolSte}. Here we have extended the analysis of phenomenological implications to the very relevant issue of structure formation. We have studied in detail the power spectrum of matter density perturbations for the FLRW type of cosmological model with running $\\Lambda$. This model is characterized by a fundamental parameter $\\nu$ associated to the running of $\\Lambda$. At the same time this parameter acts, in the present realization of the model, as a coupling between the vacuum energy and matter (mainly CDM). There are other implementations of this framework where $\\nu$ does not lead to this kind of coupling, see \\cite{Gruni}. However, in the present work we have just focused on the original RG model\\,\\cite{RGCC,CCfit}, where this coupling is present and, therefore, the cosmological term can decay into matter and vice versa. Then, under the simplifying assumptions of adiabaticity (i.e. of having no perturbations associated to entropy exchange) and of having no anisotropic stress contributions, we have derived the complete coupled set of matter density and metric perturbations for this model. The RG running of $\\Lambda$ is based on the assumption of a standard quadratic decoupling law at low energies\\,\\cite{RGCC,CCfit} and we have taken the simplest covariant form of this equation -- see (\\ref{A and B}) -- as our ansatz to derive the perturbations. This phenomenological input leads to a cosmological model with the $H^2$-dependence of the vacuum energy density\\,\\cite{nova}. It is important to emphasize that this does not mean that $\\Lambda\\sim H^2$, but rather that the (renormalization group) variation of $\\Lambda$ satisfies $\\delta\\Lambda\\sim H^2$, and hence $\\Lambda=a+b\\,H^2$ -- see Eq.\\,(\\ref{Friedmann2}). The effect produced by the running depends on a single parameter $\\nu$, which characterizes the particle mass spectrum of the quantum theory below the Planck scale. For $\\nu=0$ we retrieve the standard $\\Lambda$CDM model with time-independent cosmological constant, while the values around the ``canonical'' value $\\nu=\\nu_0\\sim 10^{-2}$ correspond to having effective particle contributions of order of the Planck mass.  The numerical analysis of the perturbations proves to be in qualitative agreement with the results of the previous estimate in Ref.\\,\\cite{OphPel}, based on restraining the amplification of the density fluctuating spectrum at the recombination era caused by the vacuum decay into CDM.  In both cases (viz. the direct calculation and the previous estimate) it is found that the ``canonical'' value of $\\nu$ leads to strong deviations from the observational data. For positive values of $\\nu$ in the range $10^{-6}\\leqslant\\nu\\leqslant 10^{-5}$, corresponding approximately to the GUT particle spectrum ($M_X \\sim (10^{-2}-10^{-3})\\,M_P$), we meet relatively small deviations of the theoretical curve from the $\\nu=0$ case. The deviations are higher the larger is $|\\nu|$. Let us recall that the upper bound on $\\nu$ obtained from the phenomenological estimate of\\,\\cite{OphPel} is $10^{-3}$ and applies only for $\\nu\\geqslant 0$. This bound is not very far away from our own most conservative result $|\\nu|<10^{-4}$, although in our case it applies for both signs of $\\nu$. Moreover, the approximate bound is much weaker than the one defined by the so-called ``safest $\\nu$ range'' -- cf. Eq.\\, (\\ref{nurange}). However, as we already emphasized, we cannot definitely exclude the $\\nu>0$ interval $10^{-6}-10^{-4}$ (and specially the subinterval $10^{-6}-10^{-5}$). Therefore, we may assert, in the most possible conservative way, that our direct bounds on $\\nu> 0$ are between one to two orders of magnitude more restrictive than those of \\cite{OphPel}, but of course they have been derived in a much more rigorous way.  Let us also remark that our exclusion of the possibility that $\\nu$ could be of the order of magnitude of the ``canonical'' value $\\nu_0\\sim 10^{-2}$ is at variance (by at least two orders of magnitude) with the estimates obtained from the approximate methods used in references \\cite{WangMeng,AlcLim}\\,\\footnote{The results of these references were presented in terms of a parameter $\\epsilon$, and they concluded that values of $\\epsilon$ at the level of $10^{-1}$ are allowed. However, $\\epsilon$ is nothing but $3$ times the parameter $\\nu$ originally defined in \\cite{RGCC}.}.  It is also interesting to compare our results with those of Ref.\\,\\cite{pavon}, where a detailed study of matter density perturbations in interacting quintessence models is addressed. These authors perform their study also under the assumption of an adiabatic regime and with no contributions from anisotropic stress. In this kind of modified quintessence models, first proposed in \\,\\cite{Amendola}, one couples matter to quintessence by means of a source function $Q$ which depends on a free parameter, called $c^2$ in Ref.\\,\\cite{pavon}. This is the only free parameter in that model, similarly as is $\\nu$ in our own framework under the specified common set of assumptions. In the models of interactive quintessence, the dark energy (DE) can also decay into matter (parallel to our case where $\\Lambda$ can also decay into matter) and as a result the amount of DE in the past is larger than in ordinary quintessence models with self-conserved DE. This corresponds, in our framework, to have $\\nu>0$ because then $\\Lambda$ becomes larger in the past. The authors of \\,\\cite{pavon} find that the effect of the $c^2$-coupling is to introduce a damping of the power spectrum of matter perturbations at low scales; that is to say, the growth of density perturbations for large wave numbers becomes smaller as compared to quintessence models with uncoupled DE to matter. In our RG framework, we observed a similar depletion of the spectrum at low scales for $\\nu>0$ as compared to $\\nu=0$ (the standard $\\Lambda$CDM model) -- cf. sections \\ref{sect:initialcond} and \\ref{sect:powerspectrum}. Moreover, we have found numerical bounds on $\\nu$ that are quantitatively similar, in fact a bit more stringent than those on $c^2$ obtained in the model of \\,\\cite{pavon}. These authors found that their power spectrum is consistent with the 2dFGRS data provided $c^2\\leq 10^{-3}$, and we find that the largest allowed value of $\\nu$ by the same LSS data is of order $10^{-4}$. The basic agreement among the different analyses suggest that the loss of power at low scales could be a general feature of dynamical models in which the dark energy decays into matter during the cosmic evolution. The latter situation is always the case in Ref. \\,\\cite{pavon} because $c^2>0$, and it is also our case for $\\nu>0$. However, in our RG framework, $\\nu$ can have any sign (depending on the dominance of bosons or fermions in the loop contributions to the running of $\\Lambda$) and so we have an opportunity to explore also the case $\\nu<0$. For this sign of $\\nu$, the running $\\Lambda$ becomes progressively smaller in the past and reaches eventually large negative values (see the detailed evolution plots of Ref.\\,\\cite{CCfit}). As a result we may expect an overproduction of structure associated to the fact that a large negative $\\Lambda$ increases the gravitational collapse. Since, however, this is not observed (cf. the plots of section \\ref{sect:powerspectrum} detecting an excess of power at low scales for $\\nu<0$), this may provide a physical explanation of why negative values of $\\nu$ are significantly more restrained. While the safest and most stringent range of values of $\\nu$ is expressed in a nutshell in Eq.\\,(\\ref{nurange}), we cannot exclude that $|\\nu|$ could be higher (mainly in the $\\nu>0$ case) once it will be possible to combine the analysis of the matter power spectrum with the corresponding CMB analysis. At the same time, we should recall that the bounds obtained on $\\nu$ in our case, and on $c^2$ in \\,\\cite{pavon}, could suffer some renormalization if we would abandon the set of simplifying hypotheses that have been made. However, our study shows at least the kind of impact to be expected on the matter power spectrum predicted by cosmological models with running $\\Lambda$ in interaction with matter. Further investigations would be needed to assess the effect in a more thorough way, but they lie beyond the scope of the present work.  Let us also finally remark that the bounds we have placed on the RG parameter $\\nu$, within our set of assumptions, can be interpreted on purely phenomenological grounds, namely as bounds on a parameter that gauges the possible degree of interaction between CDM and the vacuum energy, even if no fundamental RG model is invoked. In principle, the phenomenological implications of a non-vanishing $\\nu$ could be potentially detectable in the future supernovae observations, either through a direct measure of the running of $\\Lambda$ \\,\\cite{CCfit,AHEP03} or (more efficiently in practice) through the possibility to observe the redshift evolution of the effective equation of state parameter associated to the variable CC model \\,{\\cite{SS1,SolSte}}. However, the very study of the evolution of the density and metric perturbations, which has led us to firmly exclude the range $|\\nu|>10^{-4}$, already demonstrates the difficulty of the previous methods and at the same time emphasizes the great advantage of the density perturbations approach. Indeed, the strong sensitivity of the matter power spectrum to a variable cosmological term interacting with the CDM shows that this is perhaps the ideal observable to look at. Therefore, future data on LSS could eventually provide the clue to unravel an interaction (even if very small) between matter (essentially CDM) and vacuum, and at the same time it may hint at the possible RG physics behind it. Let us conclude by mentioning that it should be very interesting to expand our analysis to the aforementioned second model of running CC \\cite{Gruni}. In this alternate model the cosmological term is also evolving in time and redshift, but there is no energy exchange between matter and vacuum and we may expect much more freedom in choosing the value of $\\nu$. We leave this study for a future publication. \\vspace{0.3cm}  {\\bf Acknowledgments:} We are grateful to J\\'er\\^ome Martin for useful discussions, to Hrvoje \\v{S}tefan\\v{c}i\\'{c} for checking our numerical code and for interesting discussions, and finally to Javier Grande for his aid in improving the presentation of our plots and also for useful discussions. The work of J.C.F has been supported in part by CNPq (Brazil) and by CAPES/COFECUB Brazilian-French scientific cooperation. The work of I.Sh. has been supported in part by the grants from CNPq (Brazil), FAPEMIG (Minas Gerais, Brazil) and ICTP (Italy). The work of J.S. has been partially supported by MECYT and FEDER under project 2004-04582-C02-01, and also by DURSI Generalitat de Catalunya under project 2005SGR00564."
    },
    "0609/astro-ph0609050_arXiv.txt": {
        "abstract": "Published maps of red giant stars in the halo region of M31 exhibit a giant stellar stream to the south of this galaxy, as well as a giant ``shelf'' to the northeast of M31's center.  Using these maps, we find that there is a fainter shelf of comparable size on the western side as well.  By choosing appropriate structural and orbital parameters for an accreting dwarf satellite within the accurate M31 potential model of Geehan \\ea (2006), we produce a very similar structure in an $N$-body simulation.  In this scenario, the tidal stream produced at pericenter of the satellite's orbit matches the observed southern stream, while the forward continuation of this tidal stream makes up two orbital loops, broadened into fan-like structures by successive pericentric passages; these loops correspond to the NE and W shelves.  The tidal debris from the satellite also reproduces a previously-observed ``stream'' of counterrotating PNe and a related stream seen in red giant stars.  The debris pattern in our simulation resembles the shell systems detected around many elliptical galaxies, though this is the first identification of a shell system in a spiral galaxy and the first in any galaxy close enough to allow measurements of stellar velocities and relative distances.  We discuss the physics of these partial shells, highlighting the role played by spatial and velocity caustics in the observations.  We show that kinematic surveys of the tidal debris will provide a sensitive measurement of M31's halo potential, while quantifying the surface density of debris in the shelves will let us reconstruct the original mass and time of disruption of the progenitor satellite. ",
        "introduction": "Our neighboring disk galaxy M31 has turned out to be a well-equipped laboratory for research into the accretion and disruption of substructure in galactic halos.  A stellar stream of length $\\tsim 120 \\kpc$ is observed to flow from its tip south and behind M31 into M31's center \\citep{ibata01,mcconnachie03,ibata04,raja06}, while many other irregular morphological and kinematic features indicative of accretion are seen in the galactic disk and halo \\citep{ferguson02,merrett03,ibata05}. The giant southern stream is thought to be formed at the last pericentric passage of a progenitor satellite with mass $\\tsim 10^9 \\Msun$ on a highly radial orbit (\\citealp{ibata04,font06}; \\citealp{geehan06}, hereafter Paper~I; \\citealp{fardal06}, hereafter Paper~II). This stream holds out the promise of the first precise measurements of the gravitational potential in M31's halo. This is important for many reasons, including breaking the disk-halo degeneracy in the mass distribution \\citepalias{geehan06} and thereby understanding the response of the dark halo to the central baryons.  However, there are two obstacles to these measurements at present \\citepalias{fardal06}.  First, the current observational errors on the distance allow too much flexibility in the orbits to constrain the potential effectively.  Second, the stream does not follow a single orbit, but has a gradient of energy along the length of the stream, ``tilting'' the stream in phase space relative to the orbit and biasing the measurement of the potential.  The amount of this tilt depends on the location of the progenitor, which is currently unknown. Locating the progenitor would both specify the stream-orbit tilt and constrain the allowed orbital trajectories, removing both obstacles.  The ``Northeast Shelf'', a faint diffuse feature seen to the NE side of M31 in the red giant branch (RGB) star-count maps of \\citet{ferguson02}, has been suggested as a possibility for the current location of the progenitor, or at least as the forward continuation of the stream \\citep{ibata04,font06}.  Below we find that there is a similar, even fainter shelf or plateau on the other side of M31, which we call the ``Western shelf''.  In this paper, we focus on the possibility that the forward continuation of the stream makes up these NE and W shelves.  In this scenario, the shelves are a similar phenomenon to the ``shell'' systems seen around many elliptical galaxies.  These systems are believed to be produced by tidal disruption of a satellite galaxy on a nearly radial orbit \\citep[e.g.,][]{schweizer80,quinn84,hernquist88,barnes92}, exactly the phenomenon believed to be responsible for the creation of M31's southern stream, so it is not unexpected that shell-like features should be present. These coherent structures are rich in information. For example, they allow sensitive tests of the potential of the host galaxy without the need for distance information \\citep{merrifield98}, if the necessary kinematic observations can be obtained as is possible for the first time in the M31 shells.  Also, there exists a simple relation between the surface density of a shell and the rate at which stellar mass flows through the shell; when appied to the W shelf, this can be used to deduce the orbital parameters and the mass of the satellite that produced the observed debris.  Our purpose in this paper is threefold: first, to show that this paradigm of a shell debris pattern is consistent with the current observations of M31; second, to explain the basic physics that gives rise to the spatial and kinematic pattern of the debris; and third, to find observational diagnostics to further constrain the model of the debris from present and future datasets. In \\S~2, we give an overview of the complex observed structure around M31, and describe our scenario to explain a major part of this structure. In \\S~3, we explain our methods for setting up and running a simulation of this scenario.  We then compare the simulation results to the observed properties of the debris around M31, to see whether our model is plausible.  In \\S~4, we discuss the physics that governs this model of the satellite debris.  We consider several observables that can be used to confirm or rule out our scenario, and to constrain the physical properties of the stream and its progenitor.  In \\S~5, we discuss some remaining issues regarding our model, and state our conclusions.  \\nobreak ",
        "conclusions": "We have performed a simple $N$-body simulation in the gravitational potential of M31, and used this to show that an accreting satellite can explain the giant southern stream, the eastern shelf, and a similar shelf on the western side as a single continuous structure. Following the entire folded path traced by the stream and shelf particles, the structure is $\\gtrsim 230 \\kpc$ long, certainly among the largest coherent stellar structures known despite its modest total stellar mass of $\\sim 2 \\times 10^9 \\msun$.  The NE and W shelves in this model represent successive radial loops of the orbit broadened into fan shapes by passage through pericenter. The NE shelf lies closer to us than the disk of M31, and partially overlaps with it. This may have relevance to microlensing surveys which aim to detect compact objects of stellar-type mass \\citep{calchi05,dejong04}, since the lensing event probability will be boosted relative to disk self-lensing by the gap in distance between the disk and NE shelf, and by the large relative velocities between the disk and shelf stars.  The counter-rotating ``stream'' in the NE seen in both PNe and RGB stars in this view represents a continuation of the NE shelf onto the face of M31, while the apparently narrow velocity width of this stream is explained by a caustic in velocity space.  The properties of the incoming satellite are reasonable for a galaxy of its initial stellar mass, as discussed in \\S~3.1.  While we have not explored velocity distributions outside the shelf regions in detail, our impression is that the model is satisfactory in terms of where it does not put debris, as well as where it does.  Our experience with different initial conditions (\\citetalias[in][]{fardal06}) suggests that this would not necessarily be true for orbits that tried to associate the southern stream with other objects such as M32 or the Northern Spur.  We experimented with low-resolution versions of our main run, testing different orbits and satellite properties.  Our sense acquired during this exploration is that the volume of parameter space that can match current observations is large enough not to need extreme fine-tuning. The time corresponding to ``present day'' within our run is uncertain by perhaps 100 Myr, and it is not clear whether the progenitor need be disrupted or not.  As surveys of M31's disk and halo regions continue, we expect that the requirements on the model will become much more stringent.  We remind the reader that we have neglected several physical effects in computing this model that are potentially important.  These include dynamical friction, the response of the modes in the disk, rotation of the accreting satellite, dark matter in and around its stellar component, and the hydrodynamic interaction of gas within the satellite with the disk and halo of M31. In addition, the observations do not yet cover some of the most interesting areas for this model, and our comparison to the existing observations is as yet rudimentary.  We therefore regard the existing simulation as a useful illustration, rather than a definitive prediction of the debris structure.  Work on the interaction of the satellite with live dynamical components of M31 is currently underway.  Kinematic observations in the shelf regions are important for two reasons.  First, they can help clarify the nature and extent of the debris related to the giant southern stream.   The many testable predictions of our model include: \\begin{itemize} \\item The ``PNe stream'' of \\citet{merrett06} should be continuously connected to the NE shelf, since in our model it is merely a region where the kinematics in projection gives a roughly constant velocity, and not a separate structure. \\item In both shelves, the characteristic kinematic pattern of a shell should be evident in the $R_{proj}$--$v_{los}$ plane (Figure~\\ref{fig.vshelf}). \\item There should be stars with positive velocities returning to the center of M31 in the NE shelf, which connect the NE and W shelves. \\item There may be a yet weaker shelf on the E side, from debris further forward in the stream. \\end{itemize} Second, the velocity signature of the shells depend on the potential of M31. Many more spectroscopic fields will be required to allow a precise measurement of the potential.  However, unlike the case of the collimated southern stream, the line-of-sight distances need not be computed at all to probe the potential of M31, as the geometry itself selects the stars in the velocity caustic.  This removes one of the main sources of uncertainty in constraining the mass model.  Comparison of the potential derived from the shelves to that obtained from H~I rotation curves can help test the deviation of M31's mass distribution from spherical symmetry.  If confirmed as shells connected to the southern stream, the NE and W shelves would to our knowledge represent the first detection of a multiple-shell system in a late-type spiral galaxy. Shells are very common in elliptical galaxies.  137 elliptical galaxies with shells were catalogued by \\citet{malin83}, while \\citet{merrifield98} estimate that 10\\% of elliptical galaxies have shells. The reasons for this strong morphological bias have not been completely explained at present. Spirals may be destroyed and ellipticals created in the mergers that produced the shells, although this would only be plausible for merging systems much more massive than in the case discussed here \\citep[see][for a discussion of shell formation in major mergers]{hernquist92}. Alternatively, shells may dissolve faster in the non-spherical potential of the disk due to differential precession of the orbits. The difficulty of distinguishing shells from spiral features in late-type galaxies may also contribute to the lower rate of observing shells \\citep{barnes92}. If all shells around spiral galaxies were as faint as the ones in M31, they would have escaped detection, since RGB star count maps like those of \\citet{ferguson02} can be made only at close range. M31's putative shell system would also be the first in which kinematic observations have been made, allowing us to compare to and eventually constrain mass models of the host galaxy.  Theoretical studies of dark halos have established their mean mass profile to high accuracy in the absence of baryons.  Their response to the baryonic component in terms of their mass profile, ellipticity, and alignment with the disk is more difficult to simulate, making observational checks important.  However, these aspects are difficult to study in most galaxies.  The abundance in M31 of both coherent tracers of the potential, such as the tidal debris discussed in this paper, and statistical tracers such as RGB stars, satellite galaxies, and globular clusters, suggests that this galaxy can be a Rosetta Stone for understanding the properties of dark halos and their interaction with the galaxies they host."
    },
    "0609/astro-ph0609099_arXiv.txt": {
        "abstract": "We present a new interpretation of dark energy in terms of an \\textit{Abnormally Weighting Energy} (AWE). This means that dark energy does not couple to gravitation in the same way as ordinary matter, yielding a violation of the weak and strong equivalence principles on cosmological scales. The resulting cosmological mechanism accounts for the Hubble diagram of type Ia supernovae in terms of both cosmic acceleration and variation of the gravitational constant while still accounting for the present tests of general relativity. This explanation allows to build dark energy models (i) without violation of the strong energy condition $p<-\\rho c^2/3$ (ii) with non-negligible direct couplings to gravitation and (iii) natural convergence mechanism toward general relativity. ",
        "introduction": "Maybe the most helpful astrophysical evidence for determining the nature of dark energy is the dimmed magnitude of type Ia supernovae (see Astier et al. 2006 and references therein). This tell us that either the recent cosmic expansion has accelerated or (and) these objects are not standard candles, which could have happened if, for instance, fundamental constants of physics have recently varied. This is the first explanation that is retained in the so-called concordance model $\\Lambda CDM$ in which a positive cosmological constant constitutes the dark energy sector of cosmology. This convenient explanation does not require to give up neither general relativity (GR) nor the cosmological principle and furthermore allows to account for dark energy effects with only one additional parameter. However, the cosmological constant is usually interpreted as the stress-energy term associated to the vacuum state which has non-vanishing energy due to quantum fluctuations. Doing so, one recovers the old problem of the overestimation of the cosmological constant. Indeed, if one assumes that the cut-off energy scale at which GR is no longer valid is the Planck length, then the vacuum energy density acting as a cosmological constant is $\\rho_{\\Lambda}\\approx m_{Pl}^4\\approx 10^{76} GeV^4$ (in geometrical units where $\\hbar=c=1$, $G=m_{Pl}^{-2}$). Which mechanism can be so precise to make such a huge quantity vanishing? This constitutes the so-called \\textit{fine-tuning problem} of the cosmological constant. What is worst now with the astrophysical evidences for the cosmological constant, is that the vacuum energy density is indeed not vanishing but is of cosmological order of magnitude: $\\rho_{\\Lambda}\\approx \\rho_{c}(z=0)\\approx 10^{-47} GeV^4\\cdot$. We have therefore a huge disagreement between theory and experiment of more than $120$ orders of magnitude! Why is the vacuum energy density, something that should be set once for all in the universe and should depend on quantum mechanics, is of the order of the critical density today? Such a value suggests a cosmological mechanism behind the value of the cosmological constant... Even worst, this means that we live in a very particular epoch in the history of the universe : the time when the cosmological constant, whose value has \\textit{a priori} nothing to do with cosmology if it is fixed by quantum considerations, starts dominating the energy content of the universe. This constitutes the so-called \\textit{coincidence problem} associated to the cosmological constant. However, one should keep in mind that this interpretation of dark energy in terms of \\textit{weighting} quantum fluctuations is obviously a quantum gravity problem. One can therefore expect that the final solution to this problem might well be quite different from the simple and maybe rather naive cosmological constant. An interesting way-out of these problems (Padmanabhan 2005) is to consider that dark energy is not related to the absolute value of the energy of the quantum ground state but rather to fluctuations around it. For instance, one can links the dark energy density to the size of the particle (Hubble) horizon which could solve the coincidence problem. However, the resulting mechanism can no longer be interpreted as a cosmological constant as the horizon is not an absolute, constant, notion...\\\\ \\\\ Another widespread interpretation of dark energy is that of quintessence, a \"\\textit{varying cosmological constant}\" which provide a cosmological mechanism aimed to solve coincidence (and in some cases fine-tuning) problems. This most of time consists on considering dark energy as a self-interacting scalar field which couples mostly to ordinary matter indirectly, i.e. through purely gravitational couplings. This means that quintessence does only modify cosmic expansion and does not alter conventional physics. While rolling down its self-interacting potential, this scalar field might produce sufficiently negative pressures ($p<-\\rho c^2/3$) to provide the necessary cosmic acceleration. In the cases of tracking potentials, there exists attractors in the scalar field dynamics which ensure a solution to both fine-tuning and coincidence problems.\\\\ \\\\ Let us now consider what happens if dark energy couples directly to ordinary matter, which is obviously not possible for the cosmological constant. According to its \"\\textit{darkness}\", quintessence is a neutral field and the non-gravitational interactions it could have with matter will only affects the coupling constants. Such a non-minimally coupled scalar field, sometimes called \"\\textit{extended quintessence}\" therefore yields a violation of the equivalence principle. In the case where the scalar field affects the couplings to the metric and therefore makes the gravitational constant varying, we face the famous \\textit{scalar-tensor} theories of gravitation. However, if we consider that the $70\\%$ of missing energy is the whole contribution of such a non-minimally coupled scalar field, then it could be difficult to match the present tests of GR without assuming its direct couplings to be \\textit{extremely} weak. Even worst, as soon as these couplings are non-vanishing, the endless domination of a non-minimally coupled scalar field acting as dark energy might inescapably lead to a disastrous violation of the equivalence principle in the future. Indeed, an unbounded scalar field might drive the coupling constants to values arbitrary far from GR. Therefore, if dark energy is not completely dark and if it does not rule temporarily the cosmic expansion, then GR itself could be a coincidence of the present universe... In this way, the non-gravitational couplings of dark energy to matter are extremely important and should be considered carefully. \\\\ \\\\ Another possibility of such couplings is that \"dark\" energy violates the weak equivalence principle (WEP) on cosmological scales. This means that it does not couple to gravitation as usual matter does and weights abnormally. Doing so, its related gravitational energy will be felt differently by other types of matter, therefore violating also the strong equivalence principle (SEP). Under this assumption, we have built a plausible dark energy mechanism (i) without violation of the strong energy condition (SEC) $p<-\\rho c^2/3$ (ii) with non-negligible direct couplings to gravitation and (iii) with a natural convergence mechanism toward general relativity (F\\\"uzfa \\& Alimi 2006a hereafter AWEI ; F\\\"uzfa \\& Alimi 2006b hereafter AWEII). ",
        "conclusions": "The original AWE hypothesis we presented here assumes that dark energy violates the WEP on cosmological scales, which alters the well-known convergence mechanism of scalar-tensor theories toward GR during the matter-dominated era. We have shown in AWEII that this assumption allows to account for the Hubble diagram of type Ia supernovae through cosmic acceleration and time variation of the gravitational constant, account for stringent constraints on the equivalence principle today, and does not require too negative pressures ($p\\ge -\\rho/3$). Furthermore, this assumption is well-motivated as a generalization of scalar-tensor theories for which the WEP is usually assumed. AWE can also represent an interesting alternative to usual dark energy models in which this extra gravitational energy content is extremely weakly coupled to matter therefore producing very low variation of fundamental constants and violates the SEC. At least, the AWE assumption can help us to test these key features of dark energy. As well, this hypothesis is also expected to be well constrained by post-recombination on the cosmic microwave background and large-scale structure formation due to the modified recent cosmic expansion and the rather important variation of $G$ it could predict. It should also modify the universality of free fall, especially on cosmological scales where the amount of AWE is not negligible. All these effects therefore constitute very important clues for our quest to determine whether some processes in the universe are not ruled by the equivalence principle and how far Einstein was right in postulating it. For all these reasons, we claim that the AWE hypothesis should deserve a particular attention in future works as it goes well beyond the scope of finding a valuable explanation of dark energy but directly questions the deep nature of gravitation itself. \\begin{figure} \\begin{center} \\begin{tabular}{ccc} \\includegraphics[scale=0.3]{fuzfa_fig1a.eps} & \\includegraphics[scale=0.3]{fuzfa_fig1b.eps} \\includegraphics[scale=0.3]{fuzfa_fig1c.eps} \\end{tabular} \\end{center} \\caption{(left) Hubble diagram of SNLS 1st year data set with the best fit $\\Lambda CDM$ flat model (solid line, $\\Omega_m(a_0)=0.26$, $\\bar{\\chi}^2/dof=1.03$) and the AWE model (dash-dotted line, $\\bar{\\chi}^2/dof=1.09$) ($H_0=70 km/s/Mpc$); (center) cosmological evolutions of $G_N$ (in units of bare $G$), the cosmic acceleration factor and (right) the energy densities with the redshift} \\label{fig1} \\end{figure}"
    },
    "0609/astro-ph0609320_arXiv.txt": {
        "abstract": "Microwave, submillimetre-wave, and far-infrared phased arrays are of considerable importance for astronomy. We consider the behaviour imaging phased arrays and interferometric phased arrays from a functional perspective. It is shown that the average powers, field correlations, power fluctuations, and correlations between power fluctuations at the output ports of an imaging or interferometric phased array can be found once the synthesised reception patterns are known. The reception patterns do not have to be orthogonal or even linearly independent. It is shown that the operation of phased arrays is intimately related to the mathematical theory of frames, and that the theory of frames can be used to determine the degree to which any class of intensity or field distribution can be reconstructed unambiguously from the complex amplitudes of the travelling waves at the output ports. The theory can be used to set up a likelihood function that can, through Fisher information, be used to determine the degree to which a phased array can be used to recover the parameters of a parameterised source. For example, it would be possible to explore the way in which a system, perhaps interferometric, might observe two widely separated regions of the sky simultaneously. ",
        "introduction": "There is considerable interest in developing phased arrays for radio astronomy. Projects include the Square Kilometer Array (SKA), the Low Frequency Array (LOFAR), the Electronic Multibeam Radio Astronomy Concept (EMBRACE), and the Karoo Array Telescope (KAT) \\cite{A,B,C}. All of these projects are aimed constructing phased arrays for microwave astronomy, but as technological capability improves, phased arrays will eventually be constructed for submillimetre-wave and far-infrared astronomy \\cite{D,E}.  Two types of phased array are of interest: (i) imaging phased arrays, where an array of coherent receivers is connected to a beam-forming network such that synthesised beams can be created and swept across the sky; (ii) interferometric phased arrays, where the individual antennas of an aperture synthesis interferometer are equipped with phased arrays such that fringes are formed within the synthesised beams. In this way it is possible to extend the field of view, to observe completely different regions of the sky simultaneously, to steer the field of view electronically, and to observe spatial frequencies that are not available to an interferometer because the baselines cannot be made smaller than the diameters of the individual antennas.  It is important to recognise that the synthesised beams of a phased array need not be orthogonal, and may even be linearly dependent. Non-orthogonality may be built into a system intentionally as a way of increasing the fidelity with which an image can be reconstructed, or it may arise inadvertently as a consequence of RF coupling and post-processing cross-talk. In some situations, say in the case of interacting planar antennas, it may not even be clear how to distinguish one basis antenna from another, even before the beam-forming network has been connected.  In this paper, we show that the only information that is needed to determine the average powers, the correlations between the complex travelling wave amplitudes, the fluctuations in power, and the correlations between the fluctuations in power at the output ports of a phased array, or between the output ports of phased arrays on different antennas, is the synthesised beams. It is not necessary to know anything about the internal construction of the arrays or the beam forming networks. Beam patterns may be taken from electromagnetic simulations or experimental data. In the case of interferometric phased arrays, the arrays on the individual antennas do not have to be the same.  The ability to assess the behaviour of a system simply from the synthesised beam patterns separates the process of choosing the best beams for a given application from the process of understanding how to realise the beams in practice. It also suggests important techniques for simulating phased arrays, and for analysing experimental data. ",
        "conclusions": "We have studied the functional behaviour of imaging phased arrays and interferometric phased arrays, and shown that their operation is closely related to the mathematical theory of frames. In order to calculate the behaviour of an imaging phased array, or an interferometric phased array, it is only necessary to know the synthesised reception patterns, which may be non-orthogonal and linearly dependent. It is not necessary to know anything about the internal construction of the array itself. As a consequence, data can be taken from experimental measurements or from electromagnetic simulations. The theory of frames allows one to assess, in a straightforward manner, whether the outputs of a phased array contain sufficient information to allow a field or intensity distribution to be reconstructed in an unambiguous way.  Our model also allows straightforward calculation of quantities such as the correlations in the fluctuations at the output ports of phased arrays. The theory of interferometric phased arrays is almost identical to the theory of multimode bolometric interferometers, and therefore, recently developed techniques for modelling bolometric interferometers can be applied to phased arrays also: including quantum statistics. The work opens up the important possibility of constructing likelihood functions that enable the covariance matrices of source-parameter estimators to be determined. Thus, for example, one could explore the possibility of enhancing source reconstruction by packing in more and more overlapping synthesised beams into a region, or widely separated regions, of finite size. Any enhancement of the accuracy with which source parameters can be recovered, will be related, and to some extent determined, by the degree to which the beam patterns of the array form a frame with respect to all possible incoming field distributions.  In a later paper, we shall use the concepts described here, and the numerical techniques reported previously, to simulate and assess the behaviour of interferometric phased arrays when different optical systems and beam-forming networks are used.  \\begin{appendix}"
    },
    "0609/astro-ph0609116_arXiv.txt": {
        "abstract": "{ {\\sl ROSAT} observations found no convincing evidence for X-ray sources located in local dwarf spheroidal galaxies (dSph).  Now with more sensitive instruments on board {\\sl Chandra} and {\\xmm} we can reach fainter luminosity levels. We report on an observation of the Sagittarius (Sgr) dSph made using {\\sl Chandra} and an observation of the Carina (Car) dSph made using {\\xmm}.  Our observations are sensitive to sources with X-ray luminosities in the $0.1-10$~keV band of $\\sim1\\times10^{32}$ \\ergss and 3$\\times10^{34}$ \\ergss for the Sgr and Car fields respectively.  We have identified a total of 80 sources in the Sgr field and 53 sources in the Car field.  Although the source numbers are roughly consistent with the expected number of background AGN, we found a small fraction of X-ray sources which were soft and could be located in the host dSph.  Follow-up optical/IR observations may help to identify their optical counterparts and hence determine their nature. ",
        "introduction": "Dwarf galaxies are low luminosity galaxies and make up a large fraction of galaxies at the present epoch (Marzke \\& da Costa 1997). Because of their faintness, the best studied systems reside in the Local Group.  The Milky Way has 12 known satellite galaxies.  With the exception of the LMC and SMC, which are irregulars, all of them are dwarf spheroidals (dSphs).  DSphs are generally gas poor and have low metallicity.  Their star formation activity has long since ceased, and their stellar populations are old.  (See Mateo 1998 for a review of Local Group dwarfs.)  While the Local Group dSphs have been well studied in the optical and infra-red, their X-ray properties are yet to be established.  DSphs were seldom observed in the X-ray band, in the belief that they do not contain substantial number of detectable X-ray sources.  A {\\sl ROSAT} observation of the Fornax dSph (which is at a distance of 138 kpc) found 19 sources in the field (Gizis, Mould \\& Djorgovski 1993). Since this number was consistent with the expected number of background AGN, the sources were not thought to be stellar binary sources located in the dSph.  A recent {\\sl Chandra} observation of the Sculptor dSph (at a distance of 79 kpc) identified 5 sources which were likely to be located in that galaxy (Maccarone et al.\\ 2005). They have inferred X-ray luminosities $L_{\\rm x} \\sim 6-90 \\times 10^{33}$~\\ergss, which are well below $10^{37}-10^{38}$~\\ergss, typical luminosities of bright Galactic X-ray binaries.  From these studies we may conclude that dSphs in the Local Group do not have a significant number of bright X-ray sources (see Zang \\& Meurs 2001).  The number of binary X-ray sources in a galaxy is determined by the formation rate of the sources and the rate at which the sources cease to be X-ray active.  Also, it depends on the duration of the X-ray duty cycles of the sources.  Observations show that X-ray sources are more abundant in starburst galaxies and interacting galaxies than in the `milder' galaxies (see e.g.\\ Fabbiano 2003), and the brightness distribution of X-ray sources vary with their host environments even within the same galaxy.  For instance in spiral galaxies, e.g. M81 (Tennant et al.\\ 2001; Swartz et al.\\ 2003), sources in the galactic disks have power-law luminosity functions, while sources in galaxy bulges tend to have broken power-law like or exponential cut-off type luminosity functions.  Now there is strong evidence that X-ray source populations in galaxies and their brightness distribution are dependent on the star formation activity of the source host environment in recent epochs.  This phenomenon can be explained by a birth-death process as that described in Wu (2001).  The lack of very bright sources in old stellar populations can be explained by the fact that X-ray sources have finite life-spans and bright sources which require rapid mass transfer are short-lived.  If the mass transfer rate in a binary is sufficiently low or if the mass-transfer active phase in its duty cycle is sufficiently short, the system may survive a very long time as an X-ray source.  The source may exhibit alternating active and faint states, and so it may not be easily identified in a single snap-shot observation.  In the birth-death model presented in Wu (2001), there is also an old X-ray source population in any galaxy whose progenitors were the first-generation stars formed in the very early epoch.  These `primordial' binary X-ray sources are dim, and their mass-donor stars are low-mass stars.  As these binaries have a sustainable mass-transfer phase spanning over a period of $> 10$~Gyr, the average mass-transfer rate between cannot exceed $10^{15}$~g~s$^{-1}$.  If the accreting compact object is a black hole or a neutron star, the X-ray luminosity of the binary would be $< 10^{35}$~\\ergss.  Even when a detection limit of well below $10^{35}$~\\ergss has been achieved, the old binary X-ray sources are not easily identified in big spiral galaxies like M31 or M81 (see e.g.\\ Kong et al.\\ 2003; Swartz et al.\\ 2003), because there are hundreds of younger X-ray sources with luminosities above $10^{36}$~\\ergss.  DSphs have not had recent star formation activity and are therefore the most appropriate targets to search for the old X-ray source population.  {\\sl ROSAT} had the sensitivity to detect sources brighter than $\\sim 10^{34}$~\\ergss \\ in the nearby dSphs but were unable to detect sources with luminosities of $10^{32}$~\\ergss.  With {\\it Chandra} and {\\sl XMM-Newton}, we can now search for sources with luminosities below $10^{32}$~\\ergss\\ in the dwarf satellites of the Milky Way.  We have made X-ray observations of the Sagittarius dSph (Sgr) with {\\sl Chandra} in 2003 and Carina dSph (Car) with {\\sl XMM-Newton} in 2004. This paper presents the results of our X-ray observations of these two galaxies.  (The {\\sl Chandra} observation of Sgr was centered near the globular cluster M54, and analysis of the cluster sources has already been presented in Ramsay \\& Wu 2006).  Sgr, at a distance of 27.4~kpc (Layden \\& Sarajedini 2000), is the second closest satellite to the Milky Way.  It extends over at least $22^{\\circ} \\times 8^{\\circ}$ in the sky, and is located on the far side of the Milky Way to the Sun (Ibata, Gilmore \\& Irwin 1994, 1995). The absorption column density along the line-of-sight is relatively high ($1.2 \\times 10^{21}$~\\pcmsq) because of the low galactic latitude of Sgr ($-14^{\\circ}$), Star formation in Sgr mostly ceased $\\sim$10 Gyr ago (Mateo 1998).  Car was discovered in plates taken using the UK Schmidt Telescope (Cannon, Hawarden \\& Tritton 1977) and is 100 kpc distant (van den Bergh 2000).  The line-of-sight absorption column to Car is $3.9 \\times10^{20}$~\\pcmsq.  Car has had a more complex star formation history than Sgr: around half of its stars were formed 7 Gyr ago, the remainder occuring at 15 and 3 Gyr ago (Hurley-Keller, Mateo \\& Nemec 1998). ",
        "conclusions": "We have performed X-ray observations of the Sgr and Car fields.  Our observations reached a sensitivity limit (unabsorbed) $\\sim 1\\times10^{32}$ \\ergss\\ for the Sgr field and $3\\times10^{34}$ \\ergss for the Car field.  We have identified a total of 80 sources (in the $0.3-8$~keV energy band) in the Sgr field (excluding the 7 sources which we have identified as being in the globular cluster M54) and 53 sources (in the $0.2-10$~keV band) in the Car field. At energies $>$2keV we found 42 sources in the Sgr field and 40 in the Car field (these numbers do not include objects which are likely to be foreground stars).  Based on the CDF-S we estimate that $55-79$ background AGN would be in the Sgr field and $16-34$ in the Car field. (The expected numbers will vary to some extent since the number of AGN varies from field to field.)  The source numbers in the two field are broadly consistent with the expected number of background AGN. However, from the analysis of X-ray colours, we have found a small number of sources which have colours which are difficult to be explained by spectral models appropriates for AGN. We therefore propose that these are candidate stellar X-ray sources. The nature of these sources may be determined by follow-up X-ray observations or by identifying their optical/IR counterparts. Both Sgr and Car have low galactic latitudes ($b=-14^{\\circ}$ and $-22^{\\circ}$ respectively), which gives rise to a number of possible optical counterparts for many of the X-ray sources. The search for the optical counterparts to the X-ray sources will be discuss in detail in a future paper.  We note that Maccarone et al.\\ (2005) found a small number of X-ray sources which are believed to belong to the Sculptor (Scu) dSph. These sources have X-ray luminosities $L_{X} = 6 - 90\\times10^{33}$\\ergss.  The optical counterparts of these sources are relatively easy to identify as Scu has a high galactic latitude ($b=-83^{\\circ}$).  The interstellar absorption to the edge of the Galaxy towards the direction of Car and Scu are similar (3.9$\\times10^{20}$ \\pcmsq and 2.0$\\times10^{20}$ \\pcmsq respectively).  If we take a naive view that the number of X-ray sources are proportional to total stellar mass of the host galaxy, then we may obtain an estimate of the number of sources to be detected in Car from the results of Scu (we acknowledge that the reality is likely to be more complex since they have different star formation histories).  The Scu dSph has a core radius of 5.8$^{'}$ and the field of view of the {\\chan} observations of Scu covered an area of $\\sim17^{'}\\times17^{'}$.  The observations therefore covered a large fraction of the galaxies mass (6.4$\\times10^{6}$\\Msun).  The field of view of the {\\xmm} observations of Car were much larger than the core radius of 8.8$^{'}$ and covered a large fraction of the galaxies mass (1.3$\\times10^{7}$\\Msun).  (These parameters are taken from Mateo 1998). The {\\xmm} observations of Car were sensitive to X-ray luminosities $>3\\times10^{34}$ \\ergss: there were 2 Scu sources with luminosities greater than this.  If we assume that the number of X-ray sources scales with the total stellar mass of the galaxy there will be roughly 2 sources with $L_{X}>3\\times10^{34}$ \\ergss\\ in Car.  This suggests that some of the 3 soft X-ray sources in the Car field could actually be located in the Car dSph.  We observed only a small fraction of the sky coverage of the Sgr dSph -- roughly 0.05\\%.  The total mass of the Sgr dSph is $\\sim10^{9}$ \\Msun (Ibata Gilmore \\& Irwin 1995), but we sampled a fraction of the galaxy with a total mass of $\\sim5\\times10^{5}$ \\Msun.  If we scale the number of X-ray sources with mass, 1 X-ray source in Sgr would imply a total number $4 \\times10^{5}$ of X-ray sources with $L_{X} > 10^{32}$ \\ergss\\ in the Milky Way. This number is obviously rather uncertain for various reasons, and what fraction of these sources can be detected in snap-shot is unclear given that a substantial fraction of the sources are transients or variable sources.  Birth-death models (e.g.\\ Wu 2001) predict a primordial X-ray source population in all galaxies.  These sources are remnants of the first generation stars formed with the galaxy. Suppose that the slope of the luminosity function of these source is similar to the faint end of the luminosity function of the bright bulge sources in M81 (see Tennant et al.\\ 2001), we would set an upper limit that fewer than 100 primordial X-ray sources with luminosity exceeding $\\sim 10^{36}$\\ergss\\ are in large spiral galaxies similar to M81 or the Milky Way.  However, the luminosity function of the primordial X-ray sources would show an exponential cut-off (a consequence of aging effects), beyond which the source counts drop rapidly.  The number of primordial sources with such luminosities should in fact be much fewer since fainter sources can survive for a long time.  Our observation of Sgr was sensitive down to $L_{X}\\sim 1 \\times 10^{32}$\\ergss, which is sufficient to detect most of the persistent bright X-ray binaries and moderately bright CVs.  If the soft sources which we found belong to the host dSph, then they may be candidate primordial X-ray sources."
    },
    "0609/astro-ph0609102_arXiv.txt": {
        "abstract": "{The high energies of protostellar jets, implied by recent observations of X-rays from such flows, came very much as a surprise. Inferred shock velocities are considerably higher than what was previously known, hence putting even larger energy demands on the driving sources of the jets. The statistics of X-ray emitting jets are still poor, yet a few cases exist which seem to imply a correlation between the presence of \\heilong\\ emission and X-ray radiation in a given source.} {This tentative correlation needs confirmation and explanation. If the jet regions of \\heilong\\ emission are closely associated with those producing X-rays, high resolution infared spectroscopy can be used to observationally study the velocity fields in the hot plasma regions of the jets. This would provide the necessary evidence to test and further develop theoretical models of intermediately fast ($> 500 - 1500$\\,\\kms) interstellar shock waves.} {The HH\\,154 jet flow from the embedded protostellar binary L\\,1551 IRS\\,5 provides a case study, since adequate IR and X-ray spectroscopic data are in existence. The thermal X-ray spectrum is fed into a photoionization code to compute, in particular, the line emission of He\\,I and H\\,I and to account for the observed unusual line intensity ratios.} {The advanced model is capable of accounting for most observables, but shows also major weaknesses. It seems not unlikely that these could, in principle, be overcome by a time dependent hydrodynamical calculation with self-consistent cooling. However, such sophisticated model development is decisively beyond the scope of the present work.} {Continued X-ray observations, coordinated with simultaneous high resolution infrared spectroscopy, are highly desirable.} ",
        "introduction": "The radiation from HH-objects and collimated jet flows is a manifestation of interstellar shock waves, the typical velocities of which were believed to be $V_{\\rm s} \\ll 500$\\,\\kms\\ \\citep[e.g.,][]{hartigan1987,raga1996,reipurth2001}. Unexpectedly, a number of such objects from embedded protostars have recently been detected at X-ray energies \\citep[with~reference~to~HH\\,1/2,~HH\\,154,~Ser\\,SMM1,~HH\\,80/81,~TKH\\,8,~HH\\,158, HH\\,210,~respectively]{pravdo2001,favata2002,preibisch2003,pravdo2004,tsujimoto2004,gudel2005,grosso2006}.  The X-ray spectra have typically a thermal spectral distribution, characteristic of temperatures of several million degrees. This was very surprising, since implied shock velocities ($\\gg 300$\\,\\kms) would be very much higher than what had previously been inferred from spectroscopic observations at longer wavelengths. This indicated that highly significant jet material had escaped detection. As a consequence, there is a potential risk that energy requirements on the jet source, being proportional to the square of the jet velocity, had been seriously underestimated. Proposed energy extraction mechanisms from the protostar-disk system might need to be revised in the light of these new findings.  As in active stars, the {\\it continuous} X-ray emission may be accompanied by \\heilong\\ {\\it line} radiation \\citep[e.g.,][]{zirin1982}. If that is the case, this line would be a valuable tool to probe the velocity fields of the hottest plasma regions in the jet. Also, in the infrared, extinction problems would be considerably reduced compared to shorter wavelengths and, for consistency checks, recombination lines from hydrogen, i.e. \\pbetalam\\ and \\pgammalam, could also be observed simultaneously in the photometric J-band.  In Table\\,\\ref{XHeI}, observations of \\heilong\\ and X-rays toward various flows have been compiled. Evidently, many objects lack the \\heilong\\ line, but reveal a rich \\molh\\ spectrum, which is consistent with non-dissociative shocks of low velocity, of the order of 10\\,\\kms. This is true also for HH\\,1G and the Ser\\,SMM1 jet, toward which both \\heilong\\ and X-ray observations have been performed, but neither of which was detected. HH\\,2H was detected in X-rays but has apparently not been observed spectroscopically in the 1\\,\\um\\ region. On the other hand, the jet flows HH\\,158 (DG\\,Tau A) and HH\\,154 (L\\,1551 IRS\\,5) were clearly detected in both \\heilong\\ and at X-ray wavelengths, lending support to the above proposition.  It is thus the primary objective of this {\\it Research Note} to stimulate further X-ray observations for the investigation of the high energy regions in protostellar jets, alongside with coordinated high resolution near infrared spectroscopy to determine Doppler velocities and line shapes. These data can then be compared with predictions from theoretical shock models.  \\begin{table} \\begin{flushleft} \\caption{\\label{XHeI}Protostellar jets with X-ray data and/or NIR spectra} \\begin{tabular}{l ccc l} \\hline \\noalign{\\smallskip} Jet flow& \\heilong                & X-rays              & Reference & Comment \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip}  HH 158  & \\phantom{1}$\\times$     & \\phantom{1}$\\times$ & 1, 2      & [Fe\\,II]\\\\ HH 154  & \\phantom{1}$\\times$     & \\phantom{1}$\\times$ & 3, 4, 5   & [Fe\\,II] \\\\ HH 210  &                         & \\phantom{1}$\\times$ & 6         &   \\\\ HH 1G   & $-\\times$               & $-\\times$           & 7, 8      & \\molh\\ + [Fe\\,II] \\\\ HH 2H   &                         & \\phantom{1}$\\times$ & 8         &          \\\\ TKH 8   &                         & \\phantom{1}$\\times$ & 9         &          \\\\ HH 43   & $-\\times$               &                     & 10        & \\molh    \\\\ HH 24A  & $-\\times$               &                     & 11         & \\molh\\ + [Fe\\,II] \\\\ HH 25   & $-\\times$               &                     & 11        & \\molh\\   \\\\ HH 26   & $-\\times$               &                     & 11        & \\molh    \\\\ HH 72   & $-\\times$               &                     & 11        & \\molh    \\\\ Vel IRS17 & $-\\times$             &                     & 12        & \\molh    \\\\ HH 320  & $-\\times$               &                     & 11        & \\molh    \\\\ HH 321  & $-\\times$               &                     & 11        & \\molh    \\\\ HH 80   &                         & \\phantom{1}$\\times$ & 13        &          \\\\ HH 81   &                         & \\phantom{1}$\\times$ & 13        &          \\\\ Ser SMM1& $-\\times$               &  $-\\times$          & 14, 15    & \\molh    \\\\ HH 99-B0& \\phantom{1}$\\times$     &                     & 16        & \\molh\\ + [Fe\\,II]  \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{flushleft} Notes to the table: $\\times =$ detected, $-\\times$ means observed but not detected. \\\\ \\molh\\ indicates the presence of ro-vibrationally excited \\molh\\ lines; [Fe\\,II] the presence of [Fe\\,II] lines; \\molh\\ + [Fe\\,II] the presence of both species along the same line of sight/inside spectrograph slit. When possible, HH-nomenclature follows that of \\citet{reipurth1999}. \\\\ References: (1) \\citet{takami2002}, (2) \\citet{gudel2005}, (3) \\citet{itoh2000}, (4) \\citet{liseau2005}, (5) \\citet{favata2002}, (6) \\citet{grosso2006}, (7) \\citet{nisini2005}, (8) \\citet{pravdo2001}, (9) \\citet{tsujimoto2004}, (10) \\citet{giannini2002}, (11) \\citet{giannini2004}, (12) \\citet{giannini2005}, (13) \\citet{pravdo2004}, (14) Liseau (unpublished), (15) \\citet{preibisch2003}, (16) \\citet{mccoey2004}. \\end{table}  \\begin{figure}[t] \\resizebox{\\hsize}{!}{ \\rotatebox{270}{\\includegraphics{Jspectrum.ps}} } \\caption{The J-band spectrum of \\citet{itoh2000} containing the Paschen lines, with the estimated level of the rms noise indicated. The observations are displayed as the black histogram, whereas the smooth red line refers to the fit of the [Fe\\,II] spectrum \\citep[for~$A_{\\rm V} = 5.7$~mag,~][]{liseau2005}. The wavelength positions of the recombination lines \\heilong, \\pgammalam\\ and \\pbetalam\\ are indicated above and below the spectrum. This applies also to the wavelengths of the transitions to the inverted ground state of [S\\,I], i.e. $^1{\\rm D}_2 - ^3\\!{\\rm P}_2$\\,1.08212\\,\\um, $^1{\\rm D}_2 - ^3\\!{\\rm P}_1$\\,1.13058\\,\\um\\ and $^1{\\rm D}_2 - ^3{\\rm P}_0$ at 1.15376\\,\\um, respectively. } \\label{IR_lines} \\end{figure}  To date, the best available data are those of the HH 154 flow and we shall first examine the accumulated evidence for this object (Sect.\\,2). Some numerical model calculations will be presented and discussed in Sect.\\,3. In that section, we will also conclude with some suggestions for future work. ",
        "conclusions": "The overall observational evidence supports the view that the X-ray emission originates from fast shocks in the HH\\,154 jet, with velocities in excess of 500\\,\\kms\\ \\citep{favata2002,bonito2004,liseau2005,favata2006}. A reasonable scenario features a hypersonic jet, ramming into molecular cloud material which leads to the complete destruction of molecules and grains and the ionization of the gas, which is heated to high temperatures. This hot plasma is radiatively cooled by X-rays, which are capable of ionizing the neutral gas ahead of the shock, altering its state of excitation and potentially also its chemistry. In order to arrive at quantitative estimates of the associated helium and hydrogen recombination line emission, in particular, we have run photoionization models for this scenario.  The \\heilong\\ line originates from levels which are more than 20\\,eV above the ground (Fig.\\,\\ref{HeI_levels}). Evidently, the observed X-rays are sufficientally energetic to photoionize helium (see Fig.\\,\\ref{xrays}) resulting in He\\,1.0830 emission upon recombination. But is it possible to explain the fact that the observed line intensity ratio \\heilong/\\pgammalam\\,$>> 1$? Such inverted ratios are generally not revealed by models of jet shocks \\citep[e.g.][]{hartigan1987}, but have been observed to develop in supernova ejecta \\citep[e.g.][]{meikle1993,pozzo2004}.  For the model calculations we used the program Cloudy \\citep[version~96~last~described~by][]{ferland1998}, which determines the temperature structure for a gas of a certain geometry and density, including the transfer of radiation. For the calculations described here, the shape of the X-ray input spectrum is that of a blackbody (Fig.\\,\\ref{xrays}), scaled by the observed X-ray luminosity, i.e. the inner radius is $r_{\\rm in}=\\sqrt{L_{\\rm X}/4 \\pi \\sigma T^4}$. The outer radius of the spherical cloud was initially set by the observed X-ray emission measure and the assumed constant density, i.e. $r_{\\rm out} \\sim (3\\,EM/4 \\pi)^{1/3} n(\\rm H)^{-2/3}$. Solar abundances were assumed for the chemical composition of the gas, hence $n({\\rm He})/n({\\rm H})=0.1$.  The resulting model for $n({\\rm H})= 10^5$\\,\\cmthree, which is in fair agreement with the observations, is depicted in Fig.\\,\\ref{structure}. In particular, the ratios \\heilong/\\pgammalam\\ and \\heilong/\\pbetalam\\ are 13 and 7, respectively, i.e. both much larger than unity as has been observed (Table\\,\\ref{fluxes}). According to the model, with ${\\rm e}^-$-collision coefficients from \\citet{bray2000}, the excitation of the \\heilong\\ line is dominated by collisions, so that in this case, the `10830 recombination line' is not really due to He recombination (compare the $x_{\\rm He\\,II}$ and $x_{\\rm He\\,I}$ regions in the figure). For purely thermal broadening, the line would be marginally optically thick, i.e. $\\tau=1.5$ at the center of the line. The flux of the reddened 2.06\\,\\um\\ line (Fig.\\,\\ref{HeI_levels}) is predicted not to exceed the 10\\% level of that in the 1.0830\\,\\um\\ line.  A hint to the nature of the He\\,I excitation is potentially also provided by the \\heishort\\ line, originating from the state just above that of the \\heilong\\ line (see Fig.\\,\\ref{HeI_levels}). The line intensity ratio \\heishort/\\heilong\\ of the model is 0.3, i.e. the predicted (reddened) \\heishort\\ flux is $4 \\times 10^{-16}$\\,\\ecs. This is in agreement with the observation by \\citet{cohen1985}, which resulted in an upper limit, i.e. $F < 6.5 \\times 10^{-16}$\\,\\ecs\\ (the units in that paper seem erroneous). Given the information provided by \\citet{cohen1985}, it seems not likely, however, that this measurement actually refers to the location of the X-ray source. In addition, the X-ray emission is known to be variable on short time scales \\citep{favata2006}.  The observed \\heilong\\ luminosity (Sect.\\,2.1.4) is a sizable fraction of the X-ray luminosity and it comes hardly as a surprise that the model falls short with regard to $L_{1.0830\\,\\mu{\\rm m}}$. This may not be attributable solely to source variability. In principle, an ad hoc increase of the emitting volume could artificially cure this luminosity problem: for instance, shock models would lead to typically much larger cooling lengths (\\powten{15} $-$ \\powten{16}\\,cm) than what is indicated in Fig.\\,\\ref{structure} ($< 10^{14}$\\,cm). The required increase by an order of magnitude seems consistent with the models of X-ray emitting bow shocks of \\citet{raga2002}. Also, as suggested by the referee, a jet with a time variability in the ejection velocity could increase the emitting volume (hence the total luminosity). This volume increase is produced when two successive working surfaces collide with each other.  However, such an ad hoc procedure would lack any physical justification. Therefore, rather than trying `to fix' this poor model, one would like to see a proper time-dependent hydrodynamical calculation for the high velocity jets, accounting for the shocks and the cooling in a self-consistent manner. This is also warranted by the various time scales of the steady state photoionization model (e.g., for thermal equilibrium and hydrogen recombination), which are all very much longer than the dynamical time scale of less than one year.  On the observational side, one would like to see concerted efforts of X-ray observations and simultaneous intermediate to high resolution ($R \\ge {\\rm a~few} \\times 10^3$) infrared spectroscopy.  \\begin{figure}[t] \\resizebox{\\hsize}{!}{ \\rotatebox{00}{\\includegraphics{ion_struc.ps}} } \\caption{Helium ionization structure of gas with constant density $n({\\rm H})=10^5$\\,\\cmthree, being illuminated by a 4\\,MK point source. The run of the temperature $T_{\\rm e}$ (in units of \\powten{6}\\,K) is shown in red, that of the electron fraction $x_{\\rm e}$ in green. The fractional H\\,I density is shown by the black broken line, whereas those of He\\,I, He\\,II and He\\,III in blue, cyan and magenta, respectively. Tests with the previous version Cloudy\\,90 gave essentially identical results, except that the temperature structure is smoother. The hump due to `ionization jumps' for a number of species is absent, indicating that this feature does not have any major effect on the results of interest. } \\label{structure} \\end{figure}"
    },
    "0609/astro-ph0609334_arXiv.txt": {
        "abstract": "Results are presented from a detailed analysis of optical and X-ray observations of moderate-redshift galaxy clusters from the Canadian Network for Observational Cosmology (CNOC) subsample of the EMSS.  The combination of extensive optical and deep X-ray observations of these clusters make them ideal candidates for multiwavelength mass comparison studies.  X-ray surface brightness profiles of 14 clusters with $0.17<z<0.55$ are constructed from Chandra observations and fit to single and double $\\beta$ models.  Spatially resolved temperature analysis is performed, indicating that five of the clusters in this sample exhibit temperature gradients within their inner $60-200$ kpc.  Integrated spectra extracted within $\\rm{R}_{2500}$ provide temperature, abundance, and luminosity information.  Under assumptions of hydrostatic equilibrium and spherical symmetry, we derive gas and total masses within $\\rm{R}_{2500}$ and $\\rm{R}_{200}$.  We find an average gas mass fraction of $f_{gas}(\\rm{R}_{200})=0.136\\pm{0.004}~h_{70}^{-3/2}$, resulting in $\\Omega_m=0.28\\pm{0.01}$ (formal error).  We also derive dynamical masses for these clusters to $\\rm{R}_{200}$.  We find no systematic bias between X-ray and dynamical methods across the sample, with an average $\\rm{M}_{dyn}/\\rm{M}_{X-ray} = 0.97\\pm{0.05}$.  We also compare X-ray masses to weak lensing mass estimates of a subset of our sample, resulting in a weighted average of $\\rm{M}_{lens}/\\rm{M}_{X-ray}$ of $0.99\\pm{0.07}$.  We investigate X-ray scaling relationships and find powerlaw slopes which are slightly steeper than the predictions of self-similar models, with an $E(z)^{-1}$ $L_X$-$T_X$ slope of $2.4\\pm{0.2}$ and an $E(z)~\\rm{M}_{2500}$-$T_X$ slope of $1.7\\pm{0.1}$.  Relationships between red-sequence optical richness ($\\rm{B}_{gc,red}$) and global cluster X-ray properties ($T_X$, $L_X$ and $\\rm{M}_{2500}$) are also examined and fitted. ",
        "introduction": "}   Clusters of galaxies are valuable cosmological probes.  As the largest gravitationally bound objects in the universe, they play a key role in the tracing and modeling of large scale structure formation and evolution \\citep[e.g.,][]{voit,bahcall}.  Because of our rapidly expanding sample of known clusters, finding efficient means of estimating cluster properties is a highly desirable goal.  The contribution of cluster studies to the field of observational cosmology hinges on our ability to accurately estimate cluster masses.  In particular, through the determination of both gas mass and total mass, cluster analysis can lead to estimations of the cosmological mass density, $\\Omega_m$, while accurate measurement of the evolution of the cluster mass function provides important constraints on both the normalization of the matter power spectrum, $\\sigma_8$ and $w$, the dark energy equation of state~\\citep[e.g.,][]{levine02,eke}.  In addition, gaining an understanding of the evolution of X-ray scaling relationships (such as $L_X$-$T_X$ and M-$T_X$) with redshift provides an important contribution to our ability to accurately model the evolution of large-scale structure in the universe~\\citep[e.g.,][]{voit}.  Intermediate redshift ($0.2 < z < 0.6$) clusters are well situated for the study of cluster properties.  They are compact enough to be observed without tedious mosaicing, and they are present in statistically significant numbers.  Intermediate redshift clusters are also luminous enough to permit detailed investigation, with the potential for placing strong contraints on cluster evolution.  Three frequently applied approaches to estimating cluster masses are gravitational lens modeling, optical spectroscopic measurements of the cluster galaxy velocity dispersion, and characterizing X-ray emission as a means of tracing the underlying potential well of the cluster.  Each of these methods uses different observations and assumptions, which can be tested through their direct comparison.  An alternate optical approach to the efficient estimation of cluster masses involves the use of optical richness.  This method relies on the assumption that galaxy light is a reliable tracer of total cluster mass, and requires calibration via other methods.  Lensing mass estimates test both the assumption of hydrostatic equilibrium and our knowledge of the mass distribution in clusters.  Because they probe all of the projected mass along the line of sight, which may include additional mass concentrations, they are susceptible to overestimation of cluster mass~\\citep{cen, metzler01}.  Dynamical mass estimates work under the assumption that velocity dispersion is directly related to the underlying gravitational potential of the cluster.  Pitfalls of this technique also include the danger of overestimating the cluster mass in cases where substructure or mergers drive up velocity dispersion measurements~\\citep{bird}.  The hot, diffuse intracluster medium (ICM), which is observed in the X-ray, should be a direct tracer of a cluster's underlying potential well.  Under assumptions of isothermality and hydrostatic equilibrium, the surface brightness of a cluster of galaxies can provide information on gas density as well as total gravitating mass.  Factors that influence the accuracy of X-ray mass estimates are temperature gradients, substructure, and mergers, which can compromise the previously stated assumptions~\\citep[e.g.,][]{balland}.  One of the main objectives of this work is to identify, through direct comparison, any systematic biases in these methods with the ultimate goal of determining a robust calibration between optical richness and cluster mass.  Optical richness measurements are easily available due to the fact that their estimation requires very little observing time. Therefore an accurate calibration of the relationship between mass and optical richness would allow us to determine the masses of large samples of clusters in a highly efficient manner, providing strong constraints on the evolution of the cluster mass function and, consequently, key cosmological parameters.  In this paper we present a detailed analysis of high resolution Chandra X-ray observations of 14 CNOC~\\citep{yee96} clusters at z $\\sim 0.3$.  In Sections ~\\ref{s:obs}-\\ref{s:spect} we probe the temperature, metallicity, morphology, and surface brightness of the hot ICM present in each cluster.  From this initial analysis, we derive mass estimates (Section~\\ref{s:mass}) which are compared to dynamical and weak lensing mass estimates (Section~\\ref{s:masscomp}), and then we examine the X-ray scaling laws of our sample (Section~\\ref{s:laws}).  Finally, in Section~\\ref{s:or}, we use our results to calibrate relationships between red-sequence optical richness ($\\rm{B}_{gc,red}$) and global cluster X-ray properties ($T_X$, $L_X$ and $\\rm{M}_{2500}$).  Unless otherwise noted, this paper assumes a cosmology of $\\rm{H}_0=70~\\rm{km}~\\rm{s}^{-1}~\\rm{Mpc}^{-1}$, $\\Omega_m=0.3$, and $\\Omega_{\\Lambda}=0.7$ ",
        "conclusions": "}  We have presented a comprehensive analysis of Chandra observations of 14 medium redshift ($0.1709<z<0.5466$) clusters of galaxies from the CNOC subsample of the EMSS (Table~\\ref{table1}). Imaging analysis has provided information on the relative quiescence of each cluster (Figure~\\ref{fig1}).  The spatial resolution of Chandra has allowed us to determine the surface brightness profiles of these clusters down to $1\\arcsec$ scales (Figure~\\ref{fig2}).  This has enabled us to obtain precise models of cluster emission (Table~\\ref{table2}), except in the case of MS0451+02 which displays a considerable amount of substructure.  Five clusters in our sample (Abell 2390, MS0440+02, MS0839+29, MS1358+62, and MS1455+22) exhibit excess emission in their cores, indicating the presence of cool gas.  The surface brightnesses of these five clusters are better fit with the inclusion of a second $\\beta$ component to model the excess core emission (Table~\\ref{table3}).  Chandra's $0.5\\arcsec$ spatial resolution has enabled us to study the radial distribution of temperatures on small scales ($\\sim7-100$ kpc) in the inner 150-600 kpc of eleven of these objects (Figure~\\ref{fig3}).  While nine clusters in our sample have temperature profiles which are consistent with isothermality, the five cooling core clusters (listed above) show clear temperature gradients in their innermost $\\sim60-200$ kpc.  The temperature profiles of these five clusters are consistent with the ``universal'' temperature profile of \\citet{allen01}.  The energy resolution of the instrument has provided well constrained spectral analyses of the central $\\rm{R}_{2500}$ ($\\sim350-760$) kpc of each object.  Temperatures obtained initially from 300 kpc radius extraction regions began an iterative spectral extraction and fitting process culminating in the determination of robust integrated temperatures, luminosities, and abundances within $\\rm{R}_{2500}$ (Tables~\\ref{table4} and~\\ref{table5}).  Cooling core corrected spectra were used to determine global temperatures for the clusters in our sample which exhibit significant cooling.  These temperatures were then employed in the spectral determination of integrated luminosities for these objects.  Overall, our sample displays temperatures in the range of $3.4-12.1$ keV, abundances which range from $0.1-0.8$ times the solar value, and unabsorbed bolometric luminosities within $\\rm{R}_{2500}$ which span $4.9-56.1 \\times 10^{44}~\\rm{erg}~\\rm{s}^{-1}$.  Cluster gas and total masses within $\\rm{R}_{2500}$ and $\\rm{R}_{200}$ were estimated using the outcomes of spectral and surface brightness fitting, resulting in virialized cluster masses of $2.7-21 \\times 10^{14}~h_{70}^{-1}\\msun$ (Table~\\ref{table6}) and respective cluster gas masses of $0.13-2.9 \\times 10^{14}~h_{70}^{-1}\\msun$.  A weighted average of gas mass fractions gives $f_{gas}=0.136\\pm{0.004}~h_{70}^{-3/2}$, resulting in $\\Omega_m=0.28\\pm{0.01}$ (68\\% confidence, formal error), this value is in good agreement with~\\citet{spergel2}.  Dynamical masses within $\\rm{R}_{200}$ were calculated via the Jeans Equation using velocity dispersions taken from~\\citet{borgani} (Table~\\ref{table7}).  Comparisons between X-ray masses and dynamical masses result in a weighted average of $\\rm{M}_{\\rm{dyn}}/\\rm{M}_{\\rm{X-ray}}=0.97\\pm{0.05}$, indicating good agreement between between these two methods, however a fair amount of scatter is evident ($\\sim 30\\%$; Figure~\\ref{fig5}).  Dynamical masses are noticeably larger in the case of clusters which exhibit significant substructure (MS1006+12 and MS1008-12), a factor which may be responsible for an overestimation of velocity dispersions~\\citep{bird}.  Weak lensing masses within $500~h_{100}^{-1}$ kpc were obtained for seven of the clusters in our sample.  X-ray masses were calculated within this region and compared to the lensing results (Table~\\ref{table8}).  Although the distribution appears somehwat asymmetric (Figure~\\ref{fig6}), a weighted average gives $\\rm{M}_{\\rm{lens}}/\\rm{M}_{\\rm{X-ray}}=0.99\\pm{0.07}$, with a scatter of $\\sim 30\\%$.  X-ray scaling laws for this sample were investigated in a manner similar to~\\citet{allen01} and~\\citet{ettori04}, taking cosmological factors into account.  The best fitting $L_X-T_X$ relationship for our sample results in an intercept of $0.13^{+0.08}_{-0.05}$ and a slope of $2.4\\pm{0.2}$ (Figure~\\ref{fig7}).  While this slope is lower than some estimates~\\citep{arnaud99,ettori04}, it is consistent with both~\\citet{allen01} and~\\citet{ettori02}.  Moderate negative evolution is indicated by this index being greater than that predicted by self-similar models~\\citep{voit}.  The Mass-$T_X$ relationship for these clusters also exhibits a somewhat steeper slope than expected, at $1.7\\pm{0.1}$, with an intercept of $1.3^{+0.3}_{-0.2} \\times 10^{13}$ (Figure~\\ref{fig8}).  The best fitting scaling laws for our sample (listed above) were combined with Equation~\\ref{eq:txbgc}, with the ultimate goal of calibrating relationships between red-sequence optical richness ($\\rm{B}_{gc,red}$) and global cluster parameters ($T_X$, $L_X$ and $\\rm{M}_{2500}$).  We find that $T_X$ scales relatively well to $\\rm{B}_{gc,red}$ with a $40\\%$ scatter (Figure~\\ref{fig9}).  $\\rm{M}_{2500}$ shows a scatter of $\\sim70\\%$ (Figure~\\ref{fig11}), which is consistent with the $T_X$-$\\rm{B}_{gc,red}$ scatter given that ${\\rm{M}}\\propto T^{3/2}$.  The $L_X$-$\\rm{B}_{gc,red}$ relationship exhibits a significantly larger scatter at $\\sim 200\\%$ (Figure~\\ref{fig10}).  Our results indicate that $\\rm{B}_{gc,red}$ does exhibit intial promise as a mass indicator.  Accurate calibration of a relationship between optical richness and cluster mass will require the use of additional clusters which posess both well-constrained X-ray temperatures and optical richness measurements.  Overall we find that multiple cluster mass estimators - dynamics, weak lensing and X-ray observations, along with optical richness in the cluster red sequence - are converging for this sample of well-studied clusters.  While individual correlations still have significant scatter, there is little evidence for large systematic bias in any of these methods.  Cluster characteristics which might be considered problematic for one or more techniques (cluster substructure, merging, and/or the presence of a cool core) appear to perturb these relationships relatively little, as long as high quality data are obtained and the analysis is tuned to correct for these.  In particular the correlation between cluster temperature and optical richness, the most easily obtained of the mass estimators, is promising.  A remaining concern is that X-ray selected clusters may not prove to be typical of all massive clusters at these redshifts.  Additional checks involving clusters covering a broader range of redshift and selection technique (e.g., SZ, optical, weak lensing) will be necessary to solidify our understanding of the most reliable and efficient methods of cluster mass estimation."
    },
    "0609/astro-ph0609272_arXiv.txt": {
        "abstract": "We present a short general overview of the main features of exotic models of neutron stars, focusing on the structural and dynamical predictions derived from them. In particular, we discuss the presence of ``normal'' quark matter and Color-Flavor Locked (CFL) states, including their possible self-bound versions, and mention some different proposals emerging from the study of QCD microphysics. A connection with actual observed data is the main goal to be addressed at this talk and along the meeting. It is demonstrated that exotic equations of state are {\\it not} soft if the vacuum contributions are large enough, and argued that recent measurements of high pulsar masses ($M \\geq \\, 2 \\, M_{\\odot}$) create problems for {\\it hadronic} models in which hyperons should be present. ",
        "introduction": "The quest of the internal composition of compact stars has been going on for decades, in close connection with the work in nuclear and particle physics. Although deemed in some sense ``simpler\" than magnetospheric phenomena, dealing with matter at the extreme conditions inside pulsars/neutron stars has never been easy. As 2006, the consensus about the nature of matter at several times the nuclear saturation density is weak, if anything. Several phases/components of the nuclear fluid have been proposed and studied, but decisive evidence for or against them is hard to obtain. In fact, and as discussed several times during this meeting, advances on the observational side has allowed one for the first time to probe key macroscopic properties of compact stars (masses, radii and a few others) that reflect the internal composition indirectly. However, a few remarks on this last statement are in order: on the one hand the precision attained by measurements has improved greatly but {\\it not} to the point far beyond any suspicion, and on the other hand, supranuclear components with poorly constrained parameters (coupling constants, vacuum expectation values, etc.) will not be ruled out, even by very precise observations. Therefore, work is needed on the theoretical and terrestrial laboratories as well.  Overall, I believe it is fair to state that the bulk structural properties are well-known in the subnuclear domain (more strictly, below the neutron drip density $\\sim \\, 10^{11} \\, g \\, cm^{-3}$). However, important questions involving magnetic fields still remain (van Adelsberg et al. 2005). This is quite important to settle since it is where the star surface is seen by experiments due to electromagnetic radiation (see discussions about spectral lines and related topics in this meeting). There is also important information in timing irregularities, most notably glitches, believed to originate at the inner crust (although the ``conventional wisdom'' has been recently challenged, see the contribution by B. Link in this meeting and references therein) However, deep below the stellar crust it is increasingly difficult to construct a clear picture of the composition and therefore of the stellar structure. Condensates ($\\pi^{-}$, $K^{-}$, etc) are still possible depending on microphysics and drastically alter the thermal and dynamical behavior of the star. Fundamental degrees of freedom (i.e. quarks) may also constitute the main part of the core, and even most of the star if they happen to be of the ``self-bound'' type (i.e. do not decay back to ordinary nuclear species once formed). Recently, a lot of attention has been paid to paired quark matter in a variety of phases still being studied. A recent summary of these matters can be found in Weber (2005). Hereafter we shall mainly concentrate on the self-bound phases because they bring the greatest modifications to the structure and they remain viable alternatives to the supranuclear matter composition. ",
        "conclusions": "We end this brief exposition about some features of self-bound models by saying that understanding of the {\\it vacuum} is the real clue for advances in dense matter physics. This is not unlike other fields of physics, like the well-known crisis in cosmology about what the quantum vacuum should be and what actually is (Freedman \\& Turner 2003). The same vacuum issue, but related to matter well above the saturation density, looks even more formidable, and its understanding should solve in the wash the issues of the existence of self-bound states and the features of the resulting EOS. It is also important to remark again that there is no hadronic model devoid of problems with the high mass end: either they ignore hyperons or are solved in a mean field approach or some other scheme with its own problems and questions. There is, however, the possibility of an extreme stiffness of the equation of state, such as the hyperons do not appear at all, because the relatively low density at the center. The good news is that we can now foresee actual tests of the stiffness of the EOS in the near future.  As stated, this stiffness is very important to establish, because the very introduction of $\\Lambda$ particles and other hyperons would then call for very exotic interactions among them at least. Otherwise the resulting EOS become so soft that measured masses around $2 M_{\\odot}$ happen to lie well above the maximum masses of the respective theoretical sequences. This is why strongly repulsive interactions would be required if hyperons appear inside compact stars. High masses may be {\\it pointing} towards the exotica rather than excluding them (\\\"Ozel 2006, Alford, Blaschke, Drago et al. 2006).  To conclude, we would like to quote an inspiring sentence from the English literature that may (or may not) be related to these topics, in which a bit of fantasy is always hidden \\medskip  {\\it Horatio, there are more things in Heaven and Earth}  {\\it Than are dreamt of in your philosophy}  \\hfill{Hamlet, Act I, Scene V}"
    },
    "0609/astro-ph0609791_arXiv.txt": {
        "abstract": "Understanding the acceleration of the universe and its cause is one of the key problems in physics and cosmology today, and is best studied using a variety of mutually complementary approaches. Daly and Djorgovski (2003, 2004) proposed a model independent approach to determine the expansion and acceleration history of the universe and a number of important physical parameters of the dark energy as functions of redshift  directly from the data.  Here, we apply the method to explicitly determine the first and second derivatives of the coordinate distance with respect to redshift, $y^{\\prime}$ and $y^{\\prime \\prime}$, and combine them to solve for the kinetic and potential energy density of the dark energy as functions of redshift, $K(z)$ and $V(z)$.  A data set of 228 supernova and 20 radio galaxy measurements with redshifts from zero to 1.79 is used for this study. Values of $y^{\\prime}$ and $y^{\\prime \\prime}$ are combined to study the dimensionless acceleration rate of the universe as a function of redshift, $q(z)$. The only assumptions underlying our determination of $q(z)$ are that the universe is described by a Robertson-Walker (RW) metric and is spatially flat. We find that the universe is accelerating today, and was decelerating in the recent past.  The transition from acceleration to deceleration occurs at a redshift of about $z_T = 0.42 \\pm {}^{0.08}_{0.06}$. Values of $y^{\\prime}$ and $y^{\\prime \\prime}$ are combined to determine $K(z)$ and $V(z)$.  These are shown to be consistent with the values expected in a standard Lambda Cold Dark Matter (LCDM) model. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609758_arXiv.txt": {
        "abstract": "We present the results of a spectral study of the soft X-ray emission (0.2$-$2.5 keV) from low-latitude (`disk') regions of Jupiter. The data were obtained during two observing campaigns with {\\it XMM-Newton} in April and November 2003. While the level of the emission remained approximately the same between April and the first half of the November observation, the second part of the latter shows an enhancement by about 40\\% in the 0.2$-$2.5 keV flux. A very similar, and apparently correlated increase, in time and scale, was observed in the solar X-ray and EUV flux.  The months of October and November 2003 saw a period of particularly intense solar activity, which appears reflected in the behaviour of the soft X-rays from Jupiter's disk. The X-ray spectra, from the {\\it XMM-Newton} EPIC CCD cameras, are all well fitted by a coronal model with temperatures in the range 0.4$-$0.5 keV, with additional line emission from Mg XI (1.35 keV) and Si XIII (1.86 keV): these are characteristic lines of solar X-ray spectra at maximum activity and during flares.  The {\\it XMM-Newton} observations lend further support to the theory that Jupiter's disk X-ray emission is controlled by the Sun, and may be produced in large part by scattering, elastic and fluorescent, of solar X-rays in the upper atmosphere of the planet. ",
        "introduction": "\\label{}  The current generation of X-ray observatories, with their greatly improved spatial resolution ({\\it Chandra}) and sensitivity ({\\it XMM-Newton}), coupled to moderate (CCD) to high (gratings) spectral resolution, have made it feasible for the first time to study solar system objects in detail. X-ray observations of planets, satellites and comets, coupled with solar X-ray studies and solar wind and magnetospheric measurements {\\it in situ}, are being used to extend our understanding of the photon and particle processes taking place all over the solar system, of the interactions between the Sun and solar system bodies, and ultimately of the effects solar activity may have on our own Earth (for a review see Bhardwaj et al., 2006b, this issue).  Jupiter has a particularly complex magnetospheric environment, which is governed by the fast rotation of the giant planet, and by the presence of Io and its dense plasma torus. This made it an interesting target since the earliest attempts of solar system X-ray studies: Jupiter was first detected at X-ray energies with the {\\it Einstein} observatory \\citep{Metz}; later studies with {\\it ROSAT} (Waite et al., 1994, 1997) established the presence of two distinct types of X-ray emission from the planet: `auroral', from regions near the magnetic poles, and `disk' emission, from lower latitudes. The present paper focuses on the disk emission; recent detailed studies of the Jovian X-ray aurorae can be found in Branduardi-Raymont et al. (2004, 2006a,b) and \\citet{els05}, and references therein.  The first attempt to explain the origin of Jupiter's disk emission \\citep{Waite97} was based on energetic ion precipitation, the same mechanism invoked to account for the auroral emissions also seen by {\\it ROSAT}. However, a general decrease in the overall X-ray brightness of Jupiter over the years 1994-1996 was found to be coincident with a similar decay in solar activity, as characterised by the solar 10.7 cm flux; the X-ray flux from the Jovian disk also revealed an interesting dependence on local time, with the X-ray bright limb coinciding with the bright visible limb \\citep{gla98}. These two facts combined suggested that the planet's disk emission may be controlled by the Sun to some extent. Indeed \\citet{mau00} were able to demonstrate that scattering of solar X-rays in Jupiter's atmosphere may well play a role in generating its equatorial emission. They modelled this component of solar origin with a combination of elastic scattering by atmospheric neutrals, and, to a smaller degree, fluorescent scattering of C K-shell X-rays on methane molecules.  {\\it Chandra} High Resolution Camera (HRC) observations of Jupiter in Dec. 2000 \\citep{gla02} not only revealed surprising characteristics in its auroral emission, but gave us the sharpest X-ray view of the planet yet, clearly separating the bright auroral and the essentially uniform disk emissions (see Figure 28 of Bhardwaj et al., 2006b, this issue). However, no information on the spectral characteristics of the emissions could be obtained from the HRC data. Jupiter was observed again in Feb. 2003 with the {\\it Chandra} Advanced CCD Imaging Spectrometer (ACIS), and the HRC, and this time spectra were acquired of both auroral \\citep{els05} and disk \\citep{bhar06a} emissions. Recently \\citet{cra06} have produced calculations of the scattering albedo for soft X-rays from the outer planets and have compared them against the ACIS spectra of Jupiter's disk. The conclusion is that indeed the soft X-ray emissions of Jupiter (and Saturn) can largely be explained by scattering and fluorescence of solar X-rays.  A study of Jupiter's soft X-ray emission from the first {\\it XMM-Newton} observation of the planet in Apr. 2003 \\citep{bra04} clearly demonstrated the different characteristics of the auroral and disk spectra: while charge exchange was found most likely to be responsible for the auroral soft X-ray production, a coronal (i.e. optically thin collisional plasma) model best fitted the disk emission, giving support to the solar X-ray scattering hypothesis. In the following we describe the results of a second {\\it XMM-Newton} observation, carried out in Nov. 2003 (a preliminary account of this can be found in Branduardi-Raymont et al., 2006a), and we compare them with those from a re-analysis of the Apr. 2003 data. ",
        "conclusions": "\\label{}  The results of the spectral analysis reported in the previous section reinforce the idea that Jupiter's low-latitude disk X-ray emission, away from the auroral zones, is predominantly the consequence of the Sun's X-ray irradiation and of scattering in the upper atmosphere of the planet \\citep{cra06}: the {\\it XMM-Newton} spectra are very well fitted by a coronal model with a temperature which is very similar to that observed at solar maximum for the Sun's X-rays. The Mg XI and Si XIII line emissions, also observed in Jupiter, are considered the signatures of flaring episodes on the Sun: the Jovian Mg XI flux was strongest in Nov. 2003, rev. 0727, when Jupiter's disk flux was highest and flares were observed in both solar and Jupiter's X-ray lightcurves (see Bhardwaj et al., 2005).  However, with the available data we cannot exclude the possibility that other processes may play a role in this scenario: recent results \\citep{bhar06a} have confirmed the existence of a correlation between low-latitude X-ray intensity and surface magnetic field strength which was originally reported in {\\it ROSAT} data \\citep{gla98}, and which seemed to contradict other evidence supporting a solar connection. This would suggest that some phenomenon of magnetospheric origin is also at work at low latitudes on Jupiter."
    },
    "0609/astro-ph0609514_arXiv.txt": {
        "abstract": "{} {To check whether the polar angle distribution of QSOs around nearby spiral galaxies is isotropic or not.} {A statistical analysis of the polar angle distribution of large samples of QSOs from the SDSS survey and Monte Carlo simulations to calculate their significance are carried out.} {There is a clear excess of QSOs near the minor axis with respect to the major axis of nearby edge-on spiral galaxies, significant at a level 3.5$\\sigma $ up to angular distances of $\\sim 3^\\circ$ (or $\\sim 1.7$ Mpc) from the centre of each galaxy. The significance is increased to 3.9$\\sigma $ with the $z>0.5$ QSOs, and it reaches 4.8$\\sigma $ if we include galaxies whose circles of radius 3 degrees are covered by the SDSS in more than 98\\% (instead of 100\\%) of the area.} {Gravitational lensing in the halo of nearby galaxies or extinction seem insufficient to explain the observed anisotropic distribution of QSOs. The anisotropic distribution agrees qualitatively with the predictions of Arp's models, which claim that QSOs are ejected by galaxies along the rotation axis, although Arp's prediction give a distance of the QSOs $\\sim 3$ times smaller than that found here. In any case, a chance fluctuation, although highly improbable, might be a possibility rather than a true anisotropy, and the present results should be corroborated by other groups and samples, so we prefer to consider it as just  a first tentative detection.} ",
        "introduction": "The first hints of a possible relationship between nearby galaxies and high redshift QSOs came from Arp (1966, 1967), who observed radio sources across active galaxies that were identified as QSOs. Since the pioneering work by Burbidge et al.\\ (1971), several groups have demonstrated the existence of an angular correlation between samples of QSOs and low redshift galaxies (e.g.\\ the reviews of Guimaraes 2005 and Burbidge 2001 respectively of orthodox and heterodox approaches). This correlation extends  to angular scales of $\\sim 1$ degree. Although weak gravitational lensing by dark matter has been proposed to be the cause of these correlations, many authors have found this hypothesis insufficient to explain the correlations (Ben\\'\\i tez et al.\\ 2001; Gazta\\~naga 2003; Nollenberg \\& Williams 2005). The recent analysis carried out by Scranton et al.\\ (2005), who used Sloan Digital Sky Survey (SDSS) photometric data containing $\\sim 2\\times 10^5$ quasars and $\\sim 1.3\\times 10^6$ galaxies, showed that the amplitude and sign of the angular correlation function between both sets of objects depend on the magnitude limit considered for the sample of QSOs. From this correlation, Scranton et al.\\ (2005) proposed an  ad hoc  halo distribution function compatible with a cross-correlation of very small amplitude ($\\omega_{\\rm GQ}<0.04$) of faint galaxies with QSO candidates selected photometrically (5\\% of this sample are not QSOs; Richards et al.\\ 2004). It is small because the mean separation among galaxies is small and any positive correlation of QSOs around a galaxy is diluted with the contamination of many other QSOs belonging to other galaxies. This still seems insufficient to solve the most important problem of the correlation found between QSOs and the {\\it nearest and brightest} galaxies (Kovner 1989); for instance, for the high amplitude angular correlation found by Chu et al.\\ (1984) $\\omega _{\\rm GQ}\\sim 5$. Although  incompleteness could be responsible for these correlations, it is a matter that requires further study.  The existence of a correlation between samples of objects with different redshifts has been advocated by the supporters of non-cosmological redshifts (e.g.\\ Burbidge 1999) as clear evidence of physical association and as proof that QSOs are being ejected by low redshift galaxies. One variant of the model (e.g.\\ Arp 1998a, ch. 3; Arp 1999a) assumes that QSOs are ejected along the rotation axis of the parent galaxy, decreasing in redshift as they move outwards and eventually becoming normal galaxies.  There have been claims of anisotropy in the distribution of QSOs around nearby galaxies. For instance, there are configurations of QSOs aligned along the minor axis of a central Seyfert whose probability of being accidental is $10^{-9}$ to $10^{-10}$ (Arp 1998b, 1999a). Arp \\& Hazard (1980) have reported three QSOs in a straight line, together with another set of three QSOs  in a straight line, a configuration that is very difficult to explain as random. Arp (1999b) finds that the probability of having six out of six QSOs aligned within $\\pm 15^\\circ $ of the minor axis of NGC 5985 to be only $10^{-8}$ to $10^{-9}$. Arp \\& Russell (2001) found that the bright radio quasars 3C37 and 3C39 are paired across the centroid of the disturbed galaxy pair NGC 470/474 in an arrangement with a probability of only $2\\times 10^{-9}$ of being accidental. There is anisotropy in the radio QSO distribution at high flux densities (Shastri \\& Gopal 1983): the number of QSOs on one side of the M33 region is far greater ($\\sim 11\\sigma $) than that of the diametrically opposite region. This alignment of QSOs with the parent galaxy would involve conservation of momentum in their ejection in the heterodox Arp/Burbidge/Narlikar scenario (Narlikar \\& Das 1980).  In this paper we undertake the first study of the polar angle distribution of QSOs around bright  nearby galaxies in statistical terms using large samples of both types of objects. With the release of the SDSS survey, a systematic search for this effect is now possible. The motivation of this research is to test one of the predictions given by the above authors concerning the ejection model. The predictions are as follows: if the QSOs were background objects, no preferential orientation in principle should be found; however, if the QSOs were ejected by the galaxies, as Arp proposed (or alternatively a gravitational effect were predominant in the rotation axis of a galaxy), a higher concentration of objects near the rotational axis of the galaxy would be expected.  Section~2 presents the data and sample selection; the analysis and statistical tests are presented in Sections~3 and 4, Section~5 discusses some details of the method to measure the anisotropy, and in Section 6 we interpret the results. ",
        "conclusions": ""
    },
    "0609/astro-ph0609208_arXiv.txt": {
        "abstract": "Nonrotating stars that end their lives with masses $140 \\msun \\leq M_\\star \\leq 260\\,\\msun$ should explode as pair-production supernovae (\\snggo).  Here I review the physical properties of these objects as well as the prospects for them to be constrained observationally.  In very massive stars, much of the pressure support comes from the radiation field, meaning that they are loosely bound, and that $(d \\lg p/d \\lg \\rho)_{\\rm adiabatic}$ near the center is  close to the minimum value necessary for stability.  Near the end of $C/O$ burning, the central temperature  increases to the point that photons begin to be converted into electron-positron pairs, softening the equation of state below this critical value.   The result is a runaway  collapse, followed by  explosive burning that completely obliterates  the loosely-bound star.  While these explosions can be up to 100 times more energetic that core collapse and Type Ia supernovae, their peak luminosities are only slightly greater.  However, due both to copious Ni$^{56}$ production and hydrogen  recombination, they are brighter much longer, and remain observable for $\\approx$ 1 year.  Since metal enrichment is a local process, PPSNe should occur in pockets of metal-free gas over a broad range of  redshifts, greatly enhancing their detectability, and distributing their nucleosyntehtic products about the Milky Way.   This means that measurements of the abundances of metal-free stars should be thought of as directly constraining these objects.  It also means that ongoing supernova searches, which limit the contribution of very massive stars to $\\lsim 1\\,$\\% of the total star formation rate density out to $z \\approx 2,$ already provide weak constraints for PPSN models. A survey with the NIRCam instrument on  JWST, on the other hand, would be able to extend these limits to $z \\approx 10.$ Observing a $0.3$ deg$^2$ patch of sky for $\\approx 1$ week per year for three consecutive years, such a program would either detect or rule out the existence of these remarkable objects. ",
        "introduction": "Pair-production supernovae (PPSNe) are the uniquely calculable result of nonrotating stars that end their lives in the $140-260 \\msun$ mass range (Heger \\& Woosley 2002, hereafter HW02). Their collapse and explosion result from an instability that generally occurs whenever the central temperature and density of star moves within a well-defined regime (Barkat, Rakavy, \\& Sack 1967). While this instability arises irrespective of the metallicity of the progenitor star, PPSNe are in expected only in primordial environment, and there are three main reasons for this association.  Firstly, in the present metal-rich universe, it appears that stars this large are never assembled, as supported by a wide range of observations.  Figer (2005) carried out a detailed study of the $Z \\approx Z_\\odot$ Arches cluster, which is large ($M_\\star > 10^4 \\msun$), young ($\\tau$ = 2.0 - 2.5 Myrs), and at a well-determined distance, making it ideal for such studies.  No stars more massive than  130 $\\msun$ were found, although more  than 18 were expected. A similar $\\approx 150 \\msun$ limit was found in the lower metallicity cluster R136 in the Large Magellanic Cloud (Weidener \\& Kroupa 2003), and from a  grab-bag of clusters compiled by Oey \\& Clarke (2005).  However, there are good theoretical reasons to believe that the situation may have been very different  under primordial conditions. In this case the primary coolant at low temperatures is molecular hydrogen, which starts to be populated according to local thermodynamic equilibrium  (LTE) at a typical density and temperature of $\\approx 10^4$ cm$^{-3}$ and 100 K. As the Jeans  mass under these conditions is  $\\approx 10^3 \\msun$ the fragmentation of primordial molecular clouds may have been biased towards the formation of stars with very high masses (Nakamura \\& Umemura 1999; Abel Bryan \\& Norman 2000,  Schneider \\etal 2002; Tan \\& McKee 2004).  Finally, as very massive, radiatively supported stars are only loosely bound they tend to drive large winds. However, these winds are primarily line-driven and scale with metallicity as $Z^{1/2}$ or faster  (Kudritzki 2000; Vink \\etal 2001; Kudritzki 2002).  As long as another mechanism did not act to generate significant mass loss in primordial stars (\\eg Smith \\& Owocki 2006)  this raises the real possibility that they may have not only been born, but have ended the lives in the mass range necessary to drive PPSNe.  This conference proceeding summarizes both the underlying physics and the prospects for observation of these most powerful of astrophysical explosions, and it is structured as follows: In \\S2  I describe the pair-production instability in detail, and how it eventually leads to stellar disruption.  In \\S3 I discuss the post-explosion physics of these objects, and how it effects their luminosity and temperature evolution. In \\S4 I discuss the optical light-curves of PPSNe, and contrast them with  lightcuves of SNe Type-II and Ia.  \\S5 discusses the redshift evolution of metal-free stars, and its implications for PPSN environments.  In \\S6 these estimates are used to determine the feasibility of present and future PPSN. In \\S 7 and \\S 8 I discuss the likely distribution of the descendents of metal-free stars in the Galaxy, and I close with a short summary in \\S9. ",
        "conclusions": "Theoretically, PPSNe are the simplest of all supernovae.  Driven by a well-understood dynamical instability, and leading to complete stellar disruption, they are the uniquely  calculable result of nonrotating stars than end their lives in  the $140-260 \\msun$ mass range (Heger \\& Woosley 2002).   The issue is only when and where such stars existed.  In the present, enriched, universe the observed upper mass limit  of  forming stars, and the rate of mass loss in O stars argue strongly against these objects.  However, in the primordial high-redshift universe things are likely to have been very different. The typical fragmentation mass under these conditions is $\\approx 1000 M_\\odot$ and stellar winds, at least of the line driven type observed today, are expected to be negligible.   This raises the real possibility that in the metal-free universe a large fraction of stars generated PPSNe.  As metal-enrichment is an intrinsically local process that proceeds over an extended redshift range, at each redshift, pockets of metal-free star formation are naturally confined to the lowest-mass galaxies, which are small enough not to be clustered near areas  of previous star formation.  As such faint galaxies are difficult to detect and even more difficult to confirm as metal-free, the hosts of \\sngg could easily be lurking at the limits of present-day galaxy surveys.  In S05 we showed that the most important factorsfor modeling PPSNe lightcurves are the mass of the progenitor star and the efficiency of dredge-up of carbon from the core into the envelope.  In general, increasing the mass leads to greater $^{56}$Ni production,  which boots the late time SN luminosity.   Mixing, on the other hand, has two major effects: it leads to a red giant phase that increases the early-time SN luminosity and it decreases the mass of the He core, consequently leading to a somewhat smaller mass of $^{56}$Ni.  Despite these uncertainties, \\sngg in general can be characterized by: (1) peak magnitudes that are brighter than Type II SNe and comparable or slightly brighter than typical SNe Type Ia;  (2) very long decay times $\\sim 1$ year, which result from the  large initial radii and large masses of material involved in the explosion;  and (3)  the presence of hydrogen lines, which are caused by the  outer envelope.  The S05 lightcurves also allowed us to calculate the number of \\sngg detectable in current and planned supernova searches.  Here the long lifetimes help to keep a substantial number of \\sngg  visible at any given time, meaning that ongoing SN searches  should be able to limit the contribution of VMS to  $\\lsim 1\\,$\\% of  the total star formation rate density out to a redshift of 2, unless both mixing and $^{56}$Ni production are absent for all \\snggo. Such constraints already place meaningful  limits on the  cosmological models.  The impact of future NIR searches is even more promising, as the majority of the PPSN light is emitted at restframe wavelengths longward of $\\approx 8000$ \\AA.  In particular JWST surveys with NIRCam have the potential to  place fantastic constraints on PPSNe out to  $z \\approx 10.$  In this case the best approach will be a $\\approx 0.3$ deg$^2$ survey made up of $\\approx 100$ NIRCam pointings with $\\approx 10^4$ sec integrations in two or three bands, with emphasis on the redder colors.     Furthermore this field should be revisited with a cadence of roughly once per year on  three occasions. Although  this would require about 1-2 weeks of dedicated time each year, clearly this program could be carried out in the context of a more general deep field study, with a much broader set of science goals.  Closer to home, we have also studied the final distribution of the elements synthesized in primordial stars.   Despite the large uncertainties involved, all models generically predict significant PopIII star formation in what is now the Galactic halo. Thus, if they   have sufficiently long lifetimes, a significant number of stars formed in  initially primordial star clusters should be found in ongoing surveys for metal-poor halo stars. This is both good and bad news for PPSNe models.  While it implies a real lower mass limit for the PopIII IMF  of at least $0.8 \\msun,$ it also suggests that the lack of an odd-even effect in the observed  abundance ratios of metal-poor stars should be taken as evidence that a significant number of PopIII stars ended their lives as more usual core-collapse supernovae.  However, definitive limits on PPSNe will only come from space based NIR  surveys.  If they result in detections, they will open a new window on star formation and the history of cosmological chemical enrichment.  If they result in upper limits, they will place exquisite constraints on the presence of PPSN forming in primordial environments above the atomic cooing limit, but leave open the question of their formation in the first ``minihalos'' collapsing at extremely high redshift. While this would remain a possibility, the absence of detections from space-based searches would limit PPSN to the most remote and undetectable corner of the universe.  Undaunted theorists might still wish to discuss them in workshops on the deepest depths of the cosmological dark ages.  Observers may be reminded that ``he who wishes to lie, should put the evidence far away'' (Livio 2006)."
    },
    "0609/astro-ph0609722_arXiv.txt": {
        "abstract": "{Carbon is one of the most abundant metals in the universe because of its synthesis in the fundamental 3$\\alpha$ reaction. The knowledge of carbon abundances in different environments is one key ingredient to our understanding of stellar and galactochemical evolution. Studies of luminous OB-type stars allow us to address both topics even in galaxies beyond our own. Unfortunately the history of carbon abundance determinations from these objects in the last three decades is one of limited success. Analyses of the strong and weak line spectra of C\\,{\\sc ii} as well as C\\,{\\sc iii} tend to be largely discrepant. We present results of quantitative spectral analyses based on a sophisticated model atom for non-LTE line-formation calculations of C\\,{\\sc ii-iv}. As a first application, carbon abundances in a sample of B-type dwarfs and giants in nearby associations and in the field are determined. Consistency is finally achieved for all measurable lines from the three ionization stages. This includes in particular the notorious C\\,{\\sc ii} $\\lambda\\lambda$4267 and 6578/6582~{\\AA}~features which are highly important for abundance determinations of fast-rotating and extragalactic objects. The long-standing problem of carbon line formation can now be regarded as solved, with the previous difficulties related to the use of inaccurate atomic data and stellar parameters. A highly homogeneous and slightly sub-solar present-day carbon abundance from young stars in the solar vicinity of $\\log$ C/H\\,+\\,12\\,=\\,8.32$\\pm$0.04 is derived. }  \\FullConference{International Symposium on Nuclear Astrophysics --- Nuclei in the Cosmos --- IX\\\\ June 25-30 2006\\\\ CERN, Geneva, Switzerland}  \\begin{document} ",
        "introduction": " ",
        "conclusions": "A highly consistent mean carbon abundance of $\\varepsilon$(C)\\,$=$ 8.32$\\pm$0.04 is derived from the six sample stars, which provides a tight estimate to the present-day C abundance from young stars in the solar vicinity. The atmospheric composition appears to be unaffected by chemical mixing in the course of stellar evolution, i.e. we find no trend of C abundances with evolutionary age. For comparison, adopting results from \\cite{k92} one derives a mean $\\varepsilon$(C)\\,$=$\\,8.19$\\pm$0.12 from the same six stars, implying a considerable systematic shift and a significantly increased statistical scatter. More objects need to be analysed in order to confirm such homogeneous present-day (slightly) sub-solar carbon abundances -- considering as references $\\varepsilon$(C)$_{\\odot}$\\,=\\,8.39$\\pm$0.05~\\cite{ags05} or  $\\varepsilon$(C)$_{\\odot}$\\,=\\,8.52$\\pm$0.06~\\cite{gs98} -- in nearby associations (HR\\,1861: Ori OB1; $\\tau$\\,Sco, HR\\,5285: Sco-Cen) and in the field (the other stars). Despite the small sample size analysed so far, our highly accurate results indicate that the large scatter found for carbon abundances in early-type stars in previous work could be a consequence of systematic uncertainties introduced by the choice of inappropriate atomic data and stellar parameters."
    },
    "0609/astro-ph0609178_arXiv.txt": {
        "abstract": "We present a detailed investigation into the recent star formation histories of \\numl\\/ Luminous Red Galaxies (LRGs) based on the \\hdel (4101\\AA), and \\OII\\/ (3727\\AA) lines and the D4000 index. LRGs are luminous (L$>3$L*), galaxies which have been selected to have photometric properties consistent with an old, passively evolving stellar population. For this study we utilise LRGs from the recently completed 2dF-SDSS LRG and QSO survey (2SLAQ). Equivalent widths of the \\hdel\\/ and \\OII\\/ lines are measured and used to define three spectral types, those with only strong \\hdel\\/ absorption (k+a), those with strong \\OII\\/ in emission (em) and those with both (em+a). All other LRGs are considered to have passive star formation histories. The vast majority of LRGs are found to be passive ($\\sim$80 per cent), however significant numbers of k+a (2.7 per cent), em+a (1.2 per cent) and em LRGs (8.6 per cent) are identified. An investigation into the redshift dependence of the fractions is also performed. A sample of SDSS MAIN galaxies with colours and luminosities consistent with the 2SLAQ LRGs is selected to provide a low redshift comparison. While the em and em+a fractions are consistent with the low redshift SDSS sample, the fraction of k+a LRGs is found to increase significantly with redshift. This result is interpreted as an indication of an increasing amount of recent star formation activity in LRGs with redshift. By considering the expected life time of the k+a phase, the number of LRGs which will undergo a k+a phase can be estimated. A crude comparison of this estimate with the predictions from semi-analytic models of galaxy formation shows that the predicted level of k+a and em+a activity is not sufficient to reconcile the predicted mass growth for massive early-types in a hierarchical merging scenario. ",
        "introduction": "The widely accepted paradigm for galaxy formation is the hierarchical merger model which dictates that galaxies form via continuous mergers with smaller objects (e.g. \\nocite{cole1994}{Cole} {et~al.} 1994). This scenario is a natural consequence of the $\\Lambda$ cold dark matter ($\\Lambda$CDM) model of structure formation, which has been highly successful in explaining both the observed structure from the large redshift surveys (2dF Galaxy Redshift Survey (2dFGRS); \\nocite{Colless01}{Colless} {et~al.} 2001, Sloan Digital Sky Survey (SDSS); \\nocite{stoughton2002}{Stoughton} {et~al.} 2002), and the anisotropy in the cosmic microwave background \\nocite{Sper03}({Spergel} {et~al.} 2003). In the context of this model one would naively expect the most massive galaxies to also be the youngest as they have been `built up' slowly via mergers. However traditionally early-types are thought to have formed early and rapidly, with a vast amount of observational evidence points supporting this conclusion. The homogeneous nature of massive early-type galaxy properties, such as small scatter in the colour-magnitude relation (CMR) \\nocite{Vis77,Bower92,Terl2001,aragon93,stanford95,ellis97,stanford98}({Visvanathan} \\& {Sandage} 1977; {Bower}, {Lucey} \\& {Ellis} 1992; {Terlevich}, {Caldwell}, \\&  {Bower} 2001; {Aragon-Salamanca} {et~al.} 1993; {Stanford}, {Eisenhardt}, \\&  {Dickinson} 1995; {Ellis} {et~al.} 1997; {Stanford}, {Eisenhardt}, \\&  {Dickinson} 1998, among others) and fundamental plane (FP) \\nocite{Dress1987,Burstein1997}({Dressler} {et~al.} 1987; {Burstein} {et~al.} 1997), and the lack of evolution in these relations with redshift \\nocite{Kodama1998,Vandokkum2003}({Kodama} {et~al.} 1998; {van Dokkum} \\& {Stanford} 2003), is consistent with a scenario in which most, if not all massive early-types, formed at high redshift (at least $z>2$) and have experienced only passive evolution of their stellar populations since.   While this could simply be a result of selection bias \\nocite{dokkum2001}({van Dokkum} \\& {Franx} 2001), the lack of direct evidence of strong evolution in the population of early-type galaxies since $z\\sim$1 remains a serious problem for the hierarchical model of galaxy formation. In recent years much work has gone into reconciling these differences. While theoretical models have moved towards consistency with the CMR and FP \\nocite{Kauffmann98, delucia2005}({Kauffmann} \\& {Charlot} 1998; {De Lucia} {et~al.} 2005, among others), observational studies based on the spectroscopic and morphological properties of early-types have found significant populations which show evidence for recent star formation and/or interactions. A morphologically based study by \\nocite{michard2004}{Michard} \\& {Prugniel} (2004) found a significant number ($\\sim30$ per cent) of nearby ($z\\sim0.01$) early-type galaxies showed evidence of recent interactions/mergers, although 40 per cent of these showed no evidence for a young stellar population. A similar study of 86 nearby bulge-dominated red galaxies by \\nocite{vandokkum2005} {van Dokkum} {et~al.} (2005) found 71 per cent had evidence of tidal interactions. However given that both of these studies were undertaken at low redshift the volume sampled was too small for significant numbers of LRG analogues to be present. Moving to higher redshifts, \\nocite{Willis2002}{Willis} {et~al.} (2002) found evidence for current star formation, in the form of detection of the \\OII\\/ emission line, in over 20 per cent of a sample of 415 luminous field early-types at $z\\sim0.3$. While these studies show that evolution is occurring in at least some of the early-type population, it remains unclear if all early-types have experienced evolution of this type. Another contentious point is whether early-type galaxy acquire mass predominately in the form of star forming gas or already formed stars via so called `dry mergers' (\\nocite{bell2005} Bell  et al. 2005).  In this study we look for evidence of recent star formation in a sample of Luminous Red Galaxies (LRGs). LRGs are luminous ($L>3L^*$) early-type galaxies, analogous to bright cluster galaxies (BCG), selected via their red colours which are broadly consistent with a passively evolving stellar population. These galaxies are ideal candidates for the purposes of this study as they have been shown to have relatively homogeneous spectral properties \\nocite{eis01} ({Eisenstein} {et~al.} 2001), enabling any deviation due to recent star formation episodes to be readily identified. They are also the most massive galaxies, which, if the hierarchical merger model is to be correct, must also have the most vigorous merger histories. Thus a sufficiently large sample of LRGs provides the perfect test bed for current theories of galaxy formation/evolution.   Here we probe the star formation histories of our LRG sample via their spectral properties, in particular the \\OII\\/(3727\\AA) emission and \\hdel\\/(4101\\AA) absorption lines and the D4000 index. LRGs are selected from both the SDSS and 2dF-SDSS LRG and QSO (2SLAQ) surveys, resulting in a sample of \\numl\\/ LRGs covering a redshift range of $0.2<z<0.7$. Details of the sample selection are presented in Section \\ref{sec:sam}. Section \\ref{sec:meth} describes the analysis and modeling performed, with the subsequent results presented in Section \\ref{sec:results}. The implications and conclusions to be drawn from these results are then presented in Sections \\ref{sec:disc}, and \\ref{sec:conc}, respectively.  Where necessary, we have assumed a flat cosmology with $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$ and H$_0=70$ kms$^{-1}$Mpc$^{-1}$. ",
        "conclusions": "\\label{sec:conc}  We have determined the recent star formation histories of a sample of \\numl\\/ LRGs based on the equivalent widths of the \\hdel\\/ and \\OII\\/ lines. While the majority ($>$80 per cent) show the spectral properties of an old, passively evolving, stellar population, a significant number of LRGs show evidence for recent and/or ongoing star formation in the form of k+a (2.7 per cent), em+a (1.2 per cent) or em LRGs (8.6 per cent). By dividing the sample into 2 redshift subsamples from $0.45<z<0.55$ and $0.55<z<0.65$, and comparing to a $z\\sim0.15$ sample selected from SDSS, it is observed that the fraction of k+a LRGs increases with redshift as $(1+z)^{2.8\\pm0.7}$.  Spectral synthesis models, utilising the code of \\nocite{Bruz03}{Bruzual} \\& {Charlot} (2003), suggest that the k/em+a LRGs could originate from passive LRGs which undergo a starburst in which material equivalent to $\\sim 1$ per cent of the stellar mass is consumed in new star formation.  Several origins for k+a and em+a LRGs are considered, including cooling flows and mergers, however identification of the exact formation mechanism, or even if k+a and em+a LRGs share a common origin, is not possible with the current data. It is clear that k+a and em+a activity represents a signpost of recent evolution in LRGs, however by considering the life time of the k+a/em+a phase we estimate that only 8 per cent of LRGs will experience a k+a or em+a phase between $z=0.8$ and the present. By comparing this with the semi-analytic galaxy formation models of De Lucia et~al. (2005) we conclude that k+a/em+a activity cannot be responsible for the predicted mass growth of massive ellipticals since $z=0.8$."
    },
    "0609/astro-ph0609664_arXiv.txt": {
        "abstract": "{We present a Chandra analysis of the X-ray spectra of 56 clusters of galaxies at $z\\,\\ga 0.3$, which cover a temperature range of $3\\,\\la kT\\,\\la 15$~keV.}  {Our analysis is aimed at measuring the iron abundance in the intra-cluster medium (ICM) out to the highest redshift probed to date.}  {We made use of combined spectral analysis performed over five redshift bins at $0.3\\la z \\la 1.3$ to estimate the average emission weighted iron abundance. We applied non-parametric statistics to assess correlations between temperature, metallicity, and redshift.}  {We find that the emission-weighted iron abundance measured within $(0.15-0.3)\\,R_{vir}$ in clusters below 5~keV is, on average, a factor of $\\sim2$ higher than in hotter clusters, following $Z(T)\\simeq 0.88\\,T^{-0.47}\\,Z_\\odot$, which confirms the trend seen in local samples.  We also find a constant average iron abundance $Z_{Fe}\\simeq 0.25\\,Z_\\odot$ as a function of redshift, but only for clusters at $z\\ga0.5$.  The emission-weighted iron abundance is significantly higher ($Z_{Fe}\\simeq0.4\\,Z_\\odot$) in the redshift range $z\\simeq0.3-0.5$, approaching the value measured locally in the inner $0.15\\,R_{vir}$ radii for a mix of cool-core and non cool-core clusters in the redshift range $0.1<z<0.3$.  The decrease in metallicity with redshift can be parametrized by a power law of the form $\\sim(1+z)^{-1.25}$.  We tested our results against selection effects and the possible evolution in the occurrence of metallicity and temperature gradients in our sample, and we do not find any evidence of a significant bias associated to these effects.}  {The observed evolution implies that the average iron content of the ICM at the present epoch is a factor of $\\sim2$ larger than at $z\\simeq 1.2$.  We confirm that the ICM is already significantly enriched ($Z_{Fe}\\simeq0.25 \\,Z_\\odot$) at a look-back time of 9~Gyr. Our data provide significant constraints on the time scales and physical processes that drive the chemical enrichment of the ICM.} ",
        "introduction": "Clusters of galaxies are the largest virialized structures in the Universe arising from the gravitational collapse of rare high peaks of primordial density perturbations \\citep[e.g.][]{pee93, col95, pea99, rbn02, v05}.  As a result of adiabatic compression and shocks generated by supersonic motion during shell crossing and virialization, a hot thin gas permeating the cluster's gravitational potential well is formed. Typically this gas, which is enriched with metals ejected form Supernovae (SNe) explosions through subsequent episodes of star formation \\citep[e.g.][]{mat88}, reaches temperatures of several $10^7$ K and therefore emits mainly via thermal bremsstrahlung in the X-rays. At such temperatures most of the elements are either fully ionized or are in a high ionization state. Strong emission lines may originate by collisional excitation of K- and L-shell transitions in highly ionized elements, such as H- and He-like iron, oxygen, silicon or sulfur. In the isothermal approximation, the line intensities depend on the abundances of heavy elements, while the continuum intensity is mainly due to hydrogen and helium.  Therefore the equivalent width of a line, under the reasonable assumption of collisional equilibrium, gives a direct measurement of the abundance of the corresponding element.  The chemical enrichment of the intra-cluster medium (ICM) is an unambiguous signature of star formation in cluster galaxies.  A knowledge of the history of the ICM metal enrichment is also necessary in understanding the mode and epoch of cluster formation and the thermodynamic evolution of the cluster baryons in their hot and cold phases. In this respect, measuring the properties of the ICM at high redshift is important for constraining the physical processes involved in the diffusion of energy and metals within clusters. In particular, the chemical evolution of the ICM sets constraints on the SNe rate in the cluster galaxies \\citep[see][]{ett05}. The SNe explosions are, in fact, the main contributor to the metal enrichment and are also expected to provide a source of ICM heating \\citep[e.g.][]{pip02, tor04}.  In addition, much emphasis has recently been given to the evolution of global scaling relations for the ICM, such as the luminosity-temperature relation and the entropy-temperature relation \\citep{hol02,vik02,ett04,pra06,maugh06}. Such scaling relations represent a further signature of galaxy formation and super-massive black hole activity, whose relative contribution to the ICM energetic is under investigation.  \\citet[ hereafter Paper~I]{toz03} measured for the first time the average iron abundance in the ICM of hot clusters of galaxies out to redshift $z\\simeq1.3$. Their findings, which mostly probed the redshift range $0.5<z<1.3$ (with 90\\% of their sample in this range), suggested that the mean iron content of the ICM is approximately constant with a value $Z_{Fe}\\simeq 0.3\\,Z_\\odot$\\footnote{Solar abundance values are set to those provided by \\citet{and89}; in particular, the solar abundance of iron atoms relative to hydrogen is $4.68\\times 10^{-5}$. See Sect.~3.1 for discussion of the solar iron abundance.} for clusters with temperatures $kT>5$~keV. When comparing their results at high redshift with the local values of the iron abundance, they referred to the typical value $Z_{Fe}\\simeq 0.3\\,Z_\\odot$ available from the literature at that time \\citep{ren97, deg01}.  Therefore, they concluded that no evolution of iron abundance with redshift was found out to the probed redshift.  More recently, the determination of an average local value of the ICM iron content has become a harder task due to the different metal distribution inside cool-core and non cool-core clusters \\citep{deg04}. The former show a central peak of iron abundance with $Z_{Fe}\\simeq0.6-0.8\\,Z_\\odot$ and a plateau at $Z_{Fe}\\simeq 0.3\\,Z_\\odot$ in the outer regions \\citep[see][]{tam04, vik05}, and the latter a somewhat lower value ($Z_{Fe}\\simeq0.2-0.3\\,Z_\\odot$) constant with radius. The iron-rich cores typically have a size of 100~kpc \\citep[see][]{deg04, vik05}. Therefore, it is important to take into account any effects due to different physical apertures when comparing different samples from the literature.  Studies of local cluster samples have also found an increase in the iron abundance in clusters with temperatures $\\la5$~keV \\citep[see][]{arn92, fin01, bag05}, whose physical interpretation is still a matter of debate.  Single observations of high-z ($z>1$) clusters confirmed that $Z_{Fe}\\simeq 0.3\\,Z_\\odot$ or higher is common in the ICM \\citep{ros04, has04}. This implies that the last episode of star formation in clusters of galaxies must have taken place at earlier epochs in order to significantly enrich the diffuse medium with metals. Therefore, the study of the iron abundance at high redshift is expected to place strong constraints on the star formation history of cluster galaxies and on its effects on the thermodynamics of the ICM.  Here we present a significantly improved analysis compared to our previous work by substantially extending the sample (56 clusters instead of the 19 presented in Paper~I) and by using the most recent {\\em Chandra} calibrations in order to have an up-to-date data reduction at the time of writing. This increase in statistics allows us to investigate the relation between the iron abundance and global temperature of the ICM at high redshift and to derive a more robust measurement of the cosmic evolution of the average iron abundance in the ICM, which is then compared with predictions from the cosmic star formation rate.  The plan of the paper is as follows. In Sect.~2 we describe the data reduction procedure. In Sect.~3 we describe our spectral analysis (Sect.~3.1) and present the main results, subdivided as follows: single source analysis (Sect.~3.2), correlation between iron abundance and temperature (Sect.~3.3), evolution of the average iron abundance as a function of redshift (Sect.~3.4), and proper comparison with the {\\em local} iron abundance (Sect.~3.5). In Sect.~4 we discuss the implications of our findings and in Sect.~5 we summarize our conclusions. In the Appendix, we describe in detail the spectral simulations performed to investigate the possible spectral-fitting biases due to unresolved temperature and metallicity gradients in our sample. We adopt a cosmological model with $H_0=70$ km/s/Mpc, $\\Omega_M=0.3$, and $\\Omega_\\Lambda=0.7$ throughout.  Quoted confidence intervals are $1\\sigma$ unless otherwise stated. ",
        "conclusions": "We have presented the spectral analysis of 56 clusters of galaxies at intermediate-to-high redshifts observed by {\\em Chandra} and XMM-{\\em Newton}. This work improves our first analysis aimed at tracing the evolution of the iron content of the ICM out to $z\\ga1$ (Paper~I), by substantially extending the sample. The main results of our work can be summarized as follows:  \\begin{itemize}  \\item We determine the average ICM iron abundance with a $\\sim20$\\% uncertainty at $z>1$ ($Z_{Fe}=0.27\\pm0.05\\,Z_\\odot$), thus confirming the presence of a significant amount of iron in high-z clusters. $Z_{Fe}$ is constant above $z\\simeq0.5$, the largest variations being measured at lower redshifts.  \\item We find a significantly higher average iron abundance in clusters with $kT<5$~keV, in agreement with trends measured in local samples. For $kT>3$~keV, $Z_{Fe}$ scales with temperature as $Z_{Fe}(T)\\simeq0.88\\,T^{-0.47}$.  \\item We find significant evidence of a decrease in $Z_{Fe}$ as a function of redshift, which can be parametrized by a power law $\\langle Z_{Fe}\\rangle \\simeq Z_{Fe}(0)\\,(1+z)^{-\\alpha_z}$, with $Z_{Fe}(0)\\simeq0.54 \\pm 0.04$ and $\\alpha_z\\simeq1.25 \\pm 0.15$. This implies an evolution of more than a factor of 2 from $z=0.4$ to $z=1.3$.  \\end{itemize}  We carefully checked that the extrapolation towards $z \\simeq 0.2$ of the measured trend, pointing to $Z_{Fe} \\simeq 0.5\\, Z_\\odot$, is consistent with the values measured within a radius $r= 0.15\\, R_{vir}$ in local samples including a mix of cool-core and non cool-core clusters. We also investigated whether the observed evolution is driven by a negative evolution in the occurrence of cool-core clusters with strong metallicity gradients towards the center, but we do not find any clear evidence of this effect. We note, however, that a proper investigation of the thermal and chemical properties of the central regions of high-z clusters is necessary to confirm whether the observed evolution by a factor of $\\sim2$ between $z=0.4$ and $z=1.3$ is due entirely to physical processes associated with the production and release of iron into the ICM, or partially associated with a redistribution of metals connected to the evolution of cool cores.  Precise measurements of the metal content of clusters over large look-back times provide a useful fossil record for the past star formation history of cluster baryons.  A significant iron abundance in the ICM up to $z\\simeq 1.2$ is consistent with a peak in star formation for proto-cluster regions occurring at redshift $z\\simeq4-5$. On the other hand, a positive evolution of $Z_{Fe}$ with cosmic time in the last 5~Gyrs is expected on the basis of the observed cosmic star formation rate for a set of chemical enrichment models.  Our data provide further constraints on the chemical evolution of cosmic baryons in the hot diffuse and cold phases."
    },
    "0609/astro-ph0609571_arXiv.txt": {
        "abstract": "XEUS is a large area telescope aiming to rise X-ray Astronomy to the level of Optical Astronomy in terms of collecting areas. It will be based on two satellites, locked on a formation flight, one with the optics, one with the focal plane. The present design of the focal plane foresees, as an auxiliary instrument, the inclusion of a Polarimeter based on a Micropattern Chamber. We show how such a device is capable to solve open problems on many classes of High Energy Astrophysics objects and to use X-ray sources as a laboratory for a substantial progress on Fundamental Physics. ",
        "introduction": "\\label{sect:intro}  %  If we look back to the history of X-ray Astronomy we can say that after the first pioneering activity with rockets the sequence UHURU, Einstein, Chandra, following a path tracked by R. Giacconi, has been the backbone of this activity, that nowadays has arrived to a turning point. No major mission is presently established and ideas about how such a mission should be designed are not clear. Chandra has shown the great capabilities of an optics at sub-arcsecond level. Both Chandra and Newton have opened the window of high resolution spectroscopy. In some sense these two missions have brought X-ray Astronomy much closer to Optical Astronomy, but when we convert the effective areas of these telescopes into equivalent diameters we see that from the collection capability, they are equivalent to medium level optical amatorial telescopes, although equipped with an extremely performing focal plane instrumentations.  These rose the need of a new generation telescope with a collecting area two orders of magnitude higher, while still preserving an angular resolution at arcsecond level, which is imperative to study crowded deep fields avoiding source confusion. The availability of International Space Station suggested that such a telescope could be assembled in orbit, so bypassing the limits of launch vehicles. The fog that is folding the future of ISS and the general decline of funds for Space based Astronomical Research, have suggested a less ambitious solution. The present design is based on a telescope of completely new technology, drastically reducing the weight to area ratio, with an effective area of the order of 5 m$^{2}$ and a focal length of 35 m, harbored on one satellite. A second satellite will be launched with a complex focal plane instrumentation locked to the first in a formation flight. The ensemble should be put in a L2 orbit.  The focal plane would be based on two major instruments, a micro-calorimeter array based on TES technology and a Wide Field Imager based on APS technology. These two instruments will be the drivers for the study of remote Universe, explore the formation of large scale structures, the role of black holes and their interaction with galaxies. But they will provide as well the capability to study deep Physics of more bright objects. To this purpose the XEUS team is foreseeing the possibility to include more instrumentations, mainly devoted to a more restrict population of objects, but capable to address questions of extreme interest. These auxiliary instrumentations to be implemented if compatible with available resources and on the basis of a trade-off of the science, include a photoelectric polarimeter named XPOL. ",
        "conclusions": "\\label{sect:conclusions} The advent of the Micro Pattern Gas Detector, a finely subdivided pixel gas chamber detector readout by large area and small pixel size CMOS ASIC  allows to image the track of the photoelectron and derives  the polarization of the X-ray photon together with energy position and time. XPOL, as auxiliary instrument at the focus of the large mirror area of XEUS telescope, will be decisive in many of the scientific goal of the mission being therefore the high throughput no more laking tool of X-ray Astronomy."
    },
    "0609/astro-ph0609092_arXiv.txt": {
        "abstract": "The high radio-flux brightness temperature of the recently discovered class of sources known as Rotating RAdio Transients (RRATs) motivates detailed study in the X-ray band.  The source \\rrat\\ has a (large) error region which includes the frequently observed low-mass X-ray binary transient \\aql.  We describe analyses of historical X-ray data, searching for X-ray phenomena (sources, behaviors), finding no sources or behaviors which may unequivocally be associated with \\rrat.  We put forward a candidate X-ray counterpart to \\rrat, discovered in a \\chandra\\ observation in Feb 2001, which fades by a factor $>5$ prior to April 2004. Archival {\\em ROSAT} observations detect the source as early as Oct 1992, with a flux comparable to the Feb 2001 {\\em Chandra} detection. The X-ray flux and optical ($F_X/F_R>12$) and near infra-red ($F_X/F_J>35$) limits, as well as the X-ray flux itself, are consistent with an AGN origin, unrelated to \\rrat.  Searches for msec X-ray bursts found no evidence for such a signal, and we place the first observational upper-limit on the X-ray to radio flux ratio of RRAT bursts: $\\fxfr<6 \\tee{-11} \\cgsflux\\ {\\rm mJy}^{-1}$.  The upper-limit on the X-ray burst flux (corresponding to $<2.2\\tee{37} (d/3.3\\,\\kpc)^2 \\cgslum$, 2-10 keV) requires a limit on the spectral energy density power-law slope of $\\alpha<-0.3$ between the radio and X-ray bands. We place a limit on the time-average X-ray burst luminosity, associated with radio bursts, of $\\leq3.4\\tee{30}\\,(d/3.3\\,\\kpc)^2 \\cgslum$. Future definitive association between \\cxo\\ and \\rrat\\ could be made with a better radio localization of \\rrat; or discovery of msec X-ray bursts associated with the radio bursts.  Adaptive optic infra-red observations of the crowded source field of \\cxo\\ will be required to find a NIR counterpart, which may reveal its source class. ",
        "introduction": "Recently, a new observational class of neutron stars has been proposed \\citep{mclaughlin06}, comprising eleven sources defined by the phenomenon of fast (2-30 ms) radio bursts, which repeat on timescales of 4\\,min - 3\\,hr. Through period folding, periodicities were claimed to be found in the range of $P=$0.4-7\\,s for ten of the eleven sources, which suggested association with rotating neutron stars.  In three of the eleven sources, period derivatives were claimed, with one suggesting a magnetic field strength of 5\\tee{13} G -- a possible magnetar.  Because the $P-\\dot{P}$ analysis remains unpublished and undescribed, we regard association of the radio-bursting population with neutron stars to be an unsupported hypothesis; nonetheless, we adopt the hypothesis and nomenclature of referring to the sources as Rotating RAdio Transients (RRATs), and consider their association with neutron stars to be the leading hypothesis.  The proposed class was recently made more intriguing by the discovery of transient pulsed radio emission from a known magnetar, XTE~J1810-197 \\citep{camilo06}, with peak flux density $>$1 Jy and 89\\ppm5\\% polarization at 8.4 GHz, with period $P\\sim5.54\\, {\\rm s}$. The radio pulsation was not present prior to the 2003 discovery outburst of this magnetar \\citep{ibrahim04}, detected serendipitously from X-ray pulsations during observations of a nearby source; sensitive X-ray observations found a fading two-temperature X-ray source, with luminosity $\\approx\\ee{34-35}$ \\cgslum\\ on timescales of 300 and 900 days (for two different temperature components; \\citealt{gotthelf05}).  Moreover, day-to-day fluctuations in the radio peak flux were observed, of $\\approx$2. This is interpreted to imply the appearance of the radio pulsations is associated with the reconfiguration of the magnetar magnetic field, which takes place associated with magnetar bursting activity \\citep{woods01}.  The relationship between the observed process and RRATs is unclear, however observation of radio and X-ray activity associated with transient magnetic field reconfiguration may be applicable to the RRATs, in light of the observed transient nature of their radio emission.  An X-ray counterpart has been reported for RRAT J1819-1458, discovered serendipitously as CXOU J181934.1$-$145804 \\citep{reynolds06}.  This X-ray source was observed with absorption corrected flux 2\\tee{-12} \\cgsflux\\ (0.5-8 keV, with an order of magnitude uncertainty due to uncertainty in the absorption) detected within the 5\\arcsec\\ $\\times$ 32\\arcsec error ellipse of the RRAT.  Comparison with the source density of such X-ray sources find the probability of finding such a bright source by chance $<$\\ee{-4}, compared with the {\\em ASCA} Galactic Plane Survey \\citep{sugizaki01}, at the best-fit unabsorbed flux of 2\\tee{-12} \\cgsflux.  The X-ray spectrum is soft ($kT$=120\\ppm40 eV), which supports classification of a thermal, cooling isolated neutron star (INS) -- further supported by the absence of a positionally coincident optical or infrared counterpart. The most convincing evidence for association CXOU J181934.1$-$145804 with RRAT J1819-1458 is the low probability of finding such a bright X-ray source in the RRAT positional error circle.  As Reynolds \\etal\\ point out, the observed soft spectrum in light of the high Galactic column density (\\nh=1.6\\tee{22} \\perval{cm}{-2}) in the direction of this source strongly supports interpretation of the X-ray source as within the Galaxy, further supporting an neutron star interpretation.  A second source, \\rrat, has an error ellipse of 15\\arcmin$\\times$7\\arcmin (1$\\sigma$), which includes the prolific transient neutron star low-mass X-ray binary Aql X-1 \\citep{liu01}. The derived dispersion measure distance \\citep{mclaughlin06} is $d=$3.3 kpc, modestly but perhaps not significantly closer than the typical 4-6 kpc distance adopted for Aql X-1.  Four radio bursts at $\\nu=$1.4 GHz, of half-width $w_{50}$=2 msec were observed in 13 hours of observing -- the least active of the tabulated RRATs -- with peak fluxes of $S_\\nu$=250 mJy.  If we take $w_{50}$ to be the light crossing time of the emission region, this implies a source brightness temperature (e.g., \\citealt{guedel02}):  \\begin{equation} T_b = 4\\times10^{22} \\left(\\frac{S_\\nu}{250\\, {\\rm mJy}}\\right) \\left( \\frac{ 1.4\\, {\\rm GHz}}{\\nu}\\right)^2 \\left( \\frac{2\\,{\\rm ms}}{w_{\\rm 50}} \\right)^2 \\left( \\frac{d}{3.3\\, {\\rm kpc}} \\right)^2 \\, {\\rm K } \\end{equation}  \\noindent  This is significantly above the brightness temperature limit for inverse Compton scattering ($T_B\\sim\\ee{12} K$) which likely requires a coherent emission mechanism (for this, and other RRATs), possibly accompanied by X-ray emission.  The historical X-ray coverage of Aql X-1 permits study of this field (although not covering the entire error box of the RRAT).  In this paper, we report X-ray studies of the field of \\rrat.  We describe a candidate X-ray counterpart (\\cxo) to the \\rrat, detected serendipitously in an historical {\\em Chandra} observation of Aql X-1 (\\S~\\ref{sec:obs}).  In \\S~\\ref{sec:pca}, we describe a search for msec X-ray bursts from the RRAT, deriving the first limit on the X-ray to radio flux ratio for the RRAT bursts.  In \\S~\\ref{sec:con} we summarize and discuss our results. ",
        "conclusions": "\\label{sec:con}  The historical X-ray coverage of the error box of \\rrat\\ exists due to the source's close proximity to Aql X-1, a transient LMXB.  Other than Aql X-1 itself, we find one unidentified X-ray source in the {\\em Chandra} observations, which has the following properties: moderately hard spectrum ($\\alpha$=1 photon power-law index); intensity variability by a factor of $\\approxgt$5 over timescales of \\approxlt yrs.  We find no evidence of pulsations, but with weak limits (\\approxlt 42\\% rms) in the 0.0011-1.133 Hz range.  The observed X-ray flux gives a value of $F_X/F_R>12$, and $F_X/F_J> 35$, limits which are consistent with an AGN, as well as with the much higher such ratios of clusters, CV/low-mass-X-ray binaries, and isolated neutron stars.  Based on the density of AGN, the probability of finding one such object in the $15\\arcmin\\times7\\arcmin$ error box of \\rrat\\ is $\\sim$0.3 \\citep{hasinger01}, which is sufficiently large as to permit an explanation for this source as simply an unrelated AGN in the source field.  We thus do not rule out an AGN origin for this X-ray source   Though we regard the suggestion that RRATs are neutron stars to be the leading hypothesis, we would not exclude the possibility that one of the radio-bursting sources could be of a different origin.  Based on the fading X-ray behavior, we put forward \\cxo\\  as a candidate X-ray counterpart to \\rrat, warranting further study. Both \\cxo\\ and Aql X-1 fall within the error ellipse of \\rrat. No new activity was detected from Aql X-1.  Observations which would securely identify this object as the counterpart include a 1\\arcsec\\ radio localization of \\rrat, securing positional coincidence; or detection of fast ($\\sim$\\msec) X-ray bursts, which would be associated with the radio-burst phenomena.  The limit we place on the X-ray/radio flux ratio of the RRAT bursts, while the first such limit for this phenomenon, does not seem to exclude any known transient phenomena; the limit can be used to plan future, more sensitive searches."
    },
    "0609/astro-ph0609747_arXiv.txt": {
        "abstract": " ",
        "introduction": "The currently favoured Cold Dark Matter (CDM) model of galaxy formation predicts many more Dark Matter (DM) halos than are observable as galaxies. Amongst other speculations the existence of dark galaxies offers one solution (Verde et al., 2002). It is possible that the currently observable galaxies are just some fraction of a larger and more varied population of DM halos. So, in addition to the known galaxies there maybe DM halos with no baryons and/or those with baryons, but no stars (Hawkins, 1997, Jimenez et al., 1997). Without star formation the latter dark galaxies should remain un-enriched in heavy elements and contain proportions of ionised, atomic and molecular hydrogen commensurate with their gas densities and temperatures. Given that these dark galaxies contain baryons they may be accessible to observation by methods tried and tested on their more visible cousins.  The exsistence of objects that contain HI, but no stars, has been known for many years. For example High Velocity Clouds, tidal tails and clouds close to optically bright galaxies, but none of these have the characteristics that might lead them to be described as galaxies. To count as a dark galaxy a source requires HI emission over galactic scales (greater than a few kpc) with velocity widths and profiles commensurate with a dynamically stable system (velocities greater than a few tens of km s$^{-1}$).  The idea that there are galaxies only detectable by their gas content is also not new and there have been in the past blind HI surveys that might have discovered such objects if indeed they exist. Henning (1995) carried out a blind HI survey using the NRAO 300ft telescope covering 183 sq deg of sky. Most of the survey area was within the Zone of Avoidence (ZOA) thus making a convincing HI detection of a dark galaxy difficult. The Arecibo HI Sky Survey (AHISS, Zwaan et al. 1997, Zwaan 2000) covered just 13 sq deg and it appears that each detection has an associated optical counterpart. The Arecibo Slice Survey (ASS, Schneider et al. 1998, Spitzak and Schneider 1998) covered 55 sq deg. All but one of the detected galaxies has been associated with an optical source. The Arecibo Dual-Beam Survey (ADBS, Rosenberg and Schneider 2000) covered an area of 430 sq deg producing 11 objects that had no optical counterparts on the POSSI/II photographic plates. The HI Jodrell All Sky Survey (HIJASS, Lang et al. 2003) has covered about 1100 sq deg, but at the moment lacks the optical follow up to decide on the dark galaxy issue. These surveys covered small areas compared to the recent HIPASS observations (the catalogue of observations is known as HICAT, Meyer et al. 2004). The HIPASS has now covered the southern sky, $\\approx$ 21,000 sq deg, using a multi-beam instrument (HIDEEP is a smaller deeper sub-set of HIPASS). One apparent result to come from this work is that all HI sources seem to be associated with an optical counterpart (Doyle et al. 2005) and on the face of it there seems to be little scope for the existence of dark galaxies. This has important implications: for example with regard to the efficiency with which gas is constrained to remain within DM haloes and the close physical link between the availability of HI and the formation of stars.  In table 1 we list the areas covered $(A)$, channel velocity widths ($\\delta V$) and rms noise ($\\sigma_{rms}$) of the above surveys. We will return to what each might and might not have expected to detect later. In table 1 we have also included the Arecibo Legacy Fast ALFA survey (ALFALFA survey - Giovanelli et al., 2005) which is currently being carried out at the Arecibo observatory and the Arecibo Galactic Environments Survey (AGES - www.astro.cardiff.ac.uk/groups/galaxies/ages.html) which is due to start in January 2006.  \\begin{table} \\begin{center} \\begin{tabular}{l|c|c|c|c} Survey  & $A$        &  $\\delta V$    & $\\sigma_{rms}$    &  $V^{Max}$ \\\\ & (sq deg) & (km s$^{-1}$)  & (mJy beam$^{-1}$) &   (km s$^{-1}$)  \\\\ \\hline Henning &    183   &      22        &     3.4           &  6800 \\\\ AHISS   &  13      &      16        &     0.8           &  7400  \\\\ ASS     &  34      &      16        &     2.0           &  8340   \\\\ ADBS    &  430     &      34        &     3.5           &  7977   \\\\ HIJASS  &  1,115   &      13        &     14.0          &  10000  \\\\ HIPASS  &  21,000  &      13        &     13.0          &  12700  \\\\ HIDEEP  &  32      &      13        &     3.4           &  12700  \\\\ VIRGOHI &  32      &      13        &     4.0           &  10000   \\\\ ALFALFA &  10,000  &      5         &     2.0           &  18000  \\\\ AGES    &  200     &      5         &     0.5           &  18000 \\\\ \\end{tabular} \\end{center} \\caption{A comparison of survey area (A), velocity channel widths ($\\delta V$), rms noise ($\\sigma_{rms}$) and maximum velocity ($V^{Max}$) for previous and future  (ALFALFA and AGES) blind HI surveys.} \\end{table}  In 2004 we used the Jodrell Bank multi-beam instrument to carry out a deep survey of part of the Virgo cluster (Davies et al., 2004). The Virgo survey (VIRGOHI) reduced the $\\approx 14$ mJy per channel of the HIJASS survey to $\\approx 4$ mJy (Table 1) enabling us to look for both lower mass and lower column density hydrogen than would be detected in the normal survey. The area observed was also one in which we know there is a large number of nearby optically detected galaxies. Davies et al. identified four objects that had HI detections but no obvious optical counterparts. One of these was subsequently rejected as noise, a second was later identified, after obtaining new deep CCD data, with a very faint dwarf galaxy. The third object has subsequently been associated with debris stripped from NGC4388 (Oosterloo and van Gorkom, 2005). The fourth object, VIRGOHI21, remains as a strong candidate for a dark galaxy (Minchin et al. 2005, Minchin et al. 2006). VIRGOHI21 is extended over a region of about 10 kpc and has a HI velocity width of about 200 km s$^{-1}$ - according to the HI data it is a typical galaxy, but there is no optical counterpart down to very low surface brightness limits.  In this paper we first of all want to re-address the issue of whether dark galaxies can exist within our current models of galaxy formation and particularly examine the recent contention by Taylor and Webster (2005) that almost all galaxies containing HI will form stars. Secondly we want to see if a dark galaxy with the characteristics of VIRGOHI21 is a possiblity. Thirdly we want to discuss whether the existence of dark galaxies is compatable with previous blind HI surveys. Verde et al. (2002) have also produced models very similar to those discussed in this paper and predict the existence of dark galaxies, we concentrate on the detectability of these galaxies by 21cm surveys. ",
        "conclusions": "In this paper we have described a model of the galaxy population of the Universe. We have used this model to firstly infer that dark galaxies are a possible outcome of the galaxy formation process. Secondly, we have shown that the simulations predict the existence of dark galaxies including galaxies with similar properties to the candidate dark galaxy VIRGOHI21. Thirdly, we have described how previous HI surveys have placed few if any constraints on the types of dark galaxies produced by our simulation. None of these surveys have been able to rule out the existence of dark galaxies because: \\begin{enumerate} \\item Dark galaxies are rare compared to galaxies with stars. \\item Dark galaxies can have the same ranges of size and mass as do optical galaxies, but a larger fraction of them are small and of low mass. This makes them difficult to find. \\item Previous surveys have not had sufficient sensitivity to low mass objects or good enough  velocity resolution to identify the majority of dark galaxies. \\item If dark galaxies cluster like galaxies with stars then it is easy to confuse a dark galaxy detection with an incorrect optical source in the same neighbourhood. \\end{enumerate}  Future surveys using the Arecibo multi-beam instrument ALFA should reveal many hundreds of dark galaxies - if they exist."
    },
    "0609/astro-ph0609437_arXiv.txt": {
        "abstract": "Results are presented of early X-ray afterglow observations of GRB~060105 by {\\it Swift} and {\\it Suzaku}. The bright, long gamma-ray burst GRB~060105 triggered the {\\it Swift} Burst Alert Telescope (BAT) at 06:49:28 on 5 January 2006. The {\\it Swift} X-Ray Telescope (XRT) started pointed observations 87 s after the trigger ($T_0$).  The {\\it Suzaku} team commenced a pre-planned target of opportunity observation at 19 ks (5.3 hr) after the {\\it Swift} trigger as the first{\\it Suzaku} attempt at a rapid response to a new gamma-ray burst. The X-ray flux faded during the observations from $6.8 \\times 10^{-9}$ erg~s$^{-1}$cm$^{-2}$ (at $T_0 + 87$ s with the {\\it Swift}/XRT) to $1.5 \\times 10^{-13}$ erg~s$^{-1}$cm$^{-2}$ (at $T_0 + 94$--$101$ ks with the {\\it Suzaku} X-ray Imaging Spectrometer (XIS)) in the 2--10 keV energy band. The afterglow exhibited four phases of decay consisting of different decay indices during the {\\it Swift}/XRT and {\\it Suzaku}/XIS observations. The inferred flux and spectral shapes with both instruments are fully consistent within the statistics, although there are hardening in the first shallow decay phase and softening in very steep decay phases in the end of the observations. Following the prompt emission and successive very steep decay, a shallow decay was observed from $T_0+187$ s to $T_0+1287$ s. After an observation gap during $T_0 +(1.5$--$3)$ ks, an extremely early steep decay was observed in $T_0+(4$--$30)$ ks. The lightcurve flattened again at $T_0+30$ ks, and another steep decay followed from $T_0+50$ ks to the end of observations. Both steep decays exhibited decay indices $\\alpha \\sim$ 2.3 -- 2.4. This lightcurve behavior can be explained as a steep decay resulting from side expansion of a jet, while the two flattening portions suggest energy injection and refreshed shock passage, respectively. This very early break, if it is a jet break, is the earliest case among X-ray afterglow observations, suggesting a very narrow jet whose opening angle is well below $1^\\circ$. The unique {\\it Suzaku}/XIS data allow us to set very tight upper limits on line emission or absorption in this GRB. For the reported pseudo-redshift of $z=4.0\\pm1.3$ the upper limit on the iron line equivalent width is 50 eV. ",
        "introduction": "\\label{sec:intro} Gamma-ray bursts (GRB) are thought to be caused by a sudden release of a large amount of energy ($\\sim 10^{51}$ erg for long GRBs), on the scale of a single star. The prompt gamma-ray emission is attributed to internal shocks in a highly relativistic outflow (\\cite{rees94}; \\cite{sari97}) with a Lorentz factor of over 100 (see \\cite{piran99}, and references therein). After this early phase, relativistic ejecta sweep up a sufficient amount of external medium and are decelerated to cause external shocks. A highly relativistic forward shock propagates into the circumstellar medium to produce the afterglow which can be observed across a wide energy band.  Before {\\it Swift} (\\cite{gehrels04}), X-ray afterglow observations typically started several hours after the burst onset, and they showed a smooth single power-law decay $\\sim t^{-1}$. In contrast, optical afterglow lightcurves often showed an achromatic steepening to $\\sim t^{-2}$, which is attributed to the sideways expansion of a narrow jet (\\cite{rhoads99}; \\cite{sph99}). The advent of {\\it Swift} has opened the earlier time window from $10^2$ -- $10^4$ s with the onboard X-Ray Telescope (XRT; \\cite{burrows05a}). The {\\it early} --- the first few hours --- afterglow evolution is a probe not only of the early radiative energy losses from the external shock and energy injection, but also of the chemical abundance and density profile of the circumburst medium.  Because {\\it Swift} observes X-ray afterglows much more rapidly and so has revealed a fundamental discrepancy from the simple power-law decay (e.g. \\cite{nousek06}). These data show the limitations of a simple ``fire-ball model'' assuming a delta function-like ``sudden'' energy injection with self-similar evolution, and implies hidden physical processes. Zhang et al. (2006) interpreted the Nousek et al. (2006) canonical afterglow lightcurve as having four phases: (i) very steep decay ($F_\\nu (t) \\propto t^{-\\alpha}; \\alpha \\sim 3- 5$); (ii) (very) shallow decay ($\\alpha \\sim 0.5$); and (iii) somewhat steeper (or normal) decay ($\\alpha \\sim 1)$, where $F_\\nu (t)$, $t$, and $\\alpha$ are the X-ray flux, time from the burst, and decay index, respectively. In addition to those, (iv) jetlike (steep) decay ($\\alpha \\sim 2$); and temporal X-ray flares are observed in many GRBs (e.g. \\cite{burrows05b}).  Each transition suggests an evolution of the physical state, such as a start or end of energy injection, re-acceleration by a reverse or forward shock, deceleration of a structured or uniform jet, and/or spatial distribution of physical parameters. However, no significant spectral transition has been observed. Precise spectroscopy, utilizing a high sensitivity X-ray telescope, is the key to determining any spectral evolution at the phase transitions, which will enable us to understand what physics controls phase transitions in GRB jets. Observations with high sensitivity and high energy resolution also examine the question of line features reported by some authors, but which have not been confirmed by statistically strong detections (\\cite{piro99}; \\cite{yoshida01}; \\cite{reeves02}; \\cite{butler03}; \\cite{watson03}; and \\cite{sako05}). X-ray spectral data also provide limits on the environment in the very near vicinity around the gamma-ray burst and its progenitor.  We utilized the fifth Japanese X-ray observatory, {\\it Suzaku} (\\cite{suzaku06}), to conduct a follow-up observation of GRB 060105. Aiming to observe the X-ray afterglow in the fastest possible time, we organized a Target of Opportunity (ToO) team to watch and respond quickly to {\\it Swift} notices with {\\it Suzaku} during the initial performance verification phase --- the Science Working Group time. ",
        "conclusions": "We observed the early X-ray afterglow from the GRB~060105 with {\\it Swift} and {\\it Suzaku}. The X-ray lightcurve obtained with {\\it Swift}/XRT and {\\it Suzaku}/XIS exhibited an early steep decay from 0.04 day and lasted for 0.3 day. After a flattening, the lightcurve shows a steep decay again, and the slope coincides with what we saw in the first steep decay. The X-ray spectral shape, energy index and absorption column are consistent with being constant during the observations, although it infers gradual flattening and steepening through the shallow decay to successive steep decay phase. The decay indices for the two steep decays suggest a jetlike decay, and phase transitions without spectral changes suggest early jet break, which, however, requires a very narrow jet and the GRB to be an outlier of the Ghirlanda relation. We also present a very strict upper limit on the possible atomic spectral lines with the precise spectroscopy with the {\\it Suzaku}/XIS and the {\\it Swift}/XRT."
    },
    "0609/astro-ph0609351_arXiv.txt": {
        "abstract": "A new method is presented for the construction of a natural continuous wavelet transform on the sphere. It incorporates the analysis and synthesis with the same wavelet and the definition of translations and dilations on the sphere through the spherical harmonic coefficients. We construct a couple of wavelets as an extension of the flat {\\it Mexican Hat Wavelet} to the sphere and we apply them to the detection of sources on the sphere. We remark that no projections are used with this methodology. ",
        "introduction": "Multiscaling analysis techniques dealing with the analysis/synthesis of nD-images defined on intervals of $R^n$ have been applied in many fields of physics in the last 15 years.  For instance, in the case $n = 1$ one has electronics and audio signals, in the case $n = 2$ one has optical or infrared images whereas for $n = 3$ one deals with fluid dynamics or the large-scale structure of the universe as 3D-images. However, there are data given on other manifolds like the circle $S_1$ (e. g. scanning along circles the microwave sky) and the sphere $S_2$ (e. g. geophysics). In this paper, we are interested in data distributed on the sphere. Trivially, for the study of local properties (e. g. detection of objects) one can project on the tangent plane at any point on the sphere to make this type of analysis but when global properties are taken into account the curvature of the sphere can not be neglected.  A first approach to deal with these global properties is to make some global projection of all the points of the sphere. The stereographic projection has been recently used dealing with the continuous wavelet transform. In this case, to get the wavelet coefficient at any point on the sphere, one projects from the opposite point to the local tangent plane. [1] have made a connection to group theory. The translations and dilations in the wavelet have their definition on the plane. Clearly, such a projection does not take into account the topological structure of the sphere. Some applications to cosmology, in particular the study of anisotropies of the cosmic microwave background radiation have been done by some authors ([3],[7],[10]) using the projection of the {\\it Mexican hat wavelet}. A drawback of such projection is the obvious deformation of the pixels and wavelets near the projection pole. We remark that the synthesis can be done in terms of another biorthogonal wavelet [11].  Another approach uses some analyzing wavelet functions that are defined in terms of spherical harmonics [5] with a definition of the dilation operator and conditions on the wavelets in such a way to get a synthesis formula. The drawback of such methodology is: the dilations do not satisfy the appropriate flat limit in general. Also some examples of wavelet functions are poorly localized (e. g. Abel-Poison wavelets).  A different approach assumes from the beginning discrete wavelets incorporating tensor product approaches in polar coordinates, then the two poles are singular points regarding approximation/stability properties ([4],[6]). Another approach is adapted to arbitrary point systems or triangulations on the spheres, then there is no efficient tool as fast wavelet algorithms. In the approach by [8] basis are defined on a quasi-uniform icosahedral triangulation on the sphere allowing for a fast algorithm. However, biorthogonal wavelets are needed and a lifting scheme for the multiresolution is applied avoiding the concepts of translations and dilations and also it is not clear whether the construction leads to a stable $L_2 (S_2)$ basis. Haar-type wavelets have been developed using different pixel combinations ([2],[7],[9]).  The first case uses the lifting scheme weighting for the area of the pixels whereas in the other two cases an equal area pixelization is used but the Haar-type transform is only applied on regions of the sphere covering only $\\frac{1}{12}$ of the total area. Clearly, with any pixelization the symmetry on the sphere is lost.  In this paper we will consider a continuous approach, we will introduce a methodology that incorporates the analysis and synthesis of any function defined on the sphere $S_2$ using the same circularly-symmetric wavelet and also we will introduce the generalization of the translations/dilations. In this sense we follow Freeden's approach working with spherical harmonics. Examples will be given that have the appropriate flat limit. Finally, the application to the detection of a spot is given, studying the concentration of the wavelet coefficients. ",
        "conclusions": "We have developed a constructive wavelet approach on the sphere without any projection from the plane.  It is a continuous transform that allows the analysis and synthesis of any function defined on the sphere and incorporates the concepts of translation and dilation as generalizations of the elementary ones defined on the plane. It is a compensated filter that conserves the energy of any function.  We have considered some natural generalizations of the plane Mexican hat wavelet and we have applied them to the detection of a big spot. The conclusion is that one of the wavelets ($\\psi_1$) concentrates more the information than the other one."
    },
    "0609/astro-ph0609021_arXiv.txt": {
        "abstract": "% This work focuses on the properties of dusty tori in active galactic nuclei (AGN) derived from the comparison of SDSS type 1 quasars with mid-Infrared (MIR) counterparts and a new, detailed torus model. The infrared data were taken by the Spitzer Wide-area InfraRed Extragalactic (SWIRE) Survey. Basic model parameters are constraint, such as the density law of the graphite and silicate grains, the torus size and its opening angle. A whole variety of optical depths is supported. The favoured models are those with decreasing density with distance from the centre, while there is no clear tendency as to the covering factor, i.e. small, medium and large covering factors are almost equally distributed. Based on the models that better describe the observed SEDs, properties such as the accretion luminosity, the mass of dust, the inner to outer radius ratio and the hydrogen column density are computed. The properties of the tori, as derived fitting the observed SEDs, are independent of the redshift, once observational biases are taken into account. ",
        "introduction": "%  A detailed description of the torus model is given in \\cite{fritz06}; here we briefly present the main characteristics. The adopted torus geometry is a {\\it flared disk}, i.e. a sphere with the polar cones removed. Its size is defined by its outer radius, R$_{\\rm out}$, and its opening angle. The dust components that dominate both the absorption and the emission of radiation are graphite and silicate. The inner radius, R$_{\\rm in}$, depends both on the sublimation temperature of the dust grains (1500 and 1000 K, for graphite and silicate, respectively) and on the accretion luminosity. Since silicate grains have a lower sublimation temperature than graphite grains, the innermost regions of the torus only consist of graphite. The adopted absorption and scattering coefficients are those by \\cite{laor93} for dust grains of different dimensions, weighted with the standard MRN distribution \\citep{mathis77}. The gas density within the torus is modeled in such a way to allow a gradient along both the radial and the angular coordinates.  The central source is assumed to be point-like and its emission isotropic. Its spectral energy distribution is defined by means of a composition of power laws with different values for the spectra index in the UV, optical and IR. The $\\Lambda$-iteration method is used to solve the radiative transfer equation. A geometrical grid is defined along the three spatial coordinates, and the main physical quantities (dust density and temperatures, electromagnetic emission, optical depth, etc.) are computed with respect to the center of the volume elements defined by the grid. The global SED is computed at different angles of the line-of-sight with respect to the torus' equatorial plane and includes three contributions: emission from the AGN, thermal emission and scattering emission by dust in each volume element. An example of an emitted spectrum is given in Fig. \\ref{figMod}.  \\begin{figure}[!ht] \\includegraphics[angle=-90, width=11cm]{hatziminaoglou_e_fig1.ps} \\caption{Emission spectra as a function of wavelength for 10 different lines-of-sight inclinations from ${\\Psi} = 0^\\circ$ (edge-on; lower curve) to ${\\Psi} = 90^\\circ$ (face-on, upper curve) at regular steps of $10^\\circ$, for a geometrical configuration with $\\tau(9.7)=10.0$, R$_{\\rm out}$/R$_{\\rm in}=300$, torus opening angle $\\Theta=140^\\circ$ and density law $\\rho\\left(r,\\theta \\right) \\propto r^{-1} e^{-6 |cos(\\theta)|}$, where $r$ is the distance from the centre and $\\theta$ the angle from the equatorial plane.} \\label{figMod} \\end{figure} ",
        "conclusions": "This work in progress provides useful insights into the properties of dusty tori around AGN. Fitting the observed optical-to-MIR SEDs to a series of models, we showed that in most of the cases the density decreases with distance from the centre and can vary with the angle from the equator. Tori properties such as the outer radii or the mass of dust are independent of the redshift and only depend on the luminosity of the objects. The {\\it equatorial}  hydrogen column density distribution is in agreement with what is expected from the Unified Scheme."
    },
    "0609/astro-ph0609810_arXiv.txt": {
        "abstract": "{Summary of Joint Discussion 13 of the 26th IAU General Assembly, Prague.} ",
        "introduction": " ",
        "conclusions": "The next decade will see tremendous progress in our understanding of the formation of the Milky Way and the Local Group. Homogeneous large scale photometric, radial velocity and astrometric surveys, dedicated high-resolution spectroscopic follow-up, and theoretical preparation are critical. The talks and the discussion showed that the complementarity of surveys from ground and from space can probably be exploited further with the Gaia mission as the centerpiece, and that plans for this are underway. There is much to look forward to."
    },
    "0609/astro-ph0609217_arXiv.txt": {
        "abstract": "We report the diffuse X-ray emissions from the Sgr A and B regions observed with $Suzaku$.  From the Sgr A region, we found many K-shell transition lines of iron and nickel. The brightest are K$\\alpha$ lines from FeI, FeXXV and FeXXVI at 6.4~keV, 6.7~keV and 6.9~keV.  In addition, K$\\alpha$ lines of NiI and NiXXVII, K$\\beta$ of FeI, FeXXV and FeXXVI, and K$\\gamma$ of FeXXV and FeXXVI are detected for the first time. The center energy of K$\\alpha$ of FeXXV favors collisional excitation as the origin for this line emission. The ionization temperature determined from the flux ratio of K$\\alpha$ of FeXXV and FeXXVI is similar to the electron temperature determined from the flux ratio of K$\\alpha$ and K$\\beta$ of FeXXV, which are in the range of 5$-$7 keV. Consequently, the Galactic Center diffuse X-rays (GCDX) are consistent with emission from a plasma nearly in ionization equilibrium.  The radio complex Sgr B region also exhibits K$\\alpha$ lines of FeI, FeXXV and FeXXVI.  The 6.7~keV line (FeXXV) map exhibits a local excess at $(l,~~b) = (0^\\circ.612,~~ 0^\\circ.01)$, and could be a new young SNR. The 6.4~keV image is clumpy with local excesses near Sgr B2 and at $(l,~~b) = (0^\\circ.74, ~~-0^\\circ.09)$. Like Sgr B2, this latter excess may be another X-ray reflection Nebulae (XRN). ",
        "introduction": "The $Astro E2$ satellite was launched on 10th July in 2005 and was named $Suzaku$ \\cite{1} after an imaginary red bird in the ancient oriental mythology.  The X-ray Imaging Spectrometer (XIS) \\cite{2} is one of the major instruments on $Suzaku$. The XIS consists of four sets of X-ray CCD cameras, placed on the focal planes of the X-ray Telescope (XRT) \\cite{3}. Three XIS sensors contain front-illuminated (FI) CCDs, and the other has a back-illuminated (BI) CCD \\cite{2}.  The other focal plane instrument is the micro-calorimeter (XRS) \\cite{4}, which had extremely high energy resolution. The XRS was cooled down to 60 mK, and its overall energy resolution was originally 7 eV at 5.9 keV. These temperature and energy resolution were on-orbit world records. The XRS was able to resolve, e. g., the iron He-like triplets, and hence should have opened a new window of X-ray astronomy. However, regretfully, all the liquid He evaporated very quickly.  We nevertheless still have the XIS, which have the best performances among space CCD cameras. The background level for diffuse sources, in particular at the energy of 5--10 keV range is nearly one order of magnitude lower than $Chandra$ and $XMM$. The effective area and energy resolution are also comparable to or better than $XMM$. Therefore the XIS on $Suzaku$ is best suited for observations of the diffuse X-rays from the Galactic Center (GC). Early $Suzaku$ results will appear soon in the refereed journal (\\cite{5}, \\cite{6}). This paper reports the summary and essence of these papers. ",
        "conclusions": ""
    },
    "0609/astro-ph0609729_arXiv.txt": {
        "abstract": "It has been suggested that Narrow-Line Seyfert 1 (NLS1) galaxies are evolutionarily young objects, powered by the accretion of gas onto central black holes that are significantly lower in mass than those found in typical broad-line Seyferts.  We explore this hypothesis through the analysis of high-spatial resolution, near-IR imaging data obtained in {\\it J} and {\\it $K^{\\prime}$} for a sample of 11 NLS1s. Surface brightness profiles are separated into their constituent components using two-dimensional decomposition techniques.  By employing the correlation between black-hole mass and host galaxy bulge luminosity, calibrated for near-IR wavelengths using 2MASS data, we determine the mean black-hole mass for our sample to be, in solar units, $\\langle{\\rm log}(\\MBH)\\rangle$ = 7.9.  Using the correlation between the size of the broad-line region and the monochromatic continuum luminosity, we obtain black-hole mass estimates under the assumption that the emission-line gas is in virial equilibrium.  The mean black-hole mass derived from this relation is $\\langle{\\rm log}(\\MBH)\\rangle$ = 6.4.  It is found that the estimates obtained from the black-hole mass-bulge luminosity relation are systematically one full order of magnitude larger than those derived from the black-hole mass-broad-line region radius relation.  We explore possible causes for this discrepancy in $\\MBH$ estimates and the ramifications for our understanding of the role played by NLS1s in AGN evolution.  Because numerical simulations constrain the start of the AGN duty cycle to a time shortly after a significant gravitational interaction, we examine the morphology and near-IR bulge colors of the NLS1 sample for evidence of recent encounters.  The mean bulge color is found to be $\\langle({\\it J}-{\\it K_s})\\rangle$ = +1.85 $\\pm$ 0.58, which is redder than that of both a matched sample of non-active galaxies and published estimates for broad-line Seyferts.  The source of the unusual bulge colors may be an excess of flux, peaking at around 2.2~\\micron, that has been detected near the centers of some NLS1s such as Mrk 1239.  No evidence is found for light asymmetries or an extra stellar component that would indicate NLS1s are young objects. Finally, we postulate that there may be some interesting lines of circumstantial evidence suggesting that secular processes may be relevant in the evolution of NLS1s. ",
        "introduction": "}  The connection between the evolution of galaxies and that of the supermassive black holes that are known to lurk at their centers is a key issue in contemporary extragalactic studies.  It is clear that Active Galactic Nuclei (AGN) represent an important, transient phase in the development of galaxies at all redshifts \\citep{hae93}.  The source of intense, non-thermal emission generated in the cores of AGN is widely accepted to be the accretion of material onto otherwise dormant black holes \\citep{ree84, pet00a}.  However we may not yet fully appreciate the diversity of mechanisms that create the non-equilibrium conditions leading to accretion.  Significant changes in gravitational potential can be produced by interactions between galaxies of comparable mass or by the consumption of dwarf satellite galaxies \\citep{ost93}.  It also may be possible to trigger accretion through secular processes within the host galaxy such as bar-driven tidal instabilities, generating episodic bursts of nuclear activity \\citep{kor04}.  An evolutionary link between black holes and their host galaxies is strongly suggested on the basis of correlations between the black-hole mass, $\\MBH$, and various galaxy parameters.  A good correlation has been found between $\\MBH$ and the central stellar velocity dispersion, $\\sigma_s$, of the spheroid component of the host galaxy \\citep[e.g.,][]{fer00, geb00, gre05a}.  The intrinsic scatter in this relation may be more significant at low black-hole masses, however \\citep{ben05}.  It has also been shown that $\\MBH$ is a constant fraction of the mass of the host spheroid (or, in the case of ellipticals, the total mass of the galaxy), with $\\MBH$/$M_{sph} \\simeq$ 0.0015 \\citep{mag98, mar03, har04}.  Under the assumption of a constant mass-to-light ({\\it M}/{\\it L}) ratio, this relation is frequently expressed in terms of the host bulge luminosity \\citep[e.g.][]{lao01, mcl01}.  While early studies found the scatter in this mass-luminosity relation to be substantial, recent progress in the amount and quality of available data has led to an improvement in the precision of the best-fit parameters \\citep{mar03, don06}. \\citet{mar03} also demonstrated the wavelength dependence of the scatter in the relation, with near-IR observations tracing the mass distribution better than optical data, generating a tighter correlation.  There is a general consensus in the literature that AGN exhibit these correlations independent of distance or luminosity, suggesting that they are a natural product of galaxy-black hole co-evolution \\citep{kor95, bar05}.  Yet there remains much debate as to how Narrow-Line Seyfert 1 galaxies, which constitute approximately 15\\% of the total low-redshift Seyfert population \\citep{wil02}, fit into this picture.  NLS1s are a subset of AGN exhibiting an overall spectral signature consistent with a Type 1 Seyfert, but with remarkably narrow permitted emission lines.  More specifically, NLS1s are classified based on an observed broad component of H$\\beta$ emission with FWHM $<$ 2000 km~s$^{-1}$, strong $\\Fe$ emission, relatively weak [$\\Ox$], and steeper soft x-ray continuum slopes than broad-line Seyfert 1s \\citep{ost85, bol96}.  Through a Principal Component analysis, \\citet{bor92} found that NLS1s dominate the extreme end of the ``Eigenvector 1'' anti-correlation between the $\\Fe$ and [$\\Ox$] emission-line strengths.  This relation is believed to be driven by the ratio of the luminosity of the source relative to its Eddington luminosity, $L/L_{Edd}$ \\citep{sul00}.  Clearly, NLS1s occupy an important portion of AGN parameter space.  To account for the extreme properties of NLS1 galaxies, it has been conjectured that their central engines are black holes of relatively modest mass (typically of order 10$^6$ {\\it M}$_{\\sun}$) \\citep{bol96, mat00}.  This hypothesis can explain the narrow emission lines observed in these objects, under the assumption that the broad H$\\beta$-emitting region is gravitationally dominated by the black hole.  If the black holes in NLS1s are statistically smaller than in broad-line Seyfert galaxies, elevated accretion rates would be required to account for the relatively normal observed luminosities \\citep{pou95}.  Furthermore, it has been suggested that these objects represent an early stage in AGN evolution, during which the black-hole mass increases rapidly in a period of intense accretion \\citep{mat00}. As the black hole grows in this scenario, the accretion rate gradually decreases and the NLS1 matures into a ``normal'' broad-line Seyfert \\citep{gru04a}.  It remains unclear whether the aforementioned correlations with $\\MBH$ can be applied to NLS1s.  Several authors suggest that NLS1s fall below the mass-luminosity correlation of other galaxies, resulting in smaller estimates of $\\MBH$ for a given bulge luminosity \\citep{mat01,wan02}.  It has also been suggested that intense star formation activity could lead to overestimates of the bulge luminosity or, equivalently, underestimates of the ({\\it M}/{\\it L}) ratio \\citep{nel04}.  The validity of the $\\MBH$-$\\sigma_s$ correlation for NLS1s has also been questioned.  It has been demonstrated by \\citet{bar05} that AGN with small black holes follow the same relation as galaxies with more typical values for $\\MBH$.  Both \\citet{wan01} and \\citet{bot04} estimate the velocity dispersion for NLS1 galaxy samples from measurements of [$\\Ox$], and find that the $\\MBH$-$\\sigma_s$ correlation is the same as for broad-line Seyferts. Conversely, a departure from the standard $\\MBH$-$\\sigma_s$ relation for NLS1s was reported by \\citet{mat01}, \\citet{gru04a}, and \\citet{bia04}.  Each of these studies used [$\\Ox$] as a surrogate for the velocity dispersion.  It is interesting to note that different conclusions are reached in these studies even though the samples are similar.  A major thrust of this study is to compare black-hole masses derived from the well-known correlation between $\\MBH$ and the host galaxy luminosity for the NLS1 subset, hereafter the mass-luminosity correlation, with other mass estimates.  To this end, we examine a sample of 11 NLS1s using high-spatial resolution, near-IR imaging data.  The luminosities of the bulge components in {\\it J} and {\\it $K^{\\prime}$} are measured independently for each target.  Black-hole masses are obtained by utilizing the mass-luminosity relation, calibrated to near-IR wavelengths.  This work builds on previous studies by applying two-dimensional profile decomposition techniques to near-IR data of excellent quality, providing an accurate determination of the photometric properties of the host bulges.  The results are compared to values of $\\MBH$ derived via the correlation between the radius of the broad-line region and the optical continuum luminosity, which has been calibrated using reverberation mapping mass estimates \\citep{kas00, kas05}.  This comparison is used to examine the role of NLS1s in the evolution of AGN.  In \\S~2 we outline the sample selection and describe the observations. The data reduction process, including the decomposition procedure, is included in \\S~3.  The results and a discussion of their implications for the NLS1 paradigm are presented in \\S~4 and \\S~5, respectively. Conclusions are provided in \\S~6.  Throughout this work, we adopt cosmological parameters of $\\Omega_{\\Lambda}$ = 0.7, $\\Omega_m$ = 0.3 and a Hubble constant of H$_{\\rm o}$ = 75 km~s$^{-1}$~Mpc$^{-1}$. ",
        "conclusions": "}  We have studied a sample of NLS1 galaxies in order to explore the hypothesis that such objects are evolutionarily young AGN possessing relatively low-mass black holes.  Galaxy profiles were separated into their constituent components by applying the two-dimensional decomposition algorithm GALFIT to high-spatial resolution, near-IR imaging data.  Black-hole masses were derived using the correlation between $\\MBH$ and host bulge luminosity, calibrated in {\\it J} and {\\it $K^{\\prime}$} using a sample of bright elliptical galaxies with black-hole masses obtained by resolving the sphere of influence.  The scatter in the mass-luminosity relation was thus minimized by restricting the calibration sample to objects with secure $\\MBH$ estimates.  Furthermore, the near-IR form of the relation is less sensitive to attenuation by dust than at optical wavelengths.  The unweighted mean black-hole mass determined from this relation is, in solar units, $\\langle{\\rm log}(\\MBH)\\rangle$ = 7.9, in line with typical values for broad-line Seyferts.  We also derived masses for seven of the objects in our sample from the correlation between the radius of the broad-line region and the monochromatic continuum luminosity at 5100 \\AA~under the assumption that the emission-line gas is virialized and using optical spectra taken from the literature. The mean mass calculated from this relation is $\\langle{\\rm log}(\\MBH)\\rangle$ = 6.4, more consistent with the premise that the black holes in NLS1s are less massive than those found in their broad-line counterparts.  It is clear that there is a significant discrepancy between the results obtained from these correlations, with the mass-luminosity relation yielding values of $\\MBH$ that are on average more than one order of magnitude larger than those obtained from the mass-radius relation.  Formal uncertainties in the calculated masses are typically around 0.25 dex, and thus are insufficient to account for the conflicting results.  Our findings suggest that (at least) one of the correlations cannot be extended to the NLS1 subset in the same form as used for other galaxies.  It has been conjectured that NLS1s may not follow the standard mass-luminosity relation due to contamination by circumnuclear star formation.  We have argued that this extra component would not significantly affect our data.  Furthermore, we noted that the resolution of our imaging data provided excellent spatial over-sampling of the surface brightness profiles, producing highly accurate bulge luminosities.  The resulting $\\MBH$ estimates were calculated by interpolating the relation within a region known to be valid for both non-active galaxies and other AGN.  Numerical simulations of galaxy interactions constrain the high-accretion rate period that may define the start of the AGN duty cycle to a period shortly after a significant gravitational perturbation of the host galaxy.  To further investigate this possibility, we compared the NLS1 bulge ({\\it J}-{\\it K$_s$}) color indices to a matched sample of non-active galaxies.  We determined the mean bulge color for the NLS1 sample to be $\\langle({\\it J}-{\\it K_s})\\rangle$ = +1.85 $\\pm$ 0.58.  For the non-active galaxy sample, the mean color is $\\langle({\\it J}-{\\it K_s})\\rangle$ = +1.10 $\\pm$ 0.04.  Estimates of the bulge colors for broad-line Seyfert 1 galaxies were taken from the literature.  A fair comparison to our sample is complicated by the large scatter in the published results and the potential for misidentification of NLS1s as Seyfert 2 galaxies.  The NLS1 images were also examined for light asymmetries or tidal features that may reflect a recent gravitational interaction.  We find no evidence for any significant morphological asymmetries or extra stellar components that would suggest these galaxies are evolutionarily young.  We have, however, identified several lines of circumstantial evidence suggesting that nuclear fueling in NLS1s may be strongly influenced by secular processes within the host galaxy.  The bulge light profiles obtained from our two-dimensional decompositions are all nearly exponential in shape, a key characteristic of secular evolution. NLS1s may not follow the standard $\\MBH$-$\\sigma_s$ relation, which has been shown to be a natural product of galaxy mergers, while an enhanced fraction of stellar bar structures in NLS1s may provide a natural source for fueling the nucleus without invoking a recent gravitational interaction.  It is thus fair to acknowledge that we do not yet completely understand the evolution of this important subset of low-redshift active galaxies."
    },
    "0609/astro-ph0609203_arXiv.txt": {
        "abstract": "We present recent measurements of the orbital motion in the binaries, DF Tau and ZZ Tau, and in the triples, Elias 12, T Tau, and V853 Oph.  We observed these systems with the Fine Guidance Sensors on the {\\it Hubble Space Telescope} and with adaptive optics imaging at the W. M. Keck and Gemini North Observatories.  Based on our measurements and those presented in the literature, we perform preliminary orbital analyses for DF Tau, ZZ Tau, Elias 12 Na-Nb, and T Tau Sa-Sb.  Because the orbital coverage in most of these systems does not yet span a sufficient portion of the orbit, we are not able to find definitive orbit solutions.  By using a Monte Carlo search technique, we explored the orbital parameter space allowed by the current set of data available for each binary.  We constructed weighted distributions for the total mass of the binaries derived from a large sample of possible orbits that fit the data.  These mass distributions show that the total mass is already well-defined.  We compute total mass estimates of 0.78$^{+0.25}_{-0.15} M_\\odot$, 0.66$^{+0.15}_{-0.11} M_\\odot$, 1.13$^{+0.36}_{-0.09} M_\\odot$, and 4.13$^{+1.58}_{-0.97} M_\\odot$ for DF Tau, ZZ Tau, Elias 12 Na-Nb, and T Tau Sa-Sb respectively, using a distance of 140 pc.  For Elias 12 Na-Nb, where the orbital coverage spans $\\sim 164^\\circ$, we compute a preliminary orbit solution with a period of $\\sim$ 9-12 years.  By including an earlier lunar occultation measurement, we also find a likely orbit solution for ZZ Tau, with a period of $\\sim$~32~years.   With additional measurements to continue mapping the orbits, the derived dynamical masses will be useful in constraining the theoretical tracks of pre-main sequence evolution. ",
        "introduction": "Determining masses and ages of pre-main sequence (PMS) stars is important for understanding the physical processes of star formation.  For clusters of young stars, these parameters are commonly determined from a star's location relative to evolutionary tracks on a color-magnitude diagram.  However, different sets of theoretical tracks yield discrepant results, particularly for low-mass (M $< 1 M_{\\odot}$) young stars, where the predicted masses and ages can differ by factors as large as 2-3 \\citep{simon01,simon05}.  These discrepancies result mostly from differences in the choice of atmospheric opacities and the treatment of interior convection.  Calibration of theoretical evolutionary tracks requires comparison with high precision dynamical mass measurements.  Methods used to measure dynamical masses of pre-main sequence stars include: (1) observations of eclipsing double-lined spectroscopic binaries, where analysis of the radial velocity variations and the eclipse light curve yield the individual masses and radii of the component stars \\citep{popper87,casey98,covino00,alencar03,stassun04}; (2) mapping the Keplerian rotation of circumstellar disks to determine the mass of the central star \\citep{simon00,dutrey03}; and (3) combining spectroscopic and spatially-resolved observations of binary star orbits to determine the component stellar masses and distance to the system \\citep{steffen01,boden05}.  Currently, there are only a small number of low-mass PMS stars accessible to each of these techniques.  Observing eclipsing binaries, for instance, requires the orbit be inclined by $\\sim 90^{\\circ}$ along our line of sight, whereas single stars with rotating circumstellar disks must have disks large enough to detect and resolve with millimeter-wave interferometers.  Moreover, visual and spectroscopic detections of young binaries tend to sample systems distributed in different regions of orbital parameter space.  PMS spectroscopic binaries typically have periods measured on the order of a few years or less \\citep{mathieu94}, corresponding to radial velocities amplitudes on the order of tens of km~s$^{-1}$.  Meanwhile, instrumental angular resolution limits the detection of visual binaries with similarly short periods.  At the distance of the Taurus and Ophiuchus star forming regions (140 pc), the resolutions achieved through speckle interferometry and adaptive optics imaging at large 8-10 meter class telescopes or by the {\\it Hubble Space Telescope} ({\\it HST}) are capable of resolving visual binaries with periods on the order of decades.  The current generation of long-baseline optical/IR interferometers brings visual binary measurements \\citep{boden05,schaefer05} further into the orbital parameter space of PMS spectroscopic binaries.  Through a statistical comparison based on a sample of PMS dynamical mass measurements, only 9 of which are below 1~$M_{\\odot}$, \\citet{hillenbrand04} showed that for masses above $\\sim 1.2 M_{\\odot}$, the evolutionary models reproduce the measured masses within 1$\\sigma$ on average.  However, for masses below $1.2 M_{\\odot}$, most of the evolutionary models systematically underpredict stellar masses by 10-30\\%, with the models of \\citet{baraffe98} providing the most consistent results.  The goal of the work presented in this paper is to increase the number of low-mass PMS stars with dynamical mass measurements, thereby strengthening the tests of the evolutionary tracks.  In this paper, we present results from our ongoing program \\citep{simon96,schaefer03} to map visual orbits of low-mass PMS binaries using the Fine Guidance Sensors on {\\it HST} and adaptive optics imaging at the W. M. Keck and the Gemini North Observatories.  We present recent measurements of the separation of the young multiple systems DF Tau, ZZ Tau, Elias 12, T Tau, and V853 Oph.  For ZZ Tau and the close binary (Elias 12 Na-Nb) in the Elias 12 triple, we present preliminary orbit solutions.  For the remainder of the systems, it is not possible to determine a reliable set of orbital elements because our measurements do not yet span a significant fraction of the orbital period.  However, even with the limited orbital coverage, useful mass estimates can be determined based on the observed curvature in the measured orbit.  In this paper, we refine the statistical technique developed in \\citet{schaefer03} for determining the total mass of a visual binary when less than half of the orbit is mapped.  We have improved the statistical analysis by incorporating a Monte Carlo search to sample the orbital parameter space and by weighting each orbital solution by its $\\chi^2$ probability when analyzing the distributions of masses produced by the search.  We present our analysis of the orbital solutions in \\S~\\ref{sect.orbit} and discuss the dynamical masses computed from these results in \\S~\\ref{sect.mass}.  We conclude with a comparison of our dynamical mass estimates with predictions from evolutionary tracks and a brief discussion on the variability observed in the multiple systems. ",
        "conclusions": "\\subsection{Comments on Component Masses}  DF Tau is the only binary in our sample for which there is a spectral type and luminosity measured for both the primary and secondary components.  Using the evolutionary tracks computed by \\citet[DM97]{dantona97}, \\citet[BCAH98]{baraffe98}, \\citet[PS99]{palla99}, and \\citet[SDF00]{siess00} and the luminosities and spectral types listed in Table~\\ref{tab.targlist}, we compute the predicted total mass of the DF Tau binary for each set of tracks.  To convert the spectral types to effective temperatures, we use the temperature scale conversion of \\citet{kenyon95}.  The age of the DF Tau binary predicted from the different sets of evolutionary tracks ranges from 0.5 to 2 Myr.  In Figure~\\ref{fig.compmass}, we compare the masses predicted from the theoretical tracks with our dynamical mass estimate of 0.78$^{+0.25}_{-0.15} M_\\odot $ ~measured for DF Tau.  The error bars for the predicted masses are derived from an uncertainty of $\\pm1$~subclass in the spectral type determination.  The predicted masses from each of the tracks agree with our dynamical estimate within $\\sim 1.5\\sigma$.  Based on the current uncertainties, it is premature to make a qualitative assessment of the evolutionary tracks based only on the total mass estimate of DF Tau.  We are working on a program to determine the spectral types and luminosities of the close companions in ZZ Tau and Elias 12 N.  For an age of 1~Myr, the M1 spectral type of T Tau Sb \\citep{duchene02} corresponds to a predicted mass of 0.4-0.7~$M_\\odot$, depending on the choice of theoretical tracks.  Combined with the total dynamical mass measurement of 4.13$^{+1.58}_{-0.97} M_\\odot$ for the T Tau S binary, this range implies a predicted mass of 2.4-5.3~$M_\\odot$ for T Tau Sa.  Thus, T Tau Sa may be a Herbig AeBe star caught at a very early phase of its evolution.  With an estimated mass of 2.1~$M_{\\odot}$ for T Tau N \\citep{white01}, our dynamical results suggest that T Tau Sa is the most massive component in the system.  This is consistent with the photometric and spectroscopic findings in the literature \\citep{koresko97,beck04,duchene05} and the orbit fits based on radio and infrared data \\citep{johnston04, duchene06}.  \\subsection{Variability}  The two recent FGS flux measurements for DF Tau presented in Table~\\ref{tab.fgscounts} are consistent with the results in \\citet{schaefer03}, where we found that the flux of the DF Tau primary varies while the secondary remains constant.  In Figures~\\ref{fig.el12mag} and~\\ref{fig.v853ophmag}, we plot the changes in the FGS counts measured for the components in the Elias 12 triple and for V853 Oph A-B.  As with DF Tau, it appears that the primaries, Elias 12 S and V853 Oph A, show the most variability in the system.  It is interesting that in all three systems, the primaries display the activity associated with accretion from a disk.  In Figure~\\ref{fig.ttmag}, we show the cumulative magnitude variations observed for the components of T Tau, based on our measurements presented in Table~\\ref{tab.ttaumag} and measurements reported in the literature \\citep{koresko00,kohler00,duchene02,furlan03,beck04,duchene05,mayama06}.  These results show that the large variability of T Tau Sa continues in the recent measurements, with variations of $\\sim$ 2 mag taking place over the course of a few months.  \\citet{duchene05} speculate that this variability may be caused by a clumpy, nearly edge-on circumstellar disk rotating around the central star."
    },
    "0609/gr-qc0609072_arXiv.txt": {
        "abstract": "We present new results from our test of Lorentz invariance, which compares two orthogonal cryogenic sapphire microwave oscillators rotating in the lab. We have now acquired over 1 year of data, allowing us to avoid the short data set approximation (less than 1 year) that assumes no cancelation occurs between the $\\tilde{\\kappa}_{e-}$ and $\\tilde{\\kappa}_{o+}$ parameters from the photon sector of the standard model extension. Thus, we are able to place independent limits on all eight $\\tilde{\\kappa}_{e-}$ and $\\tilde{\\kappa}_{o+}$ parameters. Our results represents up to a factor of 10 improvement over previous non rotating measurements (which independently constrained 7 parameters), and is a slight improvement (except for $\\tilde{\\kappa}_{e-}^{ZZ}$) over results from previous rotating experiments that assumed the short data set approximation. Also, an analysis in the Robertson-Mansouri-Sexl framework allows us to place a new limit on the isotropy parameter $P_{MM}=\\delta-\\beta+\\frac{1}{2}$ of $9.4(8.1)\\times10^{-11}$, an improvement of a factor of 2. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609345_arXiv.txt": {
        "abstract": "{It has recently been claimed that analysis of Greenwich sunspot data over 120 years reveals that sunspot activity clusters around two longitudes separated by 180\\degr (``active longitudes'') with clearly defined differential rotation during activity cycles. In previous work we demonstrated that such effects can be observed in synthetic data without such features, as an artefact of the method of analysis.} {In the present work we extend this critical examination of methodology to the actual Greenwich sunspot data and also consider newly proposed methods of analysis claiming to confirm the original identification of active longitudes.} {We performed fits of different kinematic frames onto the actual sunspot data. Firstly, a cell-counting statistic was used to analyse a comoving system of frames and show that such frames extract useful information from the data. Secondly, to check the claim of century-scale persistent active longitudes in a contramoving frame system, we made a comprehensive exploration of parameter space following the original methodology as closely as possible.} {Our analysis revealed that values obtained for the parameters of differential rotation are not stable across different methods of analysis proposed to track persistent active longitudes. Also, despite a very thorough search in parameter space, we were unable to reproduce results claiming to reveal the century-persistent active longitudes. Previous parameter space exploration has been restricted to frames whose latitudinal profile is opposite to solar surface differential rotation. Relaxing this restriction we found that the highest values of nonaxisymmetry occur for frames comoving with the solar surface flow. Further analysis indicates that even these solutions are the result of purely statistical fluctuations.} {We can therefore say that strong and well substantiated evidence for an essential and century-scale persistent nonaxisymmetry in the sunspot distribution does not exist.} ",
        "introduction": "\\label{Section:introduction} \\subsection{Nonaxisymmetry in late-type stars} It is known that solar active regions may exist at the same position (such as recurrent sunspot groups) during several solar rotations, a new active region appearing at another position (e.g. Tuominen \\& Virtanen~\\cite{TV87}, Pulkkinen \\& Tuominen~\\cite{PT98}). The sunspots reflect intense concentrations of magnetic field and thus occur only at certain discrete locations on the solar surface. For related discussion on the observed magnetogram recordings of the full-disk field see e.g. Stenflo~\\cite{Stenflo91}. Consequently, different activity indicators tend to be statistically correlated for a long time spans and a statistical analysis of the corresponding phase distributions must take this into account.  The theory of the global magnetic field (mean-field dynamo theory) has been developed over the last 50 years, starting from a linear theory (Parker~\\cite{Parker55}, Steenbeck \\& Krause~\\cite{SK69}, for a detailed review see Krause and R\\\"adler~\\cite{KR80}), followed by a nonlinear one, where the stability of the solutions determines the symmetry properties of the field (Krause \\& Meinel~\\cite{KM88}, Brandenburg et al.~\\cite{Betal89}), and including a nonaxisymmetric solution as a possibility in the case of rapid rotation (e.g R\\\"adler et al.~\\cite{Retal90}, Moss, Tuominen \\& Brandenburg~\\cite{MTB91}). More recent models including hydrodynamics show that when the rotation velocity increases, the field geometry changes from antisymmetric with respect to the equatorial plane and axisymmmetric with respect to the longitude, to nonaxisymmetric with respect to the longitude (e.g. Tuominen et al.~\\cite{TBKR99}, Tuominen, Berdyugina \\& Korpi~\\cite{Tuominen2002}).  Recently it has been claimed that long-term, i.e. persistent for at least 120 years, nonaxisymmetric structures can be detected in the solar sunspot data using new methods of analysis that take into account the possibility that such structures follow the differential surface rotation. This is an important claim and in the present paper we investigate it in detail because the existence of persistent ``active longitudes'' (nonaxisymmetric structures) in old slowly rotating late-type stars, such as the Sun, has important theoretical implications.  \\subsection{Nonaxisymmetry in solar data} Berdyugina \\& Usoskin~\\cite{BU} (hence BU) analysed the longitudinal distribution of regions of sunspot formation in the Greenwich sunspot data for a period of 120 years. By applying phase corrections, corresponding to the differential rotation of regions of activity formation, they claimed to identify persistent (during 120 years) active longitude belts which are separated by $180\\degr$ and which drift smoothly along longitudes. The essence of their method is to apply a longitude shift correction for each Carrington rotation to represent a different rotation velocity for the active longitudes as the sunspot activity wave moves from high to low latitudes in each solar cycle. As a response to the first paper (BU) Pelt, Tuominen \\& Brooke ~\\cite{PTB} (PTB) demonstrated that large part (or all) of the persistent smooth migration of the active longitudes can be the result of a particular method of data treatment in BU. Essentially it was shown that diagrams, similar to these used as a proof in BU, can be obtained even from random data. The criticisms of the paper were not directed at the observation that there can be active longitudes which persist over time, but towards the claim that such longitudes maintain consistent migration on a century time-scale.  In recent papers Usoskin, Berdyugina \\& Poutanen~\\cite{UBP} (UBP) and Berdyugina et al.~\\cite{BMSU} (BMSU) responded to this critique with the claim that they can confirm the original BU results using much simpler constructions which do not use the data processing methods criticized in PTB. In particular they consider that the new analysis confirms that the active longitudes maintain consistent smooth migration throughout the 120 year period of the sunspot data analysed. The fact that this result was produced by two different methods of analysis of the data is considered to confirm the reality of the phenomenon.  As will be discussed in length below, the migration models given in the first paper (BU) significantly differ from those obtained as a result of the analysis in UBP and represented in the data analysis part of BMSU. This is most readily seen in the cumulative frame kinematics: the BU frame makes about 28 extra rotations compared to the Carrington frame, contrasted with the UBP values of about $\\approx 11$ rotations for the Northern hemisphere and about $\\approx 14$ for the Southern hemisphere.  The particular key feature here is that in BU the Northern and Southern frames keep coherently in phase whereas in UBP the Southern frame lags with respect to the Northern one. This indicates the need for a very thorough analysis of the methods of analysis used in constructing the differentially rotating frames considered to ``freeze'' the longitudes of these nonaxisymmetic features.  This paper is structured as follows: in section~\\ref{Section:data} we describe the data and the background to our analysis. In section~\\ref{Section:cellcounting} we demonstrate that kinematic frames which follow the rotation of sunpots exist and can clarify the basic properties of the longitudinal distribution of sunspots. In section~\\ref{Section:method} we present an examination of parameter space using the methods used in UBP but increasing the resolution and the range of the parameter space search. In section~\\ref{Section:discussion} we discuss these results and introduce clarificatory statistical computations to check whether bimodal distributions could be statistical fluctuations. Finally section~\\ref{Section:conclusions} presents our conclusions. ",
        "conclusions": "\\label{Section:conclusions} In the recent papers BU, UBP and BMSU the authors claim that, hidden in the sunspot distribution, there are two persistent active longitudes which migrate according to a certain differential rotation law. After presenting a critique of the BU results in the earlier paper (PTB), here we have tried to check the results obtained in UBP and the data analysis part of BMSU. Our results can be summarized as follows: \\begin{itemize} \\item The contramoving frames constructed by BU  differ significantly from the frames found using mathematical optimization techniques in UBP and BMSU. In the first case the frame does approximately 28 extra rotations (compared with the Carrington frame) and in the new papers the number of extra rotations is around 10-11 for the Northern hemisphere and 13-14 for the Southern hemisphere. Consequently the claim in UBP that the new analysis confirms the results obtained in the previous study is not well founded. \\item In the two latest papers authors re-iterate the claim presented in BU of the presence of the flip-flop effect with a mean period of 3.7 years. This result was obtained from (and now presumably incorrect) analysis in the first paper (BU). \\item {Our analysis demonstrated that even for fixed angular velocity contramoving frames it is quite easy to get particular longitude distributions with relatively high values of nonaxisymmetry. This result is quite robust and does not depend on the particular implementation of the UBP method (smoothing of the shift curve, sampling rates in the parameter space etc.). Consequently all the claims in UBP about $\\Gamma = 0.02-0.03$ level fluctuations in the case of $B=0$ are incorrect. The parameter space is full of parameter combinations which can give a rather high values for nonaxisymmetry.} \\item Our re-implemantion of the UBP method gave fundamentally different results from those of the original authors. We can only conjecture about the possible reasons for the discrepancies. The pre- or postselection of data points, parameter undersampling in search procedures, incorrect phase reduction etc. - these all can be considered as possible sources of differences. \\item {From inspection of the strongly fluctuating merit function curves and additional Monte-Carlo type computations for parameter space slices, we conclude that the frames obtained from the optimization procedure are pure fluctuations.} \\item We found that the parameter values obtained as results in UBP can be seen already in the distributions connected to the frame shift curve itself (which does not depend on longitudes of sunspots at all).  \\item Our analysis does not indicate that short-lived features concentrated in certain longitude are totally missing in the sunspot distribution, even though evidence for century scale persistent active longitudes can not be found. The simple cell-counting method investigated here merely as a demonstration tool may well be used to estimate real coherence times of these structures. \\item {Whichever method used, the preferred solutions occurred at positive/negative values of the differential rotation parameter $B$ in the comoving/contramoving frame systems, respectively, i.e. the frames had a tendency to follow a solar-type differential rotation pattern.}  \\end{itemize} From this we conclude that the results obtained in BU, UBP and BMSU (data analysis part) are inconsistent, and evidence for well established persistent activity migration is still lacking. In other words we are inclined to think that the enigmatic face of the Sun is still hidden and that the patterns which the authors of the three papers saw have a human origin."
    },
    "0609/astro-ph0609459_arXiv.txt": {
        "abstract": "We present results from a systematic study of the effects of initial parameters on three-dimensional thermonuclear supernova models. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609490_arXiv.txt": {
        "abstract": "The propagation of high-energy cosmic rays in the Galaxy is investigated. Solutions of a diffusion model are combined with numerically calculated trajectories of particles. The resulting escape path length and interaction path length are presented and energy spectra obtained at Earth are discussed. It is shown that the energy spectra for heavy elements should be flatter as compared to light ones due to nuclear interactions during the propagation process. The obtained propagation properties of ultra-heavy elements indicate that these elements could play an important role for the explanation of the second knee in the cosmic-ray energy spectrum around 400~PeV. ",
        "introduction": "The explanation of the origin of super-high energy cosmic rays is one of the unsolved problems in astrophysics. The energy spectra at the sources are not identical to the observed spectra at Earth. Hence, studying the sources is closely related to investigations of cosmic-ray propagation processes in the Galaxy.  For the latter, a detailed knowledge of the structure of the magnetic fields is important. Unfortunately, the configuration of the galactic magnetic field remains an open question --- different models can explain the experimental data \\cite{berezinsky,ruzmaikin,ptuskin,gorchakov}.  How cosmic rays are accelerated to extremely high energies is another unsolved problem.  Although the popular model of cosmic-ray acceleration by shock waves in the expanding shells of supernovae (see e.g. \\cite{ellison,berezhko,sveshnikova}) is almost accepted as \"standard theory,\" there are still serious difficulties.  Furthermore, the question about other acceleration mechanisms is not quite solved, and such mechanisms could lead to different cosmic-ray energy spectra at the sources \\cite{berezinsky}.  Different concepts can be verified, calculating the primary cosmic-ray energy spectrum, making assumptions on the density of cosmic-ray sources, the energy spectrum at the sources, and the configuration of the galactic magnetic fields. The diffusion model may be used in the energy range $E<10^{17}$~eV, where the energy spectrum is obtained using the diffusion equation for the density of cosmic rays in the Galaxy. At higher energies this model ceases to be valid, and it becomes necessary to carry out numerical calculations of particle trajectories for the propagation in the magnetic fields.  This method works best for the highest energy particles, since the time required for the calculations is inversely proportional to the particle energy.  Therefore, a calculation of the cosmic-ray spectrum in the energy range $10^{14}-10^{19}$~eV has been performed in a combined approach: solutions of a diffusion model are used at low energies and particle trajectories are numerically integrated at high energies.  In \\sref{assect} the basic assumptions for the diffusion model will be described. The results obtained with the propagation model are presented in the subsequent sections.  The calculated propagation path length and interaction probability of cosmic rays will be discussed in \\sref{pathsect} and \\sref{intersect}, respectively. Finally, the energy spectra are presented in \\sref{specsect}, followed by a discussion of the results (\\sref{discussion}). ",
        "conclusions": "\\label{discussion}  The energy spectra for individual elements measured at the Earth with GeV and TeV energies can be described by power laws $dN/dE\\propto E^\\gamma$ with values for the spectral index $\\gamma$ in the range -2.46 to -2.95 for elements from hydrogen to nickel \\cite{wiebel,pg}.  The measurements seem to indicate that the steepness of the energy spectra at Earth depends on the mass of the nuclei, heavier elements seem to have flatter spectra.  At higher energies in the PeV domain the measured spectra are compatible with the assumption of a \\knee for individual elements at about $Z\\cdot 4.5$~PeV \\cite{pg,aspenreview}.  The energy spectrum of cosmic rays at their source $Q(E)$ is related to the observed values at Earth $N(E)$ as \\begin{equation} \\label{sourceeq} N(E) \\propto Q(E) \\left( \\frac{1}{\\lambda_{esc}(E)} + \\frac{1}{\\lambda_{int}(E)}\\right)^{-1} \\end{equation} with the escape path length $\\lambda_{esc}$ and the interaction length $\\lambda_{int}$. Values for the former are presented in \\fref{pathlength} and for the latter in \\fref{lambdaint}. The relation between $N(E)$ and $Q(E)$ is governed by the absolute values of $\\lambda_{esc}$ and $\\lambda_{int}$ as well as their respective energy dependence.  The interaction length $\\lambda_{int}$ is almost independent of energy, the values for $\\varepsilon$ in \\eref{lambdainteq} are smaller than 0.05, see \\fref{lambdaslope}. On the other hand, the propagation path length $\\lambda_{esc}$ decreases as function of energy as $\\lambda_{esc}\\propto E^{-\\delta}$, with values between $\\delta=0.6$ for Leaky Box models and $\\delta=0.2$ for the diffusion model described in this work (see \\sref{pathsect}).  In the \"standard picture\" of galactic cosmic rays usually $\\lambda_{int}>\\lambda_{esc}$ is assumed with an energy independent interaction length and an escape path length $\\lambda_{esc}=\\lambda_{lb}\\propto E^{-0.6}$. Thus, the relation between $N(E)$ and $Q(E)$ is dominated by the energy dependence of the escape path length. To explain the spectrum observed at Earth $N(E)\\propto E^{-2.7}$ the spectrum at the sources has to be $Q(E)\\propto E^{-2.1}$.  Recent measurements by the HESS experiment of the TeV gamma ray flux from the shell type supernova remnant RX J1713.7-3946 yield a spectral index $\\gamma=-2.19\\pm0.09\\pm0.15$ \\cite{hesssnr}.  For the first time, spectra have been derived directly at a potential cosmic-ray source.  The results are compatible with a nonlinear kinetic theory of cosmic-ray acceleration in supernova remnants and imply that this supernova remnant is an effective source of nuclear cosmic rays, where about 10\\% of the mechanical explosion energy are converted into nuclear cosmic rays \\cite{voelkrxj1713}.  These investigations show that the spectrum at the source is $Q(E)\\propto E^{-2.2}$, a value relatively close to the \"standard model\". The escape path length required to obtain the spectra observed at Earth should be $\\propto E^{-0.5}$, i.e. $\\delta$ is 0.1 smaller than in the \"standard picture\".  Two questions remain open:  Do all potential cosmic-ray sources exhibit a similar spectrum? And, what fraction of galactic cosmic rays is accelerated in supernova remnants with spectra like RX J1713?  \\begin{figure*}[t] \\centering \\includegraphics[width=0.9\\textwidth]{prop.eps} \\caption{Energy spectra at Earth for elements with nuclear charge $Z$ as indicated. The dashed lines represent spectra according to the \\modell, the solid lines are expected from the diffusion model discussed, see text.  Two solid lines are shown in each panel, representing an estimate for the uncertainties.  \\label{prop}} \\end{figure*}  Considering the spectra of nuclei observed at Earth, the second term in \\eref{sourceeq} has to be taken into account.  The values of $\\lambda_{int}$ for protons are at all energies larger than the escape path length $\\lambda_{dif}$. Hence, for protons the escape from the Galaxy is the dominant process influencing the shape of the observed energy spectrum.  For iron nuclei at low energies hadronic interactions are dominating ($\\lambda_{int}\\approx2.5~\\mbox{\\gcm2} < \\lambda_{dif}$) and leakage from the Galaxy becomes important only at energies approaching the iron \\knee ($26\\cdot \\hat{E}_p$).  Finally, for elements heavier than iron, the interaction path length is smaller than the escape path length at all energies, except above the respective \\knees. For these elements the propagation path length exceeds the interaction path length by about an order of magnitude at low energies.  Hence, nuclear interaction processes are dominant for the shape of the observed spectrum.  Taking this effect into account, it is expected that the energy spectra for heavy elements should be flatter as compared to light nuclei.  The spectral index for iron nuclei should be about 0.13 smaller than the corresponding value for protons.  Direct measurements indicate indeed that the spectra of light elements are flatter as compared to heavy elements \\cite{pg}, e.g. the values for protons $\\gamma_p=2.71$ and iron $\\gamma_{Fe}=2.59$ differ as expected. This more or less obvious effect has been pointed out earlier in several articles, see e.g. \\cite{biermann,jones}. However, in the present work we extend the considerations to a wider energy range and ultra-heavy nuclei using contemporary estimates of the cross sections.  The \\modell extrapolates results from direct measurements in the GeV and TeV regime to higher energies \\cite{pg}. In the model fall-offs for individual elements proportional to the nuclear charge are assumed at energies of $E_Z=Z\\cdot \\hat E_p$, with the value $\\hat E_p=4.5$~PeV for protons. The \\knee is caused by the subsequential fall-offs for all elements, starting with protons. The second \\knee in the all-particle spectrum around 400~PeV is interpreted as the end of the galactic component, where the heaviest elements (actinides) reach their fall-off energies ($92\\cdot \\hat E_p\\approx414$~PeV). At this energy, elements heavier than iron are expected to contribute with about 40\\% to the all-particle flux \\cite{aspenphen}.  The model uses an empirical relation for the spectral index. The spectral slope decreases as function of nuclear charge.  In the following, this behavior is investigated quantitatively with the propagation model discussed above.  As source composition, the abundances of elements from hydrogen to uranium as measured in the solar system \\cite{cameron} have been weighted with $Z^{3.2}$. This choice is arbitrary to a certain extent, but may be motivated by a higher efficiency in the injection or acceleration processes for nuclei with high charge numbers. The abundances are scaled with a factor which is identical for all elements to obtain approximately the absolute values as expected at the Earth according to the \\modell.  At the source, a power law $\\propto E^{-2.5}$ has been assumed for all elements with a \\knee, caused e.g.  by the maximum energy attained during the acceleration, at $Z\\cdot4.5$~PeV, with a power law index $-3.5$ above the respective \\knee.  Using the derived propagation path length and interaction length, the amount of interacting particles has been determined.  Secondary products generated in spallation processes are taken into account, assuming that the energy per nucleon is conserved in these reactions. They are added to the corresponding spectra with smaller $Z$. The spectra thus obtained are compared to spectra according to the \\modell in \\fref{prop}.  Two features should be noted: The absolute fluxes at Earth are predicted quite well, especially when considering that only a simple scaling law has been introduced for the abundances at the sources, starting with the composition in the solar system.  More important for the present discussion is the shape of the spectra. As expected, the shape of the proton spectrum is not influenced by the (few) interactions during propagation and the difference of the spectral index at the source and at Earth $\\gamma=-2.71$ \\cite{pg} can be explained by the energy dependence of the escape path length $\\propto E^{-0.2}$.  On the other hand, it can be recognized that due to nuclear interactions the spectra for heavier elements are indeed flatter. The slopes obtained with the simple approach for the CNO, silicon, and iron groups agree well with the steepness as expected from the \\modell.  For heavy elements at low energies secondary products generated in spallation processes play an important role for the shape of the spectrum. At low energies many nuclei interact due to the large escape path length and the small interaction length, thus, the spectra of nuclei without any interaction deviate from power laws. However, the spallation products of heavier elements at higher energies compensate the effect and the resulting spectra are again approximately power laws, as can be seen in \\fref{prop}."
    },
    "0609/astro-ph0609173_arXiv.txt": {
        "abstract": "We review the fireball shock model of gamma-ray burst prompt and early afterglow emission in the light of rapid follow-up measurements made and enabled by the multi-wavelength \\swift\\ satellite. These observations are leading to a reappraisal and expansion of the previous standard view of the GRB and its fireball. New information on the behavior of the burst and afterglow on minutes to hour timescales has led, among other results, to the discovery and follow-up of short GRB afterglows, the opening up of the $z\\simg 6$ redshift range, and the first prompt multi-wavelength observations of a long GRB-supernova.  We discuss the salient observational results and some associated theoretical issues. ",
        "introduction": "\\label{sec:obs}  NASA's \\swift\\ mission \\cite{gcg+04} has enabled fundamental insights into the physics of gamma-ray bursts thanks to two new capabilities: First, the greater sensitivity of its Burst Alert Detector \\cite{bbc+05} (BAT; energy range 20--150 keV) in comparison to the preceding \\bepposax\\ and \\hete\\ missions \\cite{band06}; and second, its ability to slew in less than 100 seconds to the burst direction determined by the BAT, which allows it to position its much higher-angular resolution \\xray\\ (XRT, few-arcsec) and UV-Optical (UVOT, sub-arcsec) detectors \\cite{bhn+05,rhm+04} for observations of the prompt and early afterglow emission.   As of July~2006, over 150 bursts had been detected by BAT, at an average rate of 2 bursts detected per week. Of these, roughly 90\\% were followed promptly with the XRT within $350$ s from the trigger, and about half within 100~s \\cite{brg+05}, while $\\sim 30\\%$ were detected with the UVOT \\cite{rsf+05}.  This resulted in over 30 new redshift determinations.  Eleven definitive short GRBs were detected, of which seven had detected \\xray\\ afterglows (including three with optical afterglows, and two with radio), and five had proposed redshifts.   These \\swift\\ observations, complemented by the continuing operations of the \\hete\\ and \\integral\\ missions, have brought the total number of redshift determinations to over 60 since \\bepposax\\ enabled the first one in 1997. The redshifts based on \\swift\\ have a median $z\\sim 2.8$ \\cite{jlf+06,bmb+06}, which is a factor $\\simg 2$ higher than the median of those previously culled via \\bepposax\\ and \\hete\\ \\cite{bkf+05}. This can be credited largely to the prompt $\\sim$arcsec positions from XRT and UVOT, which make rapid ground-based observations possible at a stage when the afterglow is still bright. The highest \\swift-enabled redshift so far was that of GRB\\,050904, from 8-meter class Subaru spectroscopy, $z=6.29$ \\cite{kka+05} (the pre-\\swift\\ record was $z=4.5$ for the IPN-localized GRB\\,000131 \\cite{ahp+00}). The relative paucity of UVOT detections versus XRT detections may be ascribed in part to this higher median redshift (and correspondingly reduced median flux), and in part to the (expected) increased dust extinction at the shorter rest-frame wavelengths implied for any given observing band \\cite{rsf+05}, although the issue is still the subject of debate.   {\\it BAT light curves:} The BAT is provided with an array of triggering algorithms which provide greater sensitivity than those of previous spacecraft, including ``imaging'' triggers which, through analysis of the counts integrated over 128-s to 512-s timescales, enable the detection of faint, slow-rising events.  In a peculiar irony of GRB studies, these image triggers have made possible the discovery of both the farthest, highly time-dilated burst GRB\\,050904, and the nearest, faint and slow-rising burst GRB\\,060218 (see below for further discussion).  For some of the bursts which fall in the ``long'' category ($\\gamma$-ray duration $t_\\gamma \\simg 2$ s), the BAT has detected faint soft gamma-ray tails which extend the duration by a factor up to two beyond what the previous Burst and Transient Source Experiment (BATSE) could have detected \\cite{gehrels06}. Such extended soft tails have been found also in some ``short'' bursts, normally defined as having a hard spectrum and duration $t_\\gamma \\siml 2$ s (see below).   {\\it XRT light curves:} Striking new insights into burst and afterglow physics have come from the detailed \\xray\\ light curves, starting on average 100 seconds after the $\\gamma$-ray trigger, that result from the prompt XRT observations of BAT-detected bursts.  These observations suggest a canonical \\xray\\ afterglow (Fig.~\\ref{fig:xrt-lc}; \\cite{nkg+06,zfd+06}) with one or more of the following stages (note that the numerical subscript enumerates each stage): (1) An initial steep decay $F_X \\propto t^{-\\alpha_1}$ with temporal index $3 \\siml \\alpha_1 \\siml 5$, and an energy spectrum $F_\\nu \\propto \\nu^{-\\beta_1}$ with energy spectral index $1 \\siml \\beta_1 \\siml 2$ (photon index $0\\siml \\Gamma\\siml 1$), extending up to a time $300\\,\\s \\siml t_1 \\siml 500\\,\\s$; (2) A subsequent flatter decay $F_X \\propto t^{-\\alpha_2}$ with $0.2 \\siml \\alpha_2 \\siml 0.8$ and energy index $0.7 \\siml \\beta_2 \\siml 1.2$, at times $10^3\\,\\s \\siml t_2 \\siml 10^4\\,\\s$ (in some cases interspersed with flares, see \\cite{zfd+06}; (3) A ``normal'' decay $F_X \\propto t^{-\\alpha_3}$ with $1.1 \\siml \\alpha_3 \\siml 1.7$ and $0.7 \\siml \\beta_3 \\siml 1.2$ (generally unchanged from the previous stage, i.e.\\ $\\beta_3\\approx \\beta_2$), up to a time $t_3 \\sim 10^5\\,\\s$, or in some cases longer; (4) In some cases, a steeper \\xray\\ decay at the end, $F_X \\propto t^{-\\alpha_4}$, with $2\\siml \\alpha_4 \\siml 3$, after $t_4\\sim 10^5 \\,\\s$, resembling what is expected from jet breaks; (5) In about half the afterglows, one or more \\xray\\ flares or bumps are observed in the light curve, sometimes starting as early as 100 s after trigger, and sometimes as late as $10^5\\,\\s$.  The energy in these flares ranges from a percent up to a value comparable to the prompt emission (in GRB\\,050502b).  The rise and decay times of these flares is unusually steep, depending on the reference time $t_0$, behaving as $(t-t_0)^{\\pm \\alpha_{\\rm fl}}$ with $3 \\siml \\alpha_{\\rm fl} \\siml 6$, and energy indices which can be also steeper than during the smooth decay portions. The flux level after the flare usually decays to the value extrapolated from the value before the flare rise.  \\begin{figure}[tbh] \\centerline{\\epsfxsize=5.in \\epsfbox{xrt_genlc.eps}} \\caption{Schematic features seen in early \\xray\\ afterglows detected with the \\swift\\ XRT instrument (e.g.\\ \\cite{zfd+06,nkg+06}); see text. \\label{fig:xrt-lc}} \\end{figure}   {\\it Very high-redshift bursts:} Another major advance achieved by \\swift\\ was the detection of the long burst GRB\\,050904, which broke the $z>6$ redshift barrier. This burst was very bright, both in its prompt $\\gamma$-ray emission ($\\Eiso\\sim 10^{54}$ erg) and in its \\xray\\ afterglow. Prompt ground-based optical/IR upper limits and a $J$-band detection suggested a photometric redshift $z>6$ \\cite{hnr+06}. Spectroscopic confirmation with the 8.2 m Subaru telescope gave $z=6.29$ \\cite{kka+05}.  There are several striking features to this burst. One is its enormous \\xray\\ brightness, exceeding for a full day the \\xray\\ brightness of the most distant quasar known to-date, SDSS~J0130+0524 -- and exceeding it by a factor of $10^5$ in the first minutes \\cite{wrh+06}. The implications for the use of GRBs as a tool for probing the IGM are thought-provoking. Another notable feature was its extremely variable \\xray\\ light curve, showing many large amplitude flares throughout the first day (Fig.~\\ref{fig:lc050904}). A third exciting feature is its brief, very bright IR flash \\cite{bad+06}, comparable in brightness to the famous $m_V\\sim 9$ optical flash in GRB\\,990123.  \\begin{figure}[tbh] \\centerline{\\epsfxsize=5.in \\epsfbox{lc050904.eps}} \\caption{Prompt and afterglow lightcurves for the redshift $z=6.29$ GRB\\,050904 at hard \\xray\\ (BAT, 50~keV), \\xray\\ (XRT, 5~keV), and near-infrared ($J$-band, 1.1~\\micron; and $I$-band, 9500\\,\\AA) frequencies.  A bright early flare is seen in the IR and \\xray; in the \\xray\\ this is followed by extensive flaring activity for more than a day after the burst.  Power-law decay indices measured in the IR ($\\alpha_{\\rm o}$) are indicated, where available; a likely jet break is seen at $t\\approx 2.3$\\,days \\cite{tac+05}.  Data from \\cite{hnr+06,bad+06,cmc+06,tac+05}.} \\label{fig:lc050904} \\end{figure}   {\\it Short bursts:} The third major advance from \\swift\\ was the discovery and localization of short GRB afterglows. As of July~2006, eleven definitive short bursts have been localized by \\swift, and \\hete\\ has discovered two short bursts with subsequent afterglow emission; in addition, the short burst GRB\\,051103 was localized by the IPN to a region intersecting the nearby galaxies M81 and M82.  In seven of the \\swift\\ short bursts and two \\hete\\ bursts an \\xray\\ afterglow was detected and observed, with GRBs 050709, 050724, 051221A, 060313, and 060121 (as well as the candidate short burst GRB\\,051227) also showing optical afterglows, and GRBs 050724 and 051221A being detected in the radio \\cite{berger06}.  These are the first afterglows detected for short bursts, and yielded the first host galaxy identifications.  The hosts are of early type (ellipticals) in roughly half the cases, and otherwise appear to be dwarf irregular galaxies -- albeit with evidence for old as well as young stellar populations \\cite{sbk+06}. The redshifts of five of them are in the range 0.15 to 0.5, while GRB\\,050813 was initially thought to have $z=0.7$ \\cite{pbc+06} but may be associated with a galaxy cluster at $z\\approx 1.8$ \\cite{berger06}, and the redshift of the \\hete\\ GRB\\,060121 is estimated to be $z\\approx 4.6$ \\cite{dcg+06}. The median $z$ (as yet unaffected by these two relatively high redshifts) is $z_{\\rm med}=0.26$, which gives a $(1+z)$ that is 1/3 that of the long bursts. While there is evidence for star formation in roughly half the host galaxies, overall the host properties correspond to expectations for a progenitor population of neutron star or neutron star-black hole binary systems, the most often-discussed progenitor candidates.  Such hosts would also be compatible with other progenitors involving old compact stars.  The first short burst afterglow seen by \\swift\\ was that of GRB\\,050509B, a low luminosity ($\\Eiso\\sim 2\\times 10^{48}$ erg) burst with a simple power-law \\xray\\ afterglow which could only be followed for $\\sim 10^4$ s \\cite{gso+05}.  The subsequent afterglows of GRB\\,050724 ($\\Eiso\\sim 3\\times 10^{50}$ erg; \\cite{bcb+05}) and GRB\\,051221A ($\\Eiso\\sim 1.5\\times 10^{51}$ erg; \\cite{sbk+06}) were brighter, and could be followed in \\mbox{X-rays} for at least $10^5$ s \\cite{gbp+06,bgc+06}.  In particular, the \\xray\\ afterglow of GRB\\,050724 is remarkable in that it exhibits a typical \\xray\\ light curve as observed from long GRBs by \\swift\\ (apart from the absence of a slow-decay phase).  It shows a fast early decay, with significant \\xray\\ flares, one at 100~s and another at $3\\times 10^4$~s. The first flare has the same fluence as the prompt emission, while the late flare has roughly 10\\% of that. The interpretation of this activity poses interesting challenges, as discussed below.   {\\it The GRB-SN connection:} The fourth major advance from \\swift\\ was its observation, with BAT, XRT and UVOT, of an unusually long ($\\sim 2000$ s), soft burst, GRB\\,060218 \\cite{cmb+06}, soon found to be associated with SN2006aj, a nearby ($z=0.033$) type Ic supernova \\cite{msg+06,mha+06,pmm+06,skn+06}.  The optical light curve of this supernova peaked earlier than for most other known supernovae, and the time of core collapse is constrained to within a day of the GRB trigger.  This was the first time that a connected GRB and supernova event was observed starting in the first $\\sim$100~s in \\xray\\ and UV/optical light, and the results are of great interest. The early \\xray\\ light curve shows a slow rise and plateau followed by a drop after $\\sim 10^3$ s; the spectrum is initially dominated by a power-law component, with an increasing thermal component that dominates after $\\approx$3000~s.  Perhaps the most interesting interpretation involves shock break-out of a semi-relativistic component in a WR progenitor wind \\cite{cmb+06}, although see also \\cite{fp06}.  After this a more conventional \\xray\\ power-law decay follows, and a UV component peaking at a later time can be interpreted as due to the slower supernova envelope shock.  Finally, a radio-bright component is seen that presents an interesting comparison to other local and GRB-related type Ic supernovae \\cite{skn+06}.  A distinct GRB/SN detection based on \\swift\\ afterglow observations is that associated with GRB\\,050525A \\cite{dmb+06}.   {\\it Prompt optical observations:} A final major advance enabled by \\swift\\ has been a wealth of information about the prompt and early-time optical emission of GRBs. Thanks to \\swift\\ alerts and an array of ground-based robotic telescopes (as well as the UVOT), the detection of optical counterparts during and shortly after ($<$100~s) the gamma-ray trigger is becoming almost routine.  The \\swift\\ era began with the detection in optical \\cite{vww+05} and near-infrared \\cite{bbs+05} wavelengths of a bright flash associated with GRB\\,041219A (localized in real time by \\integral).  Apart from exhibiting the first optical flash with brightness (after correcting for Galactic extinction) similar to GRB\\,990123, this burst showed the first evidence for correlated low- and high-energy emission, intepreted as optical emission from the same internal shocks that produced the GRB \\cite{vww+05,bbs+05}.  A similar level of correlated optical emission was later observed from GRB\\,050820A \\cite{vww+06}; however in this case it was accompanied by an uncorrelated component that connected smoothly to the later afterglow.  Prompt optical emission that extrapolates smoothly to the later afterglow seems to be the more usual case; however, bright flares consistent with a reverse-shock interpretation are seen on occasion \\cite{jpk+06,kgs+06}.  The most exciting single observation is probably the $I=14.1$\\,mag optical flash from GRB\\,050904 at $z=6.3$ observed with TAROT \\cite{bad+06}, which rises to its peak after the end of gamma-ray emission and is brighter, and decays more quickly than, the later afterglow -- all properties consistent with an origin in this burst's reverse shock. ",
        "conclusions": ""
    },
    "0609/astro-ph0609615_arXiv.txt": {
        "abstract": "Within the cosmic framework clusters of galaxies are relatively young objects. Many of them have recently experienced major mergers. Here we investigate an equal mass merging event at $z\\approx0.6$ resulting in a dark matter haloe of $\\sim2.2\\times10^{14}\\hMsol$ at $z=0$. The merging process is covered by 270 outputs of a high resolution cosmological N-body simulation performed with the ART (adaptive refinement tree) code. Some $2\\Gyr$s elapse between the first peri-centre passage of the progenitor cores and their final coalescence. During that phase the cores experience six peri-centre passages with minimal distances declining from $\\sim30$ to $\\sim2\\hkpc$. The time intervals between the peri-centre passages continuously decrease from $9$ to $1\\times10^8\\yr$s. We follow the mean density, the velocity dispersion and the entropy of the two progenitors within a set of fixed proper radii (25, 50, 100, 250, 500, $1000\\hkpc$). During the peri-centre passages we find sharp peaks of the mean densities within these radii, which exceed the sum of the corresponding progenitor densities. In addition to the intermixing of the merging haloes, the densities increase due to contraction caused by the momentary deepening of the potential well. The velocity dispersions also peak during peri-centre passages. Within the fixed proper radii the entropy of the most massive progenitor after the merger settles close to its pre-merger values. At the end of the oscillatory relaxation phase the material originating from the less concentrated of the two equal mass progenitors is deposited at larger radii and shows a slightly more radially anisotropic velocity dispersion compared to the material coming from the more concentrated progenitor. Every peri-centre passage is accompanied by a substantial drop of the central potential well. We briefly discuss the possibility that AGN outbursts are triggered by the periodically changing potential. ",
        "introduction": "The current cosmological model (e.g.~\\citealt{Spergel06}) assumes the gravitating matter to be composed of $~\\sim 15$ per cent baryonic and $~\\sim 85$ per cent dark matter (DM). The former interacts electromagnetically, the latter is not directly visible. By comparing lensing observations with X-ray maps of the merging cluster 1E0657-558 \\cite{Clowe06} find strong evidence for unseen matter concentrations that are more massive than the baryonic component. These observations emphasise the complexity of the interactions between dark and baryonic matter during mayor merging events. Based on hydrodynamic simulations of  head-on mergers between model galaxy clusters \\cite{Pearce94} find a damping of the core oscillations by the gaseous component. The situation may become even more complex if the collision of gas-rich systems can trigger star formation and outflows from the central active galactic nucleus (AGN). Multi wave band observations by \\cite{Conselice01} reveal that this may actually the case for the central galaxy NGC 1275 in the Perseus Cluster.  Using high resolution gas dynamical simulations \\cite{Nagai06} demonstrate that the X-ray observations of unrelaxed clusters cannot accurately measure the density and temperature gradients in the hot baryonic intra cluster medium (ICM). In addition some potentially relevant physical processes, such as AGN bubbles, magnetic fields, and cosmic rays, are not well understood yet and are usually excluded from such simulations. Here we do not aim to model the behaviour of the gas directly. Our analysis is based on a pure N-body simulation. However, since the collisionless component dominates the gravitational potential, we hope to gain some complementary insight into the violent processes during a mayor merger of galaxy clusters. The merging event is covered by a set of 270 simulation outputs which allows us to precisely measure even highly variable features during the violent relaxation phase. Our goal is twofold: (1) we provide a precise time table of events during the core merging phase. The dynamics may be slightly altered by the presence of collisional material, however the general behaviour is not expected to change profoundly. (2) we investigate the evolution of the central potential to achieve some basic understanding of the gravitational forces which the baryons are exposed to. Ultimately we wish to uncover interactions between the DM and the baryonic components which result in observable phenomena.  The paper is organised as follows. \\S~\\ref{sec:simda} introduces the simulation and data preparation. In \\S~\\ref{sec:res} we present the evolution of the density, the velocity dispersion and the entropy of the DM within spheres of fixed proper radii. Additionally, the evolution of the central gravitational potential is discussed. Finally, \\S~\\ref{sec:sum} provides a summary of the main results and a discussion focusing on the potential interplay between the dark and the baryonic component. ",
        "conclusions": "\\label{sec:sum} We have analysed the merging of a cluster-sized dark matter haloe focusing on the evolution of the density and the velocity dispersion within fixed proper radii. This approach is complementary to the more common analysis based on the scaling relations for dark matter haloes. Whereas the latter emphasises the self similar nature of gravitationally collapsed systems, the former provides insight into the processes at fixed scales and may help to understand the cause of the self similarity. Our main results are: (1) About $3\\Gyr$s elapse from the first encounter of the progenitor outskirts to the final coalescence of their central cores. Within  the $\\sim2\\Gyr$s between the first peri-centre passage and the final core merging six peri-centre passages take place. During that process the periods decrease by a factor of ten from $\\sim10^9\\yr$s to $\\sim10^8\\yr$s. (2) The central densities of the merging progenitors peak at each peri-centre passage and are most pronounced at the first one. The density enhancement is caused by the overlapping volumes of the subhaloes but also by the contraction due to the deepening of the potential well during a peri-centre passage. (3) After the merger is completed ($z\\lesssim0.4$) the densities and the velocity dispersions within the fixed proper radii keep roughly constant. The increase for larger radii at $z\\lesssim0.1$ is associated with an ongoing intermediate merger as discussed in \\S~\\ref{sec:simda}. In agreement with \\cite{RomanoDiaz06} we find that the properties of the dark matter haloe change rapidly during merging events but remain constant during phases of passive evolution. (4) The ranking of the progenitor concentrations determines the final distribution of the particles. Less concentrated haloes get more dispersed and show higher radially anisotropic velocity dispersions.  Finally, we briefly discuss the impacts of the oscillatory relaxation process of the collisionless matter on the baryonic component. It is of interest to consider the possibility that enhanced AGN activity in the cluster-centred galaxies may be stimulated by the large oscillations of the haloe potentials during a major cluster merger.  We note that the period of relaxation oscillations for our merging cluster, $\\sim 10^8$ yrs, is similar to the duty cycle observed in AGNs \\citep{Voit05}, although only a small fraction of AGNs reside in massive clusters.  Nevertheless, to explore this idea, we performed a simple gasdynamical calculation to determine how much additional radiative cooling occurs when the merging potentials are deepened, compressing the hot cluster gas.  For this purpose we set up a traditional quasi-steady state cluster cooling flow in a potential of mass $3 \\times 10^{14}$ with typical concentration $c = 6.74$.  The gas temperature at large radii, $3.5$ keV, was chosen according to the observed cluster $M - T$ relation \\citep{Arnaud05} and the initial (hydrostatic) gas density was adjusted until the baryon mass matched the cosmic value at the virial radius.  The relaxed steady state cooling flow that results closely resembles the observations of \\cite{Vikhlinin06} and has a central cooling rate of ${\\dot M} = 43$ $M_{\\odot}$ yr$^{-1}$ (assuming abundance 0.4 solar).  The potential was then set into sinusoidal oscillations with period $0.6$ Gyr and an amplitude similar to that of Figure 10 by sinusoidally varying the concentration between $c$ and $0.43c$. During these oscillations, the cooling rate ${\\dot M}(t)$ near the centre of the flow varied with an amplitude of only 7 per cent, much less than the unperturbed steady state value which is at least $\\sim10$ times larger than allowed by X-ray observations \\citep{Peterson01}.  We conclude that the additional radiative cooling stimulated by the oscillating merger potentials is negligible.  Perhaps a more fruitful approach along these lines is the possibility that galaxies (and large black holes) orbiting about the cluster core can be processed by dynamical friction into the central galaxy/black hole during the relaxation process.  The acceleration due to dynamical friction $(du/dt)_{df} \\propto m \\rho/u^2$ depends on the mass $m$ and the velocity $u$ of the orbiting object and the density of the ambient material. By setting the velocity equal to the mean dispersion $u=\\sigma$, we see that the time scale for dynamical friction $t_{df} = u/(du/dt)_{df} \\propto \\sigma^3/\\rho = K^{2/3}$. During apo-centre passages when the dark halo entropy passes a minimum, objects orbiting near the cluster-centred galaxy may efficiently reduce their angular momentum, increasing the chance of a capture by the central black hole during the subsequent contraction phase. It is conceivable that such events, stimulated by the global merger, could result in enhanced AGN activity with a period imprinted by the dark matter oscillations."
    },
    "0609/astro-ph0609162.txt": {
        "abstract": "% {Metal-poor globular clusters (GCs) can provide a probe of the earliest epoch of star formation in the Universe, being the oldest stellar systems observable. In addition, young and intermediate-age low-metallicity GCs are present in external galaxies. Nevertheless, inferring their evolutionary status by using integrated properties may suffer from large \\emph{intrinsic} uncertainty caused by the discrete nature of stars in stellar systems, especially in the case of faint objects. } % {In this paper, we evaluate the \\emph{intrinsic} uncertainty (due to statistical effects) affecting the integrated colours and mass--to--light ratios as a function of the cluster integrated visual magnitude ($M_V^{tot}$), which represents a quantity directly measured. We investigate the case of metal-poor single-burst stellar populations with age from a few million years to a likely upper value for the Galactic globular cluster ages ($\\sim$15 Gyr).} % {Our approach is based on Monte Carlo techniques for randomly generating stars distributed according to the cluster's mass function. } % {Integrated colours and mass--to--light ratios in different photometric bands are checked to be in good agreement with the observational values of low-metallicity Galactic clusters; the effect of different assumptions on the Horizontal Branch (HB) morphology is shown to be not relevant, at least for the photometric bands explored here. We present integrated colours and mass--to--light ratios as a function of age for different assumptions on the cluster total $V$ magnitude. We find that the \\emph{intrinsic} uncertainty cannot be neglected. In particular, in models with $M_V^{tot}=-4$ the broad-band colours show an \\emph{intrinsic} uncertainty so high as to prevent precise age evaluation of the cluster. The effects of different assumptions on the initial mass function and on the minimum mass for which carbon burning is ignited on both integrated colours and mass--to--light ratios are also analyzed. Finally, the present predictions are compared with recent results available in the literature, showing in some cases non-negligible differences.} % {} ",
        "introduction": "A key issue in astronomy is to determine ages and chemical compositions of those stars and stellar systems that are needed to reconstruct the formation and evolution of galaxies. Toward this goal, the analysis of the stellar clusters population in external galaxies and in the Milky Way is fundamental for tracing the history of the parent galaxy \\citep[see, e.g.][]{Cote+98,West+04}. Contrarily to the situation in the Galaxy, in which massive ($\\approx 10^5 \\div 10^6$ stars) stellar clusters are old ($\\gsim$10 Gyr) and metal poor ($[Fe/H]\\lsim -$0.7), observations indicate that in the Magellanic Clouds, in the Local Group galaxies, and even in galaxies beyond the Local Group, massive stellar clusters spread a large range of both metallicity and age; thus they are considered the main indicators of stellar formation events in the galaxies history \\citep{Larsen00, Matteucci+02, Harris03}.  Hereinafter we will use the term ``globular clusters'' (GCs) to signify massive clusters of whatever age and chemical composition.  GCs have the advantage to be bright objects, and then to be easily observed beyond the Local Group. In addition, they are thought to be simple stellar populations (SSPs), consisting of a gravitationally bound group of stars born at nearly the same moment, and with a nearly identical chemical composition.  The integrated broad-band colours, line indices, and mass-to-light ratios, we observe from those systems, are the unique observational tools to understand their properties. Since the pioneering work by Tinsley (1972), different groups have developed population synthesis models, e.g. \\citep[][and references therein]{Brocato+90b, Charlot&Bruzual91, Buzzoni93, Bressan+94, Worthey94, Maraston98, Kurth+99, Brocato+00, Vazdekis99, Girardi+00, Anders+03, Bruzual&Charlot03, Maraston05} in order to interpret such observables.  Besides the \\emph{systematic} uncertainties due to different sets of stellar evolutionary tracks and different spectral libraries--used to transform the models from luminosity and effective temperature to observable quantities-- \\citep[see e.g.][]{Charlot+96, Maraston98, Brocato+00, Bruzual&Charlot03, Yi03}, broad-band colours may suffer from large \\emph{intrinsic} fluctuations caused by the discrete nature of the number of stars in the system.  The first studies were carried out in the optical by \\citet{Barbaro&Bertelli77} - for population I clusters - and \\citet{Chiosi+88}- for intermediate ($Z$=0.001) and solar metallicity -, while \\citet{Santos&Frogel97} analyzed the case for near infrared (NIR) bands. Among other results, the quoted authors concluded that it is necessary to include stochastic effects when deriving ages and metallicities from integrated broad--band colours.  Most studies normalized their theoretical predictions to the total number of stars or total mass in the cluster, while \\citet{Brocato+99,Brocato+00} derived the mean broad-band colours and the corresponding dispersions as a function of the cluster visual magnitude ($M_V^{tot}$) for selected values of ages. This approach directly links theory and observations and it is crucial when the cluster age and metallicity are inferred from the observed broad--band colours, especially for clusters at the faint end of the GC luminosity function.  In this paper we extend the investigations by \\citet{Brocato+99} and \\citet{Brocato+00} and we analyze stochastic effects not only on broad--band colours, but also on mass--to--light ratios as a function of the adopted cluster visual magnitude, $M_V^{tot}$.  We provide a comprehensive study on this problem by investigating the time--evolution of both integrated colours and mass--to--light ratios for a fine grid of stellar ages and for three different values for $M_V^{tot}$. The analysis is carried out using new single-burst low-metallicity models ($Z$=0.0002) based on the updated stellar models database by \\citet{Cariulo+04}.  We choose this metallicity because the metal--poor GCs can provide a probe of the earliest epoch of star formation in the Universe, being the oldest stellar systems observable. In addition, young and intermediate age low metallicity GCs are present in external galaxies (Larsen \\& Richtler 1999, 2000).  For these reasons, and also because an analysis of young metal--poor clusters gives an idea on how old GCs appeared when they formed, we explore a wide range of ages (7.7 $\\leq \\log [age (yr)] \\leq 10.3$). We note that we did not take into account redshift effects thus our calculations can be used only for objects with a redshift lower than about 0.1.  The results are compared with a sample of low-metallicity clusters in the Galaxy with different HB morphology, chosen as prototypes of the old stellar populations studied in this work. The present theoretical predictions are also compared with recent results available in the literature showing in some cases non-negligible differences.  We also discuss the influence of the adopted Initial Mass Function (IMF) on integrated colours and mass--to--light ratios and the effect of changing the maximum mass ($M_{up}$) for which carbon burning is not ignited, due to effects of the degenerate pressure and neutrino energy losses in the core.  The assumption of a fixed $M_V^{tot}$ can lead to a peculiar behaviour when varying the shape of the IMF, since an adjustment of the total number of stars might be required to keep the $M_V^{tot}$ value fixed.  The layout of the paper is the following. In Section~\\ref{section:thecode} the ingredients of the stellar population synthesis code are outlined, together with a brief description of the method adopted to derive the integrated quantities. In Section~\\ref{section:GGC} we show the comparison of the theoretical results with selected observations of galactic globular clusters.  Then, we discuss the uncertainties affecting integrated colours (Section~\\ref{section:colours}) and mass--to--light ratios (Section~\\ref{section:ML}), together with a comparison with previous works. ",
        "conclusions": "\\label{section:conclusion}  In this work we have analyzed the intrinsic uncertainties due to stochastic effects on integrated colours and mass--to--light ratios of metal-poor stellar clusters as a function of the total visual magnitude.  The calculations are performed for three different values of $M_V^{tot}$ and for a fine grid of stellar ages. Statistical errors are shown to be crucial, especially in the extreme case $M_V^{tot}=-4$ they are so high as to prevent precise quantitative evaluations for all ages. Calculations are done in the standard UBVRIJHK photometric passbands and in the Hubble Space Telescope bands (WFPC2 and NICMOS systems).  We checked the consistency of our models on observational properties of three metal poor clusters, namely M68, M15, and M30, which differ mainly in the absolute visual magnitude, and HB morphology.  For each cluster we were able to reproduce both the features of the observed CMD and the integrated colours, showing that the HB morphology does not influence the photometric indices taken into account.  A comparisons with recent results available in the literature shows, in some cases, non-negligible differences due to the large variety of prescriptions used in the model calculations.  The uncertainties in the results on both colours and mass--to--light ratios due to the still present uncertainty on the IMF shape have been quantitatively estimated; while the colours fluctuations remains within theoretical uncertainties the ${\\cal M}/L$ ratio is more sensitive to the IMF shape.  We also showed that the influence on ${\\cal M}/L$ of the adopted value for the minimum mass for which carbon burning is ignite is quite negligible."
    },
    "0609/astro-ph0609423_arXiv.txt": {
        "abstract": "Amongst the sources seen in very high gamma-rays several are associated with Pulsar Wind Nebulae (``TeV plerions''). The study of hard X-ray/soft gamma-ray emission is providing an % important insight into the energetic particle population present in these objects. The unpulsed emission from pulsar/pulsar wind nebula systems in the energy range accessible to the \\inte\\ satellite is mainly synchrotron emission from energetic and fast cooling electrons close to their acceleration site.  Our analyses of public \\inte \\ data of known TeV plerions detected by ground based Cherenkov telescopes indicate a deeper link between these TeV plerions and \\inte \\ detected pulsar wind nebulae. The newly discovered TeV plerion in the northern wing of the Kookaburra region (G313.3+0.6 powered by the middle aged \\psrkooka) is found to have a previously unknown \\inte \\ counterpart which is besides the \\vela \\ the only middle aged pulsar detected with \\inte. We do not find an \\inte \\ counterpart of the TeV plerion associated with the X-ray PWN ``Rabbit'' G313.3+0.1 which is possibly powered by a young pulsar. ",
        "introduction": "\\label{intro} Rotation powered pulsar wind nebulae (RPWN) contain an isolated neutron star that drives a relativistic wind of particles into the ambient medium fueling an extended non-thermal emission region.  \\\\ The rotational energy of the neutron star is partially transformed into a highly relativistic wind of particles that is driven~ into the ambient medium which~ (for middle aged RPWN) is potentially modified by the reverse shock of the supernova remnant ejecta \\cite{Blondin01}.  A relativistic standing shock forms at a distance of typically $10^{16}-10^{17}$~cm to the central object and is believed to be the site of particle acceleration. Among the possible ways of accelerating particles, the well-known Fermi acceleration mechanism has been considered to be a good candidate to explain the broad band power law distribution of the pair plasma in the downstream region. ",
        "conclusions": "\\label{summary}  Our preliminary study of TeV plerions with \\inte \\ indicates that some of the pulars associated with these TeV plerions are \\inte sources including young and middle aged RPWN. The \\inte \\ observations are complementary to the TeV observations as the unpulsed hard X-rays are presumably tracing the on-going particle acceleration at the relativistic standing shock. In the frame of ion induced acceleration, the observed hard X-ray spectrum constrains the bulk Lorenz factor of the relativistic wind and its composition."
    },
    "0609/astro-ph0609109_arXiv.txt": {
        "abstract": "We consider a sphere of 7.5 Mpc radius, which contains 355 galaxies with accurately measured distances, to detect the nearest empty volumes. Using a simple void detection algorithm, we found six large (mini)voids in Aquila, Eridanus, Leo, Vela, Cepheus and Octans, each of more than 30 Mpc$^3$. Besides them, 24 middle-size \"bubbles\" of more than 5 Mpc$^3$ volume are detected, as well as 52 small \"pores\". The six largest minivoids occupy 58\\% of the considered volume. Addition of the bubbles and pores to them increases the total empty volume up to 75\\% and 81\\%, respectively. The detected local voids look like oblong potatoes with typical axial ratios b/a = 0.75 and c/a = 0.62 (in the triaxial ellipsoide approximation). Being arranged by the size of their volume, local voids follow power law of volumes-rankes dependence. A correlation Gamma-function of the Local Volume galaxies follows a power low with a formally calculated fractal dimension D = 1.5. We found that galaxies surrounding the local minivoids do not differ significantly from other nearby galaxies on their luminosity, but have appreciably higher hydrogen mass-to-luminosity ratio and also higher star formation rate. We recognize an effect of local expansion of typical minivoid to be $\\Delta H/H_0\\sim (25\\pm15)$\\% ",
        "introduction": "With the discovery of the first cosmic voids in the constellation Bootis (Kirshner et al., 1981) and in other sky regions (Joveer et al., 1978, Gregory and Tompson, 1978, de Lapparent et al., 1986, Rood, 1988), a new branch of extragalactic astronomy --- ``voidology'' was formed. The concept of voids as basic elements of large-scale structure of the Universe proved to be required in modern cosmological models. Zeldovich et al., 1982 considered giant volumes between superclusters which are almost empty of visible objects in the framework of structure formation theories. Ghigna et al. (1994) found that CDM models produce smaller void-probability function (VPF) then the observational one calculated on the volume-limited sub-sample of the Perseus-Pisces survey. Sheth and van de Weygaert (2004) presented a model for the distribution of void sizes and it's evolution in the context of hierarchical scenarios of gravitational structure formation. Furlanetto and Piran (2005) produced a straightforward analytical model of the void distribution and of galaxy populations within voids. Patiri et al. (2005) compared statistical analysis of voids in 2dFGRS with that of cosmological $N$-body simulations and found that galaxies in voids are not randomly distributed but form structures like filaments. Icke (1984) on the basis of numerical studies concluded that voids beetwen filamets are likely to be roughly spherical. Regoes and Geller (1991) found that in their models of structure formation in two and three dimensions, suitable initial conditions lead to cellular structure with faceted voids similar to those observed in redshift surveys. Icke and Van de Weygaert (1991), van de Weygaert (1994) introduced geometrical model - Voronoi foam (tessellations) for the skeleton of cosmic mass on scales $10 - 200$~Mpc that includes voids as building blocks of large scale structure. Sahni et al. (1994) applied adhesion approximation to study the formation and evolution of voids in the universe. Peebles (2001) gave detailed review of modern status in theory and observations of voids.  Cosmic voids are usually considered to have typical dimensions of $\\sim$(20--50) Mpc, but the depth of voids, $\\delta\\rho/<\\rho>$, is the parameter which is still in question. Numerical modeling of the large-scale structure in $\\Lambda$CDM models (Van de Weygaert and van Kampen, 1993, Gottlober et al., 2003, Sheth et al., 2004, Colberg et al., 2005) reproduces cosmic voids of different sizes and density contrast without conflicting with their typical observational parameters.  It is evident that the most reliable characteristics of voids can be obtained by investigating the nearest ones. In his Catalog and Atlas of Nearby Galaxies, Tully (1988) noted the presence in the Local Supercluster (LSC) of the so-called Local void which begins directly from the boundaries of the Local Group and extends in the direction of North Pole of the LSC by $\\sim$20 Mpc. The Local void looks practically free from galaxies. Karachentseva et al. (1998) undertook a search for new dwarf galaxies in the region of Tully void, but failed to find in it dwarf objects with luminosities down to $10^{-3}$ of the luminosity of the Milky Way. Recently Iwata et al. (2005) have measured velocities and distances of galaxies in a vicinity of the Local void and derived some evidence of its expansion. Studying the distribution of nearby galaxies in a volume of radius 10 Mpc, Karachentsev (1994) noted the presence in it of small voids of different sizes completely free from galaxies. As new galaxies were discovered in this Local Volume and their individual distances were specified, the foamlike pattern structure of the Local Volume became more and more clear. This feature may prove to be universal, being a fractal continuation of the spectra of voids to small scales. At the present time, the sample of galaxies with distances less than 10 Mpc numbers about 450 galaxies. For half of them the distances have been measured to an accuracy as high as 8--10\\% (Karachentsev et al., 2004=CNG). The study of properties of voids in the most nearby sample, restricted by the distance, has an obvious advantage since we see here dwarf systems down to their minimum size/luminosity. This gives us unique possibility to define voids as regions empty of any type galaxies. According to scenarios of hierarchically evolving structure in the Universe voids are expected to contain some mass. But in this context our definition seems not to be unphysical because inside our voids we have certain amount of haloes of low masses that don't contain galaxies for different reasons. The absence of the effect of ``God's fingers'' in the Local Volume because of the virial motions of galaxies simplifies the analysis of the shape and orientation of nearby voids.  Below we undertake the first attempt to map voids in the Local Volume on the basis of quantitative reception of their finding. Here we present a list of nearby minivoids, maps of their distribution over the sky, and also discuss their shapes, orientations and characteristics of galaxies neighboring with voids. ",
        "conclusions": "When identifying voids in the Local Volume, we were based on the assumption that minivoids, bubbles and pores are completely free from galaxies. This assumption, important in theoretical models (Peebles, 2001), can be tested in subsequent observations. The \"blind\" survey of the whole southern sky made in the HI line of neutral hydrogen at the 64-m Parks radio telescope led to detection of new dwarf galaxies, in particular, in the Zone of Avoidance of the Milky Way (Kilborn et al., 2002, Staveley-Smith et al., 1998). A similar survey of the northern sky is conducted in Jodral Bank (Boyce et al., 2001). A much more sensitive blind survey in the HI line is now conducting at the 300-m radio telescope in Arecibo (Giovanelli et al., 2005). The minivoids Aquia, Leo, Eridanus as well as six other bubbles of smaller size fall partially within the zone of the Arecibo survey, DEC. = [0, +~38$^{\\circ}$]. We expect that the results of this survey will have a chance to confirm or reject the idea of complete absence of galaxies in nearby voids. The detection of new nearby galaxies and measurement of their distances from the luminosity of red giant branch stars will allow a more detailed study of the shape and kinematics of the nearest voids. As was noted above, the Local void in Aquila, detected by Tully, extends far beyond the limits of the Local Volume. By the present time, the number of nearby galaxies with radial velocities $< 3000$ km s$^{-1}$ has reached about seven thousand. This new observational material presents a possibility of detecting and investigating voids by two orders of magnitude exceeding the volume considered here."
    },
    "0609/astro-ph0609753_arXiv.txt": {
        "abstract": "High-precision radial-velocity techniques, which enabled the detection of extrasolar planets are now sensitive to relativistic effects in the data of spectroscopic binary stars (SBs). We show how these effects can be used to derive the absolute masses of the components of eclipsing single-lined SBs and double-lined SBs from Doppler measurements alone. High-precision stellar spectroscopy can thus substantially increase the number of measured stellar masses, thereby improving the mass-radius and mass-luminosity calibrations. ",
        "introduction": "Precise radial-velocity (RV) measurements, with long-term precisions of a few meters per second, are now routinely obtained by several telescopes around the world. The most notable scientific achievement of precise RV measurements has been the detection of planets orbiting solar-type stars \\citep{MayQue1995,MarBut1996}. In this {\\it Letter} we suggest another application of high-precision RVs, namely, the detection of relativistic effects in the Doppler shifts of close spectroscopic binary stars (SBs). \\citet{KopOze1999} have already detailed the relativistic effects one expects to find in the Doppler measurements of binary stars. Here we focus on the effects that we expect to measure in SBs, and identify the additional information they provide.  The typical velocities of components of close binary stars can be as high as $150\\,\\kms$, $\\beta \\equiv v/c \\sim {\\cal O}(10^{-4})$. The classical Doppler shift formula predicts a relative wavelength shift $\\Delta \\lambda/\\lambda$ of order $\\beta$. The next order corrections are of order $\\beta^2 \\sim {\\cal O}(10^{-8})$, which translate to ${\\cal O}(1\\,\\ms)$. Terms of order $\\beta^3$ are beyond foreseen technical capabilities. We thus limit our analysis to ${\\cal O}(\\beta^2)$ effects --- the transverse Doppler shift (time dilation) and the gravitational redshift. In principle, long term monitoring may reveal higher-order secular terms, such as the relativistic periastron shift or period decay through gravitational-wave radiation, but here we focus only on the periodic effects. These can be detected during relatively short observing runs, assuming the RV measurements are precise enough.   Recently, \\citet{Zucetal2006} have shown that the very same effects should be detectable in the stellar orbits around the black hole in the Galactic Center, after a decade of observations. The context we examine here is different, and the information that can be extracted may contribute to the statistics of close binary orbits. ",
        "conclusions": "We present here an approach to extract more information from RV data of a spectroscopic binary, based on the inclusion of relativistic effects. The practical potential lies in the time-dependent parts of the relativistic terms which are closely linked to the variation of the distance between the binary components. Thus, these terms are especially useful for orbits that are eccentric enough.  The approach is mainly useful in the case of SB2s, but precise photometry or astrometry may add the required information to utilize relativity in SB1s as well. Precise photometry space missions like MOST \\citep{Waletal2003}, CoRoT \\citep{Bag2003}, and Kepler \\citep{Basetal2005} may be able to provide precise enough measurements of $\\omega$ and the inclination of eclipsing binaries from the analysis of their light curves.  In any case, it is essential that the data quality be high enough to be sensitive to variations of the required order, namely one meter per second or less. Currently, the best precision is obtained by the ESO HARPS fiber-fed echelle spectrograph, where the RV error can be as low as $1\\,\\ms$ in certain cases and maybe even less \\citep{Lovetal2005}. In the future, much better precisions can be hoped for, on instruments designed for the ``extremely large telescopes'' \\citep[e.g.,][]{Pasetal2006}.  To demonstrate the relevance of the suggested approach in real-life cases we chose to examine the SB2 12\\,Boo. Recently, \\citet{TomFek2006} published a precise solution of 12\\,Boo based on RVs obtained at the $2.1$-m telescope at the McDonald Observatory and at the Coud\\'{e} feed telescope at Kitt Peak National Observatory, with RV precisions of $0.1$--$0.2$\\,km\\,s$^{-1}$.  The system has a period of $9.6$ days, eccentricity $0.2$, and both RV semi-amplitudes are close to $70$\\,km\\,s$^{-1}$. These orbital parameters translate to an expected relativistic amplitude variation of about $10$\\,m\\,s$^{-1}$ (Eq. \\ref{allterms}).  Furthermore, with a declination of $+25\\degr$, the star is observable by HARPS. Its brightness (5th magnitude) and spectral type (F9IV) make it fairly reasonable to expect a precision of $1$\\,m\\,s$^{-1}$ with HARPS.  We used the available $24$ RVs from \\citet{TomFek2006} and augmented them with only $3$ simulated HARPS measurements (assuming errors of $1$\\,m\\,s$^{-1}$) , including the relativistic effects. For each assumed value of $\\sin i$ we produced $1000$ sets of simulated measurements and solved for the orbital elements, using Equation \\ref{siniexpression} to estimate $\\sin i$. Figure \\ref{simsini} shows the median of the derived values in solid line, and the $25\\%$ and $75\\%$ percentiles in dashed lines.  The figure demonstrates that with reasonable efforts, $\\sin i$ can be measured satisfactorily. In the worst case where $\\sin i=1$, the standard deviation of the derived inclination is $0.14$ and a few more precise measurements can reduce this value significantly. An additional advantage of this test case is that the inclination of 12\\,Boo has already been measured by interferometry and is known to be $108\\degr$ \\citep{Bodetal2005}. Thus if this test is performed, the derived $\\sin i$ can be compared to the known value.  In real observations, more than three precise measurements may be needed in order to account for differences in zero points between instruments. Furthermore, The results depend crucially on the precision of $\\omega_1'$ and $\\omega_2'$, and a relative error of $\\epsilon$ in $\\omega_2'-\\omega_1'$ will be translated to a relative error of $1.5\\epsilon$ in the absolute masses.  The few precise measurements can be scheduled to optimize the precision of those two elements \\citep[e.g.][]{For2006}.  Currently, precisions of a few meters per second are still difficult to obtain and besides using the best instruments available, there are also several limitations imposed by the star itself. Thus, early-type stars or rapid rotators, where the spectral lines are significantly broadened, do not lend themselves easily to high-precision RV measurements. Stellar oscillations and star spots are also a concern as they can cause apparent RV modulation.  In addition, analyzing SB2s with the same level of precision as SB1s has not been easy until recently, when TODCOR \\citep{ZucMaz1994} was applied successfully to high-precision spectra by several teams \\citep{Zucetal2004,Udretal2004,Kon2005}.  The effects we have examined are most useful when the orbits are eccentric and the RV amplitudes are large enough. Large RV amplitudes are usually typical to close binaries, which are expected to have undergone orbital circularization and usually have vanishing eccentricities. However, eccentricity somewhat increases the RV amplitude, and even relatively wide binaries, with high enough eccentricities, can display quite large RV amplitudes.  Care must be taken to model correctly any other effects of order $\\beta^2$ that might contaminate the data. One such effect is the light-travel-time effect. This effect can be easily approximated to the relevant order by adding the following term to $V_{R1}$ (and a corresponding one to $V_{R2}$): \\begin{equation} \\Delta V_{\\mathrm{LT}} = \\frac{K_1^2}{c}\\sin^2\\left(\\nu+\\omega\\right)\\left(1+e\\cos\\nu\\right) \\end{equation}  An effect which should be analyzed carefully is the tidal distortion of the stellar components, in particular close to periastron. This distortion may affect the spectral lines, introducing line asymmetry, which can bias the estimated Doppler shift. RV extrasolar planet surveys use the line-bisector analysis \\citep[e.g.,][]{Queetal2001} to quantify such time-dependent asymmetries. Further development of this technique may be the key to disentangle the tidal distortion and the relativistic effects.   One important application of the proposed method is to calibrate the low-mass end of the mass-luminosity relation, to better understand the stellar-substellar borderline. This mass regime is still poorly constrained, since low mass SB2s are quite rare due to the special photometric, spectroscopic and geometric requirements \\citep{Rib2006}. Large efforts are in progress to obtain accurate stellar masses in this regime, including adaptive optics, interferometry and in the future space interferometry \\citep[e.g.,][]{Henetal2005}. We propose a new, relatively accessible tool to accomplish this goal, where the only requirements are spectroscopic.  Precise RVs for low-mass SB2s were already measured by \\citet{Deletal1999}.  Using the method presented here, their absolute masses may be derived with a relatively small observational effort. No other method exists yet to derive this information purely from RV measurements."
    },
    "0609/astro-ph0609279_arXiv.txt": {
        "abstract": "We report on our studies of the physical structure of the planetary nebula (PN) NGC7662. Using (3D) Integral Field Spectroscopy we have been able to measure the electron temperature more accurately and at a larger number of radial locations than before. Here we briefly present our method by which we find a strong positive temperature gradient with increasing radius. According to hydrodynamic models a hot halo, when compared to the central star, can be the product of the passage of an ionization front (e.g.\\ Marten 1993). Such a gradient is not found in equilibrium models, and this finding -- when confirmed for other objects -- strongly advocates the use of hydrodynamic models when modeling PN halos. ",
        "introduction": "The data presented here was obtained with the PMAS instrument placed on the 3.5m telescope at Calar Alto in September 2005. In three different pointings, forming a \\emph{mosaic}, the PMAS integral field unit (IFU) covered a distance out to 40{\\arcsec} west of the central star. The wavelength coverage was 3436--5093\\,{\\AA} with a spectral resolution of 3.6\\,{\\AA}. The IFU holds $16\\times16$ separate fibers that each represents a so called \\emph{spaxel} -- covering 1{\\arcsec}$\\times$1{\\arcsec} on the sky. Any number of spaxels can be combined to create a final spectrum with an increased S/N ratio, compared to that of a single spaxel. We have extracted six bands of spaxels from the three pointings. The flux in the spectra of the consecutive bands at increasing radius is a sum of (the spatial width in arcseconds of each band is specified as a subscript): $2_{1}$, $4_{1}$, $16_{2}$, $36_{3}$, $63_{4}$, and $128_{8}$ spaxels, respectively. ",
        "conclusions": ""
    },
    "0609/astro-ph0609086_arXiv.txt": {
        "abstract": "We report the results of near infrared spectroscopy of 11 luminosity selected candidate planetary mass objects (PMOs) in the Trapezium Cluster with Gemini South/GNIRS and Gemini North/NIRI. 6 have spectral types $\\ge$M9, in agreement with expectations for PMOs. 2 have slightly earlier types, and 3 are much earlier types which are probably field stars. 4/6 sources with types $\\ge$ M9 have pseudo-continuum profiles which confirm them as low gravity cluster members. The gravity status of the other cool dwarfs is less clear but these remain candidate PMOs. The derived number fraction of PMOs with M=3-15~M$_{Jup}$ is 1-14\\%, these broad limits reflecting the uncertainty in source ages. However, the number fraction with M$<20~$M$_{Jup}$ is at least $5\\%$. These detections add significantly to the body of evidence that a planetary mass population is produced by the star formation process. ",
        "introduction": "In recent years there have been several reported detections of free floating brown dwarfs with planetary masses in very young galactic clusters (eg. Oasa, Tamura \\& Sugitani 1999; Lucas \\& Roche 2000; Zapatero Osorio et al.2000; Allers et al.2006; Jayawardhana \\& Ivanov 2006). Many of these candidates are not yet confirmed as cluster members with the low temperatures predicted by theoretical models. However, at least 1 confirmed PMO probably has M$<10$~M$_{Jup}$: Cha 1109-7734 (Luhman et al.2005). Here we adopt a planetary mass threshold of 15~M$_{Jup}$ (0.014~M$_{\\odot}$) in order to include sources on the deuterium burning limit, which is at 13-14~M$_{Jup}$ for solar abundance (Chabrier et al.2000; Burrows et al.1997) but may be marginally higher in Orion, which has a low D/H ratio (see O'Dell 2004 and references therein).  Lucas, Roche \\& Tamura (2005) (hereafter LRT) detected 33 candidate PMOs in a deep imaging survey of the Trapezium Cluster, and showed that $<$13\\% of cluster members have masses in the range 3-13~M$_{Jup}$. These masses were estimated from source luminosities and an assumed age of 1~Myr, using the isochrone of the Lyon group. Here we present spectra of a sample of candidates chosen from LRT and Lucas \\& Roche (2000) in an attempt to constrain the inital mass function (IMF) at planetary masses. In Section 2 the observations are described. In Section 3 we determine spectral types and surface gravity status and assess cluster membership. We then compare the results with the predictions of pre-main sequence models and discuss the IMF. ",
        "conclusions": ""
    },
    "0609/astro-ph0609565_arXiv.txt": {
        "abstract": "Two little explored aspects of Compton scattering of the CMB in clusters are discussed: The statistical properties of the Sunyaev-Zeldovich (S-Z) effect in the context of a non-Gaussian density fluctuation field, and the polarization patterns in a hydrodynamcially-simulated cluster. We have calculated and compared the power spectrum and cluster number counts predicted within the framework of two density fields that yield different cluster mass functions at high redshifts. This is done for the usual Press \\& Schechter mass function, which is based on a Gaussian density fluctuation field, and for a mass function based on a $\\chi^2$-distributed density field. We quantify the significant differences in the respective integrated S-Z observables in these two models.  S-Z polarization levels and patterns strongly depend on the non-uniform distributions of intracluster gas and on peculiar and internal velocities. We have therefore calculated the patterns of two polarization components that are produced when the CMB is doubly scattered in a simulated cluster. These are found to be very different than the patterns calculated based on spherical clusters with uniform structure and simplified gas distribution. ",
        "introduction": "The pioneering experimental and observational CMB work of our esteemed colleague Francesco Melchiorri has paved the way for generations of students and colleagues in Italy and elsewhere, who are continuing the mission he helped define and lead for more than 30 years. Francesco worked hard to advance experimental and observational capabilities to measure the Sunyaev-Zeldovich (S-Z) effect in nearby clusters. His expertise and leadership resulted in the successful development of MITO, the measurement of the effect in the Coma cluster (De Petris \\ea 2002), and the upgrading of MITO to a bolometer array (Lamagna \\ea 2005). As noted by Rephaeli (2005) in Francesco's obituary, his passing marks the end of the founder's era in the growth of the thriving Italian CMB community.  In light of current and future capabilities of the many upcoming S-Z experiments, we have calculated the power spectrum and cluster number counts in the context of a specific non-Gaussian density fluctuation field. As we discuss in the next section, certain observational results point to the possible need for excess power on cluster scales. The enhanced interest in non-Gaussian models motivates contrasting the integrated S-Z observables of one such viable model with those of the standard Gaussian model whose parameters are well determined.  Future sensitive mapping of the S-Z effect in individual clusters will allow measuring its polarization level. A basic description of some of the S-Z polarization patterns in clusters has already been given by Sazonov \\& Sunyaev (1999). In Section 3 we briefly summarize some of the features of two of the polarization components that are induced when the CMB is doubly scattered in a (non-idealized) cluster whose dynamical state and gas properties are deduced directly from hydrodynamical simulations. ",
        "conclusions": ""
    },
    "0609/gr-qc0609053_arXiv.txt": {
        "abstract": "We construct quasiequilibrium sequences of black hole-neutron star binaries for arbitrary mass ratios by solving the constraint equations of general relativity in the conformal thin-sandwich decomposition. We model the neutron star as a stationary polytrope satisfying the relativistic equations of hydrodynamics, and account for the black hole by imposing equilibrium boundary conditions on the surface of an excised sphere (the apparent horizon).  In this paper we focus on irrotational configurations, meaning that both the neutron star and the black hole are approximately nonspinning in an inertial frame. We present results for a binary with polytropic index $n=1$, mass ratio $M_{\\rm irr}^{\\rm BH} / M_{\\rm B}^{\\rm NS} = 5$ and neutron star compaction $M_{\\rm ADM,0}^{\\rm NS}/R_0 = 0.0879$, where $M_{\\rm irr}^{\\rm BH}$ is the irreducible mass of the black hole, $M_{\\rm B}^{\\rm NS}$ the neutron star baryon rest-mass, and $M_{\\rm ADM,0}^{\\rm NS}$ and $R_0$ the neutron star Arnowitt-Deser-Misner mass and areal radius in isolation, respectively. Our models represent valid solutions to Einstein's constraint equations and may therefore be employed as initial data for dynamical simulations of black hole-neutron star binaries. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609014_arXiv.txt": {
        "abstract": "A large position angle misalignment between the stellar bar and the distribution of dust in the late-type barred spiral \\astrobj{NGC 3488} was discovered, using mid-infrared images from the {\\it Spitzer Space Telescope} and optical images from the Sloan Digital Sky Survey (SDSS). The angle between the bar and dust patterns was measured to be 25\\degree$\\pm$2\\degree, larger than most of the misalignments found previously in barred systems based on H$\\alpha$ or H {\\sc i}/CO observations. The stellar bar is bright at optical and 3.6$\\mu$m, while the dust pattern is more prominent in the 8$\\mu$m band but also shows up in the SDSS {\\sl u} and {\\sl g}-band images, suggesting a rich interstellar medium environment harboring ongoing star formation. This angular misalignment is unlikely to have been caused by spontaneous bar formation. We suggest that the stellar bar and the dust pattern may have different formation histories, and that the large misalignment was triggered by a tidal interaction with a small companion. A statistical analysis of a large sample of nearby galaxies with archival {\\it Spitzer} data indicates that bar structure such as that seen in \\astrobj{NGC 3488} is quite rare in the local Universe. ",
        "introduction": "\\label{Sec1} Bar structure as a major non-axisymmetric feature on all scales is important in studying the morphology, mass and light distributions (e.g., Freeman 1996; Elmegreen \\& Elmegreen 1985; Elmegreen 1996; Elmegreen et al. 1996; Eskridge et al. 2000; Men{\\'e}ndez-Delmestre et al. 2007), star formation (e.g., Zurita et al. 2004; Knapen 2005; Ondrechen \\& van der Hulst 1983; Regan et al. 1996; Sheth et al. 2000), gas dynamics (e.g., Kormendy 1983; Bettoni \\& Galletta 1988; Sancisi et al. 1979; Benedict et al. 1996; Downes et al. 1996; Regan et al. 1999) and central activities (e.g., Ho et al. 1997b; Hawarden et al. 1986; Knapen et al. 2002; Sakamoto et al. 1999; Martini et al. 2003; Sheth et al. 2005) of disk galaxies. Theoretical models, including N-body and hydrodynamic simulations, generally confirm that bar formation is spontaneous and ubiquitous in disk evolution (e.g., Athanassoula 1992; Sellwood \\& Wilkinson 1993; Friedli \\& Benz 1993, 1995; Athanassoula \\& Bureau 1999). Because of the dissipative nature of the interstellar medium (ISM), the streaming motions of the molecular gas in and around bar regions can be different from the stellar orbits (Athanassoula 1992; Regan et al. 1999; Sheth et al. 2002). Due to the delayed star formation after the clouds have been triggered ($\\sim$30Myr; Vogel et al. 1988), the locations of gas/dust in galaxies can often be offset from that of young stars (e.g., Sheth et al. 2002; Phillips 1996; Martin \\& Friedli 1997). The molecular gas can be transported from galactic disk toward central region by the gravitational torques from bars (e.g., Sakamoto et al. 1999; Sheth et al. 2002, 2005), and the condensation of gas leads to subsequent circumnuclear star formation (e.g., Ho et al. 1997b; Knapen et al. 2002; Martini et al. 2003; Jogee et al. 2005; Fisher 2006).  Observationally, the gas/dust patterns can often be seen as dust lanes, atomic and molecular gas concentrations, or isophotes of H {\\sc ii} regions with active star formation (Martin \\& Friedli 1997; Sakamoto et al. 1999; Regan et al. 1999; Rand et al. 1999; Crosthwaite et al. 2000; Sheth et al. 2002, 2005). As predicted by theoretical models (Athanassoula 1992; Friedli \\& Benz 1993, 1995), there is a small position angle misalignment between the gas/dust distribution and the stellar bar, usually of a few (and up to 10) degrees, in the sense that the former is {\\it leading}. Kenney et al. (1991) found the gaseous pattern is offset from the major axis of the stellar distribution by 24\\degree $\\pm$6\\degree\\/ in M~101. Crosthwaite et al. (2000) found that the central gas distribution as indicated by H {\\sc i} map leads the stellar bar by almost 10\\degree\\/ in the late-type galaxy IC~342. Similarly, Rozas et al. (2000) identified a large sample of H {\\sc ii} regions in barred galaxy \\astrobj{NGC 3359} and showed a position angle misalignment of a few degrees exists in H$\\alpha$ and I-band images. They also pointed out that the $u$-band image of this galaxy shows a bar pattern more aligned with  H$\\alpha$, further suggesting massive star formation ``at the leading edge of the bar''. Sheth et al. (2002) found offsets between molecular gas (CO) and star formation (traced by H$\\alpha$) in bars of six nearby spirals, which were caused by the gas flow dependent star formation. Understanding the misalignment between stellar and gas/dust patterns and their formation scenarios is crucial for studying the ISM properties and star formation processes taking place in environments where gas dynamics are strongly perturbed (e.g., Regan et al. 1996; Martin \\& Friedli 1997; Sheth et al. 2000; Zurita et al. 2004), and also offers a good opportunity to study dynamical properties and secular evolution of barred galaxies (e.g., Kormendy 1983; Benedict et al. 1996; Regan et al. 1999; Kormendy \\& Kennicutt 2004; Sheth et al. 2005; Kormendy \\& Fisher 2005; Fisher 2006; Regan et al. 2006).  The {\\it Spitzer Space Telescope}'s (Werner et al. 2004) observations in the mid-infrared, with its higher sensitivity and better angular resolution than previous observations (e.g., {\\it ISO}), provide a new opportunity to study both stellar and gas/dust structures in galaxies (e.g., Pahre et al 2004; Wang et al. 2004; Cao \\& Wu 2007). In particular, the four Infrared Array Camera (IRAC; Fazio et al. 2004) bands from 3.6 to 8.0 $\\mu$m probe both stellar continuum and warm dust emissions (of the so-called polycyclic aromatic hydrocarbon, or PAH, and dust continuum emissions) with identical spatial sampling, thus enabling a powerful probe to compare non-axisymmetric features such as bar structures involving gas/dust and stellar mass. Recently, {\\it Spitzer} observations of nearby galaxies have demonstrated the importance of using mid-infrared images for studying galaxy secular evolution driven by bar instabilities (e.g., Fisher 2006; Regan et al. 2006).  In this paper, we present an analysis of data from {\\it Spitzer} and SDSS of the late-type barred spiral galaxy \\astrobj{NGC 3488}. Previous studies show that, with an estimated distance of 39.9 Mpc (at this distance, 1$''$ corresponds to $\\sim$193 parsecs) and a total infrared luminosity of $L_{\\rm TIR}$ $\\approx$ 4.6$\\times$10$^{9}$ $L_{\\rm \\odot}$ (Bell 2003), \\astrobj{NGC 3488} [Hubble type SB(s)c] has a weak bar ($\\sim$1.5 kpc), with spiral arms beginning at the bar's end but without an inner ring. This is consistent with the conventional view that bars in most late-type spirals are relatively weak (Erwin 2005; Men{\\'e}ndez-Delmestre et al. 2007), and that weak bars tend to produce a SB(s) type response (in which the spiral arms begin at the ends of the bar; Kormendy \\& Kennicutt 2004). The data reduction is presented in $\\S$2, and results on the bar structures in \\astrobj{NGC 3488} with multi-wavelengths analysis are described in $\\S$3. Possible explanations of the large misalignment between the bar and dust patterns are discussed in $\\S$4. ",
        "conclusions": "\\label{Sec4}  \\subsection{Possible Formation Scenarios of the Misalignment Between Bars in \\astrobj{NGC 3488}} Martin \\& Friedli (1997) found that the H$\\alpha$ bars (tracing young stars and H {\\sc ii} regions) and stellar bars are misaligned by up to 10\\degree \\/ among 11 barred galaxies. Friedli \\& Benz (1993, 1995), using their N-body + SPH simulations of spontaneous bar formation, showed that the H {\\sc ii} regions tend to lead the stellar bar by an angle of several degrees. In these simulations the misalignments between bars are similar to those observed, and they occur as a result of orbit crossings of gas motion, mostly at early epochs during the formation of a strong, fast-rotating bar in the absence of an inner Lindblad resonance (Martin \\& Friedli 1997). The best example is \\astrobj{NGC 3359}, in which star formation is completely absent in the galaxy nucleus (Martin \\& Friedli 1997), and the age of the bar is young ($\\sim$400 Myr; Martin \\& Roy 1995). But this is unlikely to be the case in \\astrobj{NGC 3488}, due to the fact that the age of its stellar bar ($\\sim$4 Gyr) is much older than bars in the spontaneous bar formation scenario (e.g., \\astrobj{NGC 3359}).  Alternatively, perhaps the morphology of this galaxy represents an episode shortly after the capture of a small, secondary galaxy. In such a galaxy merger scenario, the stellar and dust patterns may have different formation histories. The stellar bar may have formed previously as in a spontaneous formation scenario, and the dust pattern could be tidally induced later from a dwarf galaxy swallowed by \\astrobj{NGC 3488}. The different stellar populations in the stellar bar and the dust pattern (see \\S3.2) also support this scenario. Berentzen et al. (2003) investigated the dynamical effects of the interaction between an initially barred galaxy and a small companion, using N-body/SPH numerical simulations. They found that the interactions can produce offset bars, nuclear and circumnuclear disks, and tidal arms connected to the end of the bar. Based on their results that the fate of the stellar bar is determined by the impact position, we speculate that in \\astrobj{NGC 3488} there may have been an impact on the bar major axis when the bar was weak. In such a case, the tidal force exerted on the bar does not disrupt much the bar structure (i.e., the stellar bar survived after the impact). The nearby bright, compact object toward the northeast of the SDSS images (see Fig. 3b), identified as SDSS J110124.04+574045.1, is a plausible candidate for a dwarf galaxy that may have hit \\astrobj{NGC 3488}. Its colors are very blue ($u-g = 0.309$, $g-r = -0.222$, $r-i = -0.757$) and similar to those of irregular galaxies (Fukugita et al. 1995). And its absolute magnitude ($-12.73$, $-13.59$, $-13.93$, $-13.82$, $-13.81$ for {\\sl u, g, r, i, z}, respectively) are located at the faint end of the luminosity function of extremely low-luminosity galaxies (Blanton et al. 2005), if we assume that its distance is the same as that of \\astrobj{NGC 3488} (39.9 Mpc). However, we cannot exclude the possibility that it is a foreground white dwarf star superposed on \\astrobj{NGC 3488}, since its very blue colors are consistent with those of spectroscopically identified white dwarf stars in SDSS (Kleinman et al. 2004).  \\subsection{Frequency of Large Misalignments Between Bars Among Galaxies Revealed By {\\it Spitzer}} Large misalignments ($>$10\\degree) between stellar bars (bright at optical and 3.6$\\mu$m) and dust distributions (shown at 8$\\mu$m) are quite rare among nearby barred galaxies with archival {\\it Spitzer} data. We have examined $\\sim$50 barred spirals (SB and SAB) in the {\\it Spitzer} Infrared Nearby Galaxies Survey (SINGS, Kennicutt et al. 2003) and in the Mid-IR Hubble Atlas of Galaxies (a {\\it Spitzer} GTO program, PID: \\#69, PI: G. Fazio; see also Pahre et al. 2004), and found that {\\it none} of them shows a misalignment as large as that found in \\astrobj{NGC 3488}. Most of the galaxies only have a single old stellar bar, which is bright at IRAC 3.6 and 4.5 $\\mu$m but absent at 8$\\mu$m (e.g., \\astrobj{NGC 7080}). Others that have younger bars always show good alignment between 3.6 and 8$\\mu$m (e.g., \\astrobj{NGC 7479}, consistent with the previous result that no misalignment was observed between its H$\\alpha$ and stellar bars; Martin \\& Friedli 1997). This statistical evidence indicates that the bar structure seen in \\astrobj{NGC 3488} is quite rare among barred galaxies in the local Universe, and that the misalignment between bar and dust patterns may be a short-lived phenomenon in the evolutionary history of the galaxy. However, any firm conclusion must await a quantitative analysis of a large, well-defined, and unbiased sample of barred galaxies."
    },
    "0609/astro-ph0609825_arXiv.txt": {
        "abstract": "{} {The aim of this paper is to study the Gamma Ray Burst (GRB) environment through the analysis of the optical absorption features due to the gas surrounding the GRB. } {To this purpose we analyze high resolution spectroscopic observations (R=20000-45000, corresponding to 14 km/s at 4200\\AA\\ and 6.6 km/s at 9000\\AA) of the optical afterglow of GRB050730, obtained with UVES@VLT $\\sim 4$ hours after the GRB trigger.} {The spectrum shows that the ISM of the GRB host galaxy at z = 3.967 is complex, with at least five components contributing to the main absorption system. We detect strong {\\ion{C}{II*}}, {\\ion{Si}{II*}}, {\\ion{O}{I*}} and {\\ion{Fe}{II*}} fine structure absorption lines associated to the second and third component.} {For the first three components we derive information on the relative distance from the site of the GRB explosion.  Component 1, which has the longest wavelength, highest positive velocity shift, does not present any fine structure nor low ionization lines; it only shows very high ionization features, such as {\\ion{C}{IV}} and {\\ion{O}{VI}}, suggesting that this component is very close to the GRB site.  From the analysis of low and high ionization lines and fine structure lines, we find evidences that the distance of component 2 from the site of the GRB explosion is 10-100 times smaller than that of component 3.  We evaluated the mean metallicity of the z=3.967 system obtaining values $\\approx 10^{-2}$ of the solar metallicity or less. However, this should not be taken as representative of the circumburst medium, since the main contribution to the hydrogen column density comes from the outer regions of the galaxy while that of the other elements presumably comes from the ISM closer to the GRB site. Furthermore, difficulties in evaluating dust depletion correction can modify significantly these values.  The mean [C/Fe] ratio agrees well with that expected by single star-formation event models. Interestingly the [C/Fe] of component 2 is smaller than that of component 3, in agreement with GRB dust destruction scenarios, if component 2 is closer than component 3 to the GRB site.} ",
        "introduction": "For a few hours after their onset, Gamma Ray Bursts (GRBs) are the brightest beacons in the far Universe, offering a superb opportunity to investigate both GRB physics and high redshift galaxies.  Tens of minutes after a GRB, its optical afterglow can be as bright as magnitude 13--15.  High resolution (a few tens of km/s in the optical band), high quality (signal to noise $>10$ per resolution element) spectra can therefore be gathered, provided that the afterglow is observed on such a short time scale by 8m class telescopes.  Fiore et al. (2005) studied the UVES high resolution spectroscopy of GRB020813 and GRB021004 afterglows; they showed that the InterStellar Medium (ISM) of GRB host galaxies is complex when resolved down to a width of a few tens of km/s, with many components contributing to each main absorption system, and spanning a total velocity range of up to thousands of km/s. Velocity ranges of hundreds of km/s were found on the high resolution spectra obtained with MIKE@MagellanII and HIRES@Keck on GRB051111 and GRB050730 (Chen et al. 2005, Prochaska, Chen \\& Bloom 2006, PCB06 hereafter, Penprase et al. 2006).  The absorption systems can be divided into three broad categories. First, those associated with the GRB surrounding medium, second, those produced by the ISM of the host galaxy along the line of sight but distant enough not to be strongly modified by the GRB, and third the intergalactic matter along the line of sight. We discuss these three systems in turn.  The physical, dynamical and chemical status of the circumburst medium in the star-forming region hosting the GRB progenitor can be modified by the explosive event, through shock waves and ionizing photons. Blueshifted absorbers may be ionized and even radiatively and/or collisionally accelerated by the GRB and its afterglow. Circumstellar absorption lines in gamma-ray burst afterglow spectra can be used to determine the main properties of the star progenitors (van Marle, Langer,\\& Garcia-Segura 2005). The variability of the intensity of absorption lines can be used to determine the location and density of the absorber, following Perna \\& Loeb (1998), Boettcher, Fryer \\& Dermer (2002), Draine \\& Hao (2002), Mirabal et al. (2002), Perna \\& Lazzati (2002), Perna, Lazzati \\& Fiore (2003). The temperature and density of the absorbing gas can be also accurately determined by studying particular line ratios, namely those of a fine structure transition to its ground state.  High resolution spectra do not reveal only the circumburst matter. They can also be used to probe the ISM of the host galaxy along the line of sight.  The study of the ISM of z$\\gs1$ galaxies has so far relied upon Lyman-break galaxies (LBGs) at z$=3-4$ (see e.g. Steidel et al.  1999) and galaxies which happen to be along lines of sight to bright background quasars (QSOs).  However, LBGs are characterized by pronounced star-formation and their inferred chemical abundances may be related to these regions rather than being representative of typical high-z galaxies.  Faint metal line systems along the line of sight to quasars probe mainly galaxy haloes, rather than their bulges or discs. Taking advantage of ultra-deep {\\it Gemini} multi-object spectrograph observations, Savaglio et al. (2004, 2005) studied the ISM of a sample of faint K band selected galaxies at 1.4$<$z$<$2.0, finding MgII and FeII abundances much higher than in QSO systems and similar to those in GRB host galaxies. Such studies will hardly be extended to higher redshift with the present generation of 8m class telescopes, because of the faintness of high-z galaxies. On the other hand, GRB afterglows provide an independent tool to study the ISM of high z galaxies.  Finally, GRB afterglow high resolution spectra can be used to probe the Ly-$\\alpha$ forest and the high-z intergalactic medium. Using GRBs as remote beacons opens up the opportunity to highlight any deviation from what is already known from quasar forests.  For example the so-called ``proximity effect'' (the feedback of QSOs on their local IGM) should be much reduced for GRBs.  The need to disentangle the contribution to the absorption spectra coming from these three different categories renders high resolution spectroscopy the ideal technique to study GRB optical afterglows.  Therefore we set up a program to observe bright optical afterglows of promptly localized GRBs with UVES@VLT, taking advantage of the capabilities of the GRB-dedicated {\\it Swift} satellite. {\\it Swift} was launched in November 2004, and provides GRB positions ($2-3\\arcmin$ precision) in $<10$~s, X--ray afterglow positions ($5\\arcsec$ precision) in $<100$~s, and an optical finding chart in $<300$~s (few arcsec to sub-arcsec positional accuracy). These characteristics, together with the possibility of quick follow-up with UVES@VLT in RRM (Rapid Response Mode), offers the opportunity to obtain high resolution spectra of GRBs with delays from a few hours down to a few tens of minutes from the event.  To date, five GRBs have been observed with UVES@VLT after a {\\it Swift} trigger: GRB050730, GRB050820A, GRB050922C, GRB060418 and GRB060607A.  In this paper, we present high resolution spectroscopy of GRB050730. This GRB was detected by {\\it Swift} on 2005 July 30 19:58:23 UT (Holland et al. 2005, Sota et al. 2005) at RA(J2000) $=14^h08^m17^s.14$ and DEC(J2000) $= -03^d46^{'}17^{''}.8$. $133$ seconds later, {\\it Swift}'s XRT started to observe the X-ray afterglow in WT mode (Barthelmy et al. 2005). Chen et al. (2005) found the redshift of the GRB at $z =3.967$ with the MIKE echelle spectrograph on Magellan II, starting the observation $\\sim 4$ hours after the burst.  The redshift was confirmed by Rol et al. (2005) using the ISIS spectrograph on the William Herschel Telescope at the Observatorio del Roque de los Muchachos on La Palma.  Starling et al. (2005) show an early WHT ISIS optical spectroscopy of GRB050730, finding that the host is a low metallicity galaxy, with low dust content. PCB06 present an analysis of the MIKE high resolution spectrum of GRB050730, treating the various absorption features as single Voigt profiles. Our main goal is instead that of trying to disentangle the relative contribution of the various components (in velocity space) contributing to each absorption system, making a step forward with respect to Chen, et al. (2005) and PCB06.  We focus here on the study of the local medium surrounding the GRB, and in particular on the analysis of the fine structure absorption lines produced by the gas in the host galaxy. The comparison of such absorption features and their corresponding ground states yields information about the gas density, the temperature and the radiation field.  A detailed analysis of the intervening absorption systems and of the Ly-$\\alpha$ forest will be presented in a forthcoming paper.  The paper is organized as follows. Section $2$ presents the observations and data reduction; Section $3$ gives a brief description of the absorbing systems identified from the spectra; Section $4$ is devoted to the analysis of the fine structure lines and to the estimate of the physical quantities characterizing the host galaxy gas; in Section $5$ the results are discussed and conclusions are drawn.  \\begin{figure} \\centering \\includegraphics[angle=-90,width=9cm]{6479f1.ps} \\includegraphics[angle=-90,width=9cm]{6479f2.ps} \\includegraphics[angle=-90,width=9cm]{6479f3.ps} \\caption{The full high resolution spectrum of GRB050730. Dashed lines represent the noise spectrum. The spectrum is cut-off at the bluemost and redmost end, where the signal-to-noise ratio becomes too small. } \\label{spe1} \\end{figure} ",
        "conclusions": "In this paper we present high resolution (R=20000-45000, corresponding to 14 km/s at 4200\\AA\\ and 6.6 km/s at 9000\\AA) spectroscopy of the optical afterglow of Gamma Ray Burst GRB050730, observed by UVES@VLT $\\sim 4$ hours after the trigger.  We confirm that the redshift of the host galaxy is z=3.96764. Four intervening systems between z = 3.56 and z = 1.77 have been identified along the GRB line of sight, in addition to the main system at the redshift of the host galaxy.  For what concerns the main absorption system, the spectrum shows that the ISM of the GRB host galaxy is complex, with at least five components contributing to this main absorption system at z= 3.967. Such components are identified in this paper with progressive numbers with decreasing velocity values.  These absorption lines appear in different flavors, from neutral hydrogen (Ly-$\\alpha$, Ly-$\\beta$, Ly-$\\gamma$), to neutral metal-absorption lines ({\\ion{O}{I}}), and via low ionization lines ({\\ion{C}{II}}, {\\ion{Si}{II}}, {\\ion{Mg}{II}}, {\\ion{Al}{II}}, {\\ion{Ni}{II}}, {\\ion{Fe}{II}}, {\\ion{P}{II}}, {\\ion{S}{II}}) to high-ionization absorption features ({\\ion{C}{IV}}, {\\ion{Si}{IV}}, {\\ion{N}{V}}, {\\ion{O}{VI}}).  In addition, we detect strong {\\ion{C}{II*}}, {\\ion{Si}{II*}}, {\\ion{O}{I*}} and {\\ion{Fe}{II*}} fine structure absorption features in components 2 and 3 (no {\\ion{Fe}{II*}} in the last component). The ratio of these lines to their ground states are used to constrain the gas density and temperature, under the assumption that they are produced via collisional excitation, or alternatively by UV pumping.  If collisional excitation is at work, we find that the gas producing the fine structure features has a temperature of a few thousand K, and an electron density $n_e =10^3\\div 10^6$ cm$^{-3}$ (component 2) or $n_e =10\\div 10^2$ cm$^{-3}$ (component 3). On the other hand, if UV pumping is the process responsible for the fine structure features, the radiation field exciting the gas has an intensity of $G/G_0 = 10^5 \\div 10^6$.  We speculate that the excitation mechanisms of the fine structure lines may be different, with component 3 probably excited by collisions, and component 2 by UV pumping. This is because the observed {\\ion{O}{I*}}/{\\ion{O}{I}} ratio in component 3 would imply a similar {\\ion{Fe}{II*}}/{\\ion{Fe}{II}} ratio in the case of UV pumping, which is not observed. On the other hand, the radiation field intensities calculated from the fine structure ratios for component 2 are all in agreement with each other.  Component 1, the one with the longest vawelenght, highest positive velocity shift, does not present any fine structure or low ionization lines; it only shows very high ionization features, such as {\\ion{C}{IV}} and {\\ion{O}{VI}}, suggesting that this component is very close to the GRB site.  For the first three components we can derive information on the relative distance to the site of the GRB explosion.  Both for collisional excitation and UV pumping, component 2 appears to be 10-100 times closer to the GRB site than component 3.  As to the metallicity, for the system at z=3.967 we obtain mean values $\\approx 10^{-2}$  of the solar metallicity or less. However, this should not be taken as representative of the circumburst medium, since the main contribution to the hydrogen column density comes from the outer regions of the galaxy while that of the other elements presumably comes from the ISM closer to the GRB site and since the correction for dust depletion is highly uncertain.  It is however interesting to note that the mean [C/Fe] ratio agrees well with that expected by the single star-formation event models of Pipino and Matteucci (2004, 2006).  We could also measure the [C/Fe] of the single components 2 and 3. The [C/Fe] ratio of component 2 is -0.15$\\pm$0.13, marginally different from that of component 3, 0.54$\\pm$0.23. Although the result is significant at only 2.6$\\sigma$ confidence level, it is interesting to speculate on its origin. If the cloud responsible for the absorption component 2 is closer than the cloud producing component 3 to the GRB site, we expect that dust destruction in this cloud, due to the interaction of the GRB radiation field with the medium, is more efficient than in the cloud responsible for component 3.  Since silicates tend to be destroyed more efficiently than graphite (Perna, Lazzati \\& Fiore 2003), we expect more iron and silicon freed in the gas phase in cloud 2 than in cloud 3, as observed."
    },
    "0609/astro-ph0609222_arXiv.txt": {
        "abstract": "The Australia Telescope Compact Array (ATCA) has been used to image class I methanol masers at 9.9, 25 (a series from J=2 to J=9), 84, 95 and 104~GHz located in the vicinity of IRAS 16547-4247 (G343.12-0.06), a luminous young stellar object known to harbour a radio jet.  The detected maser emission consists of a cluster of 6 spots spread over an area of 30\\arcsec. Five spots were detected in only the 84- and 95-GHz transitions (for two spots the 84-GHz detection is marginal), while the sixth spot shows activity in all 12 observed transitions. We report the first interferometric observations of the rare 9.9- and 104-GHz masers. It is shown that the spectra contain a very narrow spike ($<0.03$~\\ks) and the brightness temperature in these two transitions exceeds $5.3\\times10^7$ and $2.0\\times10^4$~K, respectively. The three most southern maser spots show a clear association with the shocked gas traced by the H$_2$ 2.12~$\\mu$m emission associated with the radio jet and their velocities are close to that of the molecular core within which the jet is embedded. This fact supports the idea that the class~I masers reside in the interface regions of outflows. Comparison with OH masers and infrared data reveals a potential discrepancy in the expected evolutionary state.  The presence of the OH masers usually means that the source is evolved, but the infrared data suggest otherwise. The lack of any class~II methanol maser emission at 6.7~GHz in the source raises an additional question, is this source too young or too old to have a 6.7~GHz maser? We argue that both cases are possible and suggest that the evolutionary stage where the class I masers are active, may last longer and start earlier than when the class II masers are active. However, it is currently not possible to reveal the exact evolutionary status of IRAS 16547-4247. ",
        "introduction": "Methanol masers are commonly found in massive star-forming regions. They fall into 2 categories first defined by \\citet{bat88} and further considered by \\citet{vor05b}: class~II masers (of which the 6.7 GHz is the best known and usually strongest) are closely associated with infrared sources and reside in the close environment of exciting (proto-)stars. In contrast, class~I masers (e.g. 95 GHz) are usually found offset (by up to a parsec) from continuum sources \\citep*[e.g.,][]{kur04}. Theoretical calculations strongly suggest that class~I masers are pumped via collisions with the molecular hydrogen in contrast to class II masers, which have a radiative pump \\citep[e.g.,][]{cra92,vor99,vor05b}.  Class I masers are relatively poorly studied. There is direct observational evidence to suggest a relationship between these masers and outflows in a number of sources \\citep{pla90,men96}, although another scenario, which involves cloud-cloud collisions, has also been proposed \\citep*{sob92,meh96,sal02}. The common point of these two scenarios is the presence of shocks. There is observational evidence that methanol abundance is significantly increased in the shock processed regions \\citep[e.g.,][]{gib98} and the gas is expected to be heated and compressed in such regions providing more frequent collisions and, therefore, more efficient pumping.  \\begin{table*} \\caption{Dates of observations, array configurations and observed transitions} \\label{obstab} \\begin{tabular}{@{}r@{}c@{}r@{ }r@{}c@{}c@{}crrc} \\hline \\cthead{UT}&\\cthead{Array}& \\cthead{Transition}&\\cthead{Frequency}&\\cthead{Velocity}&\\cthead{Primary}&\\multicolumn{2}{c}{Synthesized beam}& \\cthead{Spectral}&\\cthead{Velocity coverage}\\\\ \\cthead{Date}&\\cthead{configuration}& \\multicolumn{2}{c}{}&\\cthead{Uncertainty}&\\cthead{beam}&\\cthead{FWHM}&\\cthead{pa}& \\cthead{Resolution}&\\\\ &&&\\cthead{(MHz)}&\\cthead{(\\ks)}&\\cthead{(arcmin)}& \\cthead{(arcsec)}&\\cthead{(\\degr)}& \\cthead{(\\ks)}&\\cthead{(\\ks)}\\\\ \\hline 2 May 2005 & 750A &$\\hphantom{1}9_{-1}-\\hphantom{1}8_{-2}$~E\\hphantom{$^+$}& \\hphantom{2}9936.202\\hphantom{0}~(4)& 0.12\\hphantom{0} &5.09 & \\multicolumn{2}{c}{no imaging}& 0.12\\hphantom{7}& $-$92, $+$28\\\\ &      &$\\hphantom{1}2_{2\\hphantom{-}}-\\hphantom{1}2_{1\\hphantom{-}}$~E\\hphantom{$^+$} & 24934.382\\hphantom{0}~(5) & 0.06\\hphantom{0} & 2.03 & 2.0$\\times$1.3 & 88 & 0.047 & $-$36, $+$41\\\\ &      &$\\hphantom{1}3_{2\\hphantom{-}}-\\hphantom{1}3_{1\\hphantom{-}}$~E\\hphantom{$^+$} & 24928.707\\hphantom{0}~(7) & 0.08\\hphantom{0} & 2.03 & 1.2$\\times$0.8 & 6 & 0.047 & $-$104, $-$27.5\\\\ &      &$\\hphantom{1}4_{2\\hphantom{-}}-\\hphantom{1}4_{1\\hphantom{-}}$~E\\hphantom{$^+$} & 24933.468\\hphantom{0}~(2) & 0.02\\hphantom{0} & 2.03 & 1.2$\\times$0.8 & 14 & 0.047 & $-$46, $+$30\\\\ &      &$\\hphantom{1}5_{2\\hphantom{-}}-\\hphantom{1}5_{1\\hphantom{-}}$~E\\hphantom{$^+$} &24959.0789~(4)& 0.005 & 2.03 & 1.2$\\times$0.7 & 22 & 0.047 & $-$55, $-$18\\\\ &      &$\\hphantom{1}6_{2\\hphantom{-}}-\\hphantom{1}6_{1\\hphantom{-}}$~E\\hphantom{$^+$} &25018.1225~(4)& 0.005 & 2.02 & 1.6$\\times$0.8 & 20 & 0.047 & $-$54, $-$18\\\\ 16 Jun 2005& 6B    &$\\hphantom{1}9_{-1}-\\hphantom{1}8_{-2}$~E\\hphantom{$^+$}& \\hphantom{2}9936.202\\hphantom{0}~(4)& 0.12\\hphantom{0} & 5.09 & 1.6$\\times$1.2& $-$1 & 0.029  & $-$70, $+$26\\\\ &      &$\\hphantom{1}6_{2\\hphantom{-}}-\\hphantom{1}6_{1\\hphantom{-}}$~E\\hphantom{$^+$} &25018.1225~(4)& 0.005 & 2.02 & 0.6$\\times$0.5 & $-$5 & 0.023 & $-$51, $-$13\\\\ &      &$\\hphantom{1}7_{2\\hphantom{-}}-\\hphantom{1}7_{1\\hphantom{-}}$~E\\hphantom{$^+$} &25124.8719~(4)& 0.005 & 2.01 & 0.7$\\times$0.5 & $-$3 & 0.023 & $-$54, $-$16\\\\ &      &$\\hphantom{1}8_{2\\hphantom{-}}-\\hphantom{1}8_{1\\hphantom{-}}$~E\\hphantom{$^+$} &25294.4165~(4)& 0.005 & 2.00 & 0.7$\\times$0.5 & $-$3 & 0.023 & $-$47, $-$10\\\\ &      &$\\hphantom{1}9_{2\\hphantom{-}}-\\hphantom{1}9_{1\\hphantom{-}}$~E\\hphantom{$^+$} &25541.3979~(4)& 0.005 & 1.98 & 0.7$\\times$0.5 & $-$4 & 0.023 & $-$47, $-$10\\\\ 18 Aug 2005&H214C  &$\\hphantom{1}6_{2\\hphantom{-}}-\\hphantom{1}6_{1\\hphantom{-}}$~E\\hphantom{$^+$} &25018.1225~(4)& 0.005 & 2.02 &\\multicolumn{2}{c}{no imaging} & 0.023 & $-$55, $-$8\\\\ &      &$\\hphantom{1}5_{-1}-\\hphantom{1}4_{0\\hphantom{-}}$~E\\hphantom{$^+$} &84521.169~(10)& 0.04\\hphantom{0}  & 0.60 & 2.7$\\times$1.9 & 88 & 0.028 & $-$41, $-$19\\\\ &      &$\\hphantom{1}8_{0\\hphantom{-}}-\\hphantom{1}7_{1\\hphantom{-}}$~A$^+$&95169.463~(10)&  0.03\\hphantom{0}  & 0.53 & 2.3$\\times$1.7 & 87 & 0.025 & $-$41, $-$22\\\\ &      &$11_{-1}-10_{-2}$~E\\hphantom{$^+$}&104300.414\\hphantom{0}~(7)& 0.02\\hphantom{0} & 0.49 & 2.1$\\times$1.5 & 88 & 0.022 & $-$39, $-$22\\\\ \\hline \\end{tabular} \\end{table*}  To further investigate the relationship of different class I maser transitions with outflows and to provide data for theoretical studies, we have made interferometric observations of 12 class I maser transitions towards IRAS~16547-4247 (G343.12-0.06) using the Australia Telescope Compact Array (ATCA). This is a significant fraction \\citep*[50\\% according to][]{mul04} of all known class~I maser transitions, the majority of which belong to the J$_2-$J$_1$~E series. We observed at least one transition from each transition series known to exhibit class~I maser activity~\\citep{vor99} and covered all known class~I maser transitions within the frequency range of existing ATCA receivers. Only the J=5 transition from the J$_2-$J$_1$~E series at 25~GHz has been previously observed towards IRAS~16547-4247 \\citep{vor05a}, although with no position measurement. In this paper we present the first detection of 7 other transitions from this series as well as the first detection of the 84-GHz maser towards this source. The other 3 transitions were detected in a number of single dish and ATCA maser line surveys \\citep[Voronkov et al. unpublished observations]{sly94,val00}.  It is worth mentioning that this source is the only example where all these class~I maser transitions have been detected together, including the relatively rare 9.9 and 104 GHz masers.  IRAS~16547-4257 has recently been the focus of attention as the most luminous young stellar object (YSO) known to harbour a radio jet.  At a distance of 2.9 kpc (Bronfman, private communication), the YSO has a bolometric luminosity of $6.2\\times10^4$ L$_{\\odot}$, equivalent to that of a single O8 Zero Age Main Sequence (ZAMS) star. The jet was first detected by \\citet{gar03} using the ATCA at 1.4, 2.5, 4.8 and 8.6 GHz and subsequently observed using the VLA at 8.5 GHz and 14.9 GHz \\citep{rod05}. This is the first reported case of a radio jet associated with a young O-type star. Located at the centre of a massive dense molecular core (approximately 10$^3$ M$_{\\odot}$), the jet is driving a highly energetic collimated bipolar outflow that extends over 1.5~pc \\citep{bro03,gar06}. ",
        "conclusions": "\\begin{enumerate} \\item The class~I methanol maser emission consists of a cluster of 6 spots spread over an area of 30\\arcsec{ }in extent. Five spots were detected in only the 84- and 95-GHz transitions (for two spots the 84-GHz detection is marginal), while the sixth spot shows activity in all 12 observed transitions. \\item The three most southern maser spots show clear association with a jet-driven molecular outflow. Their velocities are close to that of the molecular core within which the jet is embedded. This fact supports the idea that the class~I masers reside in the interface regions of outflows. \\item Comparison with OH masers and infrared data reveals a potential discrepancy in the expected evolutionary state: the presence of the OH masers usually means that the source is evolved, but the infrared data suggest otherwise. The class~II methanol masers at 6.7~GHz are very widespread and are usually treated as an inevitable phenomenon at some stage of high-mass protostellar evolution (see section~\\ref{evolution_section}). The lack of such a maser in this source (which is active in many class~I maser transitions) raises an additional question concerning its evolutionary state: whether this source is too young or too old to have a 6.7~GHz maser? We argue that both cases are possible because the class~I masers do not require an infrared source for pumping and are harder to destroy than the class~II masers, which are usually located close to the YSO. We, therefore, suggest that the evolutionary stage where the class~I masers are present may last longer than that with the class~II masers, although the exact status of this source is not clear at present. \\item We report the first interferometric observations of the rare 9.9- and 104-GHz masers. The spectra for these transitions contain a very narrow spike ($<0.03$~\\ks) which has a brightness temperature greater than $5.3\\times10^7$ and $2.0\\times10^4$~K at 9.9 and 104~GHz, respectively. \\item High spectral resolution data leads us to suggest that the rest frequency for the $9_{-1}-8_{-2}$~E transition should be refined to 9936.201$\\pm$0.001. \\end{enumerate}"
    },
    "0609/astro-ph0609708_arXiv.txt": {
        "abstract": "LS~5039 is the only X-ray binary persistently detected at TeV energies by the Cherenkov HESS telescope. It is moreover a $\\gamma$-ray emitter in the GeV and possibly MeV energy ranges. To understand important aspects of jet physics, like the magnetic field content or particle acceleration, and emission processes, such as synchrotron and inverse Compton (IC), a complete modeling of the multiwavelength data is necessary. LS~5039 has been detected along almost all the electromagnetic spectrum thanks to several radio, infrared, optical and soft X-ray detections. However, hard X-ray detections above 20~keV have been so far elusive and/or doubtful, partly due to source confusion for the poor spatial resolution of hard X-ray instruments. We report here on deep ($\\sim$300~ks) serendipitous {\\it INTEGRAL} hard X-ray observations of LS~5039, coupled with simultaneous VLA radio observations. We obtain a 20--40~keV flux of $1.1\\pm0.3$~mCrab (5.9 ($\\pm 1.6$) $\\times$ 10$^{-12}$ erg cm$^{-2}$s$^{-1}$), a 40--100~keV upper limit of 1.5~mCrab (9.5 $\\times$ 10$^{-12}$ erg cm$^{-2}$s$^{-1}$), and typical radio flux densities of $\\sim$25~mJy at 5~GHz. These hard X-ray fluxes are significantly lower than previous estimates obtained with BATSE in the same energy range but, in the lower interval, agree with extrapolation of previous {\\it RXTE} measurements. The {\\it INTEGRAL} observations also hint to a break in the spectral behavior at hard X-rays. A more sensitive characterization of the hard X-ray spectrum of LS~5039 from 20 to 100~keV could therefore constrain key aspects of the jet physics, like the relativistic particle spectrum and the magnetic field strength. Future multiwavelength observations would allow to establish whether such hard X-ray synchrotron emission is produced by the same population of relativistic electrons as those presumably producing TeV emission through IC. ",
        "introduction": "\\label{intro}  Microquasars are radio-emitting X-ray binary sources, comprising a compact object (black hole or neutron star) and a companion star, which produce jets through the accretion-ejection mechanism and generate radiation along the electromagnetic spectrum (e.g. \\cite{mirabel99}). For several aspects, they are scaled-down versions of Active Galactic Nuclei (AGN) or quasars, from which their name comes. While less numerous than AGNs, microquasars allow for a more detailed study of some of the AGN important properties thanks to both their closer locations and their faster variability.  As X-ray binaries, microquasars generally produce electromagnetic radiation mainly through accretion and emit most of their luminosity in the soft (0.1--20~keV) and hard ($>$20\\-~keV--1~MeV) X-ray domains. $\\gamma$-ray ($E>1$~MeV) emission from microquasars is rarely observed. On the other hand, a class of AGNs, blazars, are strong $\\gamma$-ray emitters, presenting little accretion disk radiation. The $\\gamma$ radiation produced in blazars is thought in general to come from IC scattering of synchrotron photons by the high-energy particles of the jet, or Synchrotron Self Compton (SSC). This emission, linked to the jet, is therefore particularly important because it gives information on the poorly known ejection process and on the jet composition and magnetic field. However, the study of the $\\gamma$ radiation is hampered by the limited knowledge available on the seed photon field that is boosted by the high energy particles. Microquasars suffer to some extent also the same problem. Moreover, the origin of hard X-ray emission in some microquasars, like LS~5039, is still under debate since accretion disk traces are not present and the jet could be the dominant component, like in blazars. Nevertheless, there is no need at present of large Lorentz factors to explain the $\\gamma$-ray variability in microquasars, unlike the case of their extragalactic relatives.  LS~5039 is a microquasar with slightly variable and extended radio emission \\cite{paredes00,paredes02}. It is a binary system with a $\\sim$3.9 days orbital period composed by an O6.5 V-type donor star and (likely) a black hole as the compact object \\cite{casares05}. LS~5039 is the only microquasar detected at TeV energies by the HESS telescope up to now \\cite{aharonian05,aharonian06} and it is also a $\\gamma$-ray emitter in the GeV range \\cite{paredes00,ribo02t} and possibly in the MeV range \\cite{strong01,collmar04}. The properties of the emission of LS~5039 are quite unlike the ones of more common microquasars like for instance GRS~1915+105 where the thermal emission from the disk dominates the total luminosity. Therefore, LS~5039 is a laboratory in which new aspects of the accretion-ejection mechanism and radiative processes can be studied. The origin of the emission has been proposed to be due either to an accretion-powered relativistic jet (e.g. \\cite{bosch05,dermer06,paredes06}) or to the interaction between the relativistic wind of a young non-accreting pulsar and the stellar wind of the donor star (\\cite{dubus06} and references therein), although the detection of a collimated jet from $\\sim$1 to 100 milliarcsecond scales renders the second scenario unlikely.  In the context of the accretion-powered model, the high energy emission is due to IC scattering in the jet of a seed photon field. The seed photons could be photons produced by synchrotron processes, which generate also the radio emission at lower energies, or by the optical-UV photons of the donor star (see Fig.~\\ref{figsed}). The hard X-ray spectrum and the flux ratio between hard X-rays and GeV-TeV emission differ depending on the dominant energy losses, either synchrotron, Thomson IC or Klein-Nishina IC. Therefore, from the hard X-ray spectral shape and comparing the hard X-rays to the GeV-TeV flux, the magnetic field and even the dominant seed photon field (e.g. stellar or synchrotron), can be inferred. A complete modeling of the emission of LS~5039 is crucial for our understanding of the properties of the X-ray and the $\\gamma$-ray emitting region(s), as well as of the jet itself. Therefore, hard X-ray emission is to be studied because it can shed light on the highest energy part of the electron spectrum. In the following we will present previous X-ray observations of LS~5039 and report serendipitous {\\it INTEGRAL} hard X-ray observations at the position of the source. This is presented altogether with simultaneous VLA observations at 5~GHz, being also discussed in the general context of the source.   \\begin{figure*} \\centering \\includegraphics[width=0.8\\textwidth]{Goldoni_fig1.eps} \\caption{Spectral energy distribution of LS~5039 from radio to TeV energies, a model adapted from \\cite{paredes06} is superimposed in which the most important spectral components are shown. The references for the black and white data points can also be found in \\cite{paredes06}. We remark that the MeV detection obtained with COMPTEL is very probably overestimated due to source confusion. Our VLA detection is represented by a violet square, while our {\\it INTEGRAL} results are represented by red filled circles (see zoom in the inset).} \\label{figsed} % \\end{figure*} ",
        "conclusions": "Our results indicate that LS~5039 is a weak hard X-ray emitter, detected at $\\simeq$1~mCrab in the 20--40~keV energy range and with a flux below 1.5~mCrab in the 40--100~keV interval. These results are in agreement with extrapolation of {\\it RXTE}/PCA data and with previous upper limits from the ISGRI source catalog. However, they are $\\sim$3 times lower than previously reported BATSE values. The {\\it INTEGRAL} measurements can be considered as an average along an orbital cycle. Therefore, the mentioned discrepancy cannot be due to orbital variability. Alternatively, it could be due to long-term variations in the hard X-ray flux of LS~5039, as observed in soft X-rays and in the equivalent width of the H$\\alpha$ line and thought to be caused by long-term variations of the wind of the primary \\cite{reig03}.  On the other hand, the radio flux density measurements indicate that the source was not experiencing a particular low level of activity. Nevertheless, we recall that radio emission variations are small and that no long-term variability has yet been detected at radio wavelengths, suggesting that the radio and X-ray emission are not correlated.  Although these results may allow to refine the jet models proposed by \\cite{paredes06}, see Fig.~\\ref{figsed}, the modeling of the radiation in the whole spectral range is not straightforward. Actually, the relatively hard observed TeV spectrum \\cite{aharonian06} and the X-ray fluxes slightly higher than those observed at TeV energies do not fit in the one-zone leptonic scenario (the particle spectrum at the energies of interest would be dominated by synchrotron losses  instead of by IC Klein Nishina ones, rendering the TeV spectrum much softer than observed). Therefore, the soft X-rays and the TeV photons are probably coming from different regions. In fact, the detection between 20--40~keV and the upper limit in the range 40--100~keV presented here hint to a change in the spectral behavior of the source. This suggests that hard X-rays could be due to synchrotron emission that is produced by the same relativistic electrons that produce TeV radiation through IC.  A longer pointed {\\it INTEGRAL} or {\\it Suzaku} observation is needed to detect any eventual orbital hard X-ray variability, and to detect the source in the 40--100~keV energy range. Only simultaneous soft/hard X-ray and TeV observations would permit to obtain enough spectral and temporal information to fully characterize the relativistic electron spectrum, the seed photon field and the magnetic field strength, and to establish whether the X-ray and the TeV emission have indeed their origin in the same population of electrons."
    },
    "0609/astro-ph0609152_arXiv.txt": {
        "abstract": "MAGIC is currently the world's largest single dish ground based imaging atmospheric Cherenkov telescope. During the first year of operation, more than 20 extragalactic sources have been observed and several of them detected. Here we present results of analyzed data, including discussion about spectral and temporal properties of the detected sources. In addition, we discuss implications of the measured energy spectra of distant sources for our knowledge of the extragalactic background light. ",
        "introduction": "\\label{intro} The search for very high energy (VHE, defined as $E\\geq 100$ GeV) $\\gamma$-ray emission from Active Galactic Nuclei (AGNs) is one of the major goals for ground-based $\\gamma$-ray astronomy. New detections open up a possibility of phenomenological studies of the physics inside the relativistic jets in AGNs, in particular of understanding both the origin of the VHE $\\gamma$-rays as well as the relations between photons of different energy (from radio to VHE).  The number of extragalactic $\\gamma$-ray sources detected by EGRET on board the CGRO with high confidence amounts to 66 \\cite{hartman}.  However, the number of AGNs reported to be VHE $\\gamma$-ray emitters (although slowly increasing) is currently just thirteen (July 2006).   Stecker et al.\\ \\cite{ste} pointed out that the attenuation of $\\gamma$-rays due to photon-photon interactions with low energy photons from the extragalactic background light (EBL) was a likely explanation of this deficit.  In fact, the redshifts of blazars detected so far above 100\\,GeV have rather low values as expected from predictions of the correlation between the gamma-ray attenuation and the redshift of the source, known as Fazio-Stecker relation \\cite{fazio,kneiske}.  All known VHE gamma-ray emitting blazars belong to the class of high-frequency-peaked BL Lacertae objects \\linebreak (HBLs, \\cite{padovani}), a subclass of blazars characterized by a low luminosity when compared with quasars and a synchrotron peak in the X-ray band. Their Spectral Energy Distribution (SED) is characterized by a second peak at very high gamma-ray energies. In synchrotron-self-Compton (SSC) models it is assumed that the observed gamma-ray peak is due to the inverse-Compton (IC) emission from the accelerated electrons up-scattering previously produced synchrotron photons to high energies \\cite{mar}. A compilation of blazars with known X-ray spectra allowing their classification as HBLs is given in \\cite{don}.  The Major Atmospheric Gamma-ray Imaging Cheren-kov (MAGIC) telescope observed a sample of X-ray bright ($F_{1\\,\\rm{keV}}$\\,$>$\\,$2\\,\\rm{\\mu Jy}$) northern HBLs at moderate  redshifts ($z$\\,$<$\\,$0.3$). The sample of candidates was chosen based on predictions from models involving an SSC \\cite{CostGhis} and hadronic \\cite{man2} origin of the $\\gamma$-rays.   The known VHE $\\gamma$-ray emitting AGNs are variable in flux in all wavebands.  Correlations between X-ray and $\\gamma$-ray \\linebreak emission have been found on time scales ranging from several 10 minutes to days (e.g.  \\cite{fossati04}) although the relationship has proven to be rather complicated. The optical-TeV correlation has yet to be studied, but the optical-GeV correlations seen in 3C~279 \\cite{hartman01} suggest that at least in some sources such correlations do exist. Using this as a guideline, the MAGIC collaboration has performed Target of Opportunity observations whenever they were alerted about sources being in a high flux state in the optical and/or X-ray band.  Here we present results of the detected extragalactic \\linebreak sources ordered by increasing redshift values: Mkn~421 \\linebreak (z=0.030), Mkn~501 (z=0.034), 1ES\\,2344+514 (z=0.044), Mkn~180 (z=0.045), 1ES\\,1959+650 (z=0.047), 1ES\\,1218+304 (z=0.182), PG 1553+113 (unknown redshift, z$>$0.09 \\cite{z_new}). The results from galactic observations are presented elsewhere in these proceedings \\cite{rico}. ",
        "conclusions": "\\label{sec:8}  We gave an overview about outstanding findings based on MAGIC observations of extragalactic objects.  The number of extragalactic VHE $\\gamma$-ray sources detected by MAGIC is currently seven.  Two of them, 1ES 1218+304 and Mkn~180 have been discovered by MAGIC. PG 1553+113 has been co-discovered with H.E.S.S. In general, the reconstructed de-absorbed spectra seem to be the harder the further away the sources are, which might be related to an observational bias. Leptonic models (e.g. \\cite{mar}) can describe the data, but there are exceptions like in case of 1ES\\,1959+650.  For detailed modeling of the sources, extensive multiwavelength campaigns including radio through optical telescopes and X-ray satellites are inevitable.  In case of PG 1553+113, we conclude that either the redshift of the source is z$<$0.42 or for the first time a source with a clear evidence of a second emission component in the VHE $\\gamma$-ray energy range has been observed.  Re-observation of the presented sources as well as the analysis of further extragalactic objects is ongoing."
    },
    "0609/astro-ph0609364_arXiv.txt": {
        "abstract": "Compact relativistic stars allow us to study the nature of matter under extreme conditions, probing regions of parameter space that are otherwise inaccessible. Nuclear theory in this regime is not well constrained: one key issue is whether neutron stars are in fact composed primarily of strange quark matter.  Distinguishing the two possibilities, however, has been difficult.  The recent detection of seismic vibrations in the aftermath of giant flares from two magnetars (highly magnetized compact stars) is a major breakthrough. The oscillations excited seem likely to involve the stellar crust, the properties of which differ dramatically for strange stars. We show that the resulting mode frequencies cannot be reconciled with the observations for reasonable magnetar parameters. Ruling out strange star models would place a strong constraint on models of dense quark matter. ",
        "introduction": "\\label{mod}  Most previous studies computing torsional shear modes of the stellar crust assume free slip over the fluid core.  But if the strong magnetic field couples crust and core together then one should instead consider global magneto-elastic perturbations. Recent work on the coupled problem, however, suggests that the vibrations most likely to be observed have very similar frequencies to those computed in the uncoupled problem \\citep{gla06, lev07}\\footnote{The recent calculation by \\citet{lev07} is particularly interesting, as it examines the time-dependence of the interaction between the crust and the core in response to an initial perturbation. The behaviour is complicated, with time-dependence that is not pure oscillatory, but results in periodic amplification at the natural crust frequencies.}.  The natural crust frequencies are therefore physically relevant, so for simplicity we neglect coupling and apply zero traction boundary conditions \\citep{car86}.  Following \\cite{pir05} we use a plane-parallel geometry with constant gravitational acceleration $g = GM/R^2$ in the vertical direction $z$, $M$ and $R$ being the mass and radius star.  The Newtonian equations of hydrostatic equilibrium determine the crust density profile. Using a slab rather than spherical geometry allows us to incorporate the magnetic field without difficulty; we assume a constant field $\\mathbf{B} = B\\hat{z}$. For pure toroidal shear modes, which are incompressible and have no vertical component of displacement, the perturbation equations for the horizontal displacement $\\xi$ reduce to   \\begin{equation} \\frac{\\left(\\mu \\xi^\\prime \\right)^\\prime}{\\rho}  + v_A^2 \\xi^{\\prime \\prime} + \\left[\\omega^2\\left(1 + \\frac{v_A^2}{c^2}\\right) - \\frac{(l^2+ l-2)\\mu}{\\rho R^2}\\right] \\xi = 0 \\end{equation} where $\\mu$ is the shear modulus and $v_A = B/(4\\pi\\rho)^{1/2}$ the Alfv\\'en speed.  Shear speed $v_s = (\\mu/\\rho)^{1/2}$.  We have assumed a periodic time dependence $\\exp(i\\omega t)$, $\\omega$ being the frequency.  We subsequently correct for gravitational redshift to obtain the observed frequency. Primes indicate derivatives with respect to $z$. We have used $\\nabla^2_\\perp \\xi = -[(l+2)(l-1)/R^2] \\xi$ to mimic a spherical geometry, $l$ being the standard angular quantum number.  The scaling with $l$ differs from that used by \\cite{pir05} but gives better agreement between slab and spherical models in the zero field limit \\citep{mcd88,sam06}.   The shear modulus $\\mu$ \\citep{str91} is  \\begin{equation} \\label{shear} \\mu = \\frac{0.1194}{1 + 0.595(\\Gamma_0/\\Gamma)^2}\\frac{n_i (Ze)^2}{a} \\end{equation} Here $Z$ is the atomic number of the ions, $n_i$ the density of ions and $a = (3/4\\pi n_i)^{1/3}$ the average inter-ion spacing. The parameter $\\Gamma = (Ze)^2/ak_B T$, where $T$ is the temperature and  $\\Gamma=\\Gamma_0 = 173$ marks the point at which the solid lattice melts to form an ocean \\citep{far93}, and we use this to determine the upper boundary of the crust.  Crust temperature in the tail of a giant flare is not well constrained:  observations set a lower limit of $10^7$ K \\citep{tie05}, but theory suggests that it could be as high as $\\sim 10^9$ K \\citep{lyu02}, so we examine the range $T = 10^7 - 10^9$ K.  We consider stellar models with $M = 1.2 - 2.5 M_\\odot$ and $R = 8-15$ km, subject to causality constraints, to cover the full range of possible strange star parameter space \\citep{pag06}. Both thin nuclear and nugget crust models are studied \\citep{alc86, jai06}  We model magnetic fields in the range $B = 10^{12} - 10^{15}$ G.  The strong field affects only the very outermost low density layers of the crust \\citep{har06}: we estimate the corrections to computed mode frequencies to be less than 1\\%.  Energetic arguments show that global crust structure, including the suspended crust model, is unaffected.  For the thin nuclear crust there is some uncertainty in the composition of neutron-rich nuclei, which will be constrained by future rare isotope experiments.  We therefore survey a range of equations of state \\citep{hae94, rus06}.  The nugget crust is composed of a lattice of strange quark nuggets embedded in a background degenerate electron gas.  The electrons contribute to the pressure, and the nuggets to the energy density. The density within the crust is given by $\\rho = x\\epsilon_0$, where $\\epsilon_0=n_{\\mathrm{quark}} \\mu_q$ is the energy density of quark matter inside nuggets and depends on the density $n_\\mathrm{quark} \\approx 1$ fm$^{-3}$ and chemical potential $\\mu_q \\approx$ 300 MeV at which stable quark matter vanishes.  $x$ is the volume fraction occupied by nuggets and is given by \\begin{equation} x= \\frac{\\mu_e^3}{3\\pi^2(n_Q - \\chi_Q \\mu_e)} \\,, \\end{equation} where $\\mu_e$ is the electron chemical potential,  $n_Q $ is the electric charge density of the quark nugget and  $\\chi_Q$ is its charge susceptibility \\citep{jai06}. The quark matter parameters are poorly known and can only be determined within the context of specific models. In the Bag model, $n_Q=m_s^2 \\mu_q/2\\pi^2$ and $\\chi_Q=\\mu_q^2/\\pi^2$. Further, requiring SQM to be absolutely stable and simultaneously requiring normal nuclei to be metastable restricts $\\mu_q$ to a narrow range centered around $\\mu_q \\sim 300 $ MeV. Thus within the Bag model, the remaining uncertainty is parametrized through the effective strange quark mass $m_s$ which we expect to be in the range $150-250$ MeV.  The spherical nugget phase occupies most of the crust; we will neglect the small region at the base of the crust containing the pasta phase (see below). In computing shear modulus (Eqn. \\ref{shear}) the quantity \\begin{equation} Z = \\frac{4\\pi}{3} f R_n^3  (n_Q - \\chi_Q \\mu_e) \\end{equation} is the charge of the nugget where $R_n = y~\\lambda_d$ is the typical nugget size, $\\lambda_d = 1/(4\\pi\\alpha \\chi_Q)^{1/2}$ being the Debye screening length and $\\alpha$ the fine structure constant.  The factor $f$ is a correction due to screening inside nuggets, $f \\approx 3(y - \\tanh y)/y^3$, where $y \\approx 1.6$ \\citep{alf06} gives $f \\approx 0.5$ for the larger droplets. The quantity $a = R_n/x^{1/3}$ is the average inter-nugget distance, and $n_i = 3/(4\\pi a^3)$ the density of nuggets.  The variation of $\\mu_e$ with depth is given by integrating the equations of hydrostatic equilibrium.  Using the limits $x=0$ and $x=1$, \\cite{jai06} followed this procedure to estimate $\\Delta R$. In reality the region with $x \\gsim 0.5$ that contains the pasta phase occupies a tiny region with thickness $\\lesssim 1$m.  Further, with increasing $x$ the free energy gain of the heterogeneous state becomes negligible and even a small surface tension can disfavor the large $x$ region. For these reasons we set $x\\simeq 0.5$ at the base of the crust. The value of $x$ at the top of the crust is set by the melting condition: for the parameters examined, it lies in the range $x \\sim 10^{-12} - 10^{-4}$.  Figure \\ref{f1} shows the variation of $v_s$ and $v_A$ with depth for example crust models (compare to Fig. 1 of \\cite{pir05}). The shear speed in the nugget crust is smaller than in the  thin-nuclear crust. This can be understood by noting that at constant pressure $v_s \\sim \\sqrt{Z^{5/3}/A}$ where $A$ denotes the baryon number. Further, and unlike in the nuclear case, both $Z$ and $Z/A$ of the nuggets decrease rapidly with depth.   \\begin{figure} \\begin{center} \\includegraphics[width=8cm, height=6cm, clip]{figure1.eps} \\end{center} \\caption{Shear and Alfv\\'en velocities ($v_s$, $v_A$) in the crust for stars with  $M = 1.4 M_\\odot$, $R = 12$ km, $T = 10^8$ K and $B = 10^{14}$ G.  Top:  Thin nuclear crust (crust thickness $\\Delta R/R$ = 5.4\\%).  Bottom:  Crust with nuggets, for $m_s = 150$ MeV ($\\Delta R/R$ = 0.7\\%) and $m_s = 250$ MeV ($\\Delta R/R$ = 6.1\\%).} \\label{f1} \\end{figure} ",
        "conclusions": "\\label{disc}  For neutron star models torsional shear modes, or their global magneto-elastic equivalents, are a good fit for the observations.  For strange stars, the situation is much more difficult.   A thin nuclear crust model can give a fundamental in the right range if stellar mass is large, but the overtone frequencies are far too high.  The nugget crust model permits a wider range of frequencies due to the uncertainty in the strange quark mass $m_s$. The lowest order $n=0$ modes could explain some of the QPOs in the range 18-150 Hz, but because the fundamental frequencies are so low it is difficult to fit a mode sequence, given the expected scaling with $l$ \\citep{sam06}. For the radial overtones there are regions of parameter space at high temperature with modes at the right frequency, but only for magnetic fields at least an order of magnitude lower than those inferred for magnetars.  At magnetar field strengths only models at the very limits of parameter space permit an overtone frequency as low as 625 Hz.  The other constraint on the nugget crust is the sensitivity of the overtones to temperature fluctuations (greater for this model than for the thin nuclear crust or neutron star crusts).  The observations indicate that the 625 Hz QPO has high coherence and lasts for several hundred seconds, a period during which temperature could vary substantially \\citep{wat06, str06}.  We conclude that the frequencies of toroidal shear modes in strange star crusts have serious difficulty explaining the QPO frequencies observed during magnetar hyperflares.  If the results of \\cite{gla06} and \\citet{lev07} hold true, this conclusion will not change greatly when coupling to core is included in the calculation.  However, these results must now be verified using more sophisticated models that include a realistic stellar geometry, field configuration, and general relativistic corrections.   There may also be other types of modes in the right frequency range \\citep{chu06},  although these should be harder to excite and detect.  We note that there are alternative non-seismic models for the QPOs, but that these models have serious difficulties (for a discussion see \\cite{wat06b}).  The clear distinction between the theoretical predictions for neutron star and strange star crust models is extremely promising.  It offers a robust, largely model independent means of distinguishing strange stars from neutron stars, something that has in the past been lacking.  This type of study should lead to rapid progress in constraining the equation of state of compact stars.   Ruling out the strange star hypothesis would directly impact the phase diagram of QCD at finite chemical potential.  It would offer a strong constraint on models of dense quark matter, indicating that the deconfinement transition is not significantly lowered by the dynamics of the strange quarks. One important implication is that multiply strange hadronic states such as the H-dibaryon in the terrestrial context are less likely. In addition, constraints on the models at finite chemical potential should be relevant for finite temperature extensions of these models which have been employed to describe heavy-ion experiments such as the Relativistic Heavy-Ion Collider (RHIC). Perhaps most importantly, ruling out strange stars directly from observations is the only way to ascertain that terrestrial matter is stable."
    },
    "0609/astro-ph0609634_arXiv.txt": {
        "abstract": "We present long-term ($\\sim 10^4M$) axisymmetric simulations of differentially rotating, magnetized neutron stars in the slow-rotation, weak magnetic field limit using a perturbative metric evolution technique.  Although this approach yields results comparable to those obtained via nonperturbative (BSSN) evolution techniques, simulations performed with the perturbative metric solver require about $1/4$ the computational resources at a given resolution. This computational efficiency enables us to observe and analyze the effects of magnetic braking and the magnetorotational instability (MRI) at very high resolution.  Our simulations demonstrate that (1) MRI is not observed unless the fastest-growing mode wavelength is resolved by $\\gtrsim 10$ gridpoints; (2) as resolution is improved, the MRI growth {\\it rate} converges, but due to the small-scale turbulent nature of MRI, the maximum growth {\\it amplitude} increases, but does not exhibit convergence, even at the highest resolution; and (3) independent of resolution, magnetic braking drives the star toward uniform rotation as energy is sapped from differential rotation by winding magnetic fields. ",
        "introduction": "In differentially rotating neutron stars, an initially weak magnetic field will be amplified by processes such as magnetic braking and the magnetorotational instability (MRI)~\\cite{MRI0,MRI}, causing a redistribution of angular momentum.  Such differentially rotating stars may arise from the merger of binary neutron stars~\\cite{BinMerg-Rasio,BinMerg-ShapShib,BinMerg-Shib}, or from collapse of massive stellar cores, even if the cores spin uniformly at the outset~\\cite{Zwerg} (see also~\\cite{liu01}).  To better understand how magnetic braking affects differentially rotating configurations, Shapiro performed a purely Newtonian, magnetohydrodynamic (MHD) calculation~\\cite{StusFirst} in which the star is idealized as a differentially rotating, infinite cylinder of homogeneous, incompressible, perfectly conducting gas (see also~\\cite{mouschpaleo}).  The magnetic field is taken to be radial initially and is allowed to evolve according to the ideal MHD (flux-freezing) equations. This calculation demonstrates that differential rotation generates a toroidal magnetic field, which reacts back on the fluid flow. Without viscous dissipation, the toroidal field energy and rotational kinetic energy in differential motion undergo periodic exchange and oscillations on the Alfv\\'en timescale. The magnitude of these oscillations, and the maximum field strength, are independent of the initial magnetic field strength; only the growth and oscillation timescale depend on the magnitude of the seed field. If viscosity is present, or if some of the Alfv\\'en waves are allowed to propagate out of the star and into an ambient plasma atmosphere, the oscillations are damped, driving the star to uniform rotation.  Cook, Shapiro, and Stephens~\\cite{JamesCookPaper} later generalized Shapiro's calculations for compressible stars. In their model, the star is idealized as a differentially rotating, infinite cylinder supported by a polytropic equation of state. They performed Newtonian MHD simulations for differentially rotating stars with various polytropic indices and different initial values of $T/|W|$, where $T$ is the rotational kinetic energy and $W$ is the gravitational potential energy. They found that when $T/|W|$ is below the upper (mass-shedding) limit for uniform rotation, $\\beta_{\\rm max}$, magnetic braking results in oscillations of the induced toroidal fields and angular velocities, and the star pulsates stably. However, when $T/|W|$ exceeds $\\beta_{\\rm max}$, their calculations suggest that the core contracts significantly and shock waves are generated, while the outer layers are ejected to large radii to form a wind or an ambient disk.  Liu and Shapiro~\\cite{LiuShap} carried out both Newtonian and general relativistic MHD (GRMHD) simulations on slowly and differentially rotating incompressible stars. They considered the situation in which $T \\ll \\mathcal{M} \\ll |W|$, where $\\mathcal{M}$ is the magnetic energy. Due to the assumptions of slow rotation and weak magnetic field, the star is well approximated as a sphere. They found that toroidal fields are generated by magnetic braking, and the toroidal fields and angular velocities oscillate independently along each poloidal field line. The incoherent oscillations on different field lines stir up turbulent-like motion on tens of Alfv\\'en timescales, a phenomenon called phase mixing (see~\\cite{spruit99} and references therein). In the presence of viscosity, the stars eventually are driven to uniform rotation, with the energy contained in the initial differential rotation going into heat.  Most recently, Duez et al.~\\cite{MHDLett,BigPaper} and Shibata et al.~\\cite{GRB2} performed GRMHD simulations on rapidly, differentially rotating magnetized neutron stars in axisymmetry using two, newly developed GRMHD codes~\\cite{MHDIntro,SS}. They found that if a star is hypermassive (i.e.\\ the star's mass is larger than the maximum mass of a uniformly rotating star), magnetic braking and MRI will eventually induce collapse to a rotating black hole surrounded by a hot, massive torus with a collimated magnetic field aligned along the spin axis. This system provides a promising central engine for short-duration gamma ray bursts~\\cite{GRB2}.  They found the behavior of nonhypermassive, differentially rotating neutron stars to be quite different.  If a star initially spins at a rate exceeding the limit for uniform rotation (``ultraspinning'' case), then instead of collapsing, such a star settles to an equilibrium state consisting of a nearly uniformly rotating core surrounded by a differentially rotating torus.  Although this torus maintains differential rotation, the angular velocity is constant along the magnetic field lines, so further magnetic braking will not occur.  However, if a magnetized star initially exhibits rapid, differential rotation at a spin below the limit for uniform rotation (``normal'' case), the star settles into a uniformly rotating configuration.  The purpose of this paper is to study the same magnetic effects examined by Duez et al.~\\cite{MHDLett,BigPaper}, but in the slow rotation, weak magnetic field limit (i.e., $\\mathcal{M} \\ll T \\ll |W|$).  Analyzing the behavior of stars with such weak, but astrophysically realistic, magnetic fields on the magnetic braking (Alfv\\'en) timescale requires simulations spanning $\\sim 10^4M$ for the models we consider.  Thus the primary challenge of this work is its exorbitant computational expense.  To overcome this difficulty, we adopt a {\\it perturbative} metric approach similar to the one developed by Hartle~\\cite{Hartle,Hartle2}, valid to first order in the angular velocity $\\Omega$.  The second computational challenge is to resolve the wavelength of the dominant MRI mode, which is small for the weak initial fields we wish to treat ($\\lambda_{\\text{MRI}} \\propto B$). Solving the metric equations via perturbation theory allows us to adopt sufficiently high spatial resolution to track MRI.  Two aspects of our perturbative approach make simulations significantly less costly than a nonperturbative, metric evolution at a given resolution and ultimately allow us to perform simulations at roughly $1/4$ the total computational cost.  First, the perturbed metric is time independent except for the $\\phi$-component of the shift, $\\beta^{\\phi}$, which gives rise to frame-dragging.  The shift $\\beta^{\\phi}$ varies with $\\Omega$ on the Alfv\\'en timescale, equivalent to many {\\it thousands} of (Courant) timesteps in a typical simulation.  Our perturbative metric solver uses a simple, ordinary differential equation (ODE) solver to compute the shift and allows us to skip many timesteps between matter evolution updates.  Second, nonperturbative metric schemes incorporate approximate asymptotic outer boundary conditions, which cause problems if the outer boundary is moved too close to the star.  The perturbed metric on the other hand depends only on quantities defined within the (spherical) star, so the outer boundary of the MHD evolution grid may be moved significantly closer to the stellar radius with significant reduction in computational expense.  The validity of our slow-rotation perturbative code is tested to $\\sim 2$ Alfv\\'en timescales ($\\sim 10^4M$).  We evolve a star with both the perturbative and the nonperturbative Baumgarte-Shapiro-Shibata-Nakamura (BSSN)~\\cite{BSSN} gravitational field scheme, and compare the results. We validate our simulations self-consistently by checking that the evolution data satisfy the slow rotation, weak magnetic field assumptions.  We study resolution-dependent MRI effects using two techniques. First, we vary the grid spacing $\\Delta$ at fixed initial magnetic field strength, and second, we vary the initial magnetic field strength ($\\propto \\lambda_{\\text{MRI}}$) at fixed $\\Delta$.  We observe MRI-driven rapid growth of the poloidal fields when $\\lambda_{\\text{MRI}}/\\Delta \\gtrsim 10$ (consistent with~\\cite{MHDLett,BigPaper}) and convergence in the growth rate if $\\lambda_{\\text{MRI}}/\\Delta \\gtrsim 25$. However, convergence in the maximum amplitude of these fields is not achieved even when $\\lambda_{\\text{MRI}}$ is resolved to $\\approx 41$ points. This is due to small-scale turbulence intrinsic to MRI and to axisymmetry.  In addition to an analysis of MRI, we also study magnetic braking. Our simulations indicate that the winding of magnetic fields due to magnetic braking saps a considerable fraction of the energy associated with differential rotation in roughly one Alfv\\'en timscale $t_A$, regardless of resolution or metric evolution technique. Once the rotation profile becomes more uniform, the magnetic fields begin to {\\it unwind}, pumping differential rotation energy back into the star.  In the above simulations, we choose a differentially rotating star with an angular velocity profile that initially decreases with increasing distance from the rotation axis.  In our final simulation, we evolve a differentially rotating star possessing an angular velocity profile that initially {\\it increases} with distance from the rotation axis.  Magnetic braking should occur both models, but MRI should not occur in the later case~\\cite{MRI,MRIrev}. Our code yields the expected result.  The remainder of this paper is structured as follows. In Section~II, magnetic winding and MRI are explained qualitatively. Section~III presents the mathematical and numerical framework for our simulations. Section~IV outlines our initial data and numerical input parameters, Section~V analyzes the validity of our perturbative metric solver, Section~VI discusses our simulation results, and Section~VII summarizes our conclusions. We adopt geometrized units in which $G=c=1$. ",
        "conclusions": "Our perturbative metric evolution algorithm yields results quite similar to those produced by the nonperturbative, BSSN-based evolution scheme.  Because the outer boundary may be moved inward and the metric updated on a physically relevant timescale, perturbative metric simulations may be performed at $\\sim 1/4$ the computational cost of BSSN metric simulations in axisymmetry for slowly rotating, weakly magnetized equilibrium stars.  However, we found that loss of accuracy did occur after magnetic fields hit the outer boundary. This problem could be efficiently solved in future work by extending the outer boundary further from the star.  The stars we study in this paper are weakly magnetized ($\\mathcal{M}/|W| \\sim 10^{-6}$) and slowly rotating ($T/|W| \\approx 0.005$).  As a result, the Alfv\\'en timescale is very long ($\\sim 5000M$), and accurate simulations spanning many Alfv\\'en times would be prohibitively expensive. In this paper, we reliably evolved such stars out to two Alfv\\'en times at moderate resolution ($R/\\Delta = 75$).  In these simulations, we found that magnetic fields wind and unwind on the Alfv\\'en timescale, resulting in a trade-off between the kinetic energy in differential rotation and magnetic energy. Since we could not perform simulations spanning many Alfv\\'en timescales while accurately resolving MRI, we can only speculate, based on the $\\alpha$-disk model and our previous work~\\cite{StusFirst,JamesCookPaper,MHDLett,LiuShap}, that the oscillations between the magnetic and kinetic energy will be damped by MRI-induced dissipative processes over many Alfv\\'en times.  In order to observe MRI in our simulations, we find that the spatial resolution must be set so that $\\lambda_{\\text{MRI}}/\\Delta \\gtrsim 10$, in agreement with the results of~\\cite{MHDLett,BigPaper}.  We have also verified that MRI is not present if the star's angular velocity initially increases with increasing distance from the rotation axis (i.e.,~the MRI is absent when $\\partial_{\\varpi} \\Omega >0$).  We find that as resolution is increased, the effects of MRI become more and more prominent.  Due to the turbulent nature of MRI, we do not achieve convergence of field amplitude, even if $\\lambda_{\\text{MRI}}/\\Delta= 32.6$ or $40.7$.  However, we found the $e$-folding time of MRI, $\\tau_{\\rm MRI}$ does converge. The numerically determined value, $\\tau_{\\rm MRI} \\approx  5.0P_c$ is consistent with the value predicted by the linearized, local Newtonian analysis ($5.72P_c$). The small difference is due in part to the fact that our star is relativistic. In addition, the linearized analysis assumes that $\\lambda_{\\rm MRI}$ is much smaller than the length scale on which the magnetic field changes (i.e.\\ $\\lambda_{\\rm MRI} \\ll B/|\\nabla B|$), which is not quite satisfied in our magnetic field configuration.  Finally, we note that the behavior of the MRI is expected to be different in a full 3D calculation because of the effect of nonaxisymmetric MRI induced by a toroidal magnetic field. Turbulence may arise and persist more readily in 3D due to the lack of symmetry.  More specifically, according to the axisymmetric anti-dynamo theorem~\\cite{hgb95,moffatt78}, sustained growth of the magnetic field energy is not possible through axisymmetric turbulence.  Thus proper treatment of MRI in differentially rotating neutron stars requires high resolution simulations performed in full 3+1 dimensions. The computational cost of such simulations with existing 3+1 metric evolution schemes has thus far been prohibitive, but with our perturbative metric solver it may be possible to perform 3+1 simulations of weakly but realistically magnetized, slowly rotating stars at a small fraction of the computational cost."
    },
    "0609/astro-ph0609258_arXiv.txt": {
        "abstract": "{  We report results from a high-resolution optical spectroscopic survey aimed to search for  nearby young associations and young stars among optical counterparts of ROSAT All-Sky Survey X-ray sources in the Southern Hemisphere. We selected 1953 late-type ($B-V~\\geq~0.6$), potentially young, optical counterparts out of a total of 9574 1RXS sources for follow-up observations. At least one high-resolution spectrum was obtained for each of 1511 targets. This paper is the first in a series presenting the results of the SACY survey. Here we describe our sample and our observations.  We describe a convergence method in the (UVW) velocity space to find associations. As an example, we discuss the validity of this method in the framework of the $\\beta$~Pic Association. ",
        "introduction": "Post-T Tauri Stars (pTTS), and its prototypical example FK Ser \\citep{Herbig73}, were introduced as a new, systematic, category by \\cite{Herbig1978} as T Tauri stars (TTS) that follow classical and weak line TTS (cTTS and wTTS) in an evolutionary sequence. The pTTS were expected to outnumber the cTTS and wTTS populations by far, but their discovery and identification remained difficult for a long time. The relation of pTTS to \"isolated\" TTS \\citep{Quast1987} (young, low-mass pre-main sequence stars found spatially far away from any apparent dark or parental molecular cloud) remained an open problem. Interestingly, Herbig's catalog of pTTS, and the suggested list of isolated TTS produced independently by the Brazilian group are actually quite similar.    An intriguing case is TW~Hya, a high Galactic latitude object, confirmed as TTS \\citep{RucinskiKrautter1983}, with a distance of at least  13$\\degr$ from the nearest  dark clouds. Its origin remained enigmatic for a long time.  A systematic search for more isolated TTS was pursued with the  optical spectroscopic Pico dos Dias Survey (PDS) among  optical counterparts of the IRAS Point Source Catalog \\citep{gregorio92, Torres1995, Torres1998}. One of the first results of the PDS was the discovery of four additional TTS  around TW~Hya  \\citep{delareza89, gregorio92}. They concluded that this group was likely a very young association relatively close to the sun.  The putative TW~Hya association was actually a very odd result. In fact, in a recent review,  \\cite{zuckermanSong04}  reminded us that until the late 1990 the  Hyades and the UMa clusters (with an age about 600~Myr and 300~Myr, respectively) were the only coeval, co-moving concentrations of stars known within 60~pc from Earth. Thus, the closest place to find  very young stars (typically about 10~Myr) was in the nearby classical star forming regions, such as Sco-Cen, Taurus, $\\rho$~Oph (at $\\sim$150~pc) and Orion somewhat farther at $\\sim$450~pc.  Later, \\cite{Kastner97} confirmed that TW~Hya forms a physical association of about 20~Myr and at a distance of 40 to 60~pc from Earth -- the TW~Hya Association (TWA), based on the similarity of the X--ray fluxes, radial velocities, astrometry (Hipparcos) and their spectroscopic characteristics. Even additional members were found later \\citep {webb99, sterzik99, zuckerman01, song03}.  Only very few good candidates for isolated TTS were found with the PDS. But while the IR-excess selection criteria  effectively finds young stellar objects embedded in their placental material or with circumstellar disks, it fails to signal older objects whose disks have been dissipated. Hence most stars  with ages between $\\ga10-70$~Myr (i.e. wTTS and pTTS) escaped the discovery by this method.  Due to the enhanced X-ray activity of young stars \\citep{Walter86} more efficient selection criteria for post and isolated TTS candidates were developed. The high sensitivity and full sky coverage of the ROSAT all-sky survey \\citep{truemper} revealed thousands of new X--rays sources projected in the direction of nearby star forming regions \\citep{Guillout94}. Ground-based spectroscopic studies showed that a large fraction of these X--rays sources were indeed wTTS  together with older pTTS and ZAMS stars. \\citep[e.g.][]{Alcala00}. Surprisingly, many of the newly found weak-line TTS were {\\it not} obviously connected to any molecular cloud region which again raised many questions about their origin \\citep{Sterzik95}.   Young, nearby associations, and in particular TWA, are of great importance to understand the local star forming history, and they reveal samples of stars that allow to study the transition phase of disk dispersal and planet formation in great detail. Therefore several groups have started to look for similar nearby young associations hidden among the ROSAT X--ray sources. Shortly after the discovery of TWA eight more young nearby associations have been announced \\citep[see \\cite{torres03} and the review by][]{zuckermanSong04}, namely,  $\\eta$~Cha \\citep[$d\\sim97$~pc and age $\\sim9$~Myr;][]{mamajek99, Lyo}; $\\epsilon$~Cha \\citep [$d\\sim115$~pc and age $\\sim5$~Myr;][] {mamajek00, feige2003}; the Horologium Association (HorA) \\citep[$d\\sim60$~pc; age $\\sim30$~Myr;][]{torres00} and the adjacent and very similar Tucana Association (TucA) simultaneously discovered by \\cite{zucketal01}; the $\\beta$~Pic Association (BPA)  \\citep[$d\\sim35$~pc; age $\\sim15$~Myr;][]{barrado98, zuckerman01}; the AB~Dor Association \\citep[$d\\sim20$~pc; age $\\sim50$~Myr;][ - calling it as AnA]{zuckerman04, torres03}; and finally two less well defined associations, the Argus Association \\citep[$d\\sim100$~pc; age $\\sim30$~Myr;][]{torres03} and the Octans Association \\citep[$d\\sim110$~pc; age $\\sim30$~Myr;][]{torres03}.   The similarity between the adjacent TucA and HorA motivated us to start SACY (Search for Associations Containing Young stars)\\footnote{Sacy is a mythological Brazilian being, a little black boy with only one leg, that frightens night travelers.} survey in 2000. Some early obsevations, obtained in 1999 for other purposes, were also included in the SACY, whose initial goal was to examine the physical relation between these two associations. Later, SACY was enlarged in order to look for other, hitherto undiscovered, nearby young associations among the optical counterparts of the ROSAT X--ray sources.  Preliminary results of the SACY  have been presented in different conference contributions \\citep[][and references therein]{torres03}, where we have suggested eleven potential nearby associations using the preliminary data (including components of the Sco-Cen Association). One of those first results was the existence of two somewhat superposed associations which we called GAYA1 and GAYA2\\footnote{GAYA stands for Great Austral Young Association.}. GAYA1 is so similar to the  Tucana/Horologium Association proposed by \\cite{zuckermanSong04} that we can adopt this designation, although some of their proposed members seem to match better with GAYA2 ($d\\sim80$~pc; age $\\sim20$~Myr), as we will present in a forthcoming paper. In fact, the present paper is the first one of a series that aims to present the results of the SACY. Here, we present the survey, the observed sample and the observing strategy (Sec. 2). In the associated catalog, we give the spectroscopic properties (radial velocity, $V\\sin(i)$, spectral type, Lithium 6708~\\AA\\ and $H_\\alpha$ equivalent widths) as well as the \\mbox{$UBV(RI)_C$} photometry  of the observed stars. Our kinematical approach is presented and exemplified with the known BPA \\citep{zuckerman01} for which we propose new members. ",
        "conclusions": "In the present paper we present the SACY, a high-resolution optical spectroscopic survey aiming to look for undiscovered nearby associations among the optical counterparts of the ROSAT X-ray sources in the Southern Hemisphere. From an initial list containing 9574 1RXS sources in the Southern Hemisphere, we find 1953 of them having counterparts with $B-V\\geq0.6$ in the Tycho-2  or Hipparcos Catalogues. During the last seven years, 1511 objects were observed at least once. Published information for 115 objects was used to complete our measurements. SACY is $\\sim80\\%$ complete in the Southern Hemisphere.   The main results of this paper are summarized below.  \\begin{itemize}  \\item {\\it Radial velocities and $V\\sin(i)$}. Accurate radial velocities ($\\sigma(V_r)\\la0.3$~km~s$^{-1}$) have been derived for all observed sample using FEROS and/or coud\\'e observations. Good quality $V\\sin(i)$ were derived from FEROS observations. Both quantities, radial velocities and $V\\sin(i)$, were compared to recent measurements obtained with CORAVEL \\citep{Nord04} showing a very good agreement.  \\item {\\it Spectral classification}. Spectral type classification was carried out using a combination of three different methods: i) comparison with the \\cite{montes97} spectral library, ii) calibrations of \\cite{tripicchio97, tripicchio99} for the Na\\,I~D and K\\,I 7699~\\AA\\ lines and, iii) the method developed by \\cite{Torres1998} in the $H_\\alpha$ region.   \\item {\\it Young stars}. Based on our spectroscopic observations Li equivalent width could be measured for the first time for most of the 1256 dwarfs observed in the survey. Using also the information  for 115 stars from the literature and based on the Li criterion we consider, 565, 282 and 524 dwarf stars as respectively younger, having a similar age to or being older than the Pleiades.  \\end{itemize}   Since parallaxes exist only for a small part of the sample, we developed a kinematical convergence method to overcome this problem allowing us to derive dynamic parallaxes. As an example, this method is applied to an initial concentration of young stars at U,V,W=-10,-16,-9~km~s$^{-1}$ which is very close to the proposed mean velocities for BPA by \\citet{zuckermanSong04}.  \\begin{itemize}  \\item {\\it Defining the BPA in (UVW)}. According to Table~\\ref{table:betapic}, the association has a mean (UVW) of ($-10.1\\pm2.1$~km~s$^{-1}$, $-15.9\\pm0.8$~km~s$^{-1}$, $-9.1\\pm0.9$~km~s$^{-1}$).   \\item {\\it New members of the BPA}. Our definition of the BPA is close to that of \\cite{zuckermanSong04}. Using our convergence method and data from the literature we verify that  most of their northern and early-type members fit well into our definition of the BPA. According to our results we propose  13 additional members, including the well known active stars V1005~Ori and V4046~Sgr. This latter is not a \"SACY star\" and was formally considered as typical isolated TTS. Our results suggest that V4046~Sgr belongs indeed to the BPA, being one of the oldest cTTS known up to now.  We propose a wide companion red dwarf star for it. We find also a faint companion for HD~164249. We propose a list of 41 {\\it bona-fide} BPA members.  \\item {\\it Observational isochrone of the BPA}. The 10~Myr isochrone of \\cite{siess00} fits part of the temperature range of the members of the BPA. We then reversed the process and used the color-magnitude diagram of the BPA to define an {\\it ad-hoc} observational isochrone as done by \\cite{zuckerman01}.  \\item {\\it Indication for expansion}. We find  a positive correlation between the X-direction and the U-component of the space velocity vector for young stars in the SACY sample. This behavior is most clearly seen in the case of the BPA. Analysis of this expansion and its implications will be discussed in a future paper.   \\item {\\it Membership probability}. We have developed a membership assessment procedure based on a predictive classification scheme to be used with the convergence method. Both the value of the $F$ and membership probability agree well.  \\end{itemize}"
    },
    "0609/astro-ph0609291_arXiv.txt": {
        "abstract": "{Mass accretion rate determinations are fundamental for an understanding of the evolution of pre-main sequence star circumstellar disks.} {Magnetospheric accretion models are used to derive values of the mass accretion rates in objects of very different properties, from brown dwarfs to intermediate-mass stars; we  test the validity of these models in the brown dwarf regime, where the stellar mass and luminosity, as well as the mass accretion rate, are much lower than in T Tauri stars.} {We have measured simultaneously two infrared hydrogen lines, \\Pab\\ and \\Brg, in a sample of 16 objects in the star-forming region \\Roph. The sample includes 7 very low mass objects and brown dwarfs and 9 T Tauri stars.} { Brown dwarfs where both lines are detected have a ratio \\Pab/\\Brg of $\\sim 2$. Larger values, $\\simgreat 3.5$, are only found among the T Tauri stars. The low line ratios in brown dwarfs indicate that the lines cannot originate in the column of gas accreting from the disk onto the star along the magnetic field lines, and we suggest that they form instead in the shocked photosphere, heated to temperatures of $\\sim 3500$ K. If so, in analogy to veiling estimates in T Tauri stars, the hydrogen infrared line fluxes may provide a  reliable measure of the accretion rate in brown dwarfs.} {} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609772_arXiv.txt": {
        "abstract": "We conducted a radial velocity survey of the Cygnus OB2 Association over a 6-year (1999~--~2005) time interval to search for massive close binaries.  During this time we obtained \\numspectra\\ spectra on \\numstars\\ OB stars to measure mean systemic radial velocities and radial velocity variations.  We spectroscopically identify \\newOBstars\\ new OB stars for the first time, the majority of which are likely to be Association members.  Spectroscopic evidence is also presented for a B3Iae classification and temperature class variation (B3~--~B8) on the order of 1 year for Cygnus OB2 No.~12.  Calculations of the initial mass function with the current spectroscopic sample yield $\\Gamma = -2.2 \\pm 0.1$.  Of the \\numstarcomp\\ stars with the most reliable data, \\varsystems\\ are probable and \\pvarsystems\\ are possible single-lined spectroscopic binaries. We also identify 3 new and 8 candidate double-lined spectroscopic binaries.  These data imply a lower limit on the massive binary fraction of \\varpercent~--~\\pvarpercent.  The calculated velocity dispersion for Cygnus OB2 is $2.44 \\pm 0.07$ \\kms, which is typical of open clusters. No runaway OB stars were found. ",
        "introduction": "Cygnus OB2 may be one of the most massive and richest associations in the Galaxy with $2600 \\pm 400$ OB cluster members \\citep[hereafter \\Kndl]{Knodl2000} and 90 to 100 O stars \\citep[\\Kndl]{Come2002}. It has been studied numerous times due to its richness, proximity (1.7 kpc --- \\citet{Han2003,MT91,Shu58,John54}), and high extinction components with $A_{V} > 5$. The total mass of the cluster is estimated to be $(4-10) \\times 10^{4} \\; M_{\\sun}$ with a central stellar density of $\\rho_{0} = 40-150 \\; M_{\\sun}~pc^{-3}$ (\\Kndl). The radius of Cyg OB2 has been estimated as large as 30 pc (\\Kndl). For the purposes of this study, we adopt a cluster core radius of $\\sim$15 pc (30\\arcmin) inferred from the photometric survey of \\citet[hereafter MT91]{MT91}.  A handful of massive binaries have been discovered in the Association. Among the OB type binaries are MT421, MT429, MT506, MT554, MT696, Cyg OB2 No.~5, and Cyg OB2 No.~8a (MT465).  The first four of these are eclipsing binaries of the Algol type \\citep{PJ98,Kaz2000}. \\citet{Rios2004} discovered that MT696 is an early contact binary (W UMa type) consisting of a late O and an early B star with a period of 1.46 days. Cyg OB2 No.~5 is possibly a triple system consisting of an O7Ianfp + Ofpe/WN9 contact binary with period of 6.6 days \\citep{Leung78} and a B0V star \\citep{Ben2001,Wal73,Boh76}. Cygnus OB2 No.~8a is a non-thermal radio emitter \\citep{bieg89} and massive binary system consisting of O6 and O5.5 stars with a period of $\\sim$22 days \\citep{Debeck2004}.  Massive close binaries (MCB) are the progenitors of some classes of energetic phenomena such as supernovae, $\\gamma$-ray bursts, and X-ray binaries \\citep{fryer99,fryer98}.  They are also laboratories in which to study the formation mechanisms for massive stars \\citep{bonnell98}. Direct and indirect evidence reveals that up to 50~--~80\\% of massive stars reside in binary systems \\citep{vanbev2003,gies87,merm2001}. Studies of the MCB frequency are numerous and include \\citet{Hill2006},\\citet{vanbev2003},\\citet{zin2003},\\citet{merm2001}, \\citet{bonnell98},\\citet{heuvel84}, and \\citet{garmany80}.  Results have been all but conclusive.  The binary frequency for Galactic B0~--~B3 stars is at least 32\\%, and it may be anywhere between 14\\% and 80\\% for OB stars in general \\citep{vanbev2003, vanbev98, merm2001}.  In addition, \\citet{vanbev2003} speculates that there may not be a standard MCB frequency for open clusters and associations. The long term goal of this study is to provide a measurement of the MCB frequency, binary star mass ratios and orbital separations of Cygnus OB2 stars.  Section 2 presents our observations and data reduction. Section 3 discusses the identification of new early type stars and our derivations of cluster quantities such as the initial mass function, visual extinction, and distance.  Section 4 describes the calculation of radial velocities. Section 5 presents the results and a measurement of the cluster velocity dispersion. Section 6 summarizes the survey findings. ",
        "conclusions": "We conducted a radial velocity survey of the Cygnus OB2 Association over a 6-year time interval to search for MCBs using spectroscopic data from the Keck, Lick, WIYN, and WIRO Observatories.  We obtained \\numspectra\\ spectra to measure radial velocities and radial velocity variations on \\numstars\\ OB stars. There were \\newOBstars\\ identified as new early types. The calculated mean distance modulus for Cyg OB2 stars is $\\sim$11.3 mag, which is in good agreement with previous estimates. Of the 146 total OB stars, analysis of the 143 provisional members yielded an IMF slope of $\\Gamma = -2.2 \\pm 0.1$. There were a number of minor spectral classification differences, including Cygnus OB2 No.~12 and No.~21. No.~12 showed evidence of a B3Iae spectral classification in at least one epoch and a temperature class variation (B3~--~B8) over one year.  We utilized two methods for determining velocity variations, including Gaussian profile fitting and cross correlation techniques. Both methods yielded similar results, where \\varsystems\\ stars had a probability $P(\\chi^2,\\nu)\\leq0.01$ and \\pvarsystems\\ stars had a probability of $0.01 < P(\\chi^2,\\nu)\\leq 0.04$. In addition, we detected \\dsbtwos\\ SB2 systems and \\sbtwos\\ possible SB2 systems (3 of which were also designated SB1). This translates to a lower limit on the massive binary frequency of \\varpercent (36 out of 120 stars) to \\pvarpercent (50 out of 120 stars). The calculated velocity dispersion for Cygnus OB2 is $2.44 \\pm 0.07$ \\kms, which is typical of open clusters, and despite the richness of the association and the number of stars surveyed, we detected no obvious OB runaways."
    },
    "0609/astro-ph0609766_arXiv.txt": {
        "abstract": "We model the electromagnetic emission signatures of massive black hole binaries (MBHBs) with an associated gas component. The method comprises numerical simulations of relativistic binaries and gas coupled with calculations of the physical properties of the emitting gas. We calculate the accretion powered UV/X-ray and H$\\alpha$ light curves and the H$\\alpha$ emission line profiles. The simulations have been carried out with a modified version of the parallel tree SPH code {\\it Gadget}. The heating, cooling, and radiative processes for the solar metallicity gas have been calculated with the photoionization code {\\it Cloudy}. We investigate gravitationally bound, sub-parsec binaries which have not yet entered the gravitational radiation phase. The results from the first set of calculations, carried out for a coplanar binary and gas disk, suggest that the outbursts in the X-ray light curve are pronounced during pericentric passages and can serve as a fingerprint for this type of binaries if periodic outbursts are a long lived signature of the binary. The H$\\alpha$ emission-line profiles also offer strong indications of a binary presence and may be used as a criterion for selection of MBHB candidates for further monitoring from existing archival data. The orbital period and mass ratio of a binary could be determined from the H$\\alpha$ light curves and profiles of carefully monitored candidates. Although systems with the orbital periods studied here are not within the frequency band of the {\\it Laser Interferometer Space Antenna} ({\\it LISA}), their discovery is important for understanding of the merger rates of MBHBs and the evolution of such binaries through the last parsec and towards the detectable gravitational wave window. ",
        "introduction": "Massive black hole binaries (MBHBs) are expected to be the most luminous sources of gravitational radiation in the universe. The simultaneous identification of the gravitational radiation and electromagnetic counterpart from MBHBs would allow unprecedented tests of the physics of massive black holes, such as accretion, and also offer an alternative method to constrain cosmological parameters. However, there are only a handful of MBHB candidates known from observations so far and none of them are expected to merge within our lifetimes. Although galaxy mergers are the natural places to look for MBHBs \\citep*{bbr,valtaoja89,milosavljevic01,yu02}, not every such galaxy forms a MBHB which will merge within the Hubble time, hence the difficulty in determining the merger rates. Practical obstacles in the direct identification arise for several reasons. Firstly, fairly high spatial resolution and accuracy in position measurements are required to resolve a binary with two active nuclei. For example, the spatial resolution required to resolve an intermediate binary with orbital separation of 1pc at the distance of $\\sim$100$\\,$Mpc is about 2$\\,$mas. With spatial resolutions currently achievable, it is easier to spot and resolve wide binaries, like NGC 6240, whereas it has been suggested that MBHBs spend a major fraction of their life time as intermediate, hard binaries \\citep{bbr}.  It follows that a better understanding of electromagnetic signatures of galactic nuclei hosting the massive binaries is essential in order to recognize the presence of MBHBs and put tighter constraints on their merger rates. ",
        "conclusions": "Since there are reasons to expect that MBHBs spend the largest fraction of their lives at sub-parsec scales \\citep{bbr} this phase naturally emerges as an optimal one for the study of the observational signatures of such binaries. Based on the first set of calculations, one portion of which is presented here, we find that X-ray outbursts should occur during pericentric passages of a co-planar binary. We suggest that periodical X-ray outbursts should persist as long as there is a supply of the nuclear gas that the MBHB can interact with. A calculation of a much larger number of binary orbits is needed in order to confirm that the periodic outbursts are a long lived signature of the binary. The current limitation of the numerical method used in this study is that the calculation of a large number of orbits along with the hydrodynamics and radiative heating and cooling becomes prohibitively expensive.  Except for the recurrent outbursts in the X-ray light curve the signature of a binary is most easily discernible in the H$\\alpha$ emission line {\\it profiles}. The irregular shape of the broad H$\\alpha$ emission line profiles can be used as a first set of criteria when searching for MBHBs. Once candidates are selected from a large spectroscopic survey (e.g., the SDSS) they can be monitored over long time intervals to look for the time-dependent signature of a binary. Based on simulation described here, for a mass ratio of 0.1, the wavelength shift in the H$\\alpha$ emission lines over the course of an orbital period should be measurable. If one can follow the regular variations of the line over a few cycles, one could constrain the properties of the binary. If the signatures of binaries are found in observations, they could be used to estimate the number of MBHBs in this evolutionary stage and whether MBHBs indeed evolve quickly through the last parsec."
    },
    "0609/astro-ph0609285_arXiv.txt": {
        "abstract": "We present high angular resolution observations toward two massive star forming regions IRAS 18264-1152 and IRAS 23151+5912 with the Plateau de Bure Interferometer (PdBI) in the SiO (J=2-1) and H$^{13}$CO$^{+}$ (J=1-0) lines and at 1.3 mm and 3.4 mm continuum, and with the Very Large Array (VLA) in the NH$_3$ (J,K)=(1,1), (2,2) lines. The NH$_3$ (1,1) and (2,2) emission is detected toward IRAS 18264-1152 only. For IRAS 18264-1152, the SiO observations reveal at least two quasi-perpendicular outflows with high collimation factors, and the most dominant feature is a redshifted jet-like outflow with very high velocities up to about ${\\Delta}v=60$ kms$^{-1}$ with respect to the systemic velocity. The very-high-velocity component (${\\Delta}v=22 - 60$ kms$^{-1}$) of this outflow is spatially offset from its high-velocity (${\\Delta}v=3 - 21$ kms$^{-1}$) component. The SiO line profiles and position-velocity characteristics of these two components suggest that this outflow can be driven by an underlying precessing jet. For IRAS 23151+5912, the bipolar but mainly blueshifted SiO outflow coincides with the inner parts of the single-dish CO outflow. In particular the quasi-parabolic shape of the blueshifted outflow coincides with the near-infrared nebulosity and is consistent with entrainment of the gas by an underlying wide-angle wind. The analysis of the molecular outflow data of the two luminous sources further support high-mass stars forming via a disk-mediated accretion process similar to low-mass stars. ",
        "introduction": "The formation process of massive stars is poorly understood. There are two competing views: One view is that high-mass stars form via infall and disk mediated accretion processes as a scaled-up version of their low-mass counterparts \\citep[e.g.,][]{Jijina96, Garay99, Norberg00, McKee02, Yorke02, McKee03, Keto03}. The other view argues that at the dense center of evolving massive clusters, competitive accretion could become a dominant process. In the extreme case, even the coalescence of low- and intermediate-mass (proto)stars may occur to form the most massive stellar objects \\citep[e.g.,][]{Bonnell98, Stahler00,Bonnell04, Bally05}.  Massive molecular outflows play a very important role in massive star formation. In addition to dissipating excess angular momentum of the infalling material as their low-mass counterparts \\citep{Shu87}, they may inject enough energy to sustain the turbulence and prevent a decrease of the virial parameter \\citep{Quillen05}, whose value may determine whether competitive accretion is effective in cluster forming clouds \\citep{Krumholz05a}. Outflows also lead to significant anisotropy in the distribution of the circumstellar material, which greatly reduces the radiation pressure experienced by the infalling material \\citep{Krumholz05b}. The collimation factor of massive molecular outflows can help to discriminate between the disk mediated accretion process and the coalescence process, since the former requires massive outflows as collimated as their low-mass counterparts and the latter predicts outflows to be far less collimated.  Surveys with single-dish telescopes find that massive molecular outflows are ubiquitous toward high-mass star forming regions, and are far more massive and energetic than their low-mass counterparts \\citep[e.g.,][]{Shepherd96, Zhang01, Beuther02c, Zhang05}. While \\citet{Shepherd96} find that massive outflows appear to be less collimated than their low-mass counterparts, \\citet{Beuther02c} argues that this may be mostly an observational artifact caused by the large distances of the target sources and the low spatial resolution of most studies. In recent years, there have been several high spatial resolution mm interferometric studies of molecular outflows toward massive star forming regions. \\citet{Shepherd99} reveal a massive outflow with a wide opening angle toward G192.16-3.82 and suggest its driven mechanism as a strong wide-angle wind; \\citet{Cesaroni99} report a highly collimated SiO jet toward IRAS 20126+4104, and \\citet{Shepherd00} reveal the same, but a larger scale (2pc) outflow in CO in a different orientation; \\citet{Hunter99} show a well collimated SiO jet toward AFGL 5142; Beuther et al. (IRAS 05358+3543, 2002a; IRAS 19410+2336, 2003; IRAS 19217+1651, 2004) resolve simple outflows previously identified with single-dish observations into multiple and well collimated outflows when observed with interferometers at high spatial resolution. They find that their kinematics are similar to those of low-mass counterparts; \\citet{Shepherd03} present multiple, overlapping massive outflows driven by at least four protostars in the region of W75\\,N. Their energetics and near-infrared observations suggest that they are not likely to be scaled-up versions of jet-driven outflows from low-mass protostars; \\citet{Gibb03} show that a precessing jet-driven flow associated with G35.2-0.7N in single-dish observations can be explained as at least two overlapping flows; \\citet{Su04} reveal a clear bipolar morphology of molecular outflows toward two luminous sources IRAS 21519+5613 and IRAS 22506+5944; \\citet{Kumar04} unveil at least four outflows toward an ultra-compact (UC) H{\\footnotesize II} region Onsala 1; \\citet{Sollins04} resolve a highly energetic bipolar SiO outflow toward UC H{\\footnotesize II} region G5.89-0.39. All of these interferometric observations reveal new morphologies of massive outflows that can not be resolved with single-dish observations and give a wealth of information about the physical process in the innermost parts of the massive star forming regions.  As the statistics of the high spatial resolution interferometric observations are still poor and some basic issues, such as collimation factors and outflow driving mechanism, are being debated, we carry out a study of massive molecular outflows with the Plateau de Bure Interferometer (PdBI) and the Very Larege Array (VLA). Here we use the shock tracer SiO (J=2-1) to image outflows toward two massive star forming regions IRAS 18264-1152 and IRAS 23151+5912 (hereafter we refer as I18264 and I23151, respectively). The dense cores and ambient gas are mapped at 1.3 and 3.4 mm continuum, and in the NH$_3$(1,1) and (2,2) and H$^{13}$CO$^+$ (J=1-0) lines. The two sources are part of a larger sample discussed in detail by \\citet{Sridharan02} and \\citet{Beuther02b, Beuther02d}. \\citet{Beuther02c} present bipolar CO outflows toward the two sources with the IRAM 30m telescope observations. Toward I23151, \\citet{Weigelt06} reveal a cone-like nebulosity, which coincides with the blueshifted CO outflow, in the near-infrared $K^{'}$ band using the bispectrum speckle interferometry method. The kinematic distance of I23151 derived from the CS velocity is 5.7 kpc. The kinematic distance of I18264, suffering from an ambiguity, is either 3.5 or 12.5 kpc \\citep{Sridharan02}. By associating the region via the near- and mid-infrared surveys 2MASS and MSX on larger scales with sources of known distance, the ambiguity was solved to be at the near distance (S. Bontemps, 2003, priv com.). On the other hand, toward the line of sight of I18264, \\citet{Sewilo04} solve the kinematic distance ambiguity of an H{\\footnotesize II} region with velocity of 50.9 kms$^{-1}$ to be at the far distance of 12.0 kpc. For the molecular cloud with velocity of 43.6 kms$^{-1}$, Sewilo et al. consider it is associated with the H{\\footnotesize II} region of 50.9 kms$^{-1}$ and adopt its far distance of 12.4 kpc. Since the distance determinations above both use indirect methods and still give ambiguous results, we give physical parameter estimations corresponding to both the far and near distances in the following sections. Based on the High Resolution (HIRES) IRAS database, \\citet{Sridharan02} estimate the luminosities of the two sources to be $10^5 L_{\\odot}$ for I23151 and $10^4 L_{\\odot}$ and $1.2\\times10^5 L_{\\odot}$ respectively for the near and far distances for I18264. As additional evidence of massive star formation and outflows, Class {\\footnotesize II} CH$_3$OH and H$_2$O maser emission toward I18264 and H$_2$O maser emission toward I23151 is observed \\citep{Beuther02d}. \\citet{Sridharan02} detect no emission at 3.6 cm down to 1 mJy for both sources, while \\citet{Zapata06} detect 1.3 cm and 3.6 cm emission toward I18264 with a better sensitivity. After a description of observations in Sect. \\ref{obs}, we show results in Sect. \\ref{result}. Discussion is presented in Sect. \\ref{dis}, and we conclude in Sect. \\ref{conclusion}. ",
        "conclusions": "\\label{conclusion} We carry out a study of SiO outflows toward the high-mass star forming regions I18264 and I23151. According to the H$^{13}$CO$^+$ observations and SiO line profiles, the SiO emission in these two sources can be caused by the interaction between the jet or wind from the central young stars and the dense ambient gas. The mass-velocity relation of the HC of the SE outflow in I18264 and the blueshifted outflows in I23151 also can be fitted by a broken power law with the slopes steeping at about 10 kms$^{-1}$ with respective to the systemic velocity. With the PV characteristic of the terminal velocities decreasing with the distance from the driving source, the low velocity ($|{\\Delta}v|<10$ kms$^{-1}$) gas probably comprises the ambient gas entrained by the high velocity ($|{\\Delta}v|>10$ kms$^{-1}$) gas.  Toward I18264, at least two quasi-perpendicular outflows with high collimation factors ($\\sim3-4$) are resolved. The VHC of the SE outflow has very high velocities up to ${\\Delta}v\\sim60$ kms$^{-1}$. It is north offset from the HC, and its downstream peak is farther away from the driving source. Comparing with the outflows in the well studied low-mass source L1157, we suggest the SE outflow can be entrained by an underlying precessing jet. The characteristics of the HC and VHC of the SE outflow in the mass-velocity and PV diagrams also support the processing jet scenario. For I23151, the blueshifted SiO emission traces a quasi-parabolic shaped outflow which can be identified in the 3.4 mm continuum and H$^{13}$CO$^+$ emission and coincide with the near-infrared nebulosity in the literature. This outflow can be interpreted as the molecular gas entrained by the underlying wide-angle wind.  The core masses estimated from the 3.4 mm continuum are 570 $M_{\\odot}$ and 7300 $M_{\\odot}$ respectively for the near and far distances for I18264 and 170 $M_{\\odot}$ for I23151. With the assumption of the momentum conservation between the outflow and the driving agent, the estimated outflow rates lead to accretion rates of $1.2\\times10^{-3}M_{\\odot}\\,yr^{-1}$ and $4.2\\times10^{-2}M_{\\odot}\\,yr^{-1}$ respectively for the near and far distances for I18264 and $1.9\\times10^{-5}M_{\\odot}\\,yr^{-1}$ for I23151. Taking into account the missing flux due to the the missing short spacings for I23151, the accretion rate in this region should be higher. Thus for both sources the accretion rates are high enough to overcome the radiation pressure from the central objects and form massive stars.  To summarize, our molecular outflow data derived from high resolution observations toward two luminous sources show morphologies and kinematics similar to those of the low-mass sources. The outflows can be interpreted by jet or wide-angle wind entrainment models. The presented data and analysis further support that massive stars up to $10^5 L_{\\odot}$ ($\\sim30 M_{\\odot}$) form via disk mediated accretion processes as low-mass stars."
    },
    "0609/astro-ph0609550_arXiv.txt": {
        "abstract": "We present $J$ and $K$-band luminosity functions (LF) for the Group Evolution Multiwavelength Study (GEMS) sample of 60 nearby groups of galaxies, with photometry from the 2MASS survey.  We find that, as seen in $B$ and $R$-band photometry of a subsample of these groups in our earlier work, the LFs of the X-ray dim groups ($L_X\\! <\\!10^{41.7}$ erg s$^{-1}$) show a depletion of galaxies of intermediate luminosity around $\\Mk\\!=\\!-23$, within a radius $0.3\\,R_{500}$ from the centres of these groups.  This feature is not seen in the X-ray brighter groups, nor in either kind of group when the LFs are determined all the way out to $R_{500}$.  We conclude that an enhanced level of star formation is not responsible for the this feature. From the faint end of the LFs, we find support for the under-abundance of low surface brightness dwarfs in the 2MASS survey. We find that for all kinds of groups, the modelling of the luminosity function, with universal forms for the LFs of galaxies of different morphological types, fails when simultaneously required to fit the $B$ and $K$-band LFs. This means that the dip-like features seen in LFs are not merely due to the varying proportions of galaxies of different morphological types among the X-ray dim and bright groups.  We argue that this support our hypothesis that this feature is due to the enhanced merging of intermediate-mass galaxies in the dynamically sluggish environment of X-ray dim groups. ",
        "introduction": "Luminosity functions of galaxies are an essential ingredient in the study of galaxy formation. Numerical models involving merging, cooling and feedback \\citep[\\eg][]{somerville99, kauff99, benson03} produce definite predictions for the luminosity function in various environments, which can now be compared with photometric observations of wide-field and deep samples of galaxies of the field and in highly clustered regions.  The galaxian luminosity function (LF) is usually modelled as the Schechter function, which drops sharply at bright magnitudes, but rises at the faint end as a power-law of slope $\\alpha$.  In cold dark matter (CDM) dominated models; dwarf galaxy haloes collapsing at $z\\!>\\!3$ have very efficient cooling, resulting in a steep slope $\\alpha \\!\\approx\\! -2$ at the faint end of the mass function \\citep[\\eg][]{evrard89,wf91,kwg93}. However, in practice, observed galaxy samples over large samples usually yield $\\alpha\\!\\approx\\! -1$ \\citep[\\eg][APM-Stromlo, $b_J$-band] {apmstromlo95}. Steeper faint-end slopes have been found \\cite[\\eg][]{tt02,2df-lf,trentham05}, and there is some indication that this is dependent on environment, and on the photometric filter used \\citep[\\eg][]{blanton05}. Ways of reconciling LFs of the luminous component with the predicted mass function of dark haloes have variously invoked effects of the physics of cooling and ionization \\citep[\\eg][]{chiu01,somerville02}.  Recently, with the help of deep photometry of the nearby Universe, the faint end of the LF has been explored in various surveys, revealing features that makes it difficult to fit a single Schechter function to the entire data \\citep[\\eg][]{ttm06,blanton05}. Indeed, many recent studies suggest that the LF of brighter galaxies ($M_B\\!<\\!-18.5$) should be modelled as a Gaussian, while that of the fainter ones be represented as a Schechter function \\citep[\\eg][]{jerjen97,mtt05}. This approach seeks to explain the peaks and dips in the LF as due to a varying mix of galaxies of different morphological types in different environments, and highlights the connection between the evolution of galaxies and their local environment.  A good example is the luminosity function of galaxies in Hickson compact groups, where a prominent deficit (``dip'') was found at intermediate luminosities \\citep{huns98}, showing that a single Schechter function would be inadequate to describe these LFs. Indeed, even in the Coma cluster, various studies \\citep[\\eg][]{tg93,lobo97,tully05} have found that optical LF of galaxies in the cluster have features not described by a Schechter function.  In a previous paper \\citep{miles04}, we presented optical ($B$ and $R$-band) LFs of a sample of 25 nearby groups of galaxies, as part of the Group Evolution Multiwavelength Study \\citep[][GEMS]{forbes06,osmond04,khos04,rm06}, for which we have X-ray luminosities (some have upper limits) from ROSAT PSPC observations.  To investigate the overall shape of the LF, we stacked together the LFs of several groups, with galaxies chosen from within $0.3\\,R_{500}$ from the centre of each group.  We found that the LF of the X-ray dim groups ($L_X \\!<\\! 10^{41.7}$ erg~s$^{-1}$) were significantly different from those of the X-ray brighter groups, showing a prominent dip, similar to Coma, at around $\\Mb \\!=\\!-18$. We interpreted this deficiency of intermediate-luminosity galaxies as evidence of rapid evolution through merging, since X-ray dim groups have lower velocity dispersion, which would encourage tidal interaction and merging by means of enhanced dynamical friction.  There have been alternative explanations suggested for the various features seen in LFs. Some have argued \\citep[\\eg][]{fergsan91,jerjen97} that the total luminosity function of a group or cluster has features due to a varying morphological mix across various systems.  Others \\citep[\\eg][]{biv95} have suggested that the apparent dip in the optical LF might actually be an enhancement of the bright end of the LF, due to excess star formation caused by tidal effects in the inner regions of the cluster or group.  We address these issues in this paper. To investigate whether the shape of the optical LF is affected by variation in the star formation properties of galaxies, we examine the near-infrared LFs of a large well-defined sample of groups of galaxies, namely the 60 groups of the GEMS study. The JHK photometry from the 2MASS survey is likely to trace evolved stellar populations and hence the total stellar content of galaxies, and thus can be better compared with galaxy formation models. It also allows us to avoid the effects of extinction that plague optical studies.  In the following section, we analyse 2MASS photometry of galaxies belonging to the 60 GEMS groups, and, by splitting them into two categories based on their X-ray luminosity, compare their $K$-band LFs with our earlier $B$-band LFs for a subset of 25 groups \\citep{miles04}.  In \\S3, we examine how the luminosity functions of X-ray bright and dim groups change with projected radial distance from group centre. We address the issue of the shape of the LF depending on the mix of various morphological types in \\S4, where we compute model luminosity functions resulting from the summation of the LFs of individual galaxy types. In the final section, we discuss the implications in the context of the evolution of galaxies in groups.  We have used $H_0=70$ \\kms Mpc$^{-1}$; $\\Omega_0\\equiv\\Omega_M=1$ throughout. ",
        "conclusions": "The main problems associated with Luminosity Functions at optical wavelengths are that they are liable to be affected by dust extinction as well as by young stellar populations, thus rendering the light less representative of the underlying matter. This is far less of a problem in the near-infrared, where one is more likely to trace the more evolved stellar populations and hence the actual stellar content of galaxies.  Near-infrared luminosities are more directly related to stellar mass, constraining both the history of the star formation and galaxy formation models. The fact that we find dip features in the luminosity functions of X-ray dim groups of galaxies in the near-infrared regime means that it is unlikely that star formation or dust extinction is responsible for the observed depletion of intermediate magnitude galaxies.  \\begin{figure*} \\centerline{ \\epsfig{file=TMfig9.ps,angle=-90,width=0.5\\hsize} \\epsfig{file=TMfig9b.ps,angle=-90,width=0.5\\hsize} } \\caption{ {\\bf [Left]} The best fit five-component LFs, one for each morphological type (Ellipticals, S0s, Spirals, Dwarf Irregulars and Dwarf Ellipticals), following the \\citet{jerjen-lf} functional forms (see Table~1), that add up to form the best model $B$-band LF. Also plotted are the observed data, averaged for 20 X-ray dim groups ($L_X<10^{41.7}$ erg~s$^{-1}$) out to 0.3$R_{500}$. Here the dashed line = Ellipticals; dash-dot = S0s; dot-dot-dot-dash = dEs; dotted line = dIrrs \\& continuous line = sum of all components). According to these fits, spirals make a negligible contribution to the morphological makeup, contrary to observations {\\bf [Right]} The predicted $K$-band LFs of the five components, using the same line styles, together with the observed LF from the 2MASS analysis in this paper. \\label{fig:jerlow}} \\end{figure*}    \\begin{figure*} \\centerline{ \\epsfig{file=TMfig10.ps,angle=-90,width=0.5\\hsize} \\epsfig{file=TMfig10b.ps,angle=-90,width=0.5\\hsize} } \\caption{{\\bf [Left]} $B$-band LFs, similar to Fig.~\\ref{fig:jerlow}, but for the X-ray bright groups ($L_X>10^{41.7}$ erg~s$^{-1}$), in our optical sample, out to 0.3$R_{500}$. {\\bf [Right]} Corresponding predicted $K$-band LFs of the individual components, together with the observed LF from the 2MASS analysis in this paper. The line style used for the various types is the same as before. In this case, Spirals (= dash-dot-dot-dot, centred on $M_B$ = -18.3) feature in the fitting process. \\label{fig:jerhigh}} \\end{figure*}   Such features have often been explained as the result of the representation in the total galaxy luminosity function of the underlying morphological makeup of groups.  Invariant type-specific LFs have been claimed \\citep[\\eg][]{jerjen-lf}. Implicit in this notion is the assumption of passive evolution at all structure scales, and thus the shape of the LF is merely due to the combination of the individual LFs of galaxies of different morphological types in varying proportions.  The deviation from these functional forms may therefore provide evidence of predominant major mergers affecting this model, and bringing into question the invariance of the LFs of individual morphological types.  Fig.~\\ref{fig:morph} shows the observed luminosity function of galaxies visually classified as early and late types.  The superposed curves are the results from the earlier fitting process, clearly showing that the comparison with real data does not work when one assumes universal luminosity functions for different morphological types. This view subscribes to the notion of passive evolution {\\ie} an initial distribution of galaxy types was imprinted in the early Universe and very little has occurred by way of galaxy transformational processes (such as major or minor galaxy mergers, ram pressure stripping, and harrassment) to affect this makeup.  This view falters when confronted with the high incidence of early-type galaxies in the relatively shallow potential wells of groups.  We argued in \\cite{miles04} that the persistent depletion of the intermediate-luminosity galaxies in the optical luminosity functions of X-ray dim groups shows that in the dynamically sluggish environment of such groups (which have low velocity disperison), dynamical friction would facilitate more rapid merging, thus depleting intermediate-luminosity galaxies to form a few giant central galaxies, resulting in the prominent dip seen in our LFs. In this work, we have shown that this effect is also seen in the near-infrared LFs, but only in the interior regions of the groups ($R<0.3\\,R_{500}$), as in the case of the optical LFs.  We also show that this effect vanishes when one goes out to $R_{500}$, an effect we had not tested in our optical CCD data due to lack of coverage. We thus show that the intermediate depletion in the LFs is a real feature, and we argue that it is due to the merging of galaxies, rather than a bright-end enhancement caused by excess star formation.    \\begin{figure*} \\centerline{ \\epsfig{file=TMfig11.ps,angle=-90,width=0.5\\hsize} \\epsfig{file=TMfig11b.ps,angle=-90,width=0.5\\hsize}} \\centerline{\\epsfig{file=TMfig11c.ps,angle=-90,width=0.5\\hsize} \\epsfig{file=TMfig11d.ps,angle=-90,width=0.5\\hsize} } \\caption{The four panels show observed $B$-band luminosity functions for the Early-type galaxies (circles) and late-type galaxies (spirals).  The top two panels are for X-ray dim groups ($L_X\\!<\\!10^{41.7}$ erg~s$^{-1}$), and the two bottom panels for X-ray bright groups ($L_X\\!>\\!10^{41.7}$ erg~s$^{-1}$), both out to 0.3$R_{500}$.   Shown for comparison are the best fit component LFs from Figs.~\\ref{fig:jerlow} and \\ref{fig:jerhigh}. {\\bf [Left]} Ellipticals = dashed line; S0s = dash-dot; dEs = dash-dot-dot-dot; continuous line = sum of all components. {\\bf [Right]} Spirals = dash-dot-dot-dot; dIrrs = dotted line. \\newline {\\bf [Top Left]} Data points: Early-type galaxy  $B$-band LF for X-ray dim groups. Curves: Early-type components fitted to total  X-ray dim group LF. \\newline {\\bf [Top Right]} Data points: Late-type galaxy  $B$-band LF for X-ray dim groups. Curves: Late-type components fitted to total X-ray dim group LF. \\newline {\\bf [Bottom Left]} Data points: Early-type galaxy  $B$-band LF for X-ray bright groups. Curves: Early-type components fitted to total  X-ray bright group LF. \\newline {\\bf [Bottom Right]} Data points: Late-type galaxy  $B$-band LF for X-ray bright groups. Curves: Late-type components fitted to total  X-ray bright group LF. This exercise shows that the morphological abundances from the best-fit model LFs do not match the observations. \\label{fig:morph}} \\end{figure*}"
    },
    "0609/astro-ph0609620_arXiv.txt": {
        "abstract": "We examine the origin and evolution of the mass-metallicity relationship (MZR, {\\it M$_*$-Z}) for galaxies using high resolution cosmological SPH + N-Body simulations that include a physically motivated description of supernovae feedback and subsequent metal enrichment.  We discriminate between two sources that may contribute to the origin of the MZR: 1) metal and baryon loss due to gas outflow, or 2) inefficient star formation at the lowest galaxy masses.  Our simulated galaxies reproduce the observed MZR in shape and normalization both at z=0 and z=2.  We find that baryon loss occurs due to UV heating before star formation turns on in galaxies with {\\it M$_{bar}$} $<$ 10$^8$ M$_{\\odot}$, but that some gas loss due to supernovae induced winds is required to subsequently reproduce the low effective chemical yield observed in low mass galaxies.  Despite this, we show that low star formation efficiencies, regulated by supernovae feedback, are primarily responsible for the lower metallicities of low mass galaxies and the overall {\\it M$_*$-Z} trend.  We find that the shape of the MZR is relatively constant with redshift, but that its normalization increases with time.  Simulations with no energy feedback from supernovae overproduce metals at low galaxy masses by rapidly transforming a large fraction of their gas into stars.  Despite the fact that our low mass galaxies have lost a majority of their baryons, they are still the most gas rich objects in our simulations due to their low star formation efficiencies. ",
        "introduction": "\\label{intro}  The observed trend that the metallicity of nearby galaxies decreases with galaxy mass has been well established in the literature \\citep[e.g.,][]{T04,Lee06}.  This trend has also recently been observed at high redshifts \\citep[][hereafter S05 and E06, respectively]{Savaglio05,Erb06}.  The gas metallicity of a galaxy depends on reprocessed material from its stars and subsequent inflow and outflow of gas.  The mass-metallicity relationship (MZR) thus provides insight into two physical mechanisms important to the evolution of galaxies: star formation efficiency and gas inflow/outflow.  In a ``closed box'' system (no inflow or outflow of material), metallicity {\\it Z} is given by \\begin{equation} Z = y\\, {\\rm ln} f_{gas}^{\\rm -1}, \\label{eq:1} \\end{equation} where {\\it y} is the nucleosynthetic metal yield produced by stars and {\\it f$_{gas}$} is the gas mass fraction (M$_{gas}$/[M$_{gas}$ + M$_{stars}$]) \\citep[e.g.,][]{Tinsley}. In general, low mass galaxies are observed to have higher gas fractions than high mass galaxies \\citep[e.g.,][]{Geha06, West05}. This observation can easily explain the origin of the {\\it M$_*$-Z} relation (MZR) in the context of the closed box model. In this scenario, low mass galaxies are gas rich because they are inefficient at turning gas into stars, likely due to their low surface densities \\citep{Kennicutt98, MK01, Verde02, JD04}.  Alternatively, SNe feedback might expel gas in galactic winds.  In this scenario the loss of metals becomes progressively more important at lower masses due to shallower halo potential wells \\citep[for a review see][]{vcb05} and leads to the observed {\\it M$_*$-Z} trend. However, the importance of supernova feedback as a function of galaxy mass has been investigated by various groups, with results that are still debated \\citep[][E06]{vcb05, Geha06, JD06}.  Recently, studies have investigated how the MZR evolves with cosmic time \\citep[S05, E06,][]{Dave06, deRossi06, kobayashi06, tkg06}.  Fully cosmological numerical experiments are one of the best tools to infer the relative importance of star formation (SF) and feedback processes in a realistic setting \\citep[e.g.,][hereafter G06]{Okamoto05, Robertson06, Governato06}.  In this Letter, we investigate the origin and evolution of the MZR in a concordance $\\Lambda$CDM cosmological setting with N-body + SPH simulations which include a physically motivated treatment of SF and SN feedback.  These simulations allow us to study the formation history of galaxies over range of masses, and thus quantify the importance of feedback and blowout versus low SF efficiencies over their lifetimes. ",
        "conclusions": "We have presented results from high resolution simulations of field galaxies in a $\\Lambda$CDM context. The simulations are carried to z=0 and include SF and a physically motivated description of the effects of SNe feedback and metal enrichment. We have shown in previous work (G06) that this approach leads to disk galaxies that fall on the TF and baryonic TF relations and reproduce the age-mass relation of disk galaxies and abundance of galaxy satellites.  Here we present a natural extension of this analysis, showing that the simulated galaxies are also in excellent agreement with the observed MZR both at z=0 and z=2.  Simulated galaxies match both the slope and normalization of the MZR over four orders of magnitude in stellar mass. We find that the slope of the MZR is roughly constant with redshift, but that its normalization increases with time.  Such evolution agrees with the high-z observations of E06, and with early theoretical predictions \\citep{Dave06, deRossi06}.  Furthermore, the mass scales of the turnovers in our simulated {\\it Z}, {\\it y$_{eff}$}, B/DM, and {\\it f$_{gas}$} versus mass relations are in good agreement with a range of previous observational and theoretical studies that find critical transitions at $\\sim$120 km/s \\citep[e.g.,][]{JD04, Keres, DB06}.  Galaxies with M$_{bar}$ $<$ 10$^{8}$ M$_{\\odot}$ (V$_{rot}$$\\sim$50 km/s), lose a significant amount of their baryons due to combined heating from the cosmic UV field and SN feedback.  Baryon loss due to SN feedback results in a {\\it y$_{eff}$}-M$_{bar}$ trend in agreement with the results from the SDSS (T04).  However, including gas that has been lost from galaxies does not alter the O abundances significantly, and does not change the overall shape of the MZR, confirming that the lower O abundances in our low mass galaxies are primarily due to their low SF efficiencies rather than directly to blowouts.  Despite strong baryon loss from their halos, we find that our cold gas fractions at low galaxy masses are in qualitative agreement with recent observational samples \\citep{Geha06}.  SN feedback plays a crucial role in lowering the SF efficiency in low mass galaxies and originating the turnover in the MZR at the observed scales.  Without energy injection from SNe to regulate SF, gas that remains in galaxies rapidly cools, forms stars, and increases the O abundance, even by z=2 (see G06 for additional shortcomings of the no feedback case). With feedback from SNe turned off, small galaxies produce too many metals too early at the expense of their cold gas content, producing {\\it M$_*$-Z} and {\\it y$_{eff}$}-M$_{bar}$ relations too flat compared to observations, in agreement with some early theoretical predictions \\citep{deRossi06, JD06}.  SN feedback prevents low mass galaxies from enriching their gas to a high metallicity at all redshifts.  In our runs, star formation in small field galaxies occurs over extended periods of time until the present, in agreement with observations \\citep{Dolphin05}."
    },
    "0609/astro-ph0609416_arXiv.txt": {
        "abstract": "Recent timing observations of the double pulsar J0737$-$3039A/B have shown that its transverse velocity is extremely low, only 10\\,km/s, and nearly in the Plane of the Galaxy.  With this new information, we rigorously re-examine the history and formation of this system, determining estimates of the pre-supernova companion mass, supernova kick and misalignment angle between the pre- and post-supernova orbital planes.  We find that the progenitor to the recently formed `B' pulsar was probably less than 2\\,$M_{\\odot}$, lending credence to suggestions that this object may not have formed in a normal supernova involving the collapse of an iron core.  At the same time, the supernova kick was likely non-zero.  A comparison to the history of the double-neutron-star binary B1534+12 suggests a range of possible parameters for the progenitors of these systems, which should be taken into account in future binary population syntheses and in predictions of the rate and spatial distribution of short gamma-ray burst events. ",
        "introduction": "The double pulsar PSR~J0737$-$3039A/B \\citep{bdp+03,lbk+04} is an outstanding laboratory for tests of general relativity and of binary evolution and supernova theories.  With its short orbital period, $P_b=2.4$~hours, and moderate eccentricity, $e=0.088$, the general-relativistic modifications to the Keplerian orbit are the largest known.  Pulsars `A' and `B' have spin periods of 22.7\\,msec and 2.7\\,s, respectively; since both are observed as radio pulsars, the mass ratio is obtained directly from pulsar timing, leading to new theory-independent constraints on strong-field gravity \\citep{lbk+04, ksm+06}.  Such double-neutron-star (DNS) systems are descended from high-mass X-ray binaries (HMXBs) in which both stars are massive enough to undergo supernova (SN) explosions \\citep[e.g.,][]{tv03,dp03b}.  In brief, the J0737$-$3039A/B binary is thought to have begun as two main-sequence stars with masses of at least 8\\,$M_{\\odot}$.  After a first mass transfer stage, the primary formed a neutron star in a core-collapse SN.  The secondary evolved, and matter was accreted by the neutron star in an HMXB phase. Eventually, the secondary's envelope enlarged to meet the neutron star, which spiraled in, ejecting the secondary's envelope.  Angular momentum transferred to the neutron star spun it up to a period of a few tens of milliseconds; it is now observed as the A pulsar.  Its envelope expelled, the helium core of the secondary remained in a circular orbit around the neutron star until a second supernova left the B pulsar. The radio lifetime of the recycled A pulsar is far longer than that of the high-magnetic-field B pulsar; this is the reason most double-neutron star systems are observed with only the recycled pulsar still active. We are fortunate to observe the double pulsar during its relative youth.  The observed binary orbital elements and the space motion of the binary preserve important details of the evolution, including the mass lost in the second supernova (and hence the mass of the exploding helium star), the size of the pre-supernova orbit, and any asymmetry in the explosion itself. The lower limit on the mass of an exploding helium core is of particular interest; evolutionary models suggest it should be around 2.1\\,$M_{\\odot}$ \\citep{nom84,hab86}.  Also of interest is the misalignment angle $\\delta$ between the first-born NS's spin axis and the post-SN orbital angular momentum, equivalent to the tilt between the pre- and post-SN orbital planes \\citep[e.g.,][]{wkk00}.  Constraints on all of the evolutionary parameters are important for estimating the detectability of DNS systems in pulsar surveys, and hence for estimates of the total DNS birthrate and population. These quantities are of wide interest for stellar evolution and nucleosynthesis, but particularly because of their importance to event rate estimates for gravitational-wave detectors \\citep[e.g.,][]{phi91,kkl+04} and as possible progenitors of short gamma-ray bursts \\citep[GRBs;][]{pac86}.  Early attempts to constrain the pre-SN parameters and the kick velocity of J0737$-$3039A/B have led to ambiguous results.  Shortly after the discovery, analyses based on the orbital elements alone \\citep{dv04,wk04} suggested that moderate kick velocities (of 60\\,km/s or more) were needed, and that the pre-SN helium star was likely low-mass and overflowing its Roche Lobe at the time of the explosion. Later analyses by Willems et al.\\ (2004, 2006) \\nocite{wkh04,wkf+06} used observational constraints on the transverse velocity of the system, first $\\sim 140\\,$km/s from scintillation measurements \\citep{rkr+04} and later $<30$\\,km/s from pulsar timing \\citep{kll+05}, to trace the motion of the binary back through the gravitational potential of the Galaxy \\citep{kg89} and derive probability density functions for the most likely kick velocity, pre-SN mass, and tilt angle.  In the most recent paper \\citep{wkf+06}, population synthesis models were also used to estimate the most likely radial velocity of the binary (which is not directly measurable); this helped to further restrict the allowed parameters. Overall, these authors favor the standard formation scenario, reasonably high kick velocities (70--180\\,km/s) and pre-SN mass $m_{2,i}$ of at least 2\\,$M_{\\odot}$, but acknowledge that for very small transverse velocities ($\\sim 10\\,$km/s) lower-mass progenitors are allowed or even favored.  Piran \\& Shaviv (2005a,b) \\nocite{ps05,ps05b} have argued, based on the assumption that the system will oscillate vertically about the Plane of the Galaxy, that its current location close to the Plane implies that it must have a low transverse velocity and have experienced a very small or no kick at birth, indicating a pre-formation B mass less than 2\\,$M_{\\odot}$ and possibly an entirely different formation mechanism, such as the electron-capture supernova proposed by \\citet{pdl+05}. ",
        "conclusions": "Because of the constraint equations used, each point in the (${\\bf V_k}$, $m_{2,i}$, $a_i$) parameter space corresponds to a point in the ($V_r$, $\\delta$, $\\Omega$) parameter space.  Thus we have effectively constructed the posterior pdfs for the (${\\bf V_k}$, $m_{2,i}$, $a_i$) parameters as well, and may marginalize to derive confidence ranges on all the parameters.  For both J0737$-$3039A/B and B1534+12, our pdfs for the physically interesting quantities are plotted in Figure~\\ref{fig:results} and confidence ranges are summarized in Table~\\ref{tab:conf}.  We recognize, as pointed out by \\citet{wkf+06}, that projection effects are neglected in these 1-d pdfs.  The most important such effect for J0737$-$3039 is the association of large-${\\bf V_k}$, large-$m_{2,i}$ and large-$\\delta$ solutions exclusively with the large $V_r$ ($\\stackrel{>}{\\sim} 100\\,$km/s) progenitors accessible only with the gaussian $V_r$ prior.  Because these imply that the current velocity makes a very small angle to our line of sight, we regard them as inherently less plausible.  \\begin{figure} \\epsfig{file=multipdf3_1534.eps,width=8cm} \\caption{Normalized posterior pdfs for DNS progenitors.  In each panel, the red and black lines represent PSR~J0737$-$3039A/B progenitors, with the red line corresponding to the gaussian $V_r$ prior with dispersion $200\\,$km/s, and black to the $V_r$ prior which is uniform in the cosine of the angle between the velocity vector and our line of sight.  The blue lines show pdfs for progenitors of PSR~B1534+12, for the gaussian $V_r$ prior with dispersion $200\\,$km/s.  Only parameter ranges with significantly non-zero pdf values are shown. Note the association of large $\\delta$ with large kicks, as expected.\\label{fig:results}} \\end{figure}  \\begin{table*} \\centering \\begin{minipage}{140mm} \\caption{Confidence ranges on progenitor parameters\\label{tab:conf}} \\begin{tabular}{@{}lccccccccc@{}} \\hline \\multicolumn{1}{c}{$V_r$ prior} & \\multicolumn{3}{c}{Median likelihood value} & \\multicolumn{3}{c}{68\\% confidence interval} & \\multicolumn{3}{c}{95\\% confidence interval} \\\\ {} & {$V_k$} & {$m_{2,i}$} & {$\\delta$} & {$V_k$} & {$m_{2,i}$} & {$\\delta$} & {$V_k$} & {$m_{2,i}$} & {$\\delta$} \\\\ {} & {(km/s)} & {($M_{\\odot}$)} & {($^{\\circ}$)} & {(km/s)} & {($M_{\\odot}$)} & {($^{\\circ}$)} & {(km/s)} & {($M_{\\odot}$)} & {($^{\\circ}$)} \\\\ \\hline \\multicolumn{10}{l}{J0737$-$3039A/B} \\\\ Gaussian & 165 & 1.80 & 12.0 & 80--305 & 1.50--2.40 & 3.0--24.5 & 45--1005 & 1.37--4.00 & 1.0--102.5 \\\\ Uniform-in-$\\cos$ & 60 & 1.45 & 3.5 & 40--80 & 1.37--1.55 & 1.5--5.5 & 20--140 & 1.37--1.80 & 0.5--11.0 \\\\ & & & & & & & & &\\\\ \\multicolumn{10}{l}{B1534+12} \\\\ Gaussian & 235 & 2.45 & 24.5 & 200--270 & 2.00--3.35 & 20.5--28.5 & 175--305 & 1.60--3.90 & 16.5--32.5 \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*}  We can draw several conclusions about the J0737$-$3039B progenitor and the supernova kick.  First, the explosion was asymmetric, but not extremely so.  The $V_k$ pdf extends down toward $0\\,$km/s for both $V_r$ priors, but a non-zero kick is strongly favoured.  That the supernova explosion was most likely asymmetric is bolstered by the B pulsar's spin-orbit misalignment observed through long-term changes in the B profile \\citep{bpm+05} and derived through modeling of the A eclipses \\citep{lt05,lyu05}.  For the gaussian $V_r$ prior, the kick tends to be directed out of the plane of the pre-SN orbit and away from the pre-SN progenitor orbital velocity.  For the uniform-direction $V_r$ prior, the kick is directed nearly randomly in the plane perpendicular to the pre-SN progenitor orbital velocity.  The only likely radial velocities are either slightly negative or else positive.  This reflects the fact that large negative velocities (relative to our current reference frame) imply that the system must have been born in the outer reaches of the Galaxy, but there are far fewer potential birth sites at large Galactic radii.  Low, even very low ($ < 10^{\\circ}$), misalignment angles are predicted between A's spin axis and the orbital angular momentum, in excellent agreement with the observed lack of profile variations \\citep{mkp+05,ksm+06}.  The immediate B progenitor was probably less than about 2\\,$M_{\\odot}$.  Since the pre-SN orbital period was comparable to the current 0.1-day value, the low mass is likely due to significant mass loss accompanying the orbital shrinkage. \\citet{dpsv02} and \\citet{dp03b} trace the histories of He-star/NS binaries, showing that suitable parameters can be obtained from a range of starting points.  A difficulty is that the lowest-mass systems are expected to undergo common-envelope evolution and spiral-in, resulting in extremely tightly-bound or even merged systems (\\citet{dp03b}, but see \\citet{ibk+03}).  Dewi \\& Pols note, however, that the relative time-scales of the spiral-in and supernova explosion are not well known, and an explosion might occur before the spiral-in is complete.  The low-mass B progenitor is consistent with either an electron-capture collapse of the ONeMg core of a low-mass He-star \\citep{pdl+05} or the collapse of a low-mass iron core.  Either might occur sufficiently fast to prevent the development of large asymmetries, resulting in the small kick \\citep{prps02,plp+04}. Estimates of the time-scales for the onset of the two types of explosions are similar \\citep{dp03b} and appear too long relative to the spiral-in time.  Better modeling of the two types of explosions may ultimately allow us to distinguish between the two on the basis of the required short time-scale.  PSR~B1534+12 presents an interesting contrast.  Our derived pdfs are fully compatible with the parameter space allowed in TDS05; in particular the kick velocity is likely in the range 200--270\\,km/s while the progenitor companion mass was in the range 2.0--3.35\\,$M_{\\odot}$.  We report the confidence intervals only for the gaussian (dispersion 200\\,km/s) $V_r$ prior in Table~\\ref{tab:conf}; the uniform-in-cos prior gives similar results. We note that the most important constraint in this case is the measured misalignment angle \\citep[$25\\pm 4^{\\circ}$;][]{sta04}, whose only model-dependence is on the reasonable assumption that the rotating vector model \\citep{rc69a} describes the linear polarization position-angle swing of this pulsar \\citep[c.f.][]{dpp05}.  Overall, the parameters of the B1534+12 progenitor fit quite well with a more `standard' evolutionary picture that involves He-star Roche-lobe overflow but not a late spiral-in phase.  It is possible that the longer orbital period and/or higher mass of the B1534+12 progenitor permitted a larger core to develop before the SN explosion, leading to a larger kick; we note that the second-formed NS in B1534+12 is 0.1\\,$M_{\\odot}$ more massive than J0737$-$3039B.  Also of interest is that \\citet{ikb06} find PSR~B1913+16, with the largest known second-formed NS mass, cannot be formed in population syntheses with small kick velocities.  The contrast between these systems urges caution in discussing expectations based on low kicks.  In the last several years, small ($\\sigma \\sim 50\\,$km/s) kicks have been proposed to be nearly universal for nascent neutron stars processed in relatively close binary systems, due to the early loss of the NS progenitor's envelope and subsequent formation of a low-mass core which in turn provides a small kick upon collapse \\citep[e.g.,][]{plp+04}.  This hypothesis has been used to explain the existence of long-period, low-eccentricity high-mass X-ray binaries \\citep{prps02}, the retention of large numbers of neutron stars in globular clusters \\citep{prps02,plp+04}, the apparent correlation between the orbital eccentricity and the recycled pulsar spin period in DNS systems \\citep{fkl+05,dpp05} and the apparent dearth of isolated mildly recycled pulsars (`failed DNSs') ejected from unbinding second SN explosions \\citep{dpp05}, as well as to predict larger numbers of DNS mergers than are inferred from the observed set of objects \\citep{prps02}.  Our derived progenitor parameters for J0737$-$3039A/B certainly imply that some fraction of DNS progenitors will be low-mass and experience small kicks, and hence that each of these explanations is plausible at some level.  However, the much larger kick needed for B1534+12, which is also a short-period (10h) binary, indicates that a range of parameters must be considered in all of these arguments.  Noting the additional rough correlation between the orbital eccentricity and the mass of the second-formed compact object in the six systems for which the masses are well-determined (PSRs J1756$-$2251, J0737$-$3039, J1141$-$6545, B2303+46, B1534+12 and B1913+16; see also \\citet{fau04}), we speculate that the magnitude of the kick may depend quite sensitively on the size of the collapsing core.  This speculation is not supported by the most recent 2-D numerical simulations of hydrodynamic instabilities during SN explosions \\citep{skjm06}, which suggest that core mass and NS velocity are not correlated, but more work will be required on 3-D simulations and on other kick models. At the same time, careful population synthesis, incorporating initial mass functions and preferably an improved understanding of the supernova vs.\\ orbital-evolution time-scales, will be required to determine whether this correlation or the (likely related) spin-period--eccentricity relation can truly be reproduced by a mass-dependent kick.  As pointed out by \\citet{pdl+05}, the low space velocity of J0737$-$3039 may provide a further probe of its long-term evolutionary history.  A competing model to the `standard' DNS evolutionary model outlined in Section~1 is the `double-core' scenario \\citep{bro95,bb98b} in which the He cores of two main-sequence stars of nearly equal mass undergo a spiral-in through the envelopes of both stars simultaneously.  Thus at the time of the first SN, the secondary star is a (low-mass) He star rather than a massive main-sequence star, and the systemic velocity after the first supernova explosion would be expected to be higher than in the standard model.  The observed low space velocity therefore makes the double-core scenario less likely. If SN kicks are in fact lower in He stars processed in binaries, this argument is slightly weakened; however we note that under the mass-dependent kick hypothesis, since A's mass is close to that of the companion in B1534+12, its kick velocity might well have been large.  Finally, our low-velocity progenitor for J0737$-$3039A/B may have implications for the nature of short GRBs.  In the popular DNS inspiral models \\citep{pac86}, a poor correlation between active star formation and GRB activity can be a natural consequence of the time delay between birth and inspiral and the space velocity of the post-supernova binary.  The few detected counterparts to short GRBs place these objects in the outskirts of a broad variety of galaxies without active current star formation \\citep[e.g.,][]{pbc+06}.  Allowing for large DNS kicks, this is consistent with the production of the progenitors during normal galactic star formation \\citep{bpb+06}.  However, if smaller DNS kicks are assumed, then it may also be necessary to produce a significant fraction of the progenitors in globular clusters in the galaxy haloes \\citep{gpm06}.  Our low velocity for J0737$-$3039A/B tends to support the latter scenario, but also suggests that similar systems could be a source of GRBs near the Earth, with possible implications for our terrestrial environment \\citep{tho95b,tmj+05}."
    },
    "0609/astro-ph0609370_arXiv.txt": {
        "abstract": "% Star clusters are observed to form in a highly compact state and with low star-formation efficiencies, and only 10~per cent of all clusters appear to survive to middle- and old-dynamical age. If the residual gas is expelled on a dynamical time the clusters disrupt. Massive clusters may then feed a hot kinematical stellar component into their host-galaxy's field population thereby thickening galactic disks, a process that theories of galaxy formation and evolution need to accommodate.  If the gas-evacuation time-scale depends on cluster mass, then a power-law embedded-cluster mass function may transform within a few dozen Myr to a mass function with a turnover near $10^5\\,M_\\odot$, thereby possibly explaining this universal empirical feature. Discordant empirical evidence on the mass function of star clusters leads to the insight that the physical processes shaping early cluster evolution remain an issue of cutting-edge research. ",
        "introduction": "\\label{pk_sec:earlyevol} The star-formation efficiency (sfe), $\\epsilon\\equiv M_{\\rm ecl}/(M_{\\rm ecl} + M_{\\rm gas})$, where $M_{\\rm ecl}, M_{\\rm gas}$ are the mass in freshly formed stars and residual gas, respectively, is $0.2 \\simless \\epsilon$ $\\simless 0.4$ \\citep{LadaLada2003} implying that the physics dominating the star-formation process on scales $<10$~pc is stellar feedback.  Within this volume, the pre-cluster cloud core contracts under self gravity thereby forming stars ever more vigorously, until feedback energy suffices to halt the process ({\\it feedback-termination}) \\citep{WeidnerKroupa2006}.  This occurs on one to a few crossing times ($\\approx 10^6$~yr), and since each proto-star needs about $10^5$~yr to accumulate about 95~per cent of its mass \\citep{WuchterlTscharnuter2003}, the assumption may be made that the embedded cluster is mostly virialised at feedback-termination. Its stellar velocity dispersion,  \\begin{equation} \\sigma \\approx \\sqrt{G\\,M_{\\rm ecl}/(\\epsilon \\, R)}, \\end{equation}  \\noindent may then reach $\\sigma=40\\,$pc/Myr if $M_{\\rm ecl} = 10^{5.5}\\,M_\\odot$ which is the case for $\\epsilon\\,R < 1$~pc. This is easily achieved since the radius of one-Myr old clusters is $R\\approx 1$~pc with a weak, if any dependence on mass \\citep{Bastian_etal2005}. Very young clusters (age~$\\simless 10$~Myr) would thus appear super-virial, i.e. with a velocity dispersion too large for the cluster mass.  The above exercise demonstrates that the possibility may be given that a {\\it hot kinematical component} could add to a galactic disk as a result of clustered star formation for reasonable physical parameters. Thickened galactic disks may result.  But this depends on  \\begin{description} \\item (i) $\\epsilon$,  \\item (ii) $R$ (cluster concentration) and  \\item (iii) the ratio of the gas-expulsion time-scale to the dynamical time of the embedded cluster, $\\tau_{\\rm gas}/t_{\\rm cross}$. \\end{description}  \\subsection{Empirical constraints} \\label{pk_ssec:emp} The first (i) of these is clearly fulfilled: $\\epsilon<40$~per cent \\citep{LadaLada2003}.  The second (ii) also appears to be fulfilled such that clusters with ages $\\simless 1$~Myr have $R\\simless1$~pc independently of their mass. Some well-studied cases are tabulated and discussed in \\cite{Kroupa2005}. Finally, the ratio $\\tau_{\\rm gas}/t_{\\rm cross}$ (iii) remains uncertain but critical.  The well-observed cases discussed in \\cite{Kroupa2005} do indicate that the removal of most of the residual gas does occur within a cluster-dynamical time, $\\tau_{\\rm gas}/t_{\\rm cross} \\simless 1$. Examples noted are the Orion Nebula Cluster (ONC) and R136 in the LMC both having significant super-virial velocity dispersions. Other examples are the Treasure-Chest cluster and the very young star-bursting clusters in the massively-interacting Antennae galaxy which appear to have HII regions expanding at velocities such that the cluster volume may be evacuated within a cluster dynamical time.  A simple calculation of the amount of energy deposited by an O~star within a cluster crossing time into its surrounding cluster-nebula also suggests it to be larger than the nebula binding energy \\citep{Kroupa2005}. Furthermore, \\cite{BastianGoodwin2006} note that many young clusters have a radial-density profile signature as expected if they are expanding rapidly.  Thus, the data suggest the ratio $\\tau_{\\rm gas}/t_{\\rm cross}$ to be near one, but much more observational work needs to be done to constrain this number. Measuring the kinematics in very young clusters would be an extremely important undertaking, because the implications of $\\tau_{\\rm gas}/t_{\\rm cross}\\simless 1$ are dramatic.  To demonstrate these implications it is now assumed that a cluster is born in a very compact state ($R\\approx 1$~pc), with a low sfe ($\\epsilon <0.4$) and $\\tau_{\\rm gas}/t_{\\rm cross}\\simless 1$. As noted in \\cite{Kroupa2005}, ``in the presence of O~stars, explosive gas expulsion may drive early cluster evolution independently of cluster mass''. ",
        "conclusions": "\\label{pk_sec:concs}  Observations show that star clusters of any mass are formed extremely compact ($\\simless 1$~pc), while older clusters appear well distended with radii of a few-to-many~pc. Very young clusters seem to be super-virial, and expanding HII shells also indicate the explosive removal of residual gas.  Star-cluster formation thus appears to be rather violent leading to a loss of a large fraction of stars and probably in many cases total cluster disruption.  As this discussion shows, the formation process of star clusters may have significant cosmological implications in that the morphology of galaxies may be shaped by the clusters born within them: the disks of galaxies are probably thickened during periods of high star-formation rates while low-mass dwarf galaxies may attain halos of stars as a result of the violent processes associated with cluster birth as a result of which stars may spread outwards with relatively high velocity dispersions.  Particular examples of such events have been noted above.  If this is true can be verified by measuring the velocity dispersions of stars in clusters younger than at most a few~Myr. This is very important to do because thickened galactic disks are often taken to be evidence for the merging of cold-dark-matter sub-structures. Naturally, in-falling dwarf galaxies will lead to thickened disks, and counter-rotating thick disks would be prime examples of such processes \\citep{YoachimDalcanton2005}, but this discourse shows that sub-pc-scale processes would probably need to be incorporated in any galaxy-evolution study.  The mass function of old globular clusters has a near-universal turnover near $10^5\\,M_\\odot$ which has so far defied an explanation, given that the mass function of very young clusters is observed to be a power-law, and that the modelling of classical disruption through two-body relaxation and tidal fields poses a challenge for evolving a power-law to the approximate log-normal form for old clusters.  Again, mass-dependent processes related to the physics of residual gas expulsion can explain such a turnover, but in this case the turnover would have to become evident in cluster populations younger than about 100~Myr.  The results of \\cite{deGrijs_etal2005} suggest this to be the case.  The results of \\cite{ZhangFall1999}, however, appear to show no such evidence for the cluster population in the Antennae galaxies: the mass function for two age groups of clusters (younger than about 10~Myr, and between 10 and 160~Myr) show the same power-law form, albeit with different normalisations as a result of the high infant mortality rate. If this is true, then orbital anisotropies in the young globular cluster populations as proposed by \\cite{FallZhang2001} would be needed to evolve a power-law mass function to the log-normal form.  It therefore appears that the broad principles of early cluster evolution remain quiet unclear.  Given that {\\it clusters are the fundamental building blocks of galaxies} \\citep{Kroupa2005} this is no satisfying state of affairs.  On the positive side, the importance of star clusters to stellar populations and other cosmological issues (merging histories of galaxies, galactic morphology) means that much more theoretical and observational effort is needed to clarify which processes act during the early life of a star cluster, and how these affect its appearance at different times."
    },
    "0609/astro-ph0609693_arXiv.txt": {
        "abstract": "We describe observations of the seventh accretion-powered millisecond pulsar, HETE~J1900.1$-$2455 made with the {\\it Rossi X-ray Timing Explorer}\\/ during the year of activity that followed its discovery in 2005 June. We detected intermittent pulsations at a peak fractional amplitude of 3\\%, but only in the first two months of the outburst. On three occasions during this time we observed an abrupt increase in the pulse amplitude, approximately coincident with the time of a thermonuclear burst, followed by a steady decrease on a timescale of $\\approx10$~d. HETE~J1900.1$-$2455 has shown the longest active period by far for any transient accretion-powered millisecond pulsar, comparable instead to the outburst cycles for other transient X-ray binaries. Since the last detection of pulsations, HETE~J1900.1$-$2455 has been indistinguishable from a low-accretion rate, non-pulsing LMXB; we hypothesize that other, presently active LMXBs may have also been detectable initially as millisecond X-ray pulsars. ",
        "introduction": "The accretion-powered millisecond pulsars (AMSPs) have emerged as a distinct class in recent years, beginning with the discovery of the pulsing nature of SAX~J1808.4$-$3658 in 1998 \\cite[]{wij98b,chak98d}. Since then, \\nmspsmo\\ more AMSPs were discovered during transient outbursts typically lasting a few weeks \\cite[see][for a review]{wij03a}. These transient sources verify the evolutionary link between low-mass X-ray binaries (LMXBs) and recycled millisecond radio pulsars \\cite[]{alpar82,rs82}. Extensive observations have revealed a number of properties which (until now) appeared characteristic of the class. The outbursts tend to be of short duration, typically a few weeks (but as long as 50~d in XTE~J1814$-$338). Pulsations are persistently detected at fractional amplitudes of typically $\\sim5$\\% rms.  Where thermonuclear bursts are present, oscillations at the pulsation frequency and roughly the same fractional amplitude are present throughout \\cite[e.g.][]{chak03a,stroh03a}.  \\src, the most recently-discovered AMSP, remarkably appears to share none of these properties. This source was discovered when a strong thermonuclear (type-I) burst was detected from a previously unknown source by {\\it HETE-II} \\cite[]{vand05a}. A subsequent {\\it Rossi X-ray Timing Explorer} (\\xte) observation of the field revealed pulsations at 377.3~Hz, confirming the bursting source as a new pulsar \\cite[]{morgan05}. The pulse amplitude early in the outburst was as low as 1.5\\% rms, fading to below the detection limit following a brief increase of the persistent flux level \\cite[]{kaaret05a}. Source activity continued throughout 2005 \\cite[]{gal05d} and 2006 (although the source was close to the sun between 2005 December and 2006 January and hence unobservable). This is the first AMSP that has shown ``quasi-persistent'' activity.  Here we present analysis of a series of observations over the 2005--06 outburst of \\src. In \\S\\ref{obs} we describe the observations and analysis procedures. Our analysis of the pulsation amplitudes and pulse shape is presented in \\S\\ref{pulsations}. We measure the variation of pulse amplitude and arrival time with energy in \\S\\ref{pulse_energy}. Finally, we discuss the implications of our measurements in \\S\\ref{sec2}. ",
        "conclusions": "\\label{sec2}  \\src\\ is the most remarkable accretion-powered millisecond pulsar discovered since the first example, SAX~J1808.4$-$3658. In several important respects, its behaviour is distinctly different from all other sources in the class. First, the source has been active for $>1$~yr, roughly an order of magnitude longer than any other AMSP. In this respect, the behaviour more closely resembles transients like MXB~1659$-$298 or KS~1731$-$260, with outbursts of 2.5 and 12.5~yr duration, respectively \\cite[e.g.][]{wij03c}. For a distance of 5~kpc \\cite[]{kawai05}, the mean X-ray flux of \\src\\ indicates a steady accretion rate of 2--3\\%~$\\dot{M}_{\\rm Edd}$. Although this is lower than the maximum reached in the outbursts of some other AMSPs, the ongoing activity means a higher time-averaged accretion rate since discovery \\cite[see also ][]{gal06b}.  The ongoing activity may be related to the stellar evolution of the mass donor. At 83.25~min, the orbital period of \\src\\ is essentially identical to the well-known lower limit of $\\approx80$~min for a main-sequence mass donor. The Roche lobe in a binary with a smaller orbital period cannot accommodate a H-rich donor, and such sources instead accrete from evolved stars or white dwarfs \\cite[]{ps81,nrj86}. An extended episode of enhanced mass accretion may occur as the binary evolves through the main-sequence orbital period limit, as the outer H-rich layers are gradually stripped away leaving only the He-rich core \\cite[see also Fig. 3 of][]{prp02}.  Second, the amplitude of the pulsations decreased systematically on a timescale of $~10$~d following several of the bursts observed early in the outburst. \\cite{kaaret05b} reported that the pulsations first became undetectable following a large flare early in the outburst, likely resulting from a temporary increase in the accretion rate. Our pulse amplitude measurements have revealed that the amplitude had been decreasing steadily since the burst on MJD~53548, suggesting that the disappearance of the pulsations after MJD~53559 may have instead been a consequence of the elapsed time since that burst, and the arrival of the flare was a coincidence. Indeed, pulsations only became detectable again following a subsequent thermonuclear burst which was detected by the PCA. It remains unclear why the burst just before the flare on MJD~53558, as well as the bursts observed later in the outburst, did not trigger similar episodes of high pulse amplitude.  Such pulse amplitude variations have not been reported in any other AMSP; we note that bursts have been detected in two others, SAX~J1808.4$-$3658 \\cite[]{chak03a} and XTE~J1814$-$314 \\cite[]{stroh03a}. It is particularly intriguing that the onset of the pulsations can apparently occur up to 2.4~hr {\\it before} or 0.5~hr {\\it after} the actual time of burst ignition (or not at all). That the pulsations may appear prior to the burst seems to rule out oscillations which are directly triggered by the bursts, although some dependence of the pulsation mechanism on the properties of the envelope is indicated.  The third aspect in which \\src\\ behaves distinctly differently from the other six known AMSPs is that the pulsations in \\src\\ were present only in the first few months of the outburst. While the amplitude of the pulsations in the other AMSPs may vary throughout the outburst, they are always detectable except at the end of the outburst when the source flux has dropped essentially to the background level. It is an open question as to what physics prevents persistent millisecond pulsations from being detected in most LMXBs containing neutron stars. That \\src\\ apparently makes a transition from persistently-pulsing (albeit at rather low amplitude) to a non-pulsing, quasi-persistent source within a timescale of a month indicates that the visibility of the pulsations is determined in the outer layers of the neutron star envelope, since this timescale is far too short to affect the % properties of the core. One possibility that has been discussed theoretically is that magnetic channeling of the flow (which is thought to give rise to the pulsations) will cease once the magnetic field is ``buried'' by the accreted material \\cite[e.g.][]{czb01,payne05}. The cessation of pulsations 70~d after the beginning of the outburst in \\src\\ suggests that only $10^{-10}\\ M_\\odot$, or equivalently a sustained rate of $\\approx2$\\%~$\\dot{M}_{\\rm Edd}$ is necessary to bury the field. We hypothesize that other, presently-active low-$\\dot{M}$ LMXBs may have also been detectable as millisecond X-ray pulsars when they first became active."
    },
    "0609/astro-ph0609146_arXiv.txt": {
        "abstract": "s{ A possible explanation of the present discrepancy between the abundance of galactic subhaloes predicted by N-Body simulations with those observed in the Local Group is presented. We study the impact of an early reionization on the baryonic component of the Universe using SPH simulations on group and galactic scales. We use a simplified model for reionization described as an instantaneous increment of the IGM temperature (with $10^5<T_{IGM}<10^6$ K) at a given redshift ($z_r \\geq 8$). We find that a key role is played by compton cooling (interaction between hot electrons and the CMB photons); at high redshift ($z_r>10$) this cooling is very efficient and it is able to counteract any heating of the gas within few Myrs ($\\approx 70$). This means that a late reionization is needed ($z_r<9$) to sufficiently reduce the number of luminous dwarf satellites around our Galaxy. For a reionization redshift $z_r=8$ and a reionization temperature of $T_{IGM} \\approx 10^5$ K we are able to reproduce the observed number of Local Group dwarf galaxies in our simulations.} ",
        "introduction": "Cold Dark Matter (CDM) models are the current paradigm for the formation of Large Scale Structure and they successfully explain most properties of the local and high-redshift Universe. One of the (few) problems of such models is their failure to reproduce the abundance of satellites with a circular velocity of $V_{c} \\approx 10-30$ km sec$^{-1}$ observed within the Local Group. High-resolution N-body simulations reveal a large number of embedded subhaloes that survive the collapse and virialization of the parent structure (Klypin et al 1999; Moore et al 1999). Such large numbers of subhaloes are in apparent discrepancy with the number of faint satellites observed in the Local Group by an order of magnitude (i.e. Macci\\`o et al 2006a). Many authors have pointed out that accretion of gas onto low-mass haloes and subsequent star formation are inefficient in the presence of a strong  photoionizing background (i.e. Quinn, Katz, \\&  Efstathiou 1996). The ultraviolet radiation and mechanical energy that preheated and reionized most of the hydrogen in the intergalactic medium (IGM), ending the ``dark ages'', was probably the result of an early generation ($z = 10-15$) of subgalactic stellar systems aided by a population of accreting black holes (Madau \\& Rees 2001). This picture, consistent with the hierarchical nature of structure formation in the Universe, is supported by the recent analysis of the third year's data from the WMAP satellite suggesting that the Universe was reionized at redshift $z_{ion} \\simeq 10$ (Spergel et al 2006). Reionization raises the entropy of the gas, preventing it from accreting onto small dark matter haloes and lengthening the cooling time of that gas which has accreted (Bullock et al (2000) see also Kravtsov, Gnedin \\& Klypin 2004). Moreover, Benson \\& Madau (2003; BM03 hereafter) suggested that winds from pregalactic starbursts and miniquasars may pollute the intergalactic medium (IGM) with metals and raise its temperature to a much higher level than expected from photoionization alone and so further inhibit the formation of early galaxies.  The detailed history of the Universe during and soon after these crucial formative stages depends on the power spectrum of the density fluctuations on small scales and on a complex network of poorly understood ``feedback'' mechanisms. For the first time, in order to elucidate the effect of all the cooling mechanisms that can act at high-redshift against the photo evaporation of gas in DM haloes we use several ``toy'' models for reionization within high-resolution hydrodynamic N-body simulations of structure formation. We particularly focus our attention on the Compton cooling (due to scattering between hot electrons and CMB photons, CC hereafter) which we argue is of critical importance for forming structures before and during reionization. ",
        "conclusions": "We have tested the impact of a reionization on the abundance of visible haloes around a Milky Way like simulated galaxy. We have used high resolution hydrodynamical simulation in which we implemented a simple treatment of reionization. In our model, cosmic reionization  is parametrized as an instantaneous increase of the temperature of the IGM. Two free parameters are in the model: $T_r$ the temperature of the IGM after reionization and $z_r$ the redshift at which it occurs. Using a first set of simulations of a group like object ($M \\approx 10^{13} \\Msun$) we inspected a wide range of parameters for our model. Our first finding is that a key role is playing by the Compton Cooling (between hot electrons of the IGM and CMB photons). If the reionization occurs at high redshift ($z>9$), the compton cooling time is very short with respect to the hubble time (cfr fig \\ref{fig:one} left panel). Thus CC is able to cool down the IGM in few Myrs ($\\approx 100$), counteracting any reionization effect. To avoid this high cooling we need a later reionization ($z_r \\approx 9-8$) as suggested by WMAP third year results in contrast with WMAP first year results ($z_r \\approx 17$) (Spergel et al 2006). Starting from these results we ran 6 high resolution ($N_{gas} \\approx 6 \\times 10^5$) hydro galaxy simulations spanning the parameter space ($z_r:T_r$). We showed that with later reionization ($z_r=8$) and a mildly high IGM temperature $T_r=5 \\times 10^5$ K (in agreement with observational data on quasar absorption lines around $z=3$, BM03) the total number of luminous satellites surviving at the present epoch is in agreement with the dozen of dwarf galaxies observed around the Milky Way. In this way it is possible to reconcile the  flat luminosity function observed within the Local Group with the steep mass function of haloes and subhaloes within the cold dark matter model.    \\begin{figure} \\begin{minipage}{0.45\\textwidth} \\begin{center} \\includegraphics[width=0.6\\textwidth]{galdm.bw.2.ps} \\end{center} \\end{minipage} \\hfill \\begin{minipage}{0.45\\textwidth} \\begin{center} \\includegraphics[width=0.6\\textwidth]{R5e5_z8.bw.2.ps} \\end{center} \\end{minipage} \\caption{Left: $G_0$ run dark matter satellites overdensity at z=0. Right: $G_3$ stellar satellites overdensity at z=0} \\label{fig:three} \\end{figure}"
    },
    "0609/astro-ph0609236_arXiv.txt": {
        "abstract": "%  We use new {\\it FUSE} data to determine the column densities of interstellar D\\,I, N\\,I, O\\,I, Fe\\,II, and \\hmol~along the HD\\,41161 and HD\\,53975 sightlines. Together with $N$(H\\,I) from the literature (derived from {\\it Copernicus} and {\\it IUE} data) we derive D/H, N/H, and O/H ratios. These high column density sightlines have both log $N$(H\\,I) $>$ 21.00 and allow us to probe gas up to $\\sim$1300 pc. In particular these sightlines allow us to determine the gas phase D/H ratio in a hydrogen column density range, log $N$(H) $>$ 20.70, where the only five measurements available in the literature yield a weighted average of D/H = (0.86 $\\pm$ 0.08)$\\times10^{-5}$. We find D/H = (2.14 $\\pm~^{0.51}_{0.43}$)$\\times10^{-5}$ along the HD\\,41161 sightline. This ratio is $\\sim$3$\\sigma$ higher than the weighted mean D/H ratio quoted above, for sightlines with log $N$(H) $>$ 20.70, while the D/H ratio for the HD\\,53975 line of sight, D/H = (1.02 $\\pm~^{0.23}_{0.20}$)$\\times10^{-5}$, agrees within the 1$\\sigma$ uncertainties. Our D/H measurement along the HD\\,41161 sightline presents the first evidence of variations of D/H at high $N$(H). Our result seems to indicate that either the long sightlines that according to the deuterium depletion model are dominated by cold undisturbed gas where deuterium would be depleted onto carbonaceous grains occur at higher $N$(H) than previously thought or that the clumping of low D/H values in the literature for the long sightlines has another explanation. Another possibility to explain the high D/H toward HD\\,41161 is local production of D\\,I which cannot be ruled out. Both of the O/H ratios derived here, (9.12 $\\pm~^{2.15}_{1.83}$)$\\times10^{-4}$ and (5.37 $\\pm~^{1.35}_{1.14}$)$\\times10^{-4}$ (for HD\\,41161 and HD\\,53975, respectively) are higher than what has been found by other authors in studies containing many sightlines, although not inconsistent, in the case of HD\\,53975, with individual measurements along some sightlines. However, O/H along HD\\,41161 is extremely high, and together with other high O/H ratios based on {\\it FUSE} data might indicate the presence of unknown systematic effects when deriving $N$(O\\,I) along high column density sightlines. Finally, we derive (N/H)$\\times10^{5}$ =  8.32 $\\pm~^{2.09}_{1.76}$ and 5.07 $\\pm~^{1.45}_{1.21}$ and (D/O)$\\times10^2$ = 2.29 $\\pm~^{0.40}_{0.35}$ and 1.91 $\\pm~^{0.51}_{0.43}$, for HD\\,41161 and HD\\,53975, respectively. In addition, the relatively high signal-to-noise ratio of the HD\\,41161 data allows us to place constraints on the $f$-values of some neutral chlorine transitions, present in the {\\it FUSE} bandpass, for which only theoretical values are available. ",
        "introduction": "\\label{intro}  Deuterium is thought to be produced in appreciable amounts only in the Big Bang \\citep{1976Natur.263..198E}. Of all the light elements produced during Big Bang Nucleosynthesis (e.g., $^3$He, $^4$He, $^7$Li), deuterium, easily destroyed in stellar interiors (astration), is the one whose relatively simple evolution makes it a particularly sensitive estimator of the baryonic density. Conversely, measurements of the cosmic microwave background, such as those performed with {\\it WMAP} and other missions, can be used to infer the primordial abundance of deuterium, (D/H)$_{\\rm P}$. Analysis of the three year {\\it WMAP} data \\citep{2006astro.ph..3449S} yields (D/H)$_{\\rm P}$ in agreement with that derived by \\citet{2003ApJS..149....1K}, (D/H)$_{\\rm P}$ = (2.78 $\\pm~^{0.44}_{0.38}$)$\\times10^{-5}$, by taking the weighted average of five D/H measurements toward QSOs in the range $z$ = 2.079--3.572. \\citet{2004ApJS..150..387S} found D/H = (2.2 $\\pm$ 0.7)$\\times10^{-5}$ for Complex C, a high-velocity low-metallicity cloud falling into our galaxy, which has presumably experienced more stellar processing than the gas seen toward QSOs.      Precise measurements of the gas-phase D/H ratio in the ISM were first performed with the {\\it Copernicus} satellite \\citep[e.g.,][]{1973ApJ...186L..95R} followed by the {\\it Hubble Space Telescope} twenty years later \\citep[{\\it HST}, e.g.,][]{1995ApJ...451..335L}, IMAPS \\citep[Interstellar Medium Absorption Profile Spectrograph,][]{1999ApJ...520..182J,2000ApJ...545..277S}, and more recently the {\\it Far Ultraviolet Spectroscopic Explorer} \\citep[{\\it FUSE}, e.g.,][and references therein]{2006ApJ...642..283O}. Variations in the D/H ratio were observed early on \\citep{1983ApJ...264..172Y} and by different satellites \\citep[e.g.,][]{1999ApJ...520..182J}. \\citet{1999NewA....4..231L} discuss some of the mechanisms that could be responsible for the observed variations, based on the D/H data available at the time. However only after the launch of {\\it FUSE} in 1999, has a large enough sample of sightlines with differing D/H ratios been accumulated, allowing the question of what causes the D/H variations to be addressed in a more systematic way. While the constancy of the D/H ratio in the Local Bubble \\citep[LB, corresponding to log $N$(H) $\\leq$ 19.2,][]{1999A&A...346..785S} seems to be a well accepted fact \\citep[(D/H)$_{\\rm LB}$ = (1.56 $\\pm$ 0.04)$\\times10^{-5}$ from the compilation by][]{2004ApJ...609..838W} two distinct explanations have emerged to explain the large scatter of gas-phase D/H measurements, in the Solar neighborhood outside of the LB. These measurements range from D/H = (0.50 $\\pm$ 0.16)$\\times10^{-5}$ toward $\\theta$~Car derived by \\citet{1992ApJS...83..261A} to D/H = (2.18 $\\pm~^{0.22}_{0.19}$)$\\times10^{-5}$ toward $\\gamma^2$~Vel derived by \\citet{2000ApJ...545..277S}.    \\citet{2003ApJ...599..297H} proposed a low value for the present-epoch deuterium abundance in the Solar neighborhood, based on indirect measurements of D/H from D/O and the constancy of the O/H ratio \\citep[O/H = (3.43 $\\pm$ 0.15)$\\times10^{-4}$ from][updated $f$-value]{1998ApJ...493..222M}. Using more measurements than \\citet{2003ApJ...599..297H}, \\citet{2006ASPC..348...47H} proposed then D/H = (0.7 $\\pm$ 0.2)$\\times10^{-5}$ for the present-epoch deuterium abundance. The high D/H ratios observed would be due to unknown systematic errors in the determination of $N$(H\\,I), unknown deuterium enrichment processes and/or local infall of deuterium-rich gas.  The possibility that deuterium could be depleted in dust grains was first proposed by \\citet{1982auva.nasa...54J}. \\citet{2004AAS...205.5704D} and \\citet{2006ASPC..348...58D} explored further the theoretical plausibility of this process leading \\citet{2004ApJ...609..838W} and \\citet{Linsky2006} to propose that the depletion of deuterium into dust grains leads to different levels of depletion along sightlines with different histories, producing the observed scatter in D/H. In this scenario, the highest measurements of the gas-phase abundance of deuterium in the Solar neighborhood place a lower limit on the total deuterium abundance, D/H $\\ge$ (2.31 $\\pm$ 0.24)$\\times10^{-5}$, which one can then compare to the primordial value through chemical evolution models. Supporting evidence for this model includes correlations of D/H with the depletion of Ti \\citep{2005ApJ...620L..39P}, Fe and Si \\citep{Linsky2006}, and with the average sightline density \\citep{2006ApJ...642..283O}.    Galactic evolution models have to reproduce a number of observables such as the metallicity distribution of late type stars, along with predictions of the evolution of deuterium from its primordial value to its present day distribution in the Galaxy, and in particular its value in the Solar neighborhood. \\citet{2006MNRAS.369..295R} have addressed the different views on the behavior of the D/H ratio by considering what value of D/H is predicted for the Solar neighborhood by chemical evolution models that satisfy the majority of the observational constraints (such as the G-dwarf metallicity distribution). They find that a modest star formation and a continuous infall of unprocessed gas is required to fit all the observational data, leading only to a modest decrease of the deuterium abundance from its primordial value. According to the models of \\citet{2006MNRAS.369..295R}, which adopt (D/H)$_{\\rm P}$ = 2.6$\\times10^{-5}$, D/H in the Solar neighborhood should be in the range of (1.4--2.0)$\\times10^{-5}$, implying astration factors $<$1.8. Low D/H ratios, as the one proposed by \\citet{2006ASPC..348...47H} are ruled out because they produced significant disagreements with other observational constraints, in particular with the G-dwarf metallicity distribution.    Testing the two explanations above for the scatter in D/H, requires more measurements of D/H and other species that might yield information about the depletion and physical conditions along the sightlines, particularly along high column density sightlines. In this work we present analyses of the D/H and other ratios along the extended sightlines to HD\\,41161 and HD\\,53975. Both of these sightlines have total hydrogen column densities greater than 10$^{21}$ cm$^{-2}$, and according to the variable deuterium depletion model such high column density sightlines should be biased by cold, not recently shocked gas, where deuterium would be depleted in grains. For the HD\\,53975 sightline we find D/H = (1.02 $\\pm~^{0.23}_{0.20}$)$\\times10^{-5}$, a value slightly higher but comparable to those found previously for similar sightlines. The high D/H ratio along the HD\\,41161 sightline, D/H = (2.14 $\\pm~^{0.51}_{0.43}$)$\\times10^{-5}$, presents the first evidence of D/H variations at high $N$(H), indicating that our understanding of the behavior of D/H is still incomplete.  This paper is organized as follows. The targets and their sightlines are presented in $\\S$\\ref{targets}. $\\S$\\ref{obs} describes the observations and data processing, while $\\S$\\ref{ana} presents the analyses. The abundance of neutral hydrogen along the two sightlines is discussed in $\\S$\\ref{hi}. The results are presented and discussed in $\\S$\\ref{results}. Constrains on the $f$-values of Cl\\,I transitions present in the spectra of HD\\,41161 are derived in $\\S$\\ref{cl_disc}. Our work is summarized in $\\S$\\ref{summary}. All uncertainties are quoted at the 1$\\sigma$ level unless noted otherwise. ",
        "conclusions": "\\label{results}  Table \\ref{ratios} summarizes several ratios along the two sightlines. Below we discuss each ratio in detail.  \\subsection{D/H ratios}  In the discussion below the D/H ratios are determined using D/H = [$N$(D) $+$ $N$(HD)]/[$N$(H\\,I) $+$ 2$\\times$$N$(\\hmol) $+$ $N$(HD)].  \\subsubsection{D/H along the HD\\,41161 sightline} \\label{highhihd41161}    Figure \\ref{dhratioplot} presents the D/H ratio as a function of log $N$(H). Open squares represent values from the literature and the three new values from \\citet{2006ApJ...642..283O} (references for all the ratios can be found in this paper), filled circles the ratios derived in this work. The Local Bubble D/H ratio and uncertainties are represented by horizontal solid and dashed lines. The vertical dashed line at log $N$(H) = 19.2 represents the edge of the Local Bubble, while the vertical dashed line at log $N$(H) = 20.7 represents the approximate position of where a new low and constant D/H regime was proposed to exist by \\citet{2004ApJ...609..838W} and \\citet{Linsky2006}.   There are five sight lines with log $N$(H) $>$ 20.7, with D/H ratios available in the literature \\citep[HD\\,198965, HD\\,191877, JL\\,9, LS\\,1274, and HD\\,90087;][]{2003ApJ...586.1094H,2004ApJ...609..838W,2005ApJ...635.1136H}. For these sightlines with 20.7 $<$ log $N$(H) $<$ 21.20, the weighted mean is D/H = (0.86 $\\pm$ 0.08)$\\times10^{-5}$. HD\\,41161 has $N$(H) similar to these sightlines, but a D/H ratio inconsistent by $\\sim$ 3$\\sigma$. To make D/H along the HD\\,41161 sightline consistent with the weighted ratio of the other 5 sightlines, one would have to increase $N$(H) by a factor of $\\sim$ 2.49, or log $N$(H\\,I) from 21.00 to $\\sim$ 21.44. To test whether D/H along this sightline could be high due to $N$(H\\,I) being underestimated we performed fits where log $N$(H\\,I) was forced to 21.44. Figure \\ref{hi21.43_hd41161} presents the fits to H\\,I Ly$\\beta$ and Ly$\\gamma$ with log $N$(H\\,I) forced to 21.44. The degree of the polynomial used to fit the continuum of these two transitions, is the same as the one used for the fit where we determined $N$(H\\,I) = 21.13 $\\pm~^{0.04}_{0.05}$. However, in this case, the continuum is forced to be artificially high to accommodate the higher $N$(H\\,I). In addition, the fit grossly overestimates the absorption on the blue side of the core of the Ly lines as shown in the two bottom panels of Figure \\ref{hi21.43_hd41161}. To further check that log $N$(H\\,I) = 21.44 overestimates the true H\\,I absorption we retrieved the {\\it IUE} Ly$\\alpha$ spectrum for the HD\\,41161 sightline from the MAST archive (no further processing was applied). Figure \\ref{hi21.44iue} presents the fit to H\\,I Ly$\\alpha$ forcing log $N$(H\\,I) = 21.44. Similarly to the fits to Ly$\\beta$ and Ly$\\gamma$ discussed above, forcing such a high value of $N$(H\\,I) on the data forces the continuum to be artificially high and overestimates the absorption near the core of the Ly$\\alpha$ line.    The $N$(H\\,I) adopted for this sightline is based on measurements performed by two different sets of authors \\citep{1982ApJS...48...95S,1994ApJ...427..274D} who obtain similar results, although both used the same {\\it IUE} dataset. In addition, fits of the Ly$\\beta$ and Ly$\\gamma$ H\\,I transitions produce $N$(H\\,I) in agreement with the value adopted in $\\S$\\ref{hd41161los}, while forcing $N$(H\\,I) to the value required to bring D/H to agree with the other low D/H measurements in the literature (discussed above) produces fits that clearly overestimate the H\\,I absorption. We conclude then that the high D/H along this sightline is unlikely due to $N$(H\\,I) being underestimated.  Another possibility to explain the high D/H along this sightline is that $N$(D\\,I) is overestimated. This is probably unlikely since saturation tends to move column densities in the opposite direction, i.e., underestimation rather than overestimation. However other factors, such as continuum placement, could produce the desired effect. A D\\,I column density of log $N$(D\\,I) = 16.00 is required to bring D/H along this sightline down to the weighted mean derived by the 5 sightlines discussed above. To test whether this $N$(D\\,I) value is consistent with the data we consider the $\\lambda$916.18 and $\\lambda$919.1 D\\,I transitions. With the AOD method we derive log $N$(D\\,I) = 16.37 $\\pm$ 0.04 from the $\\lambda$916.2 transition and our fits yield log $N$(D\\,I) = 16.43 $\\pm$ 0.03. A column density of log $N$(D\\,I) = 16.0 clearly underestimates the true D\\,I absorption, as shown in Figure \\ref{di16} where we overplot the above column density on the $\\lambda$916.2 transition.  To further check our $N$(D\\,I) measurement along this sightline we can consider the $\\lambda$919.1 D\\,I transition. This transition, not used in our analysis of this sightline, is blended with \\hmol~($J$ = 4) and according to our model slightly saturated. Using our \\hmol~fit model we carefully deblend the D\\,I and \\hmol~transitions. We use then the AOD technique to derive a lower limit to $N$(D\\,I) using the deblended D\\,I $\\lambda$919.1 transition. This exercise yields log $N$(D\\,I) $>$ 16.13, indicating that it is unlikely that $N$(D\\,I) is overestimated.  We conclude then that the high D/H ratio along this sightline is unlikely due to either underestimation or overestimation of the H\\,I or D\\,I column densities, respectively.  According to the variable deuterium depletion model proposed by \\citet{Linsky2006}, sightlines that probe gas with high column densities should be biased by cold, not recently shocked gas, where deuterium would be depleted in grains. The high D/H toward HD\\,41161 might indicate that this scenario occurs at $N$(H) higher than what was previously thought or that the clumping of low D/H values in the literature has another explanation. There is also the possibility that this sightline is dominated by gas which has been shocked and where the typical traces of a shock, such as highly ionized species (not detected along this sightline), have been erased. In this scenario the timescale to erase the traces of the shock would have to be smaller than that for D to deplete into the grains. The presence of cold \\hmol~($T_{\\rm 01}$ = 84 $\\pm$ 7 K) along this sightline does not necessarily imply that the gas is cold and undisturbed since \\hmol~is only a minor constituent of the sightline. Finally, we cannot rule out that local production of D\\,I of the type proposed by \\citet{1999ApJ...511..502M} or \\citet{2003ApJ...597...48P} is responsible for all or part of the deuterium enhancement along this sightline.    To determine the iron depletion along this sightline, D(Fe) = Log (Fe/H)$_{\\rm gas}$ - Log (Fe/H)$_{\\odot}$, we use the solar abundance ratio Log (Fe/H)$_{\\odot}$ = $-$4.55 $\\pm$ 0.05 from \\citet{2005ASPC..336...25A}. We find for this sightline, D(Fe) = $-$1.55 $\\pm$ 0.11. In a manner similar to some other sightlines with high D/H ratios, the Fe/H ratio along this sightline does not fit the trend of high Fe/H versus high D/H, displayed for most of the cases shown in Figure 3 of \\citet{Linsky2006}. There are however a few sightlines in Figure 3 of the paper mentioned above that seem to follow a trend different than most of the points, in that for the same Fe depletion a higher D/H ratio is found. The D/H and Fe/H ratios along the sightline to HD\\,41161 seem to follow this trend as well. \\citet{Linsky2006} discuss some of the reasons that could lead to this sort of behavior and argue that these sightlines could either contain a smaller percentage of carbon grains (thus fewer sites for deuterium to deplete) or could have had weak shocks pass by which would have evaporated deuterium from the grain mantles.   \\subsubsection{D/H along the HD\\,53975 sightline}  For the HD\\,53975 sightline we derive D/H = (1.02 $\\pm~^{0.23}_{0.20}$)$\\times^{-5}$ in agreement with the weighted mean D/H = (0.86 $\\pm$ 0.08)$\\times10^{-5}$, quoted above. Although the D/H ratio along the HD\\,41161 sightline indicates that D/H is variable for log $N$(H) $>$ 20.7, the D/H ratio along the HD\\,53975 sightline follows the trend displayed by the other other five sightlines in the literature with log $N$(H) $>$ 20.7. Taking into account this ratio, along with the ratios for the five sightlines mentioned above, yields a weighted mean D/H = (0.89 $\\pm$ 0.07)$\\times10^{-5}$ (with $\\chi^2_\\nu$ = 0.4 for 5 degrees of freedom; the average of the upper and lower 1$\\sigma$ uncertanties of individual ratios are used to determine weighted means in this work). The iron depletion along this sightline, D(Fe) = $-$1.75 $\\pm$ 0.09, follows the trend of low Fe/H versus low D/H displayed in Figure 3 of \\citet{Linsky2006}.  \\subsection{O/H Ratios}    The O/H ratio along the HD\\,41161 sightline, O/H = (9.12 $\\pm~^{2.15}_{1.83}$)$\\times10^{-4}$, is more than 3$\\sigma$ away from O/H = (3.43 $\\pm$ 0.15)$\\times10^{-4}$ derived for 13 lines of sight probing gas within 1500 pc (most within 500 pc), with 20.18 $\\leq$ log $N$(H\\,I) $\\leq$ 21.28 by \\citet{1998ApJ...493..222M} (updated $f$-value). For HD\\,53975, O/H = (5.37 $\\pm~^{1.35}_{1.14}$)$\\times10^{-4}$, is $\\sim$1.7$\\sigma$ away from the \\citet{1998ApJ...493..222M} O/H ratio. The {\\it FUSE}-based O/H measurements present a larger scatter than those based on {\\it HST} data \\citep[see for example][]{2003ApJ...591.1000A,2004ApJ...613.1037C}. Figure \\ref{oiscatterplot} presents a compilation of O/H measurements performed with {\\it FUSE} \\citep[from this work and the compilation by][asterisks and filled circles, respectively]{2006ApJ...642..283O} and {\\it HST} data \\citep[from][diamonds]{2004ApJ...613.1037C}. For the {\\it FUSE} measurements we consider only sightlines that probe gas beyond the Local Bubble, i. e., log $N$(H) $>$ 19.2, since previous measurements by \\citet{2005ApJ...625..232O} have shown that O/H = (3.45 $\\pm$ 0.19)$\\times10^{-4}$ in the LB. The assumption that all the {\\it HST} O/H measurements can be described by a single mean leads to the weighted mean O/H = (3.51 $\\pm$ 0.10)$\\times10^{-4}$ with $\\chi^2_{\\nu}$ = 1.35 for 51 degrees of freedom. Assuming that the {\\it FUSE} measurements can also be described by a single mean leads to the weighted mean O/H = (3.56 $\\pm$ 0.13)$\\times10^{-4}$ but with $\\chi^2_{\\nu}$ = 5.4 for 18 degrees of freedom. The scatter in the {\\it FUSE} O/H measurements is more pronounced in the range 20.0 $\\leq$ Log $N$(H) $\\leq$ 20.9, where there are only a few {\\it HST} measurements available. Hence, the scatter in the {\\it FUSE} measurements might be due to there really being more scatter for lower column density sightlines, even if the O/H values for the LB show little scatter \\citep{2005ApJ...625..232O}.  The larger \\citep[than][]{1998ApJ...493..222M} sample of sightlines probed by the work of \\citet{2004ApJ...613.1037C} contains some sightlines with O/H ratios $\\sim$6$\\times10^{-4}$. For example, HD\\,91824 ($l$ = 285.70$^\\circ$, $b$ = 0.07$^\\circ$) and HD\\,91983 ($l$ = 285.88$^\\circ$, $b$ = 0.05$^\\circ$) are both members of the Carina OB1 association (located at $d$ $\\sim$ 2.5 kpc) and have O/H = (6.91 $\\pm~^{1.21}_{0.92}$)$\\times10^{-4}$ and (5.25 $\\pm~^{1.21}_{0.98}$)$\\times10^{-4}$, respectively \\citep[][based on {\\it HST} data]{2004ApJ...613.1037C}. HD\\,90087 ($l$ = 285.16$^\\circ$, $b$ = $-$2.13$^\\circ$) is also part of the Carina OB1 association and is $\\sim$100 pc away from the two other stars. For this sightline the {\\it FUSE} derived O/H = (5.8 $\\pm$ 1.0)$\\times10^{-4}$ \\citep{2005ApJ...635.1136H} is in good agreement with the O/H ratios derived for the other members of this association.   In view of the \\citet{2004ApJ...613.1037C} measurements, the O/H ratio along the HD\\,53975 sightline does not seem to abnormally high, however the ratio along the line of sight to HD\\,41161 does. $N$(O\\,I) is determined using the only unsaturated O\\,I transition in the {\\it FUSE} bandpass, $\\lambda$974, which is heavily blended with \\hmol~$J$ = 2 and $J$ = 5. One might suspect that an erroneous $f$-value would lead to an inaccurate $N$(O\\,I). However, previous studies by \\citet{2005ApJ...635.1136H} indicate that $N$(O\\,I) derived from this transition is consistent with $N$(O\\,I) derived both from weaker and stronger O\\,I transitions. Unknown blends with stellar lines are also the unlikely cause of a seemingly high $N$(O\\,I) due to the high $v_{\\rm sini}$ of HD\\,41161. We suspect that unknown systematic effects when deriving $N$(O\\,I) from the heavily blended $\\lambda$974 O\\,I transition might be responsible for at least part of the scatter in the {\\it FUSE} based O/H measurements.   \\subsection{N/H, O/N, and D/O Ratios}  We derive N/H = (8.32 $\\pm~^{2.09}_{1.76}$)$\\times10^{-5}$ and N/H = (5.07 $\\pm~^{1.45}_{1.21}$)$\\times10^{-5}$ for the HD\\,41161 and HD\\,53975 sightlines, respectively. The first ratio is consistent, within the uncertainties, with N/H = (7.5 $\\pm$ 0.4)$\\times10^{-5}$ derived by \\citet{1997ApJ...490L.103M}. The N/H ratio for the HD\\,53975 sightline is $\\sim$1.7$\\sigma$ away from the \\citet{1997ApJ...490L.103M}.   We derive O/N = 10.96 $\\pm~^{2.10}_{1.85}$ and O/N = 11.74 $\\pm~^{3.36}_{2.80}$ along the HD\\,41161 and HD\\,53975 sightlines, respectively.  The O/N ratios for HD\\,41161 and HD\\,53975 are more than 3$\\sigma$ and 2.7$\\sigma$, respectively, away from O/N = 4.1 $\\pm$ 0.3 derived using the values of O/H and N/H determined by \\citet{1998ApJ...493..222M} and \\citet{1997ApJ...490L.103M}. This is likely a consequence of the high O/H for the two sightlines, in addition to the low N/H for the HD\\,53975 sightline.   We derive D/O = (2.29 $\\pm~^{0.40}_{0.35}$)$\\times10^{-2}$ and D/O = (1.91 $\\pm~^{0.51}_{0.43}$)$\\times10^{-2}$ along the HD\\,41161 and HD\\,53975 sightlines, respectively. Both of these ratios are consistent within the uncertainties with the weighted average D/O = (1.75 $\\pm$ 0.18)$\\times10^{-2}$ derived by \\citet{2005ApJ...635.1136H} for the distant sightlines. The low dispersion of D/O at high $N$(H) has been used by \\citet{2003ApJ...599..297H} and \\citet{2006ASPC..348...47H} as an argument against a high D/H ratio being representative of the interstellar medium in the Solar neighborhood. However, as pointed out above, the large scatter in the {\\it FUSE} derived O/H ratios seems to indicate that $N$(O\\,I) might suffer from unknown systematic errors, which would affect the derived D/O ratios. If the high value of $N$(O\\,I) measured along the HD\\,41161 sightline is indeed too high due to systematic effects related to using the $\\lambda$974 transition, then the D/O ratio along this sightline would be have to be revised upwards becoming inconsistent with the \\citet{2005ApJ...635.1136H} D/O ratio for distant sightlines. This would be consistent with our claim, from D/H, that the deuterium abundance towards HD\\,41161 is higher than that along the other log $N$(H) $>$ 20.7 sightlines.  Finally, we would like to point out that there are other sightlines in which high D/H, low D/O, and high O/H ratios are also observed, similarly to the HD\\,41161 sightline. These are Feige 110 \\citep{2002ApJS..140...37F,2005ApJ...635.1136H}, PG\\,0038$+$199 \\citep{2005ApJ...625..210W}, and LSE\\,44 \\citep{2006ApJ...638..847F}. Discussions about the possible causes of these ratios are given in the papers quoted for each sightline.     \\label{summary}  We have used {\\it FUSE} data to determine the abundance of several atomic and molecular species along the sightlines to HD\\,41161 and HD\\.53975. Together with $N$(H\\,I) from the literature, we derive D/H, N/H, O/H, and D/O ratios for these sightlines. We find that while the D/H ratio along the HD\\,53975 sightline is consistent with five previous measurements at similar $N$(H), the ratio toward HD\\,41161 is surprisingly high and presents the first evidence of D/H variations for sightlines with high $N$(H). The pattern of low D/H for sightlines with large $N$(H), pointed out by \\citet{2003ApJ...599..297H} and emphasized by \\citet{2004ApJ...609..838W} and \\citet{2005ApJ...635.1136H} is based on a small number of sightlines. Taking into account that at this high $N$(H) there is a selection bias against high D/H measurements \\citep{2006ApJ...642..283O} it is not surprising that only 1 out of 7 sightlines present high D/H ratios. As analyses of D/H along other high column density sightlines are being performed, one might expect that future measurements will also produce high D/H ratios, indicating that our understanding of the mechanisms responsible for the observed D/H pattern are still incomplete. Finally, the large number of Cl\\,I transitions detected in the spectra of HD\\,41161 allows to place significant constraints on the $f$-values of several transitions for which only theoretical values are available."
    },
    "0609/astro-ph0609831_arXiv.txt": {
        "abstract": "In the course of the formation of cosmological structures, large shock waves are generated in the intra-cluster medium. In analogy to processes in supernova remnants, these shock waves may generate a significant population of relativistic electrons which, in turn, produce observable synchrotron emission. The extended radio relics found at the periphery of several clusters and possibly also a fraction of radio halo emission may have this origin.  Here we derive an analytic expression for (i) the total radio power in the downstream region of a cosmological shock wave and (ii) the width of the radio-emitting region.  These expressions predict a spectral slope close to -1 for strong shocks. Moderate shocks, such as those produced in mergers between clusters of galaxies, lead to a somewhat steeper spectrum. Moreover, we predict an upper limit for the radio power of cosmological shocks.  Comparing our results to the radio relics in Abell\\,115, 2256, and 3667, we conclude that the magnetic field in these relics is typically at a level of $0.1\\,{\\rm\\mu G}$. Magnetic fields in the intra-cluster medium are presumably generated by the shocks themselves, this allows us to calculate the radio emission as a function of the cluster temperature. The resulting emissions agree very well with the radio power-temperature relation found for cluster halos. Finally, we show that cosmic accretion shocks generate less radio emission than merger shock waves. The latter may, however, be detected with upcoming radio telescopes. ",
        "introduction": "\\label{sec-intro}  Diffuse radio objects that extend over several Mpc but have no optical counterpart have been observed in several clusters of galaxies. These objects fall into two categories: radio halos are centrally located in the cluster, they are fairly regular in shape and show very little polarisation. In contrast, radio relics appear at the periphery of clusters, have a filamentary morphology and are often highly polarised. Abell\\,3667, for instance, shows a spectacular double relic \\citep{roettgering:97}, where the two radio objects, separated by $\\sim 4\\:{\\rm Mpc}$, straddle symmetrically the cluster X-ray emission. Similarly extended and remote relics have also been found in A115 \\citep{govoni:01}, A2256 \\citep{giovannini:99}, and A2345 \\citep{giovannini:99} to name just a few. All of these clusters also show signs of an ongoing or recent merger: A115 has a double peak in the centre in X-ray observations \\citep{shibata:99}, A2256 shows also a double peak \\citep{sun:02}, A2345 is ass umed to be a dynamically young system since it shows multiple X-ray substructures \\citep[see discussion in][]{dahle:02}, and in A3667, among other indications, a cold front has been found which is presumably related to a slightly supersonically moving substructure \\citep{vikhlinin:02}. In particular, the last example suggests that the relics are located at the shock wave that has been produced by the substructure. \\citet{roettiger:99} were able to reproduce the X-ray morphology and the location of the radio emission for this particular cluster in a merger simulation.   Even if the X-ray emission at the periphery of galaxy clusters is too faint to identify clearly shock fronts, radio relics seem to trace cosmological shock waves. Two models have been proposed for the origin of radio relics: They could either be caused by diffusive shock acceleration of electrons at the shock waves themselves \\citep{ensslin:98,miniati:01}, or they could be old radio bubbles that are compressed by the passing shock wave and thus induced to emit observable synchrotron emission again \\citep{ensslin:01,ensslin:02,hoeft:04}. The former relics are sometimes called {\\it radio gischt}, while the latter ones are known as {\\it radio phoenix} \\citep{kempner:04}. The radio relic in A85 \\citep{slee:01} may serve due to its size, morphology, and strong polarization as prototype for the radio phoenix class.   Radio halos, in contrast, do not directly trace the shock fronts. Instead, the radio morphology is similar to the overall X-ray morphology of the cluster \\citep{govoni:01b}. The radio spectrum of halos tends to steepen with radius \\citep{feretti:04,feretti:05b}. This may be regarded as evidence that the emitting electrons are accelerated by turbulence in the ICM \\citep{brunetti:04}, which is a result of a merger. However, the origin of the halos is still unclear. There is increasing evidence that halos are indeed related to cluster mergers, but it is a subject of a debate whether the relativistic electron population that is responsible for the radio emission is emitted by a central AGN, whether it is produced by strong shocks, whether it stems from the decay of relativistic protons or whether it is generated by turbulence in the downstream region of the shocks \\citep[for a review see e.\\,g.][]{feretti:05}. Interestingly, the radio power of halos scales with the cluster X-ray luminosity \\citep{feretti:05} or, equivalently, with its temperature \\citep{keshet:04}. Therefore, the non-thermal component in a galaxy cluster, if existent, seems to be related to the thermal plasma.   Numerical simulations show that the formation of the large-scale structures leads to a variety of shocks in the intra-cluster medium (ICM) and inter-galactic medium (IGM) \\citep{miniati:00, ryu:03, pfrommer:06, hoeft:06}. These shocks are crucial for the thermal state of the ICM. As matter falls into the deep potential wells of galaxy clusters, its shocks to the high temperature of the ICM of $\\sim 10^7-10^8\\:{\\rm K}$.  \\citet{ryu:03} distinguish between `internal' and `external' shocks. The former affect material that is already heated to roughly the cluster temperature, whereas the latter heat rather cold gas for the first time. Hence, internal shocks are typically weak while external shocks are very strong.   Cosmic rays with energies up to $\\sim10^{15}\\:{\\rm eV}$ are presumably accelerated at the shock fronts of supernova remnants \\citep[e.\\,g.][]{berezhko:03,vink:03}. The acceleration is believed to be caused by diffuse shock acceleration: Particles may be scattered repeatedly between the upstream and the downstream regions, separated by the shock. Each time the particle crosses the front, it gains kinetic energy \\citep[see for a review][]{axford:78,bell:78a,bell:78b,blandford:78}. Even though the physical conditions of shock waves in the ICM are somewhat different from those in supernova remnants, one would also expect diffuse shock acceleration to be efficient at cosmological shocks.   Various groups have tried to estimate the radio emission of the shock fronts by combining the probability distribution of cosmic shock fronts with models for diffuse shock acceleration and the subsequent synchrotron emission. \\citet{keshet:04} found that 10\\,\\% of the extragalactic radio background below 500\\,MHz should be caused by cosmic structure formation shocks for $\\Lambda$CDM cosmology. They assumed that the emission comes mainly from external accretion shocks, in particular those that surround massive clusters of galaxies. Essential for the radio emission is the strength of the magnetic field since the emission is synchrotron emission. Unfortunately, the magnetic field distribution in clusters of galaxies and their surroundings is difficult to determine and still poorly constrained \\citep[for a review see][]{govoni:04}.   In this paper, we compute the radio emission in the downstream region of a shock front. We assume that diffuse shock acceleration produces a relativistic electron population with a power-law distribution in the energy spectrum (\\Sec{sec-DSA}). As the plasma moves downstream, the high-energy electrons cool by inverse Compton and synchrotron losses (\\Sec{sec-cool}). The main objective of this paper is to derive an analytic expression for (i) the total radio power in the downstream region of a cosmological shock wave and (ii) the width of the radio-emitting region. In particular, we wish to explore the dependence of the power, spectral shape and spatial extent of the radio emission on the magnetic field and the shock parameters.  We apply our model to the relics observed in Abell\\,115, 2256, and 3667 and infer the magnetic field strength in the relic region (\\Sec{sec-examples}). Assuming that the magnetic field energy density is a fixed fraction of the thermal energy allows us to express the radio power as a function of the cluster temperature (\\Sec{sec-P-T}). Finally, we make predictions for the radio emission of cosmic accretion shocks (\\Sec{sec-cosmic}). ",
        "conclusions": "\\label{sec-summary}   Galaxy clusters grow either by mergers or by steady accretion, both of which produce large shock fronts. In particular, the extended radio relics ({\\it radio gischt}) in the periphery of clusters are believed to trace these fronts.  In analogy to supernova remnants, diffuse shock acceleration may produce a population of relativistic electrons that cool in the downstream region of the shock by inverse Compton and synchrotron losses. Here, we compute the total emission in the downstream region per unit area of the shock surface. The resulting analytic expression shows several interesting features: (i) the total spectrum is a power-law even if locally the spectrum may steepen towards higher frequencies. The overall slope is close to unity and depends solely on the slope of the initial electron spectrum which, in turn, is directly related to the strength of the shock. Thus, the overall slope of the radio emission may allow us to infer the Mach number of the shock. (ii) the radio emission depends almost like a step-function on the shock strength. For Mach numbers $\\lesssim 3$, only very little radiation is emitted while for ${\\cal M} \\gtrsim 10$ the emission saturates.  (iii) in the regime of weak magnetic fields, $B\\ll B_{\\rm CMB}$, the total emission is proportional to $B^{1+s/2}$ while for strong magnetic fields the total emission is virtually independent of $B$. (iv) The width of the emission region, i.\\,e. the extension perpendicular to the shock front, has a maximum extension for $B\\sim B_{\\rm CMB}$.   We analysed the relics in A\\,115, A\\,2256, and A\\,3667 and found that the shocks have Mach numbers in the range of 3 to 5. Weaker shocks would lead to practically no radio emission. Furthermore, we can infer the magnetic fields in the emission regions. We find that for three sample relics the magnetic field is in the range from 0.1 to $0.3\\:{\\rm\\mu G}$. The extension of the relics perpendicular to the shock front serves as an independent indicator for the field strength. The energy fraction stored in the magnetic field seems to be rather low, namely $\\xi_B \\sim 10^{-3}$.   Assuming that the magnetic field energy density is on average proportional to the thermal one, we derive a radio emission-temperature relation. It agrees strikingly well with that found for radio halos. Moreover, if in the centres of galaxy clusters, i.\\,e. where the radio halos reside, the magnetic fields is about a few ${\\rm\\mu G}$, the extension of the emission region perpendicular to the shock front would be several hundred kpc. Therefore, it is conceivable that at least part of the halo emission stems from shock waves which swept over the cluster centres within the last 1\\,Gyr. Using the value obtained for $\\xi_B$, we have finally estimated the emission from cosmic accretion shocks. We inferred that at 30\\,MHz a radio power of $10^{29}$ - $10^{30}\\:{\\rm erg \\, s^{-1} \\,Hz^{-1}}$ may be expected from regions where the accretion shock surface is parallel to the line-of-sight. A detailed computation of the cosmic radio background will be subject of a forthcoming study. \\\\   \\noindent {\\sc Acknowledgement } We gratefully acknowledge support by DFG grant BR 2026/2. We thank Torsten Ensslin for helpful discussions.   \\newcommand{\\aap  }{A\\&A} \\newcommand{\\araa }{ARA\\&A} \\newcommand{\\apj  }{ApJ} \\newcommand{\\apjs }{ApJS} \\newcommand{\\apjl }{ApJL} \\newcommand{\\apss }{ApSS} \\newcommand{\\aapr }{A\\&A~Rev.} \\newcommand{\\aj   }{AJ} \\newcommand{\\jgr  }{JGR} \\newcommand{\\mnras}{MNRAS} \\newcommand{\\physrep  }{PhysRep}"
    },
    "0609/astro-ph0609326_arXiv.txt": {
        "abstract": "This paper constructs an analytic framework for calculating the assembly of galactic disks from the collapse of gas within dark matter halos, with the goal of determining the resulting surface density profiles.  In this formulation of the problem, gas parcels (baryons) fall through the potentials of dark matter halos on nearly ballistic, zero energy orbits and collect in a rotating disk. The dark matter halos have a nearly universal form, as determined previously through numerical simulations. The simplest scenario starts with a gaseous sphere in slow uniform rotation, follows its subsequent collapse, and determines the surface density of the disk. This calculation is carried out for pre-collapse mass distributions of the form $M(r) \\sim r^p$ with $1 \\le p \\le 3$ and for polytropic spheres in hydrostatic equilibrium with the halo potential.  The resulting disk surface density profiles have nearly power-law forms $\\sigma \\sim \\varpi^{-q}$ (where $\\varpi$ is the radial coordinate in the disk), with well-defined edges determined by the centrifugal barrier $R_C$ -- the radius to which gas with the highest specific angular momentum falls during the collapse. This idealized scenario is generalized to include non-spherical starting states, alternate rotation profiles, and multiple accretion events (e.g., due to gas being added to the halo via merger events). This latter complication is explored in some detail and considers a log-normal distribution for the angular momenta of the pre-collapse states of the individual components. In the limit where this distribution is wide, the composite surface density approaches a universal form $\\sigma_T \\sim \\varpi^{-2}$, independent of the shape of the constituent profiles. When the angular momentum distribution has an intermediate width, however, the composite surface density attains a nearly exponential form, roughly consistent with profiles of observed galaxies. ",
        "introduction": "Although the formation of dark matter halos is rapidly becoming understood, the formation of gaseous galactic disks remains more elusive.  Disk formation is complicated and involves a number of physical processes, including the large scale collapse of the dark matter, merging of separate halos, gas cooling, and the infall of gas through the gravitational potential of the dark matter. Numerical simulations now include all of these processes and are beginning to produce realistic gaseous disks (e.g., Governato et al. 2004; Robertson et al. 2004; Keres et al. 2005).  This paper adopts a complementary approach and seeks to study the infall process --- separately and analytically --- where the starting conditions are informed by numerical simulations that have been previously carried out. This paper thus develops an analytic framework from which to study the infall process in galactic disk formation. This formalism applies to a wide range of infall scenarios and can be readily generalized to include additional physical processes. The main focus of this paper is to investigate the range of surface density profiles produced by gaseous collapse within the potential of a dark matter halo.  Although the aforementioned simulations demonstrate that the current paradigm of structure formation can produce galactic disks with reasonable properties, an analytic description of disk formation is useful in several respects. Analytic solutions provide us with a direct assessment of how the final disk surface density properties depend on starting conditions and other parameters of the problem. Along the way, we obtain an analytic description of several ancillary quantities, e.g., the velocity and density fields of the gaseous material as it falls toward the disk, as well as the hydrostatic equilibrium profiles of the pre-collapse states.  Once found, the resulting analytic solutions, both for the ancillary quantities and the resulting disk surface density profiles, can be used in a variety of other contexts.  In the standard picture of galactic disk formation (e.g., Rees \\& Ostriker 1977; Fall \\& Efstathiou 1980; White \\& Frenk 1991), the gas within a forming dark matter halo approaches a density profile that is roughly parallel to that of the dark matter.  This assumption is justified by complementary simulations of gas dynamics (from Evrard et al. 1994 to Kaufmann et al. 2006). In order for the gas to fall inward toward the nascent galactic disk, the gas must cool.  The gas cooling time depends primarily on density, which is a decreasing function of galactic radius, so the cooling time increases outward (note that metallicity, temperature, etc., also influence the cooling time). There are thus two regimes of interest. In one limit, the entire virialized region has a cooling time shorter than the age of the universe (at the formation time of the halo).  In this case, the gas cools rapidly and never approaches hydrostatic equilibrium. The gas supply is then limited by the rate at which the forming galactic structure accretes gas from larger radii. In the other limit, the cooling time is longer than the age of the universe for much of the virialized region, and infall is limited by the rate of gas cooling (White \\& Rees 1978). The flow of gas into the central regions is then analogous to cooling flows (e.g., Fabian et al. 1984). In both regimes, the gas falls inward through the gravitational potential provided (primarily) by the dark matter halo.  In this paper, we study the infall of baryonic gas moving in the potential of a dark matter halo.  Since dark matter halos approach a known nearly universal profile (e.g., Navarro et al. 1997; hereafter NFW), we consider the potential to have a given form and follow the orbits of gas parcels as they fall through the dark matter and build up a galactic disk structure.  Specifically, we consider the halo to have a fixed profile, independent of the gas, and to take the form of a Hernquist potential (Hernquist 1990; eq. [\\ref{eq:hq}]). Notice that this form is essentially the same as the more commonly used NFW profile, except that it has a steeper density profile at large radii and hence converges to a finite mass (see also Busha et al. 2005). As outlined above, the gas must cool before it falls through the halo potential toward the central regions. For the first part of this work, we describe this process as a single, overall collapse, even though the gas cools at different rates.  This complication will not affect the resulting surface density profiles of the disk as long as the central portions fall in first (as expected, since they tend to be denser and cool faster) so that infalling shells do not cross.  Notice also that we are implicitly assuming that the gas mass is small compared that of the dark matter.  After a sufficiently large quantity of baryonic mass has accumulated near the galactic center, however, the back reaction of the gas on the dark matter potential should be included.  The goal of this calculation is thus relatively modest: We want to develop an analytic transformation between starting gas distributions and the ``final'' surface density profiles of the disk. In this context, the starting gas distributions reflect the initial profiles of specific angular momentum, which depend on gas cooling and other effects that determine the starting state. The final surface density profiles (as considered here) are determined by where gas falls onto the galactic disk with no additional evolution (e.g., due to gravitational torques, which act on longer time scales). We discuss the effects of additional evolution, e.g., angular momentum redistribution, in the final section.  This work expands upon previous work concerning galactic disk formation (see, e.g., Gunn 1982, Kauffmann 1996, Dalcanton et al. 1997, Firmani \\& Avila-Reese 2000, Bullock et al. 2001, van den Bosch 2001, Thacker \\& Couchman 2001). These previous studies explore disk formation within dark matter halos with increasing levels of complexity, primarily through numerical simulations. The generic finding is that the resulting surface density profiles are more concentrated in the center than an exponential law, and typically display sharper outer boundaries.  This result depends on the starting (pre-collapse) distribution of angular momentum of the gas.  Dalcanton et al. (1997) start with a uniformly rotating sphere, Firmani \\& Avila-Reese (2000) assume that each infalling mass shell has a constant spin parameter $\\lambda$, and Bullock et al. (2001) determine the initial angular momentum distribution directly from N-body simulations. In these studies, the disk surface density profiles are the result of angular momentum conservation, although the initial distributions of angular momentum vary. One contribution of this present study is to provide a systematic exploration of possible starting states and their effect on the corresponding surface density profiles. In particular, we generalize the starting conditions to include a wide variety of initial density profiles, rotation profiles, and initial geometries. We also explicitly determine the velocity fields and density fields of the infalling material.  More significantly, this analytic framework is sufficiently robust to account for multiple accretion events (see also Firmani \\& Avila-Reese 2000). For comparison, previous numerical work (Robertson et al. 2004; Keres et al. 2005) includes more physical processes (e.g., explicit calculation of gas cooling and angular momentum transport during collapse), but does not isolate the dynamical mechanism(s) by which galactic disks are assembled.  This paper is organized as follows. We first extract the relevant orbit solutions that describe gas parcels falling through the gravitational potential of a dark matter halo (\\S 2); this calculation determines the velocity and density fields of the infalling material, as well as the locations where given parcels fall onto the disk plane. In \\S 3 we apply these results to the collapse of rotating gaseous spheres and find the resulting surface density profiles as a function of the starting density and angular momentum distributions. These results are generalized (in \\S 4) to include more complex starting geometries, including holes, removal of polar caps, and filamentary structures.  Motivated by the fact that cosmic structure is built up through a hierarchical process, \\S 5 considers the effects of multiple accretion events and finds the corresponding composite surface density profiles. The paper concludes in \\S 6 with a summary of our results and a discussion of how exponential disks can be produced. ",
        "conclusions": "\\subsection{Summary of results}  This paper has studied the assembly of galactic gaseous disks by following orbit solutions for baryonic gas parcels as they fall through the potential wells of dark matter halos. Our general analytic formulation of the problem (as presented in \\S 2) is robust and can be used in variety of contexts. In this paper, we calculate the velocity fields of the incoming gas, the density distribution of the gaseous material as it falls inward, and the surface density profiles of the resulting gaseous disks.  Our specific results are summarized below:  [1] The generic outcome for a gaseous disk produced by a single spherical collapse is a power-law surface density that is truncated at a centrifugal radius $R_C$ (see Figs. \\ref{fig:profiles}, \\ref{fig:profiles2}, and \\ref{fig:isodisk}). Here, the centrifugal radius $R_C$ is the radius to where the material with the highest (initial) specific angular momentum falls during the collapse. This result suggests that the ``fundamental units'' of galactic disk assembly are not exponential disks, but instead are truncated power-laws.  For a constant rotation rate of the initial gaseous sphere, the power-law index $q$ of the surface density distribution ($\\sigma \\propto \\varpi^{-q}$) depends on the initial mass distribution such that $q = 2 - p$, where $p$ is the power-law index of the starting mass profile ($M \\propto r^p$). For a rotation rate that varies with radius according to $\\Omega \\sim r^{-s}$, the power-law index of the resulting disk surface density is given by $q = 2 - p/(2-s)$. We also determined the density profiles appropriate for hydrostatic equilibrium of both isothermal and polytropic gas in the background potential of a dark matter halo (\\S 3.5) and followed the subsequent collapse (Figs. \\ref{fig:isodisk} and \\ref{fig:polytrope}). For polytropic equations of state with relatively large indices $n$, the resulting surface density profiles are close to exponential (except near the centrifugal barrier).  [2] In addition to determining the disk surface density profiles, this calculation provides an analytic description of the dynamics that builds up galactic disks. This description includes the velocity fields of the gas during the collapse (eq. [\\ref{eq:vfield}]), the density distribution of the infalling baryonic material (eqs. [\\ref{eq:density0}] and [\\ref{eq:density1}]), and the mapping between the starting locations in the pre-collapse cloud and the ``final'' radial positions within the disk (e.g., Figs. \\ref{fig:diagram} and \\ref{fig:conediagram}). We also obtain analytic forms for the hydrostatic equilibrium profiles of polytropic gaseous spheres in the gravitational potential of a dark matter halo (eqs. [\\ref{eq:isodense}] and [\\ref{eq:polydense}]).  [3] This collapse picture has been generalized to include less idealized starting configurations (\\S 4). The first generalization considers the starting gaseous spheres to have evacuated regions in their centers. The signature of these central holes in the starting states is to produce impoverished central regions in the corresponding gaseous disks (see Fig. \\ref{fig:radholes}). The starting states can also be considered to have limited angular extent. The removal of the polar caps from the initial states also leads to a deficit of material in the disk centers (relative to the spherical initial state) as shown in Figure \\ref{fig:polarholes}.  [4] This formulation of the collapse problem can be further generalized to include multiple accretion events, which are expected to occur due to merging of individual halos to form larger, composite structures (\\S 5). Each incoming halo provides a new supply of gas, which subsequently cools and falls inward to produce another component to the disk.  Here we consider the addition of a large number of such constituent disks, where the centrifugal radii are drawn from a log-normal distribution (consistent with the numerical finding that halo spin parameters display a log-normal distribution) and find an interesting and somewhat unexpected result: The addition of a large number of power-law disk profiles with index $q < 2$ results in a composite surface density profile of the form $\\sigma \\sim \\varpi^{-2}$ in the limit of a wide distribution for $R_j$, where this result is independent of the index $q$ (provided that $q < 2$). In the limit of a narrow distribution of $R_j$, the composite surface density is the same as the constituent surface density profiles. In the intermediate regime, however, the composite surface density can display a nearly exponential form over a range of radii (Figs. \\ref{fig:comp} and \\ref{fig:comptest}). The asymptotic forms are obtained in the limit of large numbers $N$ of constituent disks, and the approach to large $N$ is shown in Figure \\ref{fig:compvsn}. Although the large $N$ limit is approached rather slowly, this sequence shows that tens of subunits (constituent disks) are necessary to produce a smooth disk with a nearly exponential surface density profile.  \\subsection{The quest for exponential disks}  One of the defining features of observed disk galaxies is that their surface brightness distributions display a nearly exponential form. One might hope that numerical simulations of galaxy formation would naturally produce gaseous disks with exponential surface density profiles, so that star formation could take place {\\it in situ} and thereby account for the observations. Unfortunately, the production of exponential surface density profiles is not automatic. Building upon earlier work (Dalcanton et al. 1997, van den Bosch 2001, Kaufmann et al. 2006, and many others), this paper shows that the fundamental disks produced by a single collapse event have the form of a truncated power-law rather than an exponential profile (Figs. \\ref{fig:profiles} and \\ref{fig:profiles2}). The transformation between these truncated power-law forms and exponential profiles most likely involves several processes, as outlined below:  [A] The observed disk brightness profiles are close to exponential only over a limited range of radii (roughly 2 -- 20 kpc) and a correspondingly small range of scale lengths (roughly $r/\\rexp \\approx 0.5 - 5$). The inner regions depart from an exponential form due to contamination from bulge components and the outer regions are too faint to be observed. Further, as shown by Figure \\ref{fig:profiles}, the difference between a truncated power-law and an exponential is not too large over a limited range of radii. One of the largest differences between the power-law profiles and an exponential profile occurs at small radii, where the power-law profiles have cusps; this discrepancy is largely alleviated, or at least masked, by the presence of the bulge component (for $r < 2$ kpc). In addition, if the initial (pre-collapse) gas departs from a complete spherical form, either through inner evacuated regions or the removal of polar caps, then the resulting surface density profiles will not have inner cusps (Figs. \\ref{fig:radholes} and \\ref{fig:polarholes}), or they will be less pronounced.  [B] The calculation presented in this paper accounts for only the formation of the disk, not its subsequent evolution. If galactic disks are produced with sharp edges (as indicated here for the case of a single collapse event), then gravitational instabilities will grow robustly and cause the disk to spread.  Although previous authors (Dalcanton et al. 1997) have argued that the time scale for such instability is too long, the quoted (long) time scales are appropriate for disks that already have soft (e.g., exponential) edges. If disks have sufficiently hard edges, as produced here, then gravitational instabilities can grow on the dynamical time scale of the outer disk edge (Adams et al. 1989; Ostriker et al. 1992). This time scale is given by the inverse of the orbital frequency at the centrifugal radius, roughly $\\tau = R_C / v_{rot} \\approx$ 100 Myr ($R_C$/20 kpc). This time scale is still somewhat long to account for all of the necessary angular momentum redistribution, but gravitational instabilities will have time to grow and act to spread out the disk. Notice that a truncated power-law profile that spreads out will become more like an exponential profile.  [C] Exponential profiles are observed for the surface brightness distributions of galactic disks, not the surface density of the gas, so it is possible for the star formation process to help produce the observed exponential forms. A good demonstration of this effect is found in Lin \\& Pringle (1987), where initial surface density profiles roughly like those of Figure \\ref{fig:profiles} spread out under the action of disk viscosity and convert some fraction of the mass into stars.  The resulting surface density of the stars is much more like an exponential form than the starting gas profile.  [D] Finally, as shown in \\S 5, the addition of multiple disk surface density profiles can produce a composite profile that displays a nearly exponential form --- much more like an exponential than the initial (constituent) starting profiles. The ingredients necessary to produce a nearly exponential profile are a large number of constituent disks and a log-normal distribution for their centrifugal radii. This result is independent of the shape of the constituent disk profiles as long as they have effective power-law indices $q < 2$.  Since the current paradigm of structure formation indicates that galaxies are assembled through a process of hierarchical merging, multiple accretion events are highly likely. Further, numerical simulations suggest that the distribution of halo angular momentum is log-normal (Bullock et al. 2001) so that the distribution of centrifugal radii is also expected to be log-normal.  In practice, all four of these considerations (and perhaps others) will conspire to produce the observed exponential disk profiles. Individual infall events will each produce disk surface density profiles that are truncated power-laws, but multiple infall events will jointly build up a nearly exponential disk (as shown in Fig. \\ref{fig:comp}). The effects of star formation, viscous spreading, and disk evolution via gravitational torques will act to smooth out the disk and drive the surface density towards a more exponential form. Given the presence of the bulge contribution at small radii and the limited dynamic range of the observed profiles, the exponential disks in observed galaxies can be understood through the joint action of these processes.   \\newpage  \\centerline{\\bf Acknowledgment}  We would like to thank Michael Busha, Gus Evrard, and Lars Hernquist for many beneficial discussions.  This work was supported at the University of Michigan by the Michigan Center for Theoretical Physics; by NASA through the Terrestrial Planet Finder Mission (NNG04G190G) and the Astrophysics Theory Program (NNG04GK56G0); and by NSF through grants DMS01038545 and DMS0305837. Finally, we thank the referee for a comprehensive and detailed set of comments that improved the manuscript.  \\vskip1.0truein"
    },
    "0609/astro-ph0609110_arXiv.txt": {
        "abstract": "{Echelle spectra of 10 bright asteroids are presented and compared against an observed twilight spectrum and a computed Solar spectrum. Spectra covering a 2130~\\AA\\ spectral range centered on $\\lambda = 5785$~\\AA\\ are of high resolving power and high signal to noise ratio. So we focus on a comparison of detailed properties of spectral lines and not on albedo variations. It is shown that the normalized Solar and asteroid spectra are identical  except for radial velocity (RV) shifts which  can be predicted at the accuracy level of 1~m~s$^{-1}$. So asteroids are proposed as new and extremely accurate radial velocity standards.  Predicted and measured RVs of observed asteroids indeed match within limits of accuracy of the instrument. There are numerous absorption lines in the reflected Solar spectrum. This allows a direct mapping of the resolving power of a spectrograph between and along echelle spectral orders.  So asteroid spectra can be used to test wavelength calibration and resolving power  of spectrographs on the ground as well as in space, including the Gaia mission of ESA. All spectra are also given in an electronic form. ",
        "introduction": "Spectroscopic observations are usually wavelength calibrated using spectra of standard calibration lamps. Still, it is desirable to check this calibration with real objects on the sky in order to discover possible instrumental effects. These effects include: the calibration and the stellar beams not being exactly parallel when they enter the spectrograph slit; uneven illumination of the slit, which causes spectral shifts and therefore errors in the measurement of radial velocity; flexures and/or temperature instabilities (present in most of the spectrographs which are mounted directly on the telescope), which could cause significant variation of velocity and/or spectral shape if not monitored and properly accounted for.  IAU Commission 30 \"Radial velocities\" has defined a list of standard stars with accurate velocities  constant over many years. The list of Elodie-CORAVEL high-precision standard stars   (published on obswww.unige.ch/$\\sim$udry/std/std.html) contains a total of 140 stars, 38 of them with a  radial velocity error $\\lsim 50$~m~s$^{-1}$\\ (Udry, Mayor \\&\\ Queloz 1999). Most of these stars are bright ($V < 7$) so they may be unsuitable as radial velocity standards for spectrographs on  large telescopes due to saturation problems. The same is true for the now quite numerous stars which are being searched for extrasolar planets. Also, the RV residuals for some of these stars show a peculiar behaviour possibly due to yet undiscovered planetary companions, so their RV is not predictable accurately enough for calibration purposes. What is more, in some cases only variation but not the absolute value for RV has been derived and published. Nordstrom et al.\\ (2004) published radial velocities of the first 15\\%\\ of stars in the Hipparcos catalogue, part of which may be used as references if an error up to 300~m~s$^{-1}$\\ is acceptable. The sample which is being observed by the RAVE survey (Steinmetz et al.\\ 2006) is much larger and fainter, but of moderate accuracy. Most of the stars are observed only once, so some of the radial velocities may belong to (the yet undiscovered) binary stars.  Here we propose to use asteroids as objects for which radial velocity at the time of observation can be easily calculated to an extreme accuracy level. To the best of our knowledge no high resolution spectra of even the brightest asteroids exist in the literature. So we followed our initial suggestion (Zwitter \\&\\ Crifo 2003, Mignard 2003) and obtained high resolution and high signal to noise echelle spectra of 10 bright asteroids. Measurements of their radial velocities are compared to the accurate calculated values. Normalized asteroid spectra are also compared against both theoretical and observed (twilight) Solar spectra for each of the echelle orders. We show that asteroid spectra have identical spectral lines to twilight spectra. Excellent distribution of Solar spectral lines over the whole UV, optical and IR domain allows for a more accurate assessment of wavelength accuracy and resolving power than the conventional methods which use telluric lines.  The next section describes the observations and data reduction procedures. In section 3 we briefly describe the computations necessary to derive an accurate prediction of asteroid radial velocity at the time of observation. Observed asteroid and Solar spectra are compared in Sections 4 and 5, to be followed by Conclusions.  \\begin{table}[!Ht] \\tabcolsep 0.08truecm \\label{table1} \\caption{Wavelength ranges of individual spectral interference orders of the Asiago echelle spectrograph.} \\begin{tabular}{clclcclcl} \\hline \\hline &&&&&&&&\\\\ order&\\multicolumn{3}{c}{range [\\AA]}&\\ \\ \\ \\ &order&\\multicolumn{3}{c}{range [\\AA]}\\\\ &&&&&&&&\\\\ \\hline &&&&&&&&\\\\ 33&6720   &--& 6850&& 41&5410   &--& 5510\\\\ 34&6520   &--& 6650&& 42&5280   &--& 5380\\\\ 35&6330   &--& 6460&& 43&5150   &--& 5250\\\\ 36&6160   &--& 6270&& 44&5040   &--& 5130\\\\ 37&5990   &--& 6110&& 45&4930   &--& 5020\\\\ 38&5830   &--& 5940&& 46&4820   &--& 4910\\\\ 39&5680   &--& 5790&& 47&4720   &--& 4810\\\\ 40&5540   &--& 5650&& &       &  &     \\\\ &       &  &     && &       &  &     \\\\ \\hline \\end{tabular} \\end{table}    \\begin{table}[!Ht] \\tabcolsep 0.08truecm \\label{table2} \\caption{Log of asteroid observations. $N$ is the number of exposures, $t$ is the total observing time, {\\it Date} is the truncated date of observation and $V$ is the apparent V magnitude (derived from the Lowell Observatory  asteroid database). $S/N$ is the signal to noise ratio per 0.128~\\AA\\ wavelength bin at  $6070\\pm 4$~\\AA. } \\begin{tabular}{rllrccrr} \\hline \\hline &&&&&&&\\\\ \\multicolumn{2}{l}{Name}&$N$&$t$~[s]&\\ \\ &Date&V& $S/N$\\\\ &&&&&&&\\\\ \\hline &&&&&&&\\\\ 1 & Ceres   & 3 & 3000 && 2005 04 03.12 & 7.6& 188 \\\\ 2 & Pallas  & 4 & 3300 && 2005 04 02.95 & 7.3& 242 \\\\ 3 & Juno    & 3 & 3000 && 2004 09 25.81 &10.4&  81 \\\\ 4 & Vesta   & 4 & 3200 && 2004 09 26.00 & 6.3& 204 \\\\ 9 & Metis   & 3 & 3000 && 2004 09 25.45 & 9.2& 153 \\\\ 21& Lutetia & 3 & 3600 && 2004 09 26.13 &10.7&  98 \\\\ 27& Euterpe & 3 & 3600 && 2004 09 27.13 &10.2& 214 \\\\ 40& Harmonia& 3 & 3600 && 2004 09 25.09 & 9.8&  94 \\\\ 49& Pales   & 3 & 3600 && 2004 09 29.17 &11.5&  71 \\\\ 80& Sappho  & 3 & 2700 && 2004 09 28.64 &11.1&  84 \\\\ & dawn sky& 3 &  480 && 2004 09 26.20 &    & 340 \\\\ &         &   &      &&               &    &     \\\\ \\hline \\end{tabular} \\end{table}  \\begin{table*}[!Ht] \\tabcolsep 0.08truecm \\label{table3} \\caption{Radial velocities measured from individual spectral orders. Asteroid's radial velocity was determined as a mean of results for orders 33--43, skipping the 2 most deviant values (see text). } \\begin{tabular}{rlrrrrrrrrrrrrrrr} \\hline \\hline &&&&&&&&&&&&&&&&\\\\ \\multicolumn{2}{c}{Name}&\\multicolumn{15}{c}{radial velocity from echelle order}\\\\ &          &\\multicolumn{1}{c}{33}&\\multicolumn{1}{c}{34}&\\multicolumn{1}{c}{35}&\\multicolumn{1}{c}{36}& \\multicolumn{1}{c}{37}&\\multicolumn{1}{c}{38}&\\multicolumn{1}{c}{39}&\\multicolumn{1}{c}{40}& \\multicolumn{1}{c}{41}&\\multicolumn{1}{c}{42}&\\multicolumn{1}{c}{43}&\\multicolumn{1}{c}{44}& \\multicolumn{1}{c}{45}&\\multicolumn{1}{c}{46}&\\multicolumn{1}{c}{47}\\\\ \\hline &&&&&&&&&&&&&&&&\\\\ 1 &Ceres     &-12.03  &-12.31  &-12.29  &-12.36  &-12.46  &-12.42  &-12.62  &-12.49  &-13.18  &-13.45  &-13.72  &-13.32  &-13.47  &-13.73  &-14.00  \\\\%01270129 2 &Pallas    & 14.14  & 14.18  & 13.96  & 13.78  & 13.47  & 13.71  & 13.46  & 13.68  & 13.04  & 13.19  & 12.82  & 13.64  & 13.64  & 13.67  & 13.18  \\\\%01110114 3 &Juno      & 15.34  & 14.64  & 15.13  & 15.20  & 15.16  & 15.66  & 15.30  & 15.36  & 15.17  & 15.16  & 15.15  & 15.82  & 15.88  & 16.41  & 16.17  \\\\%00370039 4 &Vesta     &  9.73  &  9.69  &  9.95  &  9.89  &  9.82  & 10.14  & 10.28  & 10.56  & 10.23  & 10.24  & 10.36  & 10.84  & 10.79  & 11.12  & 11.19  \\\\%00570060 9 &Metis     &  0.73  &  0.44  &  0.80  &  0.57  &  0.54  &  0.71  &  0.76  &  0.97  &  0.64  &  0.76  &  0.87  &  1.44  &  1.56  &  1.95  &  1.77  \\\\%00520054 21&Lutetia   & -8.24  & -9.15  & -8.66  & -8.86  & -9.09  & -8.26  & -8.63  & -8.79  & -8.67  & -8.60  & -8.70  & -8.11  & -8.27  & -7.90  & -7.98  \\\\%00680070 27&Euterpe   &-19.54  &-19.80  &-19.69  &-19.78  &-19.93  &-19.38  &-19.55  &-19.47  &-19.77  &-19.59  &-19.68  &-18.95  &-18.78  &-18.51  &-18.65  \\\\%01170119 40&Harmonia  & -6.50  & -6.59  & -6.51  & -6.73  & -6.87  & -6.74  & -6.69  & -6.46  & -6.76  & -6.60  & -6.75  & -6.16  & -6.05  & -5.73  & -5.55  \\\\%00250027 49&Pales     &-13.87  &-13.34  &-13.43  &-13.76  &-13.81  &-12.92  &-13.67  &-13.52  &-13.86  &-13.53  &-13.57  &-13.13  &-12.66  &-12.52  &-12.78  \\\\%01880190 80&Sappho    &-11.01  &-10.54  &-10.46  &-10.76  &-10.76  &-10.18  &-10.48  &-10.33  &-10.73  &-10.59  &-10.57  &-10.04  & -9.78  & -9.48  & -9.61  \\\\%01500152 &twilight  & -0.12  & -0.43  & -0.35  & -0.52  & -0.66  & -0.33  & -0.51  & -0.34  & -0.83  & -0.76  & -0.86  & -0.34  & -0.28  & -0.09  & -0.21  \\\\%00820084 &&&&&&&&&&&&&&&&\\\\ \\hline \\end{tabular} \\end{table*}   \\begin{table*}[!Ht] \\tabcolsep 0.08truecm \\label{table4} \\caption{Calculated radial velocities compared to observations. Columns (3) and (4) give Sun--asteroid and asteroid to Asiago observatory distances in a.u. The next columns are velocities in km~s$^{-1}$: column (5) is radial velocity of asteroid vs.\\ Sun, column (6) is radial velocity of asteroid vs. Asiago observatory and column (7) is their sum. Column (8) gives radial velocity as derived from observations (see Table 3). } \\begin{tabular}{rllllrrrrr} \\hline \\hline &&&&&&&&&\\\\ \\multicolumn{2}{c}{Name}&\\multicolumn{1}{c}{JD(UTC)}&\\multicolumn{1}{c}{d$_\\odot$}& \\multicolumn{1}{c}{d$_{tel.}$}&\\multicolumn{1}{c}{RV$_\\odot$}&\\multicolumn{1}{c}{RV$_{tel}$}& \\multicolumn{1}{c}{RV(calc)}&\\multicolumn{1}{c}{RV(obs)}&\\multicolumn{1}{c}{O$-$C}\\\\ \\multicolumn{2}{c}{(1)}&\\multicolumn{1}{c}{(2)}&\\multicolumn{1}{c}{(3)}&\\multicolumn{1}{c}{(4)}& \\multicolumn{1}{c}{(5)}&\\multicolumn{1}{c}{(6)}&\\multicolumn{1}{c}{(7)}&\\multicolumn{1}{c}{(8)}& \\multicolumn{1}{c}{(9)}\\\\ \\hline &&&&&&&&&\\\\ 1&Ceres   & 2453463.619167&2.66095& 1.83302&  1.3092&--13.9723&--12.6631&     --12.43 & \t0.23\\\\ % 2&Pallas  & 2453463.448085&2.38562& 1.41984&  3.8052&   9.6790&  13.4842&\t 13.80 & \t0.32\\\\ % 3&Juno    & 2453274.309769&2.82011& 2.50488&--4.3485&  19.1563&  14.8078&\t 15.20 & \t0.39\\\\ % 4&Vesta   & 2453274.498461&2.38032& 1.41049&  1.7008&   8.0025&   9.7033&\t 10.02 & \t0.32\\\\ % 9&Metis   & 2453274.449896&2.30037& 1.32308&--2.3264&   3.2949&   0.9685&\t  0.70 &      --0.27\\\\ % 21&Lutetia & 2453274.628148&2.17825& 1.34490&  2.6739&--10.9357&  -8.2618&\t--8.68 &      --0.42\\\\ % 27&Euterpe & 2453275.632188&2.10125& 1.28461&--2.9481&--16.5471&--19.4952&     --19.65 &      --0.15\\\\ % 40&Harmonia& 2453273.592419&2.16246& 1.20745&  0.1823&  -6.9727&  -6.7904&\t--6.66 & \t0.13\\\\ % 49&Pales   & 2453277.665752&2.36502& 1.52997&--0.3002&--13.6541&--13.9543&     --13.60 & \t0.35\\\\ % 80&Sappho  & 2453276.636134&1.92316& 1.26508&  2.7895&--13.4288&--10.6393&     --10.56 & \t0.08\\\\ % &&&&&&&&&\\\\ \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "Observations of asteroids are convenient to measure radial velocity error as well as resolving power variation along and between the interference orders of an echelle spectrograph. They can be used to assess the accuracy of the absolute wavelength calibration on the level of m~s$^{-1}$. This is particularly valuable for wavelength calibration of slitless spectrographs, such as the radial velocity spectrograph on the forthcoming Gaia mission of ESA (Katz et al. 2004, Wilkinson et al. 2005). Rotation of asteroids with irregular shape and/or reflectance distribution over their surfaces can introduce radial velocity shifts on the m~s$^{-1}$ level. Most of the radial velocity standard stars are very bright and could be expected to have a certain degree of intrinsic spectrum variability. Asteroids, on the other hand, are observed at a wide range of magnitudes, so allowing a calibration suitable to the saturation limit of the telescope. For the calibration of the GAIA radial velocity spectrometer (RVS), only the brightest asteroids, i.e.\\ those with a not too fancy shape, will be used.  Asteroid light follows the same path through the atmosphere, telescope and spectrograph as the light from other objects on the sky. The numerous absorption lines in the reflected Solar spectrum therefore allow a detailed mapping of line broadening due to variations in the resolving power along the spectrum. Solar spectrum is not expected to change much over time, moreover its light output is being constantly monitored. So it is safe to use the observed width of most individual Solar lines reflected from an asteroid to measure the resolving power of the spectrograph.  Dedicated instruments are used to accurately map time variability in spectra of small areas on the Solar surface. But spectra reflected from asteroids offer a convenient way to measure the integrated high-resolution spectrum from the whole Solar disk using a normal telescope. So a variation of the high resolution integrated Solar spectrum with the Solar cycle could be measured."
    },
    "0609/astro-ph0609783_arXiv.txt": {
        "abstract": "%  The SPIRE Photometer Simulator reproduces the entire Herschel-SPIRE system in a modular IDL program.  Almost every aspect of the operation of SPIRE can be investigated in a systematic way to ensure that observations are performed in the most efficient way possible when Herschel flies.  This paper describes some of the work done with the Simulator to help prepare for large observing programs such as deep extra-galactic, high-redshift surveys. ",
        "introduction": "%  Expensive space missions with limited operating lifetimes, such as Herschel \\citep{pilbratt}, require efficient operation in order to obtain the best possible value for money in terms of the quantity and quality of astronomical data obtained.  Careful preparation of every aspect of the mission needs to be undertaken, including how the various instruments on board the spacecraft are operated to achieve optimum performance.  SPIRE \\citep{matt1} is one of three scientific instruments on board the Herschel spacecraft and covers the longer wavelength range from $\\sim~200 - 600~\\mu$m with both a photometer and an imaging Fourier Transform Spectrometer (FTS).  This paper describes how the operation of the photometer half of the SPIRE instrument has been optimised through the use of an instrument Simulator, and how the same Simulator can be used to prepare for large observing programmes.  \\subsection{The SPIRE Photometer Simulator}  The Simulator is an IDL coded virtual version of the photometer half of the Herschel-SPIRE system.  It incorporates as many of the physical instrumental and telescope characteristics as it is possible to include in a computationally practical and user-friendly program. Full details of the individual modules and their interaction with each other are given in \\citet{bruce1}.  Briefly, the user creates a suitable input sky -- as realistic or fantastic as desired -- for each of the three SPIRE bands.  These are fed into the Simulator where they are convolved with a representative beam profile and then 'observed' with the bolometer detector arrays.  Parameters for the observation are predefined by the user in the same way that a real observation would be planned.  The astronomical power from the sky and the background radiation from the telescope and internal instrument components are all passed into a module containing a model of the individual detectors, which calculates their response to the incident radiation.  This bolometer model also calculates and superimposes realistic noise on the output detector time-lines.  The detector time-lines are then low-pass filtered and sampled, in the same way as done by the on-board electronics, to produce output voltage time-lines.  Additionally, a pointing time-line is generated based on the observation parameters. ",
        "conclusions": "The SPIRE photometer Simulator is a very powerful tool for investigating the consequences of instrumental effects on the quality of astronomical data to be obtained with SPIRE.  Not only has it helped in the optimisation of the observing modes but it is also providing a head start in understanding the data and how it will be best processed when the real thing becomes available.  As well as simulating cosmological surveys the Simulator is also being used to simulate galactic (Sibthorpe, in preparation) and local Universe observations. Together, these simulations are central to the selection and development of the map making code to be delivered as part of the SPIRE pipeline processing package.  This process is currently underway and as launch approaches the Simulator will be used to make further simulations to help refine the performance of the map making algorithm as well as helping to develop optimised source detection routines capable of fully exploiting the SPIRE data.  Finally, being a modular program the Simulator is fully customisable and can be modified to represent any other facility of a similar design to SPIRE, or even to help design the next generation of far-IR observatories."
    },
    "0609/astro-ph0609260_arXiv.txt": {
        "abstract": "% We describe the scientific motivation for achieving photometric precision and accuracy below the 1\\% level, and we present a calibration philosophy based on using calibrated {\\it detectors} rather than celestial {\\it sources} as the fundamental metrology reference. A description of the apparatus and methodology is presented, as well as preliminary measurements of relative system throughput vs.~wavelength for the Mosaic imager at the CTIO Blanco 4m telescope. We measure the throughput of the optics, filter, and detector by comparing the flux seen by the instrument to that seen by a precisely calibrated monitor photodiode,  using a tunable laser as the illumination source. This allows us to measure the transmission properties of the system, passband by passband, with full pupil illumination of the entire optical train. These preliminary results are sufficiently promising that we intend to further pursue this technique, particularly for next-generation survey projects such as PanSTARRS and LSST. ",
        "introduction": "The challenge of photometry is to extract knowledge of the location and flux distribution of astronomical sources, based on measurements of the 2 dimensional distribution of detected photons in a focal plane. Each pixel $i$ in the detector array sees a signal $S_i$ given by   \\begin{equation} S_i = \\sum_{sources~j} \\int \\Phi_j(\\lambda) R_i(\\lambda) T(\\lambda)  A_i ~ d\\lambda, \\label{eq:psignal} \\end{equation}  \\noindent where the sum is taken over all sources (including the sky) that contribute to the flux in the pixel, $\\Phi_j(\\lambda)$ is the photon spectrum for source $j$, $R_i(\\lambda)$ is the throughput of the pixel, including the transmission of the optics and the pixel's quantum efficiency,  $T(\\lambda)$ is the optical transmission of the atmosphere, and $A_i$ is the effective aperture of the system for pixel $i$, essentially the wavelength-independent part of the instrumental response.  Traditional flat-fielding techniques attempt to extract knowledge of $\\Phi(\\lambda)$ from this array of sums of integrals by first dividing the flux in each pixel by a scalar number, $F_i^B$, the ``flat field'' for that pixel in a passband $B$. It is however clear from the above equation that this is not arithmetically correct if multiple sources with different photon spectra are contributing to the flux in the pixel. Furthermore, airmass corrections typically assume that the atmospheric transmission for a given band $B$ depends only on the secant of the zenith angle, ignoring any wavelength dependence across the passband. Astronomical instruments are currently calibrated, in practice, using celestial calibrators. At present the most popular photometric system is based on our knowledge \\citep{Hayes75} \\cite{Hayes75b} \\cite{Castelli94} of Vega. Our observational information about the spectrum of Vega is of course fundamentally based upon the terrestrial blackbody sources against which Vega was itself calibrated, by ground-based measurements. Present work, described elsewhere in this volume, uses a combination of photospheric modeling and observation to construct a synthetic spectra of spectrophotometric calibration objects. Since we don't know the distances or radii of the sources well enough to determine what their apparent fluxes should be, there is an overall multiplicative ambiguity that is (for Vega-based magnitude systems) tied to a monochromatic flux of Vega. This in turn links all celestial calibrators to the terrestrial blackbody sources that were used in the calibration of Vega.  Considerable careful effort has been expended on building a network of spectrophotometric stars across the sky \\cite{Oke83} \\cite{Colina94}  \\cite{Hamuy92}  \\cite{Hamuy94}  \\cite{Bohlin96} \\cite{Megessier95}. These stars, in conjunction with the network of secondary photometric standards \\cite{Landolt92}, form the basis for determining the instrumental throughput of different instruments, and for making the photometric transformations into a ``standard'' system  \\cite{Fukugita96}.  The current state of the art with this approach produces flux measurements with fractional uncertainties at the few percent level \\cite{SDSS04}  \\cite{Bessell99}  \\cite{Bessell05}. Considerable further progress along these lines is described in other contributions to this conference.   A variety of forefront scientific issues motivate breaking through the 1\\% barrier that has long been a limitation in the precision of ground-based photometry. One example, which has in large part motivated the work described here, is the Dark Energy problem. Type Ia supernovae are a powerful probe of the history of cosmic expansion \\cite{highz} \\cite{scp}, and they provide strong evidence for the existence of the Dark Energy. As we move from the detection to the characterization of Dark Energy, the challenge is detecting subtle signals in the Hubble diagram. In particular, measuring the equation of state parameter $w=P/\\rho$ of the Dark Energy requires confidence in photometry at or below the 1\\% level.  The fundamental measurement in supernova cosmology is a determination of apparent brightness vs.~redshift. In order to measure brightness in the same spectral region for each supernova, we must shift to redder passbands for supernovae at increasing redshifts.  Any systematic miscalibration of the photometric zeropoints in different passbands would produce a corresponding systematic distortion in the Hubble diagram. Furthermore, a detailed knowledge of the instrumental response function is required, in order to properly account for the effects of the redshifted SN spectrum as seen in the passbands used (i.e.~ to perform``K corrections'' \\cite{Peacock98}).  An alternative approach, which we intend to pursue, is to {\\it measure} the quantities $R_i(\\lambda)$, each pixel's response vs.~wavelength, and $T(\\lambda)$, the optical transmission of the atmosphere for each image.   We have described the formalism of this approach elsewhere (Stubbs \\& Tonry 2006) and we encourage the reader to consult that paper in conjunction with this one, which describes a preliminary implementation of this technique. In Stubbs \\& Tonry (2006) we point out  (for a major wide field survey which will measure many of currrent photometric standards as a matter of course) the merits of reporting all photometric survey results in the ``natural'' photometric system of the survey, rather than inflicting the systematic errors introduced by transferring measurements into a ``standard'' system.  In particular we intend to fully exploit the fact that we can obtain detectors whose response vs.~ wavelength can be far better characterized \\cite{Larason98} than the spectrum of a calibration source.  The measurement of atmospheric transmission is another essential ingredient in this approach, which is a topic of current research \\cite{Adelman96} \\cite{Granett05}, but we will not address that issue here. ",
        "conclusions": "We consider the preliminary results we obtained to be sufficiently promising to proceed with this calibration scheme. We note that the alternative approach we are pursuing does not in any way compromise or collide with reducing survey data with standard techniques. In fact comparing the results with the two methods will provide useful constraints on systematic effects.  Our near term goal is to obtain a full set of throughput calibration data for the various filters on the Mosaic instrument, by the end of 2006. We also will  \\begin{itemize}  \\item{} Develop and refine atmospheric transmission measurement techniques,  \\item{} Develop and refine self-luminous flat field screens, as described in the companion paper by Brown et al in these proceedings,  \\item{} Continue to work towards implementing these techniques for both PanSTARRS and LSST.  \\end{itemize}"
    },
    "0609/astro-ph0609056_arXiv.txt": {
        "abstract": "Globular clusters rotate significantly, and with the increasing amount of detailed morphological and kinematical data obtained in recent years on galactic globular clusters many interesting features show up. We show how our theoretical evolutionary models of rotating clusters can be used to obtain fits, which at least properly model the overall rotation and its implied kinematics in full 2D detail (dispersions, rotation velocities). Our simplified equal mass axisymmetric rotating model provides detailed two-dimensional kinematical and morphological data for star clusters. The degree of rotation is not dominant in energy, but also non-negligible for the phase space distribution function, shape and kinematics of clusters. Therefore the models are well applicable for galactic globular clusters. Since previously published papers on that matter by us made it difficult to do detailed comparisons with observations we provide a much more comprehensive and easy-to-use set of data here, which uses as entries dynamical age and flattening of observed cluster and then offers a limited range of applicable models in full detail. The method, data structure and some exemplary comparison with observations are presented. Future work will improve modelling and data base to take a central black hole, a mass spectrum and stellar evolution into account. ",
        "introduction": "Many globular clusters do show some amount of rotation, even the old ones of our own galaxy (\\citealt{meylan86,lupton87} and some ten others, K. Gebhardt, personal communication). The more accurate and star-by-star observations become available, the more important it is to properly model the effects of rotation together with evolution of globular star clusters, since typically the amount of rotational energy in clusters is not dominant, but also not negligible. It is long known that indeed flattenings of galactic globular clusters correlates with their rotation, suggesting that rotation still is important for the shape \\citep{white87}. Proper Motions of stars in $\\omega $ Cen have been measured to get a real 3D profile of stellar motions including rotation, and in fantastic improvements of observations using the Hubble Space Telescope it became even possible to measure rotation of 47 Tuc in the plane of the sky \\citep{anderson03}. While without rotation the wealth of observational data such as luminosity functions and derived mass functions, color-magnitude diagrams, and population and kinematical analysis, obtained by e.g. the Hubble Space Telescope, for extragalactic as well as Milky Way clusters (cf. e.g. \\citealt{piotto02,barmby02,pulone03,rich05,zhao05,beccari06}, to mention only the few most recently appeared papers), is balanced by a decent amount of detailed modelling (see below), theorists and observers alike seem to be appallingly puzzled by rotation in clusters, and are disappointed that (multi-mass) King model fitting does not work very well (see e.g. \\citealt{piotto99b}, and for a specific example \\citealt{mclaughlin03}, therein).  One notable exception being the early three-integral models by \\citet{lupton87} nobody seemed to care about rotation in globular cluster modelling since then for a long time. There was a thesis by Jeremy Goodman \\citeyearpar{goodman83}, whose part on rotation of clusters remained unpublished, but who was the pioneer for the main idea to use a two-dimensional approximation using the distribution function as function of energy $E$ and $z$-component of angular momentum $J_z$, and neglect the possible dependencies on a third integral $I_3$ for the time being. Unfortunately the solution of the 2D orbit averaged Fokker-Planck (henceforth FP) equation, including the self-consistent numerical determination of diffusion coefficients, is an extremely complex numerical problem, involving two levels of numerical integrations in a second order integro-differential equations, in other words we need 4D nested discretization in $E$, $J_z$, $\\varpi$, and $z$ where the latter are the spatial coordinates.  It turned out that the problem was too challenging in the mid-eighties for the then available computers. With the help of Goodman, who kindly made his PhD thesis available at an early stage, \\cite{einsel99} could re-write from scratch a new FP code and publish a sequence of rotating cluster models (pre-collapse, equal masses). These studies have been recently improved in a comprehensive effort on the FP side \\citep{kim02,kim04}, and also promising results exist which show that the neglect of the third integral does not harm too much (if flattening is not extreme), since there is fair agreement with direct $N$-body models (see some examples in \\citetalias{kim02} and \\citealt{boily00}). In a new approach (\\citealt{ardi05}) we are working to back up the amount of good $N$-body data for rotating clusters by a large number of small $N$ averaged models and a few large models, as was done for the non-rotating case by \\cite{gierszh94,giersz96} and \\cite{gierszs94}. It is clear that while all work known to the authors at this moment concentrates on self-gravitating star clusters, the improvement of our knowledge and methods in the field of rotating dense stellar systems is extremely important for galactic nuclei, too, where a central star-accreting black hole comes into the game (some stationary modelling exists, such as \\citealt{quinlan90}, \\citealt{murphy91}, \\citealt{freitag02}). Also, claims have been made that some globular clusters in the halo of our galaxy and around M31 contain massive black holes \\citep{gerssen02}, and these results have been challenged by \\cite{baumgardt03a}. However, the latest paper could only provide a very limited number of case studies, due to the enormous computing time needed even on the GRAPE computers. Therefore it is very urgent to develop reliable approximate models of rotating star clusters with black hole, which is subject of current work.  Dynamical modelling of globular clusters and other collisional stellar systems (like galactic nuclei, rich open clusters, and rich galaxy clusters) still poses a considerable challenge for both theory and computational requirements (in hardware and software). On the theoretical side the validity of certain assumptions used in statistical modelling based on the FP  and other approximations is still poorly known. Stochastic noise in a discrete $N$-body system and the impossibility to directly model realistic particle numbers with the presently available hardware, are a considerable challenge for the computational side.  A large amount of individual pairwise forces between particles needs to be explicitly calculated to properly follow relaxation effects based on the cumulative effects of small angle gravitative encounters. Therefore modelling of globular star clusters over their entire lifetime requires a star-by-star simulation approach, which is highly accurate in following stellar two- and many-body encounters. Distant two--body encounters, the backbone of quasi-stationary relaxation processes, have to be treated properly by the integration method directly, while close encounters, in order to avoid truncation errors and downgrading of the overall performance, are treated in relative and specially regularized coordinates \\citep{ks65,mikkola98}, their centers of masses being used in the main integrator. The factual world standard of such codes has been set by Aarseth and his codes NBODY$n$-codes ($0\\leq n \\leq 7$, see \\citealt{aarseth99a,aarseth99b,aarseth03}). Another integrator {\\sc Kira} aimed at high-precision has been used \\citep{mcmillan96,portegies98,takahashi00}. It uses a hierarchical binary tree to handle compact subsystems instead of regularization.  Unfortunately, the direct simulation of most dense stellar systems with star-by-star modelling is not yet possible, although recent years have seen a significant progress in both hardware and software \\citep{makino98,aarseth99a,aarseth99b,spurzem99,makino05}. Despite such progress in hardware and software even the largest useful direct $N$-body models for both globular cluster and galactic nuclei evolution have still not yet reached the realistic particle numbers ($N \\sim 5\\times 10^5$ for globular clusters, $N \\sim 10^6$ to $10^9$ for galactic nuclei). However, recent work by \\cite{baumgardt02} and \\cite{baumgardt03b} has pushed the limits of present direct modelling for the first time to some $10^5$ using either NBODY6++ on parallel computers or NBODY4 on GRAPE-6 special purpose hardware. There is a notable exception of a direct one-million body problem with a central binary black hole tackled by \\cite{hemsendorf03}. Their code, however, being very innovative for the large $N$ force calculation, still lacks the fine ingredients of treating close encounters between stars and black hole particles from the standard $N$-body codes.  Bridging the gap between direct models and the most interesting particle numbers in real systems is far from straightforward, neither by scaling \\citep{baumgardt01}, nor by theory. There are two main classes of theory: (i) FP models, which are based on the direct numerical solution of the orbit-averaged FP equation \\citep{cohn79,murphy91}, and (ii) isotropic \\citep{lyndenbell80, heggie84} and anisotropic gaseous models \\citep{gierszs94,spurzem96}, which can be thought of as a set of moment equations of the FP equation.  On the side of the direct FP models there have been two major recent developments. \\cite{takahashi95,takahashi96,takahashi97} has published  new FP models for spherically symmetric star clusters,  based on the numerical solution of the orbit-averaged 2D FP equation (solving the FP equation for the distribution $f=f(E,J^2)$ as a function of  energy and angular momentum, on an $(E,J^2)$-mesh). \\cite{drukier99} have published  results from another 2D FP code based on the original Cohn (1979) code. In such 2D FP models  anisotropy, i.e. the possible difference between radial and tangential velocity dispersions in spherical clusters, is taken into account.  Secondly, another 2D FP model has been worked out recently for the case of axisymmetric rotating star clusters \\citepalias{einsel99,kim02}. Here, the distribution function is assumed to be a function of energy $E$ and the $z$-component of angular momentum $J_z$ only; a possible dependence of the distribution function on a third integral is neglected. As in the spherically symmetric case the neglection of an integral of motion is equivalent to the assumption of isotropy, here between the velocity dispersions in the meridional plane ($r$ and $z$ directions); anisotropy between velocity dispersion in the meridional plane and that in the equatorial plane ($\\phi$-direction), however, is included. We realize that the evolutionary models provided by us for rotating globular clusters are difficult to use for direct comparisons with observations, because they are not easily analytically describable. But they are the only ones which fully cope with {\\it all} observational data available nowadays (full 3D velocity data, including velocity dispersions in $\\varpi$ and $\\phi$-direction, rotational velocity, density, all as full 2D functions of $\\varpi$ and $z$). No other evolutionary model exists so far which is able to provide this information. With the advent of our new post-collapse and multi-mass models \\citepalias{kim02,kim04} and the inclusion of stellar evolution and binaries (work in progress) we will be able to deliver even more interesting results. Already the existing $N$-body study \\citep{ardi05} shows that rotation not only accelerates the collisional evolution but also leads to an increasing binary activity in the system.  Here we present a detailed and comprehensive set of our model data, covering a fair range of rotation rates and initial concentrations of galactic globular clusters, which is aimed to enable observers to use them for comparison with their data. The basic idea is to start up with a dynamical cluster age and a present day observed flattening, and then being able to pick a number of models from our evolutionary model data base which allow a prediction (within certain ambiguity) of how this clusters started up initially. In this paper we present the basic procedure of modelling and how we create the data, and give a few exemplary data to show the capabilities of our data base which is to be found in the web ({\\tt http://www.ari.uni-heidelberg.de/clusterdata}).  Sect. 2 and 3 summarize briefly the numerical and physical procedure used for our evolving rotating cluster models, and the choice of initial conditions. Sect. 4 describes the comprehensive set of data we provide for every of our models using one example set, and Sect. 5 shows how observational data can be used to obtain the proper entry points in our data for a more detailed theory-observation comparison. ",
        "conclusions": "Most globular clusters are flattened by rotation, and show interesting features in their rotation curves. While rotation is not dominant it is also not negligible, the total amount of rotational (ordered) kinetic energy (as compared to unordered kinetic energy) can be as high as a few 10 \\% typically. Young globular clusters are usually more flattened than older ones. This was observed in the Magellanic clouds \\citep{kontizas90,hill06} and in M31 \\citep{barmby02}; and theoretically studied \\citep{frenk82,goodwin97,bekki04}. But see also \\cite{mackey05} for the MW. While observations, especially using the utmost accuracy for single star measurements in globular clusters with the Hubble Space Telescope (e.g. \\citealt{anderson03,rich05}) provide a marvelous detail of morphological and sometimes surprising kinematical information, people seem sometimes still surprised that even the most elementary feature of flattening and rotation (and its dynamical implication for rotational and dispersion velocities) are so difficult to model and will not be fitted well by multi-mass King models (\\citealt{longaretti97}; \\citetalias{kim04}). We have provided detailed evolutionary models here which make it possible to compare them easily with observations, using data first obtained by \\cite{einsel99} for the equal mass pre-collapse phase. The basic idea of our data collection, which we have discussed here in a few examples, and which is complete at the web URL {\\tt http://www.ari.uni-heidelberg.de/clusterdata} is that one can go in it using simple global observational data (dynamical age, flattening and/or concentration of a cluster), and then pick from a small number of theoretical evolved models to compare with.  2-dimensional axisymmetric FP models were used in order to follow the dynamical evolution of rotating stellar systems by using initial parameters (King dimensionless central potential parameter $W_0$, dimensionless measure of rotational energy $\\omega_0$). Our data sample is the only one available from theoretical evolutionary modelling of rotating globular clusters and is aimed to allow direct comparison with observational data, as the presented by \\cite{merritt97} and \\cite{gebhardt00}, by using our formatted data files and IDL plotting routines provided. Rotation conduces the system to a phase of strong contraction (during a mass and angular momentum loss process) and accelerates the collapse time (gravo-gyro phase; e.g. \\citealt{hachisu79,hachisu82}). The observed rotational structure in the meridional plane and $v_{\\rm rot}/\\sigma$ curves (despite of the central region of clusters like M15, discussed in Sect.\\ref{obsdata}) can be reproduced by our models. Nevertheless our models are still idealized and simplified. Physical scalings, such as time scales, may change if more effects, such of binaries, are included.  The provided database of models with theoretical cluster data will make it possible for observers to select model data to compare with \\footnote [4]{(look for  http://www.ari.uni-heidelberg.de/clusterdata/)}. Our database will in the future be extended to more realism including a stellar mass spectrum and stellar evolution of singles and binaries, which all is subject of ongoing work (\\citealt{ardi05}). Also a study of rotating clusters with a central massive star-accreting black hole is under way (Fiestas et al. 2006, in prep.), which will provide important data for galactic nuclei and to check claims that some globular clusters contain intermediate mass black holes."
    },
    "0609/astro-ph0609506_arXiv.txt": {
        "abstract": "We present spectroscopy of a transit of the exoplanet HD~189733b. By modeling the Rossiter-McLaughlin effect (the anomalous Doppler shift due to the partial eclipse of the rotating stellar surface), we find the angle between the sky projections of the stellar spin axis and orbit normal to be $\\lam$. This is the third case of a ``hot Jupiter'' for which $\\lambda$ has been measured. In all three cases $\\lambda$ is small, ruling out random orientations with 99.96\\% confidence, and suggesting that the inward migration of hot Jupiters generally preserves spin-orbit alignment. ",
        "introduction": "A primary reason to study planets of other stars is to learn how typical (or unusual) are the properties of the Solar system. For example, the nearly circular orbits of Solar system planets were once considered normal, but we now know that eccentric orbits of Jovian planets are common (see, e.g., Halbwachs, Mayor, \\& Udry 2005; or Fig.~3 of Marcy et al.~2005). Likewise, gas giants were once thought to inhabit only the far reaches of planetary systems, an assumption that was exploded by the discovery of ``hot Jupiters'' (Mayor \\& Queloz 1995, Butler et al.~1997). This inspired theoretical work on planetary migration mechanisms that can deliver Jovian planets to such tight orbits (as recently reviewed by Thommes \\& Lissauer 2005 and Papaloizou \\& Terquem 2006).  Another striking pattern in the Solar system is the close alignment between the planetary orbits and the Solar spin axis. The orbit normals of the 8 planets are within a few degrees of one another (Cox et al.~2000, p.~295), and the Earth's orbit normal is only 7~degrees from the Solar spin axis (Beck \\& Giles 2005, and references therein). Presumably this alignment dates back 5~Gyr, when the Sun and planets condensed from a single spinning disk. Whether or not this degree of alignment is universal is unknown. For hot Jupiters in particular, one might wonder whether migration enforces or perturbs spin-orbit alignment.  For exoplanets, the angle between the stellar spin axis and planetary orbit normal (as projected on the sky) can be measured via the Rossiter-McLaughlin (RM) effect: the spectral distortion observed during a transit due to stellar rotation. The planet hides some of the velocity components that usually contribute to line broadening, resulting in an ``anomalous Doppler shift'' (for the theory, see Ohta et al.~2005, Gim\\'enez 2006, or Gaudi \\& Winn~2006).  Observations of the exoplanetary RM effect have been reported for HD~209458 (Bundy \\& Marcy 2000, Queloz et al.~2000, Winn et al.~2005, Wittenmyer et al.~2005) and HD~149026 (Wolf et al.~2006). Here we report observations of the RM effect for HD~189733. This system, discovered by Bouchy et al.~(2005), consists of a K dwarf with a transiting Jovian planet ($M_P = 1.15$~$M_{\\rm Jup}$) in a 2.2~day orbit. Our observations are presented in \\S~2, our model in \\S~3, and our results in \\S~4, followed by a brief summary and discussion. ",
        "conclusions": "We have monitored the apparent Doppler shift of HD~189733 during a transit of its giant planet. By modeling the RM effect, we find that the stellar spin axis and the orbit normal are aligned to within a few degrees.  This is the third exoplanetary system (and the third hot Jupiter) for which it has been possible to measure $\\lambda$. The first system was HD~209458, for which $\\lambda=-4\\fdg4\\pm 1\\fdg4$ (Winn et al.~2005; see also Wittenmyer et al.~2005, who modeled the RM effect but required $\\lambda=0$). The second system was HD~149026 (Wolf et al.~2006), for which $\\lambda = 11\\pm 14\\arcdeg$. The small observed values of $\\lambda$ suggest that the most common end-state of the inward migration of a hot Jupiter involves a close alignment.  With only 3 systems, we cannot yet measure the distribution of $\\lambda$, but we can test the hypothesis of random orientations (i.e., a uniform distribution in $\\lambda$). The weighted mean of the measured values of $|\\lambda|$ is 2.6\\arcdeg. If we replace the measured values by random numbers drawn from a uniform distribution, the probability that the weighted mean\\footnote{Here we have assumed the uncertainty in $\\lambda$ is independent of $\\lambda$, a good approximation because all 3 systems have an intermediate impact parameter (Gaudi \\& Winn 2006).} will be this small is only 0.04\\%. Hence we rule out the hypothesis of random orientations with 99.96\\% confidence.  Winn et al.~(2005) argued that tides from the star would not ordinarily cause alignment within the star's main-sequence lifetime. There are therefore two basic possibilities: either the alignment is primordial and was not disturbed by migration, or there was a different mechanism to damp any initial or induced misalignment. Among the various theories of hot Jupiter migration, some would tend to enhance any initial misalignments, and are thereby constrained by our results. Such scenarios include planet-planet scattering followed by circularization (Rasio \\& Ford 1996, Weidenschilling \\& Marzari 1996), Kozai migration (Wu \\& Murray 2003; Eggenberger et al.~2004), and tidal capture (Gaudi~2003).  The agreement among the three systems studied to date is clear, but should not discourage future measurements. Obviously, a sample of three is only barely sufficient to draw a conclusion. And of course, the discovery of even a single example of a grossly misaligned system would be of great interest."
    },
    "0609/astro-ph0609676_arXiv.txt": {
        "abstract": "We found a very likely counterpart to the recently discovered hard X-ray source \\igr\\ in the multi-wavelength observations of the source field. The source, originally discovered in the 20-40 keV band, is now confidently detected also in the 40-80 keV band, with a flux of $(1.4 \\pm 0.4) \\times 10^{-11} \\enf$. A 5 ks \\swift\\ observation of the \\igr\\ field revealed a hard point-like source with the observed 0.5--10 keV flux of 3.4$^{+0.7}_{-0.8}\\times$10$^{-12}$ erg cm$^{-2}$ s$^{-1}$ (90\\% confidence level) at $\\alpha = 20^{\\rm h}18^{\\rm m}38\\fs 55$, $\\delta = +40^\\circ 41' 00\\farcs4$ (with a 4\\farcs2 uncertainty). The combined {\\sl Swift}-{\\sl Integral} spectrum can be described by an absorbed power-law model with photon index $\\Gamma=1.3\\pm 0.2$ and $N_{\\rm H}=6.1^{+3.2}_{-2.2}\\times 10^{22}$ cm$^{-2}$. In archival optical and infrared data we found a slightly extended and highly absorbed object at the \\swift\\ source position. There is also an extended VLA 1.4 GHz source peaked at a beam-width distance from the optical and X-ray positions. The observed morphology and multiwavelength spectra of \\igr\\ are consistent with those expected for an obscured accreting object, i.e. an AGN or a Galactic X-ray binary. The identification suggests possible connection of \\igr\\ to the bright $\\gamma$-ray source GEV J2020+4023 detected by {\\sl COS B} and {\\sl CGRO} EGRET. ",
        "introduction": "The first hard X-ray imaging observation of the Cygnus region with the IBIS-ISGRI imager (Lebrun et al.\\ 2003) aboard the {\\sl International Gamma-Ray Astrophysics Laboratory} ({\\sl INTEGRAL}; Winkler et al.\\ 2003) has revealed a new source, IGR J2018+4043 (Bykov et al.\\ 2004). The 320 ks fully coded field of view (FCFOV) observation yielded the source  flux of $(1.7 \\pm 0.4) \\times 10^{-11} \\enf$ in the 25-40 keV band. The Cygnus region is very rich in starforming sites, X-ray binaries, and supernova remnants (SNRs). \\igr\\ is located in the field of the SNR G78.2+2.1 (commonly known as $\\gamma$-Cygni SNR). A high-energy \\gr source, 2CG 078+2, in the \\gc\\ field was discovered with the {\\sl COS B} satellite (e.g., Swanenburg et al.\\ 1981). Observations with {\\sl CGRO} EGRET showed that this source (2EG\\,J2020+4026 = 3EG\\,J2020+4017 = GEV\\,J2020+4023), whose photon flux is $\\sim$ 1.2 $\\times 10^{-6}\\fl$ above 100 MeV (Sturner \\& Dermer 1995; Hartman et al.\\ 1999), is the brightest among apparently steady unidentified EGRET sources. {\\sl Whipple} \\gr\\ telescope observations (e.g., Buckley et al.\\ 1998) have established an upper limit of 2.2 $\\times 10^{-11}\\fl$ for the flux above 300 GeV, indicating a break in the spectrum above a few GeV. Bykov et al.\\ (2004) suggested that the EGRET and \\isgri\\ sources could be the $\\gamma$-ray and hard X-ray counterparts of the same source in \\gc, e.g., a Geminga-like pulsar or extended emission from particles accelerated in the supernova ejecta/shell. In particular, the interaction of the wind of the early type O9V star HD 193322, located just 7\\arcmin\\ from IGR J2018+4043, with \\gc\\ would be a plausible mechanism of particle acceleration. On the other hand, the \\isgri\\ source position within the SNR field does not necessarily imply their physical connection. An example of such a confusion is the source IGR J17204--3554 projected onto the rich Galactic starforming region NGC 6334. The most likely identification of that source is a highly obscured extragalactic object NGC 6334B (see Bassani et al.\\ 2005; Bykov et al.\\ 2006).  Given the nominal 12\\arcmin\\ angular resolution (FWHM) and about 1\\arcmin point source localization of \\isgri\\ (Lebrun et al.\\ 2003), a reliable identification of the relatively weak (4--5 $\\sigma$) \\ii\\ source requires high-resolution observations with other instruments. We report here the results of recent 1.5 Ms \\hbox{FCFOV} observations of the source with \\isgri\\ (\\S2.1) and the detection of a counterpart of IGR J2018+4043 in the 0.5--10 keV band with {\\sl Swift} XRT (\\S2.2). The accurate positioning of the {\\sl Swift} source allowed us  to find its possible radio/IR/optical counterparts in archival data (\\S2.3 and \\S2.4). Based on these multiwavelength observations, we suggest that the source is a strongly obscured accreting source, likely an AGN, and discuss the possibility that the EGRET source is indeed its $\\gamma$-ray counterpart (\\S3). ",
        "conclusions": "The {\\sl Swift}-ISGRI spectrum of J2018 shows that it is a nonthermal, highly obscured source. The hydrogen column density estimated from the {\\sl Swift} data (\\S2.2) looks too high for an object in the \\gc\\ SNR ($d\\approx 1.5$ kpc, expected $N_{\\rm H} \\sim 1$--$3\\times 10^{21}$ cm$^{-2}$; Landecker et al. 1980). J2018 is likely a more distant object, not associated with \\gc.  A number of ISGRI sources have been identified as Galactic X-ray binaries, mostly HMXBs (e.g., Filliarte \\& Chaty 2004; Tomsick et al.\\ 2006, and references therein). The J2018's X-ray spectrum (\\S2.2) and luminosity, $L_X\\sim 3\\times 10^{35}(d/{\\rm 10\\,kpc})^2$ ergs s$^{-1}$ in 1--80 keV range, as well as the magnitudes of the IR-optical counterpart (\\S2.4) do not contradict such an interpretation. The apparent $\\sim$$10\\arcsec$ extension of the optical and 2MASS counterparts could be associated with a powerful wind of an HMXB's O or B companion. Alternatively, the extended IR-optical source could be interpreted as a galaxy that hosts an AGN detected in X-rays.    An extrapolation of the X-ray spectrum of J2018 into the EGRET range reaches the flux level of GEV J2020+4023, which suggests that these sources could be X-ray and $\\gamma$-ray counterparts of the same object. An extensive dedicated search for a counterpart to 3EG J2020+4017 with \\chan\\ and the Green Bank radio telescope in a close vicinity of its nominal position failed to show any clear candidate object (Becker et al.\\ 2004; Weisskopf et al.\\ 2006). However, J2018 is located outside the fields of these searches.   The position of J2018 is 30\\farcm6 (35\\farcm5) offset from the nominal position of GEV J2020+4023 (3EG J2020+4017), exceeding the $8'$ ($9\\farcm6$) 95\\% localization accuracy in those catalogs (Lamb \\& Macomb 1997; Hartman et al.\\ 1999). However, such offsets are not unusual among EGRET sources. The position accuracy of relatively faint EGRET sources in the Galactic plane was estimated by Gehrels et al.\\ (2001) as $\\sim$$1\\degmark$. Moreover, even the very bright Vela and Geminga pulsars are located outside their EGRET 99\\% confidence contours (see Hartman et al.\\ 1999). 3EG J2020+4017 is a bright source, but it is located in a very crowded region of the plane, with enhanced, highly inhomogeneous diffuse emission from the \\gc\\ SNR and nearby molecular clouds, and at least two other EGRET sources (3EG J2021+3716 and J2016+3657) within the 3$\\degmark$ vicinity, which may have affected the source localization. In such an environment, systematic uncertainties can  increase the localization error up to $\\sim$0\\fdg5. Thus, despite the positional discrepancy, we believe that one can consider the possibility that {\\sl Swift}-ISGRI and EGRET detected the same source. Future {\\sl GLAST} (e.g., Gehrels \\& Michelson 1999) observation could resolve the issue.  \\begin{figure*} \\begin{center} \\includegraphics[width=0.95\\textwidth,angle=0]{f3.eps} \\caption{ Observed ({\\em top}) and extinction-corrected ({\\em bottom}) multiwavelength spectra of J2018. The NIR-optical and soft X-ray points in the bottom panel are corrected for absorption assuming $A_V=9.8$ (see \\S3) and $N_{\\rm H}=6.1\\times 10^{22}$ cm$^{-2}$, respectively. Note that the observational data are noncontemporaneous. The EGRET data points for 3EG\\,J2020+4017 are included, although the association with J2018  is not firmly established (see \\S3).} \\label{fig3} \\end{center} \\end{figure*}   It is also worth mentioning the $\\sim 3\\,\\sigma$ detection of a possible Fe K$_\\alpha$ feature at 6.2 keV in the {\\sl RXTE} spectrum of \\gc\\ (Bykov et al.\\ 2004). The feature, which could be attributed to a putative hidden accreting source (e.g., an AGN), does not appear in the {\\sl ASCA} spectra of the \\gc\\ regions observed by this satellite, which did not include the J2018 position, while this position was well within the {\\sl RXTE} collimator field.   To investigate the nature of J2018, we constructed its spectral energy distribution (SED; Fig.\\ \\ref{fig3}) using the \\isgri\\ and \\swift\\ observations and the archival data described above. It should be noted that the data are noncontemporaneous, which is a point of some concern because the source could be variable. The Galactic absorption ($A_V\\approx 9.8$) was accounted for in the unabsorbed NIR-optical spectrum using the reddening maps by \\citet{schleg}\\footnote{See also http://irsa.ipac.caltech.edu/applications/DUST/}. Being rather accurate for high latitude sources, the method often overestimates the reddening by a factor of up to 1.5 for strongly reddenned objects (Arce \\& Goodman 1999). The optical emission is apparently extended, but with the data available it is not possible to disentangle the core component, corresponding to the point source detected  by {\\sl Swift}, from the extended halo. Therefore, the optical points can be lower than shown in Figure 3, and the extinction corrected (with $A_V= 9.8$) optical fluxes in Figure \\ref{fig3} should be regarded as upper limits, because of a possible overestimation (up to 3 mag) in the $A_V$ used.    The hypothesis that the 1.4\\,GHz VLA excess is due to a point-like source at the \\swift\\ source position (e.g., an AGN) projected onto an extended radio source is not in conflict with the data available. Therefore, we conditionally include the VLA point, as well as the EGRET data, in Figure 3. Within a more conservative approach, they should be regarded as upper limits.   The radio through hard X-ray SED in Figure \\ref{fig3} is consistent with that expected for a nearby Seyfert AGN or a radio-galaxy (e.g., Dermer \\& Gehrels 1995). For instance, at $z=0.02$ the source luminosity below 100 keV would be about $10^{43}$ ergs s$^{-1}$.  If GEV J2020+4023 is indeed associated with J2018, then the radio through $\\gamma$-ray source could be a blazar. Given the $A_V$ uncertainties discussed above, the multiwavelength spectrum resembles the two-peak synchrotron self-Compton  SED of S5 0716+714, a blazar studied by Foschini \\etal (2006). Simultaneous multiwavelength observations are required, however, to verify the blazar hypothesis and understand the processes responsible for such a spectrum. On the other hand, the leptonic jet model for a microquasar, such as discussed by Dermer \\& B\\\"{o}ttcher (2006) for LS 5039, also provides a SED generally consistent with that given in Figure 3.  NIR and optical spectroscopy could distinguish between the two interpretations by measuring the redshift of the galaxy that hosts the putative AGN. Also, microquasars are usually variable on minutes scale, while AGNs are variable on longer time scales making a variability study conclusive.  The analysis given above allows us to conclude that the properties of the source discovered by \\sw\\ are consistent with those of \\igr. However, that does not exclude a sizable extended hard X-ray flux from the possible interaction of the \\gc\\ SNR with stellar winds as discussed by Bykov et al.\\ (2004). The extended VLA feature seen in Figure 1 can also be attributed to the SNR-wind interaction. High-resolution observations with \\chan\\ and \\xmm, as well as simultaneous multiwavelength studies, are needed to separate the components and to make a firm conclusion on the nature of IGR J2018+4043."
    },
    "0609/astro-ph0609559.txt": {
        "abstract": "% context heading (optional) {} % aims heading (mandatory) { We present an overview of a high-mass star formation region through the major (sub-)mm, and far-infrared cooling lines to gain insight into the physical conditions and the energy budget of the molecular cloud. %, and the evolution of the cloud. } % methods heading (mandatory) { We used the KOSMA 3m telescope to map the core ($10'\\times 14'$) of the Galactic star forming region DR\\,21/DR\\,21\\,(OH) in the Cygnus X region in the two fine structure lines of atomic carbon (\\CIw{} $^3\\RM{P}_1\\TO^3\\RM{P}_0$ and $^3\\RM{P}_2\\TO^3\\RM{P}_1$) and four mid-$J$ transitions of CO and $^{13}$CO, and CS $J=7\\TO6$. These observations have been combined with FCRAO $J=1\\TO0$ observations of $^{13}$CO and C$^{18}$O. Five positions, including DR21, DR21\\,(OH), and DR21\\,FIR1, were observed with the ISO/LWS grating spectrometer in the \\OI\\ 63 and 145 $\\mu$m lines, the \\CII\\ 158 $\\mu$m line, and four high-$J$ CO lines. We discuss the intensities and line ratios at these positions and apply Local Thermal Equilibrium (LTE) and non-LTE analysis methods in order to derive physical parameters such as masses, densities and temperatures. The CO line emission has been modeled up to $J=20$. } % results heading (mandatory) { From non-LTE modeling of the low- to high-$J$ CO lines we identify two gas components, a cold one at temperatures of T$_\\RM{kin}\\sim 30-40$\\,K, and one with T$_\\RM{kin}\\sim 80-150$\\,K at a local clump density of about n(H$_2$)$\\sim 10^4-10^6$\\,cm$^{-3}$. {{\\NEW While the cold quiescent component is massive containing typically more than 94\\,\\% of the mass, the warm, dense, and turbulent gas is dominated by mid- and high-$J$ CO line emission and its large line widths. %     gas is found at all five positions. %     This dense and turbulent gas is due to the large line widths very prominent in mid- and high-$J$ CO line emission. }} The medium must be clumpy with a volume-filling of a few percent. The CO lines are found to be important for the cooling of the cold molecular gas, e.g. at DR21\\,(OH). Near the outflow of the UV-heated source DR21, the gas cooling is dominated by line emission of atomic oxygen and of CO. Atomic and ionised carbon play a minor role. %     The cold and quiescent gas at DR21\\,(OH) is cooled dominantly by the CO emission lines. } % conclusions heading (optional), leave it empty if necessary {} ",
        "introduction": "Emission lines of \\CIIw{}, \\CIw{}, and CO in molecular clouds are among the most important diagnostic probes of star formation and particularly important for studying the influence of massive stars and their UV radiation on the environment. Obtaining the physical parameters of such regions %using observations of spectral lines is often done with simplified radiative transfer assumptions. %As long as only low-$J$ rotational lines of CO are observed, Local Thermal Equilibrium (LTE) and the Large Velocity Gradient (LVG) analysis have proven to yield good results for the bulk of the cold gas (see, e.g., \\citet{black2000}). However, the assumption of a homogeneous medium does not represent the typical physical conditions of a high mass star formation region. {\\NEW %Naturally, Instead, the observing beam will contain several gas components with different excitation conditions and with unknow filling factors. %Instead, the observed The observed line transitions are sensitive to gas at different density, temperature and opacity, corresponding to different regions. } % %But as soon as mid- and high-$J$ CO lines are included, %a high temperature and high density solution is obtained which does not fit well to estimates from low-$J$ only transitions, %because the transitions are sensitive to gas at different density, temperature and opacity, e.g., in different regions. %It was soon found that {\\NEW By selecting designated %By assigning dedicated %By assigning different lines as tracers of } the different parts of the gas, %(with respect to the density and temperature regime), one can obtain a rather complete coverage of excitation conditions accounting for different physical and also chemical conditions {\\NEW within the beam}. A more differenciated analysis using non-LTE radiative transfer algorithms (see \\citet{vanzadelhoff2002} for an overview) is needed for this purpose. It can provide %remedy the problem a more detailed description of the structure of the cloud including clumpiness and temperature or density gradients (e.g., \\citet{ossenkopf2001}, \\citet{williams2004}). % \\OUT{ Although optically thin thermal dust emission at mid- to far-IR wavelengths is a valuable tool to determine the column density and the temperature it may not resolve the velocity structure or give an estimate of the local density. } %  % Changed ossk Using the results of this analysis, we can try to retrieve a self-consistent picture including all major processes driving and driven by high mass star formation (e.g, with respect to turbulence or chemistry). %However, only in a self-consistent approach (e.g., including cloud dynamics or timedependent chemistry), %processes driving and driven by high mass star formation may be fully understood. %physical models solving the physical %chemical %Although, those multi-wavelength studies begin to draw a coherent picture, %the processes driving and driven by high mass star formation are still not well understood. %For instance, observations of \\HII{}-regions and Young Stellar Objects (YSOs) in the near-infrared, X-Ray, mm-continuum, and others %open very rich and active fields. %As such, {\\NEW Thus,} we obtain a detailed view of the Photon Dominated Regions (PDRs) %forming {\\NEW in} the environment of young stars. PDRs are transition regions from ionized and atomic to dense molecular gas in which physics and chemistry {\\NEW of} the gas are dominated by the UV radiation from young massive stars. % of massive star forming clouds. Processes heating and cooling the clouds are qualitatively well understood on a theoretical level \\citep[e.g.,][]{tielens1985,sternberg1989}, and in general we find a reasonable agreement between the theoretical predictions and observations of the brightest cooling lines. Still, there is an ongoing debate about the details of the energy input contributing to the heating. We find a combination of mechanisms such as stellar far-ultraviolet radiation (FUV), dynamical energy through shocks or supernovae, or an enhanced cosmic ray rate. There are also many open questions with respect to the chemical composition of these regions. The role of nonequilibrium chemistry is still not well understood. For instance Polycyclic Aromatic Hydrocarbons (PAHs) as a reservoir of fresh gas phase carbon (e.g., \\citet[][]{oka2004}, \\citet{habart2005}) may change the overal picture. %the cooling of gas phase species are important.  % \\citep[][]{bradford2003}. {\\NEW Emission from warm carbon in the form of atoms, ions or in carbonated molecules (mainly CO) provides a significant fraction of the gas cooling at sub-millimeter and far-infrared (FIR) wavelengths} %A significant fraction of the gas cooling is emitted by warm carbon species %(in molecular, atomic, or ionic form) at sub-millimeter and far-infrared (FIR) wavelengths %The widely observed high fraction of warm molecular gas (e.g., \\citet{giannini2000,giannini2001}, \\citet[][]{schneider2003}, \\citet{kramer2004}, or \\citet{bradford2003,bradford2005}). %e.g, in the starburst galaxy NGC\\,253, %It is also of importance to exploit %The carbon species -- in molecular, atomic, or ionic form -- with lines at sub-millimeter and far-infrared (FIR) wavelengths %account -- with respect to the energy balance -- for a significant fraction of the gas cooling. %Other important species are those with oxygen, e.g., \\OI, H$_2$O, OH. Finding the prevailing cooling channels among these and other species %with oxygen, (e.g., \\OI, H$_2$O, OH) may provide significant insight into the details of the chemical composition and the energy balance of PDRs. Since systematic observations of sub-mm- and FIR-line cooling lines with a good spatial resolution {\\NEW are still sparse, data to test the model predictions are still largely missing and many questions concerning the details of the chemistry and energy balance remain unanswered. By studying a well-known, massive Galactic star forming region, we aim at gaining new insight into these problems.} %are still rare, many serious gaps %exist in our understanding of these processes. % end changes ossk %In the following, we aim to gain insight into these %aspects using a well studied Galactic massive star forming region.  %%%Here, radiative transfer calculations can be a helpful tool to estimate the total line cooling power. %With the help of radiative transfer calculations, %the total line cooling power %at sub-millimeter and far-infrared (FIR) wavelengths %can be estimated. %as done, e.g., in %\\citet[][]{schneider2003}, \\citet{kramer2004}, \\citet{giannini2000,giannini2001}, %\\citet{bradford2005}, and others. % \\OUT{ the mid- and high-$J$ CO line emission was assigned to a high temperature gas component at densities of $10^5$\\,\\percc{} and higher. This component constitutes only a few percent of the total cloud mass and is thus not representative for the analysis of the (cooler) main gas component. %hinders the analysis of the main gas component. %For a better overview we need to address the following questions:  In order to address this problem we need to pin down the origin of the line emission. %{\\it (i)} CO rotational lines proved to be good tracers of the molecular gas. This was shown many times in comparison studies to other surrogates as well as to dust emission studies \\citep[e.g., ][]{goldsmith1997}. %{\\it (ii)} Collisions with H$_2$ molecules and the radiative transfer through the local radiation field dominate the CO excitation. They depend on the kinetic temperature and density structure of a cloud. Hence, consistent modeling using radiative transfer methods should tell us about the physical conditions in the cloud. } % %{\\it (iii)} % % \\OUT{ Theoretical \\citep{tielens1985,sternberg1989} and observational studies (references) all come to the same conclusion of a layered structure: the ionising photons first create a zone (at a visual extinction of A$_v\\sim$1$^m$) where hydrogen and atomic oxygen are neutral and atomic carbon ionised. Deeper penetration into in a clumpy molecular medium \\citep[ and other references]{stutzki1988} leads to an enriched zone of molecular gas where carbon first becomes neutral, and the increasing attenuation of the UV radiation leads to a warm molecular zone of high excitation (at A$_v\\sim$2--5$^m$), followed by the cloud core (A$_v\\sim$10$^m$) with carbon mostly contained in CO molecules. The brightest cooling lines in the sub-millimeter and far-infrared (FIR) wavelength regime for the gas are the \\CII\\ 158\\,\\micron{}, \\OI\\ 63\\,\\micron{} finestructure lines, the CO rotational lines, H$_2$ near infrared lines, and numerous Polycyclic Aromatic Hydrocarbon (PAH) features. They are crucial to determine the physical parameters of the PDR.  The common approach to study PDRs is to combine observed line intensities and ratios from PDR tracers with theoretical calculations in order to reproduce the observations by varying the hydrogen density and intensity of the UV radiation field. Several studies of high mass star forming regions have already been presented (references) but many lack a sufficient number of observed species and positions in order to pinpoint exactly the performance and limitations of existing PDR models. }  \\OUT{ In this paper, we present maps and single pointings of a large number of PDR tracers in the DR21 region (see below) in order to examine the physical properties of this star forming region. Particularly multi-wavelength studies enable to draw a coherent picture of the processes driving and driven by high mass star formation. Thus, large-scale surveys in \\CI\\ and CO emission obtained with, e.g., the KOSMA, Mt.Fuji, and AST/RO telescopes in conjunction with the ISO or Kuiper Airborne Observatory (and in future with Herschel) far-IR data provide a useful diagnostic of the conditions associated with active massive star formation \\citep[e.g.,][]{schneider2003,kramer2004,yamamoto2001,oka2004,zhang2001}. }    \\subsection*{The DR\\,21 region}  The DR\\,21 \\HII-region/molecular cloud is part of the Cygnus~X complex of molecular clouds. Cygnus~X is known to host a large number of \\HII{}-regions \\citep{downes1966,wendker1984} associated with molecular clouds %JSt this holds for basically every and any HII region! %JST as seen, e.g., in the CfA $^{12}$CO 1\\TO0 survey (\\citet{leung1992}). An extended (11 square degrees) map of the Cygnus X region in the $^{13}$CO 2\\TO1 line and smaller maps in the CO and $^{13}$CO 3\\TO2 lines were obtained with the KOSMA telescope and are presented in \\citet{schneider2006} (from now on \\citetalias{schneider2006}). The distance to DR\\,21 was sometimes estimated to be $\\sim3$\\,kpc, %based on arguments other than the radial velocity. based on visual interstellar extinction and radial velocity \\citep[e.g.,][]{campbell1982}. For this paper, we follow the argumentation from \\citetalias{schneider2006} based on a morphological comparison of the molecular line data with mid-IR emission and the OB associations, and favour a value of 1.7\\,kpc as distance toward DR21 and DR21\\,(OH). However, for easier comparison with the literature, we explicitly give distance depended physical properties as a function of distance.  The DR\\,21 region has been subject to numerous studies at different wavelengths. Several active star formation sites and cometary shaped \\HII{}-regions were identified in maps of radio continuum emission \\citep[e.g., ][]{dickel1983,cyganowski2003}. %The presence of non-metastable ammonia lines, as detected by Ammonia observations by \\citet[][]{mauersberger1985} were interpreted %in \\citeyear{mauersberger1985} toward DR21, was interpreted as evidence for high densities ($\\sim$$10^{5.5}$\\,\\percc{}) paired with a high degree of clumpiness. H$_2$O masers \\citep{genzel1977} and a recently obtained map of 1.3-mm continuum emission \\citep{motte2005} show that DR\\,21 belongs to a north-south orientated chain of {\\sl massive star forming complexes}. Additional indication of active star formation is also given by observations of vibrationally excited H$_2$, tracing the hot, shocked gas in the vicinity of DR\\,21 and DR21\\,(OH) (e.g., \\citet{garden1986}, {\\NEW and more recent results by \\citet{davis2006}}). H$_2$ images reveal an extended ($5'$=2.5\\,pc(D/1.7\\,kpc)) north-east to south-west orientated outflow centered on DR21. %similar to Spitzer 4.5\\,\\micron{} data \\citep{marston2004}. To the west, the outflow is highly collimated, whereas toward the east side, gas is expanding in a blister style \\citep{lane1990}. \\citet{jaffe1989} found very intense CO 7\\TO6 emission peaking on the \\HII{}-region and following the H$_2$-emission.      %________________________________________________________________ % %                          Summary (Introduction) %________________________________________________________________  In this paper, we present the results of sub-mm mapping observations with KOSMA and discuss far-IR line intensities observed with ISO. % of DR21 as described in the following section. This comprehensive data set allows to study the spatial and partly also the kinematic structure of all major cooling lines of the ISM in the vicinity of DR\\,21 and DR\\,21(OH). In Sect. 3, we analyse the line emission and continuum data toward five positions, by comparing the spatial distribution, line profiles, and the physical properties through selected line ratios and the far-IR dust continuum. In Sect. 4, we model the integrated line intensity ratios and absolute intensities with radiative transfer models. The impact of the line emission on the total cloud cooling is dicussed in Sect. 5. %, and whether PDR excitation is sufficient or other heating (e.g., shock excitation) is necessary, A summary is given in Sect. 6.      %________________________________________________________________ % %                          End Introduction %________________________________________________________________    %________________________________________________________________ % %                          Table ISO-Positions %________________________________________________________________  \\begin{table}[ht] \\caption[] {Positions of ISO observations in absolute coordinates. DR21\\,C{} is centered on the \\HII{}-region. Projected distance of DR21\\,E{} (resp. DR21\\,W{}) from position DR21\\,C{} is 0.77\\,pc (resp. 0.89\\,pc) at a distance of 1.7\\,kpc. The TDT-number is given in the second column. %The ProposerID is PCOX, JTAUBER, and EVDISHOE. %Offsets are relative to DR\\,21 $\\alpha_{(B1950)}=20^\\RM{h}39^\\RM{m}01^\\RM{s}.1$, %$\\delta_{(B1950)}=+42:19:43$. %Offsets are relative to DR\\,21 $\\alpha_{(J2000)}=20^\\RM{h}38^\\RM{m}00^\\RM{s}.0$, %$\\delta_{(J2000)}=+42^\\RM{d}00^\\RM{m}00^\\RM{s}$.  % ref 681.1 1183.0 \\label{tab-iso-positions}} \\begin{tabular}{lclcc} \\noalign{\\smallskip} \\hline \\hline \\noalign{\\smallskip} Name & TDT-No. & R.A. [J2000] & Dec. [J2000]\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} DR\\,21 E{}/East & 15200785 & $20^\\RM{h}39^\\RM{m}09^\\RM{s}.14$ &$42^\\circ20'05''.0$ \\\\%771 1205   & 90 22\\\\ %DR\\,21 East& 55500823 & \\\\%746 1178.5 &\\\\ %           & 54001323 &            &\\\\ DR\\,21 C{}/Center & 15200786 & $20^\\RM{h}39^\\RM{m}00^\\RM{s}.93$ &$42^\\circ19'42''.0$ \\\\%679 1182   & -2 -1\\\\ %DR\\,21 West& 54001470 & \\\\%601 1147   &\\\\ DR\\,21 W{}/West & 15200787 & $20^\\RM{h}38^\\RM{m}52^\\RM{s}.23$ &$42^\\circ19'00''.5$\\\\%582 1140.5 & -99 -42\\\\ DR\\,21 (OH)& 34700439 & $20^\\RM{h}39^\\RM{m}00^\\RM{s}.95$ &$42^\\circ22'37''.6$ \\\\%679 1358   &-2 175\\\\ DR\\,21 FIR1& 35500317 & $20^\\RM{h}39^\\RM{m}00^\\RM{s}.2$  &$42^\\circ24'34''.6$\\\\%671 1475   &-10 292\\\\ %           & 52800917 &            &\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{table}  %________________________________________________________________ % %                          Figures KOSMA-MAPS %________________________________________________________________  \\begin{figure*}[ht] % @ plot.graphic \\includegraphics[width=9.05cm,angle=-90]{images/out_co43_2000.eps} \\includegraphics[width=9.05cm,angle=-90]{images/out_co76_2000.eps} \\caption{Integrated intensity map of CO 4\\TO3 ({\\it left}) and CO 7\\TO6 ({\\it right}) of DR\\,21 \\& DR21\\,(OH). The integration interval is $-20$ to 30 $\\kms$. The contours for CO 4\\TO3 run from 66($3\\sigma$) to 462$\\,\\kkms$ in steps of $3\\sigma$, and from 68($2\\sigma$) to 671$\\,\\kkms$ in steps of $2\\sigma$ for CO 7\\TO6. The HPBW is indicated in the upper left corner. The positions of ISO/LWS observations (cf. Fig. \\ref{fig_spitzer}) are indicated by circles. %Black crosses are positions of radio continuum sources identified as OB stars by \\citet{roelfsema1989}. %Maps were created using a re-gridding algorithm. %The small points indicate the observed positions. The dashed white lines correspond to the cuts shown in Fig. \\ref{plot-co}. \\label{fig_map_co} } \\end{figure*} \\begin{figure*}[ht] \\includegraphics[width=9.05cm,angle=-90]{images/out_ci10_2000.eps} \\includegraphics[width=9.05cm,angle=-90]{images/out_ci21_2000.eps} \\caption{Integrated intensity map of \u00a0\\CI{} 1\\TO0 ({\\it left}) and of \\CI{} 2\\TO1 ({\\it right}). The integration interval is $-10$ to 20 $\\kms$. The contours for \\CIw{} 1\\TO0 run from 12($2\\sigma$) to 60$\\,\\kkms$ in steps of $2\\sigma$, and from 20($1\\sigma$) to 108$\\,\\kkms$ in steps of $\\sigma$ for \\CI{} 2\\TO1. \\label{fig_map_ci} } \\end{figure*}     %________________________________________________________________ % %                          Observations %________________________________________________________________ ",
        "conclusions": "%________________________________________________________________ % %                      ISO Data / UV-field %________________________________________________________________  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%   \\OUT{ The observed region is known to harbour high mass star forming regions at different evolutionary stages (references?), and the maps confirm the diversity of this region. The combination of different observations into a unique dataset allowed a deep view into this high-mass star-forming region. }  We presented the first comprehensive mapping of sub-mm lines in the whole DR21/DR21\\,(OH) region. The line emission of \\CI{} at 492\\,GHz and 809\\,GHz, CO 3\\TO2, 4\\TO3, 6\\TO5, and 7\\TO6, $^{13}$CO 3\\TO2 and 6\\TO5, and of CS 7\\TO6 %The results of our detailed line analysis are: %\\begin{enumerate} %\\item reveals contributions from different gas components: {\\it (i) } A very wide winged emission in mid-$J$ CO transitions having a line width of more than $15\\,\\kms$ and self-absorption, and {\\it (ii)} narrow (4 to $6\\,\\kms$) emission in the other lines. Based on these lines, and supported by mm- and FIR-lines from FCRAO and ISO, we confined temperature, density, masses, and volume filling for five positions toward DR21, DR21\\,(OH) and DR21\\,FIR1 using a radiative transfer code by modeling the spectral line profiles.  %\\item  %\\item   %\\item 1. {\\NEW We found that 2-component models consisting of a warm inner region and a cold envelope} are needed to explain the observed strengths of the energetically low and high lying transitions of CO and to explain the self-absorption.  2. The bulk of molecular gas at DR21 \\& DR21\\,(OH) is cold ($\\sim$$30$\\,K) and concentrated in the ridge. The clump volume filling is typically a few percent and the density within a clump is of the order $10^6$\\,\\percc{}.   3. By fitting the dust emission SED, we constrain the H$_2$ column density to 2--5 $10^{22}$\\,\\percc{} for DR21\\,E{}, W{} and toward FIR1. At DR21/DR21\\,(OH) it rises to above $10^{23}$\\,\\percc{}. The dust temperatures vary between 31--43\\,K and are close to the temperature of the cold gas.   %\\item 4. Warm gas (80--150\\,K) was found  at all five positions. %but  toward the density peaks and correlated with the outflow axis. The warm gas mass is typically a few percent of the total mass but accounts for most of the total CO emission (80\\.\\% to 94\\,\\%, mainly in the mid- and high-$J$ tranitions). DR21\\,(OH), as the coldest source, still shows a significant amount of high-$J$ CO emission.  5. CS 7\\TO6  emission can be reproduced if the local density in the envelope region is sufficiently high. Therefore, high densities are expected at DR21 and DR21\\,(OH) supporting again the low volume filling factors.  %\\item 6. Supported by the model results, we derived the total intensity of the strongest emission lines (\\OI{}, CO, \\CII{}, and \\CI{}). Line emission provides up to 2\\,\\% to the total cooling of the atomic and molecular gas and by dust grains. At the dense cores DR21 and DR21\\,(OH), the cooling efficiency is less than 0.5\\,\\%. At DR21, \\OIw{} and CO are equally strong, whereas at DR21\\,(OH) the CO lines dominate the cooling with more than 75\\,\\%. \\CIIw{} contributes generally less than 20\\,\\% and \\CIw{} not more than 1.5\\,\\% to the line cooling. The \\OI{} emission shows a large variation from $<15$\\,\\% at DR21\\,(OH) and FIR1 to $61$\\,\\% at E.  %\\item FUV field was estimated from the ISO/LWS continuum.  %\\end{enumerate}   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0609/astro-ph0609454_arXiv.txt": {
        "abstract": "{Since the discovery, with the EINSTEIN satellite, of strong X-ray emission associated with HD~93162 ($\\equiv$WR~25), this object has been predicted to be a colliding-wind binary system. However, radial-velocity variations that would prove the suspected binary nature have yet to be found. } {We spectroscopically monitored this object to investigate its possible variability to address this discordance. } {We compiled the largest available radial-velocity data set for this star to look for variations that might be due to binary motion. We derived radial velocities from spectroscopic data acquired mainly between 1994 and 2006, and searched these radial velocities for periodicities using different numerical methods. } {For the first time, periodic radial-velocity variations are detected. Our analysis definitively shows that the Wolf-Rayet star WR~25 is an eccentric binary system with a probable period of about 208 days. } {} ",
        "introduction": "Massive stars of spectral type O and Wolf-Rayet (WR) are important objects that play a crucial role in the dynamic and chemical evolution of galaxies. They are the major source of ionizing and UV radiation and, through their huge mass-loss rates, they have a strong mechanical impact on their surroundings. Despite their importance, our knowledge of these objects and of their evolution is still fragmentary. The parameters that predominantly determine the evolution of a massive star are its mass (as for any star in the HR diagram) and its mass-loss rate, although rotation could also have an important impact \\citep{mem05}. In this context, massive O+O and WR+O binaries are key objects because their binary nature allows us to determine minimum masses from the radial velocity orbital solution.  It is now well established that strong winds of massive stars in binary systems collide \\citep{1992ApJ...386..265S} generating hard X-ray emission. For O stars, this comes in addition to a softer component that is intrinsic to individual objects \\citep{brg97, san06}. The existence of intrinsic emission for WR stars is still under investigation \\citep{osk03, gos05}. As for colliding-wind binaries, the observed characteristics of the collision region may contribute to constrain the mass-loss rates and, eventually, the inclination of those systems (by studying the X-ray emission modulation during the orbital cycle). Although colliding-wind binaries are interesting objects, only a handful of them have been studied in detail, which is particularly true for WR+O systems.  The star \\object{HD~93162} \\citep[WR~25 in][]{huc01} is a bright ($V$~=~8.1) galactic Wolf-Rayet located in the Carina Nebula region. Its binary nature has been a matter of debate for many years. It has been classified as WN7 + O7 by \\citet{smi68} because the spectrum appears as a superposition of WN-type emission lines and of absorption lines corresponding to an O-type star. The morphology of the blue optical spectrum of WR~25 has been further discussed by \\citet{wal74}. He classified this star as WN6-A \\citep[see also][]{wal00} and considered that there was not sufficient information to confirm its binary nature, since absorption lines intrinsic to a WN star had been observed in the neighbouring binary \\object{HD~92740}$\\equiv$WR~22 \\citep{nie73}. The massive binary status of WR~22 with absorption and emission lines moving in phase has been confirmed by several authors \\citep{nie73,1979ApJ...228..206C,rau96,1999A&A...347..127S}. In later classification systems of WR stars \\citep[e.g.][]{huc81,smi96}, the presence of absorption lines has been noted in different ways, e.g.\\ ''+abs'' or adding an ''a'', indicating that absorption lines of unknown origin are observed in the otherwise emission-line spectrum.  Radial velocity studies of WR-type spectra with absorption lines are thus needed to shed light on the origin of these lines. In the case of WR~25, previously reported radial velocity studies have not revealed any orbital motion. From the study of 15 photographic spectra obtained during consecutive nights, \\citet{1978A&A....68...41M} concluded that WR~25 is probably a single star. \\citet{1979ApJ...228..206C} reached a similar conclusion although they noticed a larger scatter in the radial velocities of WR~25 than that observed for other stars in their study (such as \\object{HD 93131}$\\equiv$WR~24 and WR~22 besides its orbital motion).  Using EINSTEIN observations, \\citet{1982ApJ...256..530S} found that WR~25 has an abnormally high X-ray flux and, later, \\citet{1991IAUS..143..105P} suggested that this very high X-ray luminosity might be caused by colliding winds in a binary system. More recently, \\citet{raa03} analyzed {\\it{XMM-Newton}} observations and provided further evidence that the X-ray emission could be indicative of a colliding-wind binary (CWB) although its apparent stability has been, on the contrary, taken as an argument against binarity. \\citet{pol06} reported significant variability as detected with {\\it{XMM-Newton}} and discussed the possibility of periodic variations in the X-ray flux.  \\citet{1992ApJ...386..288D} found polarimetric variability in WR~25, as well as a wavelength dependence of its polarization angle. One of the two proposed explanations was that this variability may arise in a long-period binary system. The alternative explanation attributed particular properties to the intervening interstellar medium. In his recent catalog, \\citet{huc01} classified WR~25 as WN6h+O4f considering that the WR emission lines showed evidence of being diluted, although no radial-velocity variations had ever been detected.  To investigate the binary status of WR~25, the massive star research groups of Li\\`ege and La Plata independently collected high-resolution spectra of this star over the past 10 years. The acquisition of high-resolution spectra of WR stars is rather unusual due to the difficulties presented by the comparatively broad emission lines typical of these stars. However, the radial velocities (RVs) measured for each data set showed variations larger than the expected errors. No single periodicity could be found until the data obtained by the two groups were combined. In the following, we show how the analysis of the combined spectroscopic data set enabled us to reveal the binary nature of WR~25 and to derive a preliminary orbital solution for this system. ",
        "conclusions": "We have found that the RVs of the N\\,{\\sc iv} $\\lambda$4058 emission line in the spectrum of WR~25 show variations with a period of about $P$=208 days. These variations indicate that the WN star WR~25 is an eccentric binary system.  The He\\,{\\sc i} $\\lambda$5876 and $\\lambda$4471 absorption lines seem to show an anti-phase motion and should thus partly belong to the companion which is supposed to be of the OB type.  The companion seems to be fainter than the WR primary, indicating that WR~25 is perhaps another example of a young binary system where the WR component is the most massive star. The prototype of this class is WR~22 \\citep[WN7h+O9III-V,][]{huc01} which turned out to be a massive binary system  (55 M$_\\odot$ + 21 M$_\\odot$) in an 80-day period orbit \\citep{rau96, 1999A&A...347..127S}. The primary of WR~22 is most probably an example of the long-searched for case of a core H burning star resembling a WR star due to its luminosity and mass. The star is on its way to becoming a hydrogen-free WR-type star. These objects are extremely rare and thus interesting to study in detail.  WR~25 is one of the brightest WN stars in the X-ray domain, too bright to be explained by intrinsic X-ray emission. It thus was customarily classified as a putative colliding-wind binary. Fig.~\\ref{RVs} is the first proof that WR~25 is a binary as suspected from its X-ray emission. On the basis of the existing {\\it{XMM-Newton}} data, \\citet{pol06} found that the X-ray emission of WR~25 is variable although the period cannot be derived independently from the X-ray data alone. WR~25 was brighter in the X-ray domain on JD\\,2\\,452\\,842.6 which according to our ephemeris corresponds to a phase of 0.96-0.98.  The star thus seems to be brighter at periastron as predicted for colliding-wind binaries in the adiabatic regime \\citep[see e.g.][]{1992ApJ...386..265S}. Although this should be further investigated, WR~25  appears to be an example of a colliding-wind binary system.  We are currently planning to monitor the optical spectrum of WR~25 around the expected time of minimum radial velocity of the WR emission lines in order to derive a full definitive orbital solution. High-resolution high S/N ratio data are essential. These spectra will help to improve the orbital solution, and to detect the absorption lines of the companion, which for this particular phase domain should be located to the red of the WN emission line, thus avoiding confusion with the absorption components of  P-Cygni profiles from the primary."
    },
    "0609/astro-ph0609348_arXiv.txt": {
        "abstract": "Cosmic microwave background (CMB) temperature anisotropies follow a Gaussian statistical distribution in the standard inflationary model, but there are non-Gaussian contributions due to astrophysical foregrounds. The detection of the non-Gaussian signal due to extragalactic point sources and its distinction from the possible intrinsic non-Gaussianity is an issue of great importance in CMB analysis. The Mexican Hat Wavelet Family (MHWF), which has been proved very useful for the detection of extragalactic point sources, is applied here to the study of non-Gaussianity due to point sources in CMB maps. We carry out simulations of CMB maps with the characteristics of the forthcoming Planck mission at 70 and 100 GHz and filter them with the MHWF. By comparing the skewness and kurtosis of simulated maps with and without point sources, we are able to detect clearly the non-Gaussian signal due to point sources for flux limits as low as 0.4 Jy (70 GHz) and 0.3 Jy (100 GHz). The second and third members of the MHWF perform better in this respect than the Mexican Hat Wavelet (MHW1) and much better than the Daubechies 4 wavelet. We have also estimated the third order, $K_3$, and fourth order, $K_4$, cumulants produced by point sources at these Planck channels by means of a fit with the MHWF. The average relative errors with respect to the real values are below $12\\%$ for fluxes down to $0.6$ Jy (70 GHz) and $0.4$ Jy (100 GHz). The values of these cumulants allow us to distinguish between different source counts models. From the estimated cumulants and assuming a power law for the source number counts we are able to obtain the coefficients $\\alpha$ and $A$ of the differential number counts, $\\alpha=2.19\\pm 0.46$ ($\\alpha=2.26\\pm 0.19$), $A=24.3\\pm 4.3$ ($A=21.0\\pm 4.2$) for 70 (100) GHz, assuming a flux limit of 1 Jy. These results are consistent with the values obtained from simulations containing only point sources, which are: $\\alpha=2.32\\pm 0.06 $ ($\\alpha=2.35\\pm 0.16$), $A=22.1\\pm 1.5$ ($A=20.0\\pm 3.7$) for 70 (100) GHz. ",
        "introduction": "\\footnotetext{E-mail: argueso@uniovi.es}   According to the standard cosmological model, CMB temperature anisotropies are caused by the inhomogeneities in the distribution of matter and radiation present at the decoupling time. These inhomogeneities are the product of quantum fluctuations amplified in the inflationary era and follow a Gaussian statistical distribution. Therefore, the observed CMB anisotropies are a realization of a homogeneous and isotropic Gaussian random field on the sphere. Nevertheless,  some degree of non-Gaussianity can be introduced, for instance, by non standard inflation \\citep{bar04} or by topological defects \\citep{tur90,dur99}. The analysis of the first and three-year Wilkinson Microwave Anisotropy Probe (WMAP) data by the WMAP team shows that CMB temperature fluctuations are consistent with a Gaussian distribution \\citep{kom03,spe06} in agreement with the standard inflation paradigm. However, some studies have detected non-Gaussian features in the WMAP data, for instance in \\cite{vie04} and \\cite{cru05,cru06} a cold non-Gaussian spot of unknown origin was detected and studied by using the Spherical Mexican Hat Wavelet. Non-Gaussian signatures have also been found by using different methods: phase correlations \\citep{chi03,col04}, local curvature methods \\citep{han04,cab05}, multivariate analysis \\citep{din05}, Minkowski functionals \\citep{par04,eri04} and higher criticism \\citep{cay05}.  On the other hand, astrophysical foregrounds contaminate the cosmological signal and also give rise to a non-Gaussian contribution. In particular, the emission due to point sources (radio and infrared galaxies are seen as point-like objects due to their very small projected angular size in comparison with the typical experimental resolutions) is clearly non-Gaussian \\citep{kom01,kom03,arg03,pie03,her04,bar06}. In the last few years, new high frequency ($\\nu\\geq 10$ GHz) surveys of radio sources have been carried out \\citep{wal03,ric04}. All the recently published data on counts and on the emission properties of extragalactic radio sources have guided the development of new cosmological evolution models able to improve on previous ones \\citep{dun90,tof98} in predicting number counts and source statistics at high radio frequencies \\citep{dez05}. We will use this new model to carry out point source simulations.  It is very important to assess the degree of non-Gaussianity generated by point sources in order to distinguish it from the non-Gaussian signal coming from other origins. The ESA Planck satellite \\citep{man98,pug98}, which will be launched in 2008, will analyze the CMB anisotropies with unprecedented accuracy and resolution and will make possible a very precise non-Gaussianity study, for which the use of wavelets can be very suitable.  The Mexican Hat Wavelet (MHW1) has been very useful in detecting point sources \\citep{cay00,vie01a,vie01b} in CMB maps, by using the signal amplification going from real space to wavelet space. The Mexican Hat Wavelet Family (MHWF) was introduced by \\cite{gon06} as an extension of the standard MHW1, which produced higher amplification (see below) and then a better performance at point source detection. These new wavelets were also used for point source detection in Planck maps in \\cite{lop06}. The MHWF is a family of isotropic 2D wavelets obtained by applying the Laplacian operator iteratively to the 2D Gaussian.  The expression of these new wavelets in real space is \\begin{equation} \\label{cc} \\psi_{n}(x)=\\frac{(-1)^n}{ 2^{n}n!}\\triangle^{n}{\\varphi(x)} \\end{equation} \\noindent where $x$ is the radial coordinate and $\\varphi$ is the 2D Gaussian \\begin{equation} \\label{cc} \\varphi(x)=\\displaystyle\\frac{e^{-x^2/2}}{2\\pi} \\end{equation}  \\noindent   Note that $\\psi_{1}(x)$ is the standard MHW1, we call the other corresponding members of the family MHWn. Any member of the family can be written very easily in Fourier space as \\begin{equation} \\hat\\psi_{n}(k)=\\frac{k^{2n}e^{-\\frac{k^{2}}{2}}}{2^{n}n!} \\end{equation}  \\noindent \\cite{gon06} considered a point source convolved with a Gaussian instrumental beam of dispersion $\\gamma$ \\begin{equation} \\gamma=\\frac {FWHM}{\\sqrt{8\\log(2)}} \\end{equation} \\noindent where FWHM denotes the Full Width Half Maximum of the instrumental beam.  The point source was embedded in a CMB map, including noise and all the relevant foregrounds. Then, they filtered the map with the first members of the MHWF at different scales R. \\begin{equation} \\Psi_{n}(\\vec{x}; \\vec{b}, R) \\equiv \\frac{1}{R^{2}}\\psi_{n} \\Big(\\frac{|\\vec{x}-\\vec{b}|}{R}\\Big), \\end{equation}  \\noindent The amplification, $\\lambda$, is defined as the ratio between the intensity of the point source $I_w$ and the rms deviation $\\sigma_w$ in the wavelet filtered image divided by this ratio in the original image filtered with a Gaussian-shaped instrumental beam.  \\begin{equation} \\lambda=\\frac{I_{w}/\\sigma_w}{I_{g}/\\sigma_{g}} \\end{equation}  \\noindent $I_w$ can be calculated for the different members of the MHWF as  \\begin{equation} \\label{cc} I_{w}=\\displaystyle\\frac{I_g \\beta^{2n}}{(1+\\beta^{2})^{n+1}} \\end{equation}  \\noindent with $\\beta=R/\\gamma$.   It was proved that when filtering CMB maps with the first three members of the MHWF at an optimal scale, \\textit{R}, the MHW2 and MHW3 produced a higher amplification than the standard MHW1 and were hence more effective at source detection. Other higher n members of the MHWF did not yield higher amplifications.  After removing the detected point sources, there could still be a remaining non-Gaussianity due to undetected point sources. It is of great interest to characterize this non-Gaussianity in order to distinguish it from the intrinsic one.  The detection of the non-Gaussian signal due to point sources poses a similar problem to that of point source detection, these new wavelets will enhance the signal (non-Gaussian) due to point sources, thus making it possible to detect this kind of non-Gaussianity more effectively. Since the skewness and kurtosis of a Gaussian distribution are zero and three respectively, any significant deviation from these values will mean the presence of a non-Gaussian signal, so that the skewness and the kurtosis of the wavelet filtered maps can be used as estimators of the non-Gaussianity due to point sources.  Filtering CMB maps with the MHWF at optimal scales amplifies point sources and can also be used to determine the third and fourth order cumulants produced by the point sources. From these cumulants, we can also determine the parameters $A$ and $\\alpha$, if we assume that the differential number counts can be expressed as a power law $dn/ds=AS^{-\\alpha}$. In \\cite{pie03} a chi-squared fit to a power law was explored considering the moments of the CMB maps. We have considered in this paper a chi-squared fit of the third order and fourth order cumulants of the wavelet filtered maps combining several scales including the optimal ones.  The plan of the paper is as follows: in \\S 2, we calculate the cumulants of a wavelet filtered map from those of the original map, this allows us to relate the skewness and kurtosis of the filtered maps to the amplification. In \\S 3, we carry out realistic CMB simulations with and without point sources, considering the characteristics of the Planck experiment at 70 and 100 GHz and perform the non-Gaussianity detection by filtering with the MHWF. In \\S 4, we obtain the third and fourth order cumulants produced by point sources by means of a fit of the cumulants of wavelet filtered CMB maps and we also determine the parameters $\\alpha$ and $A$. Finally in \\S 5, we draw the main conclusions of the paper. ",
        "conclusions": "We have considered a natural generalization of the MHW1 on the plane, $R^2$, based on the iterative application of the Laplacian operator to the MHW1. We have called this group of wavelets the Mexican Hat Wavelet Family (MHWF) \\citep{gon06,lop06}. We have applied the MHWF to the study of the non-Gaussianity produced in CMB maps by extragalactic point sources, a very important issue for the separation and analysis of different non-Gaussian components in the CMB.  We have first derived a formula, eq. (14), which relates the cumulants of 2D point source maps filtered with the MHWF with the cumulants of the original point source maps. From this formula, we have obtained several equations, eqs. (18-23), which relate the contribution of point sources to the skewness and kurtosis of wavelet (MHWF) filtered 2D CMB maps to their contribution to the skewness and kurtosis of beam filtered CMB maps.   According to these formulas, the skewness of a wavelet filtered map is proportional to $\\lambda^{3}$, with $\\lambda$ the amplification, and the kurtosis is proportional to $\\lambda^{4}$. This means that the higher the amplification the higher the non-Gaussian signal will be. Therefore, the MHW2 and MHW3 at their optimal scales can be very suitable filters to detect the non-Gaussian signal due to point sources, since they produce a higher amplification than the MHW1 \\citep{gon06,lop06}.  We have then considered Gaussian simulations of CMB plus anisotropic noise and compared them with simulations including all the relevant components: CMB, anisotropic noise, point sources and Galactic foregrounds. The simulations are 2D maps of $12.8\\times12.8$ square degrees with the characteristics of the Planck experiment at 70 GHz and 100 GHz. We perform a Gaussianity test based on the skewness and kurtosis of the maps filtered with the MHW1, MHW2 and MHW3 at the optimal scales (those which produce the highest amplifications), we compare 1000 simulations of CMB plus noise with 1000 simulations including point sources up to different flux limits from 0.1 Jy to 1 Jy.   The power of the test is higher for the MHW2 and the MHW3 than for the standard MHW1, the kurtosis test is more powerful than the skewness test. We could detect the non-Gaussianity due to point sources down to a flux limit of 0.4 Jy (70 GHz) and 0.3 Jy (100 GHz), the powers are higher for 100 GHz than for 70 GHz. The results can be seen in Figure 2 and Table 1. We would like to point out that according to \\cite{lop06} point sources could be detected in these channels, assuming a $5\\%$ of spurious detections, down to $0.47$ Jy (70 GHz)and $0.25$ Jy (100 GHz). This means that, using the MHWF, we can see a non-Gaussian signal due to weak point sources, which are in the limit of detection of current source extraction techniques. Therefore, it is very important to characterise the non-Gaussianity due to unsubtracted point sources, in order to avoid its possible misidentification as intrinsic non-Gaussianity.   The influence of Galactic foregrounds is negligible, since when we consider simulations without point sources, but keeping the Galactic foregrounds, the powers go down below $6\\%$. We have also studied the performance of the Daubechies 4 wavelet, obtaining powers below $50\\%$ even for flux limits as high as 1 Jy. The MHWF performs much better than this wavelet (see Figure 2), since it is adapted to deal with point sources.  This method could be used for all the frequencies of the Planck experiment and could also be extended to spherical CMB images, as the ones provided by the Planck experiment.  We have also tried to estimate the third order cumulant $K_3$ and the fourth order cumulant $K_4$, produced by point sources at the 70 and 100 GHz Planck channels. We filter $12.8\\times 12.8$ square degrees CMB maps with the characteristics of the Planck experiment with the MHW1, MHW2 and MHW3 at three scales for each wavelet (including the optimal ones). The maps include CMB, anisotropic noise, point sources and other foregrounds. We calculate the third and fourth order cumulants of these wavelet filtered maps and write chi-squared functions, eqs. (24) and (25), which allow us to estimate the cumulants of the original point source maps; we have used eq. (14) to relate the cumulants of the wavelet filtered maps with those of the original maps. We have compared our estimated values with those, $K_3^{ps}$ and $K_4^{ps}$, obtained directly from simulations of the same type but including only unfiltered point sources without any other component. The results can be seen in Table 3. The average error (from ten half skies) is below $12\\%$ down to a flux limit of $0.6$ Jy (70 GHz) and $0.4$ Jy (100 GHz). The determination of these cumulants would allow us to distinguish between different source counts models (see Tables 3 and 4).  Finally, we have also studied the fit of the differential number counts at 70 and 100 GHz to a power law $dn/dS=A S^{-\\alpha}$. From the cumulants previously estimated, we calculate by using eqs. (32-34) the amplitude $A$ and slope $\\alpha$ of the differential number counts. The results at 70 GHz for a 1 Jy flux limit are: $\\alpha=2.19\\pm 0.43$, $A=24.3\\pm 4.3$. When we calculate these magnitudes from the cumulants $K_3^{ps}$ and $K_4^{ps}$, computed directly from the point source only simulations, we obtain $\\alpha^{ps}=2.32\\pm 0.06$, $A^{ps}=22.1\\pm 1.5$. These are the average and rms deviation from ten half skies. For this flux limit, the average relative error of $\\alpha$ ($A$) with respect to $\\alpha^{ps}$ ($A^{ps}$) is below $10\\%$ ($15\\%$). The dispersion in the values obtained with the MHWF fit is high, though the average is consistent with the value obtained directly from point sources down to a flux limit of 0.7 Jy, as can be seen in Table 5.  When we consider the 100 GHz channel, the values obtained from the MHWF fit for a flux limit of 1 Jy are: $\\alpha=2.26\\pm 0.19$, $A=21.0\\pm 4.2$, the results from the point source simulations are: $\\alpha^{ps}=2.35\\pm 0.16$, $A^{ps}=20.0\\pm 3.7$. The fit is better for this channel than for the 70 GHz one, having a lower dispersion. For this flux limit, the average relative error of $\\alpha$ ($A$) with respect to $\\alpha^{ps}$ ($A^{ps}$) is below $5\\%$ ($6\\%$). The results are consistent down to a flux limit of 0.7 Jy. The detailed values of the parameters $\\alpha$ and $A$ can be seen in Table 5.  In Figure 3 we plot the differential source counts of the De Zotti model, the fit to a power law by using the MHWF and the fit obtained from the point source only simulations, in both cases with a 1 Jy flux limit. As we can see in this figure, both fits are consistent with the input counts in the range between 0.1 Jy and 1 Jy."
    },
    "0609/astro-ph0609662_arXiv.txt": {
        "abstract": "%  We have measured the 3.6 $\\mu$m luminosity and mass evolution of about 1000 galaxies in 32 clusters at $0.2<z<1.25$ with a special attention to methodological issues, as emphasized in this proceeding contribution. We find that the luminosity of our galaxies evolves as an old and passively evolving population formed at high redshift without any need for additional redshift-dependent evolution. Models with a prolonged stellar mass growth are rejected by the data with high confidence.  The data also reject models in which the age of the stars is the same at all redshifts.  Similarly, the characteristic stellar mass evolves, in the last two thirds of the universe age, as expected for a stellar population formed at high redshift. Together with the old age of stellar populations derived from fundamental plane studies, our data seems to suggest that massive early-type cluster galaxies have been completely assembled at high redshift, and not only that their stars are old. The quality of the data allows us to derive the LF and mass evolution homogeneously over the whole redshift range, using a single estimator. The Schechter function describes the galaxy luminosity function well. The characteristic luminosity at $z=0.5$ is found to be 16.30 mag, with an uncertainty of 10 per cent. ",
        "introduction": "%  The luminosity function (LF) is the basic statistic  used to understand galaxy properties, giving  the relative frequency of galaxies of a given luminosity in a given volume. Since starlight at 3.6 $\\mu$m very nearly follows the Rayleigh-Jeans limit of blackbody emission for $T > 2000$ K, the colors of both early- and late-type stars are similar.  There is virtually no dust extinction at this wavelength either, since any standard extinction law predicts only a few percent of the extinction of optical wavelengths. The 3.6 $\\mu$m light therefore traces the stellar mass distribution free of dust obscuration effects (Pahre et al. 2004). Thus, a useful approach to understanding how galaxies form is to track their growing stellar mass, measured through the evolution of the 3.6 $\\mu$m LF.  Several studies found LF results consistent with the behavior of a simple, passive luminosity evolution model in which galaxies form all their stars at high redshift and thereafter passively evolve. However, seldom these studies address which models data exclude. What it actually interesting is whether some plausible scenario can be rejected by the data, because sometime a (wrong) scenario cannot be rejected simply because of the low information content of the data in hand. Furthermore, previous works rarely address the topic of sample representativity. Specifically, clusters in the sample are typical or atypical clusters? All types of clusters are included in the studied sample? Some type of cluster is over- or under- represented in the sample? The above questions are closely related to the following one: can a result found to hold on a sample of unknown representativity (the studied sample of clusters) be generalized to the parent distribution (i.e. to clusters)? Results obtained on uncontrolled sample rarely hold in general (and when they hold, they hold by good chance).  Throughout this paper we assume $\\Omega_M=0.3$, $\\Omega_\\Lambda=0.7$ and $H_0=70$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": ""
    },
    "0609/astro-ph0609381_arXiv.txt": {
        "abstract": "Recent seismic observations coming from acoustic and gravity modes clearly show that the solar standard model has reached its limits and can no longer be used to interpret satisfactorily seismic observations. In this paper, we present a review of the non-standard processes that may be added to the solar models in order to improve our understanding of the helioseismic data. We also present some results obtained when applying ``non-standard'' stellar evolution to the modelling of the Sun. ",
        "introduction": "Dynamical processes are thought to occur in stars all across the Hertzsprung-Russell diagram. The development of advanced and dedicated observational facilities have greatly improved the constraints that we can derive from observations. In particular, helioseismology, and more recently asteroseismology, have given a new insight to our understanding of solar and stellar physics in both the convective and radiative regions \\citep{Couvidat03}; \\citep{TuC04}; \\citep{JCD05}. Spectroscopic and (spectro-)polarimetric data collected for stars of different masses, metallicities and evolutionary stages, including the Sun, also demonstrate the actual presence and action of dynamical processes such as rotation and magnetic fields from direct measurements as well as from indirect evidence provided by the abundance anomalies.\\\\ However, while processes such as rotation, magnetic fields, turbulence and internal waves appear to play a crucial role in many aspects of stellar evolution, they are not well understood and their effects are seldom accounted for in stellar and solar evolution models. At the dawn of a new era for stellar physics, these processes can no longer be neglected.\\\\ The dynamical processes occurring in stellar interiors may have a direct impact on the fundamental parameters of stars (modification of the basic equations of stellar structure by rotation or magnetic fields, see for instance \\citep{KT70}). They also lead to transport of matter and/or of angular momentum. Last but not least, they operate on very different scales at which their action has to be understood if one wants to properly take them into account in the stellar modelling. Two complementary approaches have emerged:\\\\ (1) the multi-dimensional simulation of the physical process in restricted areas on dynamical time-scales \\citep{BT02};\\citep{BBT05}\\\\ (2) the incorporation of non-standard\\footnote{By non-standard, we mean all processes different from convection.} physical processes in stellar evolution codes in order to follow their integrated effects over evolutionary timescales.\\\\ In the following, we will adopt the second point of view and evaluate importance of various physical processes in the radiative stellar zones. The first approach has been essentially applied to the convective zones that evolve on dynamical scales, but few studies of radiative zones also exist (see e.g. \\citep{T05}).. ",
        "conclusions": ""
    },
    "0609/astro-ph0609038_arXiv.txt": {
        "abstract": "% Using data obtained as part of the {\\it Spitzer} Infrared Nearby Galaxies Survey \\citep[SINGS;][]{rk03} and WSRT-SINGS radio continuum survey \\citep{rb06}, we study the effects of star-formation activity on the far-infrared (FIR)--radio correlation {\\it within} galaxies. This is done by testing a phenomenological model for the correlation, which describes the radio image as a smeared version of the FIR image. We find that this description works particularly well for galaxies with higher infrared surface brightnesses, yielding best-fit smoothing scale-lengths of a few hundred parsecs, substantially shorter than those for lower surface brightness galaxies. We interpret this result to suggest that galaxies with higher disk averaged star formation rates have had a recent episode of enhanced star formation and are characterized by a higher fraction of young cosmic-ray (CR) electrons compared to galaxies with lower star formation activity. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609512_arXiv.txt": {
        "abstract": "We present far-UV to optical analyses of four hydrogen deficient central stars of planetary nebulae : BD +30 3639, NGC 40, NGC 5315 and NGC 6905.  Using the radiative transfer code CMFGEN, we determined new physical parameters and chemical abundances for these stars. The results were analyzed in the context of the [WR] $\\rightarrow$ PG 1159 evolution via the transformed radius-temperature ($R_T \\times T_*$) and HR diagrams. We found that the use of clumping increases $R_T$ by up to $\\sim 0.5$ dex relative to previous results, whilst our homogeneous models showed systematic changes of $\\sim 0.1-0.2$ dex. NGC 5315 showed itself as an odd object among the previously analyzed central stars. Its temperature ($\\sim 76kK$) is considerably lower than other early-type [WR] stars ($\\sim 120-150kK$). From our models for NGC 5315 and NGC 6905, it is unclear if early-type [WR] stars have smaller C/He mass ratios than other spectral classes, as claimed in the literature. In particular, the ratio found for NGC 6905 ($\\sim 0.8$) is in rough agreement with evolutionary calculations, and with values derived for [WCL] stars. We analyzed FUSE spectra of these stars for the first time, and identified phosphorus in the spectra of BD +30 3639, NGC 40 and NGC 5315 through the doublet transition \\ion{P}{5} $\\lambda\\lambda  1118,  1128$ ($3p\\,\\,^{2}P^{o}-3s\\,\\,^{2}S$). The iron, silicon, phosphorus, sulfur and neon abundances were analyzed in the context of the nucleosynthesis occurring in previous evolutionary phases. We found evidence for an iron deficiency in BD +30 3639 and NGC 5315, and from fits to the \\ion{Si}{4} lines we determined a solar silicon abundance for BD +30 3639 and NGC 40. For phosphorus, an oversolar abundance in the NGC 5315 model was preferred, while in the other stars a solar phosphorus abundance cannot be discarded. Regarding sulfur, we estimated upper limits for its abundance, since no conspicuous lines can be seen in the observed spectra. We found that neon is overabundant in BD +30 3639. In the other stars, neon is weak or undetectable and upper limits for its abundance were estimated. Our results are in agreement with theoretical predictions and show the usefulness of [WR] stars as testbeds for nucleosynthesis calculations in the AGB and post-AGB phases. ",
        "introduction": "Hydrogen deficient central stars of planetary nebulae (CSPN) present a challenge to stellar evolution theory: it is still unclear how a star removes its hydrogen when going from the AGB to the planetary nebula stage. One of the first theories trying to explain this phenomenon was proposed by Iben et al. (1983). In their {\\it born-again} scenario, a very late thermal pulse (VLTP) occurs in the white dwarf cooling track of the HR diagram. As a consequence, all the remaining hydrogen is burned and the star is driven to red giant dimensions, near the AGB, for a second time. Two other thermal pulses could also potentially give rise to hydrogen deficient CSPN. One is the AFTP (AGB final thermal pulse), occurring at the end of the AGB phase, and the other is the LTP (late thermal pulse), which occurs when the star is already on the post-AGB track. It seems that all three scenarios are necessary in order to explain the wide range of hydrogen deficient CSPN properties (e.g., hydrogen abundance and nebular diameters), and that there is no unique mechanism or single channel of evolution (Bl\\\"ocker 2001; Herwig 2001).  Binarity could also provide an avenue for the creation of hydrogen deficient CSPN. De Marco \\& Soker (2002), for example, raised the interesting possibility of a companion engulfed by an AGB star. In this way, extra mixing would be induced and a hydrogen deficiency would be achieved. However, the scenario is still speculative and lacks quantitative calculations.  The evolutionary sequence of hydrogen deficient CSPN is also uncertain. Generally these objects are divided into three classes : [WR], [WC]-PG 1159 and PG 1159 stars. [WR] stars show spectra very similar to Wolf-Rayet stars of Pop. I, with strong and broad emission lines. [WC]-PG 1159 stars show both absorption and weak emission lines in their optical spectra, while wind features are seen in the UV. Parthasarathy et al. (1998) claim that [WC]-PG 1159 and the {\\it weak emission line stars} ([WELS]), presented in the work of Tylenda et al. (1993), constitute the same class. PG 1159 stars show mainly absorption lines in their spectra.  Analysis of the spectra of the central stars or of their nebulae seems to indicate the following evolution : [WCL] $\\rightarrow$ [WCE] $\\rightarrow$ [WC]-PG 1159 = [WELS] $\\rightarrow$ PG 1159 (Werner \\& Heber 1991a,b; Zijlstra et al. 1994; Gorny \\& Tylenda 2000; Pe\\~na et al. 2001; Leuenhagen et al. 1996; Koesterke \\& Hamann 1997a,b; Leuenhagen \\& Hamann 1998). However, there are still fundamental questions to be addressed. A major problem is the C/He mass ratio obtained for early-type [WR] stars, which is $\\sim 0.35$ on average. This is lower than found in the other classes, and is thus in contradiction to the above evolutionary sequence. It also disagrees with evolutionary models which find this ratio to be $\\sim 1$ (Herwig 2001). A second problem comes from the work of Gorny \\& Tylenda (2000), who showed that some stars might not follow this evolution channel. Through an analysis of nebular parameters, the authors suggest that some objects that have experienced the VLTP could evolve directly to the PG 1159 stage without becoming [WC] stars. Another important problem is that the role of the [WELS] is not yet understood (Marcolino \\& de Ara\\'ujo 2003). Indeed,  Pe\\~na et al. (2003a;b) were able to show that nebular expansion velocities of [WC] stars are generally higher than for the [WELS], in contradiction to that expected from [WC] $\\rightarrow$ [WELS] evolution. Moreover, the presence of hydrogen is still questionable (De Marco \\& Barlow 2001) and can favor a specific thermal pulse model. While in the AFTP and LTP a small amount is found, in VLTP models the hydrogen is completely burnt.  So far, the majority of [WR] stars analyzed by means of non LTE expanding atmosphere models is due to the Potsdam group code (see Koesterke 2001 and references therein). Additional studies of three [WCL] stars were undertaken by De Marco \\& Crowther (1998, 1999) using an independent code (CMFGEN, Hillier \\& Miller 1998). All objects were studied without metal line-blanketing or clumping (exceptions are the recent work of Stasi\\'nska et al. 2004 and of Pe\\~na et al. 2004 regarding two stars on LMC, Herald \\& Bianchi 2004a,b, which included one [WC4] and two [WC]-PG 1159 stars and Crowther et al. 2006, who compared the [WC9] star BD +30 3639 with the Wolf-Rayet star HD 164270). These two effects have been shown to be very important when analyzing hot stars with winds (Hillier \\& Miller 1998, 1999; Gr\\\"afener et al. 2002), having an impact on important physical parameters (e.g., $T_{eff}$ and $\\dot{M}$).  With the aim of addressing some of the problems and uncertainties discussed above, we have derived physical parameters and chemical abundances of a small sample of H deficient CSPN (BD +30 3639, NGC 40, NGC 5315 and NGC 6905) using the CMFGEN code developed by Hillier \\& Miller (1998). Contrary to most of the previous works on CSPN, line-blanketing and clumping are taken into account in the models. We also rely on improved atomic data and fits to larger spectral intervals, including the FUSE region now available for some CSPN. Among the issues that we address are : the influence of line-blanketing and clumping on the evolutionary sequence; the peculiar values of the C/He mass ratio for the early-type [WR] stars and the iron abundance and its possible depletion as predicted by AGB and post-AGB evolutionary calculations.  This paper is organized as follows. The observational data is described in Section 2. In Section 3 we explain the code used (CMFGEN) and its main assumptions, while in Section 4 we present our analysis from the far-UV to the optical for each object. In Section 5 we discuss our results. Finally, in Section 6 we summarize the main points of our work. ",
        "conclusions": "\\subsection{C/He mass ratios of the early-type [WR] stars}  It is claimed that for early-type [WR] stars the C/He mass ratio ($\\beta_{C}/\\beta_{He}$) is considerably smaller than in other classes ([WCL], [WELS] and PG 1159). In Table 4 we show the average mass fractions of He, C and O for all H deficient central stars analyzed so far. The range of values measured by atmosphere models is shown for each element. This problem was first mentioned by Koesterke \\& Hamann (1997a,b) and contradicts the evolutionary sequence, because early-type [WR] stars are supposed to be descendents of the [WCL] stars. Furthermore, $\\beta_{C}/\\beta_{He}$ is expected to be about unity from stellar evolution models regardless which thermal pulse (AFTP, LTP or VLTP) is responsible for the hydrogen deficiency (Herwig 2001). It is still unclear if such a discrepancy is due to effects not taken into account in older models (e.g., line-blanketing).  The determination of $\\beta_{He}$ and $\\beta_{C}$ is usually based on the transitions \\ion{He}{2} $\\lambda 5412$ and \\ion{C}{4} $\\lambda 5471$. These lines are adopted since both are formed by recombination in a similar region of the stellar wind. In Fig. 12 we have a spectral interval comprising these two lines for the two early-type stars of our sample, NGC 5315 and NGC 6905, respectively. In order to investigate the low $\\beta_{C}/\\beta_{He}$ issue found in previous works, we plot two models along with the observations : one with $\\beta_{C}/\\beta_{He} \\sim 0.35$, which is the average value found for the early-type [WR] stars, and another with this ratio approaching unity. For NGC 5315, none of the models present a reasonable fit to these lines. We also tested lower values for $\\beta_{C}$ and $\\beta_{He}$ at the expense of an increase in $\\beta_{O}$. Although the fit showed a better quality, some oxygen lines increased considerably (e.g., \\ion{O}{3}\\ion{-}{5} $\\lambda 5595$). In the case of NGC 6905, a $\\beta_{C}/\\beta_{He}$ ratio of $\\sim 0.8$ is better than $\\sim 0.35$ (bottom panel). This lower value is essentially the same as that adopted by Koesterke \\& Hamann (1997). The ratio of $\\sim 0.8$ is more in line with values found in other classes (see Table 4) and with what is expected from evolutionary calculations. It is very important to note however, that the fits are not perfect. The determination of the continuum and a \\ion{He}{2} $\\lambda 5412$ nebular contamination (as is clear from the figure for NGC 6905) are factors that can complicate the analysis. At least from our results, it is not clear if early-type [WR] stars have smaller $\\beta_{C}/\\beta_{He}$ ratios. A larger sample of objects with better signal-noise spectra should be carefully analyzed. In doing so, we should keep in mind a commitment with other parts of the spectrum when fitting \\ion{He}{2} $\\lambda 5412$ and \\ion{C}{4} $\\lambda 5471$.  \\subsection{Silicon, phosphorus, sulfur and neon abundances}  The silicon features present in the spectra of [WR] stars deserve special attention. Stellar evolution models predict that silicon has essentially a solar mass-fraction in the H-He intershell matter in AGB stars. As a consequence, we expect to measure roughly a solar value, i.e., $\\beta_{Si} \\sim 0.07\\%$, when analyzing H-deficient CSPN (see Werner \\& Herwig 2006 and references therein). However, Leuenhagen \\& Hamann (1998) derived $\\beta_{Si} = 0.5 - 3\\%$ in a couple of [WCL] stars. These mass fractions correspond to 7 to 40 times the solar value, indicating a huge discrepancy with theory. Motivated by this, we used silicon lines in the spectra of BD +30 3639 and NGC 40 in order to check if the same high over solar abundances would be found. It turned out that the transitions \\ion{Si}{4} $\\lambda 1722$ in NGC 40 and  \\ion{Si}{4} $\\lambda 1394$, $\\lambda 1403$ and $\\lambda 4089$ in BD +30 3639 were well reproduced using a solar silicon abundance. Despite an uncertainty of approximately a factor of two in our determination, our results are in agreement with stellar evolution models. As reported by Werner \\& Herwig (2006), near solar values were also found in a few PG 1159 stars.  The atomic model for silicon used in the work of Leuenhagen \\& Hamann (1998) was relatively simple (see their table 1). Moreover, neither clumping nor metal line-blanketing was taken into account in their study. Therefore, it would be interesting to re-analyze the stars in their sample in order to check if more sophisticated models could influence the silicon abundance determination. This could confirm or nullify the conflict with stellar evolution models.   Thanks to the FUSE spectra available for all stars in our sample, we could identify phosphorus in the spectrum of BD +30 3639, NGC 40 and NGC 5315. The $1100-1130$\\AA\\, interval could only be fitted after the inclusion of \\ion{P}{5}, which gave origin to the transitions $\\lambda 1118$ and $\\lambda 1128$ ($3p\\,\\,^{2}P^{o}-3s\\,\\,^{2}S$). From nucleosynthesis calculations, phosphorus is expected to be observed with $4-25$ times the solar abundance, depending on the treatment of convective mixing (Werner \\& Herwig 2006). In fact, the NGC 5315 spectrum could be better matched using abundances in this range. On the other hand, BD +30 3639 and NGC 40 models were too insensitive to changes in the phosphorus abundance for a reliable abundance to be derived. They allow not only over solar values in the expected range but also a solar abundance. Hence, we can say that the result for NGC 5315 supports nucleosynthesis predictions, while for the other two stars the error in the abundance determination is too large to draw any conclusions.   Evolutionary calculations predict that sulfur should be observed with an abundance of $0.6-0.9$ of the solar value (Werner \\& Herwig 2006). We searched for sulfur lines in the spectra of each star of our sample. In all objects they are very weak or undetectable. From our models, we found that the following lines can be used as diagnostics for the abundance determination : \\ion{S}{5} $\\lambda 1502$, \\ion{S}{5} $\\lambda 2654$, and \\ion{S}{4} $\\lambda 3107$. For BD +30 3639 we estimate an upper limit of $\\beta_S \\lesssim \\beta_{S_{\\odot}}$, in agreement with what is expected from stellar evolution models. For NGC 40 and NGC 5315, a higher limit was adopted, namely, $\\beta_S \\lesssim 2\\beta_{S_{\\odot}}$. We could not estimate the sulfur abundance in NGC 6905. Given its high temperature, \\ion{S}{6} is fully ionized throughout its stellar wind and we currently lack \\ion{S}{7} atomic data.  Neon is expected to be overabundant in hydrogen deficient CSPN. According to Werner \\& Herwig (2006), its mass frass fraction should be $\\sim 2\\%$, i.e., about 10 times the solar value, although slightly lower or higher values are predicted depending on the initial mass of the star. In our models, neon diagnostic lines were found in the 2000-3000\\AA\\, interval. For NGC 40 and NGC 5315, the IUE spectra in this region have low resolution, limiting the accuracy of the analysis. No conspicuous neon transitions are seen. From the lines \\ion{Ne}{4} $\\lambda 2364$, \\ion{Ne}{3} $\\lambda 2553$ and \\ion{Ne}{3} $\\lambda 2678$, we estimate a rough upper limit of $\\sim 8$ times the solar neon value. For NGC 6905, by using the theoretical transitions \\ion{Ne}{6} $\\lambda 2641$ and \\ion{Ne}{6} $\\lambda 2687$ we estimate an upper limit of 5 times the solar neon value. Higher abundances make these lines much stronger than observed.  In contrast to the other stars, BD +30 3639 present neon lines in the near UV. Although the IUE spectrum in the 2000-3000\\AA\\, region does not have a good signal-noise ratio, we can identify \\ion{Ne}{3} $\\lambda 2553$ and \\ion{Ne}{3} $\\lambda 2678$ (see Fig. 3). With a solar neon abundance, this last transition is practically absent in the models. We could only have a reasonable agreement with the observed \\ion{Ne}{3} $\\lambda 2678$ line by using a neon mass fraction of $\\sim 2\\%$. Therefore, BD +30 3639 is the only case in our sample that definitely support the prediction of stellar evolution models.  \\subsection{Evolutionary sequence} \\label{sect_ev_seq}  A comparison among our results and previous works that did not include line-blanketing and clumping can be seen in Fig. 13. In this diagram, the suggested evolution is to decrease the transformed radius while the temperature increases when going from the [WCL] to the early-type [WR] stars. This would be the result of a shrinking radius and higher mass-loss. As we get near the white dwarf phase, i.e.,  shortly afterwards the [WELS] and PG 1159 stage, a wind shut down is expected (Koesterke \\& Werner 1998). The abrupt decrease in mass-loss rate increases the transformed radius. At the same time, regarding the spectra, the number of absorptions starts to overwhelm the number of emission lines. From our results, we can conclude that :  \\begin{itemize}  \\item The inclusion of clumping in [WR] models can change the $R_T \\times T_*$ diagram in Fig. 13 considerably. If clumping is present in all [WR] stars in a uniform way, i.e., if they can be described by the same filling factor $f$, the mass-losses will be diminished by a factor of $\\sim 3$ (for $f=0.1$). Thus, we will have a vertical shift of the same amount for each star because the transformed radius is inversely proportional to $\\dot{M}^{2/3}$ :  \\[  R_T = R_* \\left( \\frac {v_\\infty/2500 km\\, s^{-1}} {\\dot{M}/10^{-4}M_{\\odot}\\,yr^{-1}}\\right )^{2/3}.  \\]  Among the analyzed stars, direct evidence for clumping is shown only for BD +30 3639 and NGC 40.  \\item Compared to previous homogeneous models, the transformed radii increase to $\\sim 0.5$ dex. Without clumping, we still found a change of $\\sim 0.1-0.2$ dex.  \\item Our result for NGC 6905 suggests that the gap found between early-type [WR] stars and [WELS] could disappear if clumping is not important or present in this latter group.  \\item A far-UV to optical coverage plus the use of line-blanketing is probably the reason for the change in the temperatures of all stars regarding previous results. However, we did not see any systematic change that would allow any conclusion regarding the evolutionary sequence. We decreased the temperatures of NGC 40 and NGC 5315 and increased the temperatures of NGC 6905 and BD +30 3639. An important case to note is the one of NGC 5315, which will be discussed in the next section. This star is clearly far from the other early-type objects. An analysis of a larger sample is required to check if the same behavior can be found.  \\end{itemize}  Fig. 13 allows us to check the differences between our work and previous results. However, it is interesting to compare our results also with evolutionary tracks in the HR diagram. Indeed, this kind of comparison was one of the strongest arguments first supporting an evolutionary connection between [WR], [WELS] and PG 1159 stars (see Hamann 1997). Unfortunately, there are some drawbacks in doing this. First, the radii determined from the models are distant dependent whilst Fig. 13 is not. Hence, the determination of surface gravities can be very uncertain. Absorption lines in PG 1159 stars are used to overcome this difficulty, but the same is not possible for the [WR] stars. Another concern are the evolutionary tracks. Generally, the ones used are tailored for H-rich objects, which could be another source of uncertainty. The main reason for this is the complexities involved in getting a H-deficient object and its evolution in an appropriate way. Nevertheless, very recently, Althaus et al. (2005) made available a complete evolutionary track for a 2.7 $M_{\\odot}$ star from the main sequence up to the normal (H-rich) white dwarf stage. Further, after a very late thermal pulse (VLTP), the calculations were carried out to the H-deficient PG 1159 stage and finally, to the  DO, DB and DQ white dwarf region. Their work fits our needs and is reproduced here in our Fig. 14. The number of data points are augmented by our analysis and by new results regarding PG 1159 stars (H\\\"ugelmeyer et al. 2005). We also add to the plot all the previously analyzed Galactic [WR] stars (Koesterke 2001). All [WR] stars are assumed to have 0.6$M_{\\odot}$. Although the [WR] $\\rightarrow$ PG 1159 evolution can be inferred from Fig. 14 (based on an appropriate evolutionary track), some remarks should be made :  \\begin{itemize}  \\item There is not a smooth transition from [WCL] to the early-type [WR] stars (see Log $T_{*} > 4.9$). This reflects the long standing issue regarding the lack of stars in the subtypes [WC5-7], i.e., there are no [WR] stars with $T_{*}$ between $\\sim 80kK$ and $\\sim 110kK$;  \\item NGC 5315 ([WC4]-[WO4]) is clearly displaced from other early-type [WR] stars. The majority of early-type [WR] stars are distributed over a very narrow region near the PG 1159 stars, where $T_{*}$ $> 100kK$. In contrast, the position of NGC 5315 is practically the same as NGC 40, which is a [WCL] star. It must be kept in mind that although some early-type [WR] and PG 1159 stars occupy a region in common in this diagram (near Log g = 6), their spectra are quite distinct. In contrast with PG 1159 stars, all early-type [WR] stars present mass-loss;  \\item As already noted by Hamann (1997), the position of [WELS] in Fig. 14 is slightly before the early-type [WR] stars. This is another point showing the need for the analysis of more objects in this class.  \\end{itemize}  Until we elucidate the points described above, the detailed evolutionary sequence [WCL] $\\rightarrow$ [WCE] $\\rightarrow$ [WO] $\\rightarrow$ [WELS] $\\rightarrow$ PG 1159 must be considered with caution.  It would be very useful to have the same diagram but for the other two thermal pulse models (AFTP and LTP) and for different remnant stellar masses. This is beyond the scope of the present work.  \\subsection{The early-type [WR] star NGC 5315}  As already mentioned, the $T_{*}$ found for NGC 5315 is lower than previous determinations for other [WC4] stars, namely : IC 1747, NGC 1501 and NGC 6369 (see Koesterke 2001). However, it is important to note that they all now have different stellar classifications. Recently, because some [WCE] stars present oxygen lines more intense than carbon lines, new classification systems (Crowther et al. 1998; Acker \\& Neiner 2003) replaced [WC2-3] by [WO1-4] subtypes, changing the status of several objects (including some [WC4]) and giving a more appropriate description of the ionization of the wind. It turns out that NGC 5315 now is the only object analyzed from its class, which according to Crowther et al. (1998) is [WC4] and to Acker \\& Neiner (2003), [WO4]. As a matter of fact, NGC 5315 is clearly separated from other early-type [WR] stars in the diagrams in Figs. 13 and 14. This result should not be interpreted as the influence of the use of metal line-blanketing, and calls for an analysis of more objects in the [WC4]-[WO4] class."
    },
    "0609/astro-ph0609809_arXiv.txt": {
        "abstract": "Simulations of radio emission from extensive air showers we have published so far were performed with a Monte Carlo code using analytical parametrisations to describe the spatial, temporal, energy and angular particle distributions in the air shower. A new version of our simulation code, which replaces these parametrisations with precise, multi-dimensional histograms derived from per-shower CORSIKA simulations, is now available. The new code allows an independent selection between parametrisation and histogram for each of the relevant distributions, enabling us to study the changes arising from using a more realistic air shower model in detail. We describe the new simulation strategy, present some initial results and discuss the new possibilities. ",
        "introduction": "In the past few years, we have modelled geosynchrotron radio emission from extensive air showers with both analytic calculations \\cite{HuegeFalcke2003} and extensive Monte Carlo simulations \\cite{HuegeFalcke2005a}, the results of which have been presented and summarised in a parametrisation formula in \\cite{HuegeFalcke2005b}. These Monte Carlo simulations were based on an air shower model using analytical parametrisations (see \\cite{HuegeFalcke2003}) to describe the spatial, temporal, energy and angular distributions of the air shower particles. The new Monte Carlo code presented here replaces these analytical parametrisations with multi-dimensional histograms generated for each individual shower by the air shower simulation code CORSIKA \\cite{Heck1998}, allowing a much more precise calculation of the radio signal.  In the following, we present the new simulation strategy before discussing a few examples of the effects arising as a consequence of the transition to the realistic, histogrammed particle distributions and the elimination of the track length parameter that was present in the earlier simulations. Afterwards, we discuss some examples of how the new code can be used to explain the shape of the emitted radio pulses, followed by our conclusions. ",
        "conclusions": "We have developed a next-generation Monte Carlo code for the calculation of geosynchrotron radio emission from CORSIKA-simulated cosmic ray air showers. The transition from our earlier code to this new version is gradual, allowing to understand the changes in the radio pulses in detail. A last free parameter remaining in our earlier code has been eliminated in a natural, self-consistent way. The adopted two-step approach proves to be very powerful in many situations.  Although there are significant differences in the particle distributions calculated by CORSIKA in comparison to the parametrisations used earlier, the changes of the resulting radio pulses are not dramatic. Along the north-south observer axis, a damping of about a factor of two and a somewhat flatter frequency distribution in the shower centre arise."
    },
    "0609/astro-ph0609724_arXiv.txt": {
        "abstract": "Using data obtained for twelve galaxies as part of the {\\it Spitzer} Infrared Nearby Galaxies Survey (SINGS) and the Westerbork Synthesis Radio Telescope (WSRT)-SINGS radio continuum survey, we study how star formation activity affects the far-infrared (FIR)--radio correlation {\\it within} galaxies by testing a phenomenological model which describes the radio image as a smeared version of the FIR image. The physical basis of this description is that cosmic-ray (CR) electrons will diffuse measurably farther than the mean free path of dust-heating photons before decaying by synchrotron radiation. This description works well in general. Galaxies with higher infrared surface brightnesses have best-fit smoothing scale-lengths of a few hundred parsecs, substantially shorter than those for lower surface brightness galaxies. We interpret this result to suggest that galaxies with higher disk averaged star formation rates have had a recent episode of enhanced star formation and are characterized by a higher fraction of young CR electrons that have traveled only a few hundred parsecs from their acceleration sites in supernova remnants compared to galaxies with lower star formation activity. ",
        "introduction": "Radio continuum emission from normal galaxies arises from a combination of non-thermal (synchrotron) and thermal (free-free) processes; the former critically depends on a galaxy's magnetic field and cosmic ray (CR) electron distributions. To date, the evolution of CRs has been studied directly only in the Milky Way. In other galaxies, information has come from sampling the CR electron component via multi-frequency radio observations \\citep[e.g.][]{nd91,dlg95,ul96,ji99,rb05}. Radio continuum data alone, however, only describe the present distribution of CR electrons, without providing any information regarding the initial source distribution or propagation history.  The close spatial correlation between the thermal dust and non-thermal synchrotron emission within galaxies argues for a shared origin, presumably massive star formation \\citep{har75,gxh85}. Young massive stars, still in or near their dusty natal molecular cloud, emit photons that are re-radiated by dust grains at far-infrared (FIR) wavelengths. These stars are also the progenitors of supernovae (SNe) whose remnants (SNRs) are thought to accelerate CR electrons responsible for synchrotron radiation. Since the mean free path of dust-heating photons ($\\sim$100~pc) is far shorter than the diffusion length of CR electrons ($\\sim$1-2~kpc), \\citet{bh90} conjectured that the radio image of a galaxy should resemble a smoothed version of its infrared image. This prescription has been shown to hold for galaxies observed at the ``super resolution'' ($\\la$1$\\arcmin$) of {\\it IRAS} HIRES data \\citep{mh98}. Using {\\it Spitzer} imaging, \\citet[hereafter M06]{ejm06} tested this model on sub-kiloparsec scales within four nearby galaxies and found it applicable to first order, improving the correlation between the infrared and radio maps while leaving patterns in the residuals that mirror the structure in the galaxy's image. More recently, \\citet{ah06} has found that within the Large Magellanic Cloud (LMC), synchrotron halos around individual star-forming regions are more extended than FIR-emitting regions, corroborating this phenomenology on scales $\\ga$50~pc.  In this {\\it Letter} we explore how star formation activity affects the relative spatial distributions of FIR and radio continuum emission for a sample of 12 spiral galaxies observed as part of the {\\it Spitzer} Infrared Nearby Galaxies Survey (SINGS) legacy science project \\citep{rk03} and the Westerbork Synthesis Radio Telescope (WSRT)-SINGS radio continuum survey. This analysis builds on that of M06, though the larger sample has increased the range in star formation rates (SFRs) investigated by a factor of $\\sim$5. Such a study is only now possible due to the high spatial resolution and sensitivity of the {\\it Spitzer} Space Telescope at FIR wavelengths. ",
        "conclusions": "\\label{secdisc} We have analyzed the effect of star formation intensity on the FIR-radio correlation {\\it within} galaxies using a phenomenological image-smearing model. We find that the best-fit scale-length decreases as the star formation activity within each galaxy increases, implying a systematic decrease in  the mean distance diffused by CR electrons \\citep{hb93}. The measured diffusion scale-lengths can provide independent estimates for the ages of related star formation episodes as well as help constrain interstellar medium (ISM) parameters associated with the propagation characteristics of CR electrons.   The physical interpretation of our results indicates that CR electrons are, on average, closer to their place of origin in galaxies having higher star formation activity. Since the diffusion scale-length of CR electrons depends only on their age and ability to diffuse through a galaxy's ISM, three explanations of the results are possible.  The CR electrons may: (1) have relatively short lifetimes due to a high energy loss rate; (2) diffuse at a slower rate due to the ISM having a high density and magnetic field strength, resulting in a shorter mean free path; or (3) have been accelerated recently and be relatively young. The first two of these explanations are applicable in the case of steady-state star formation, while the third requires a recent episode of enhanced star formation. We will try to distinguish among these three explanations.   As CR electrons propagate through the ISM, they lose energy due to synchrotron, inverse-Compton (IC) scattering, bremsstrahlung, ionization, and adiabatic expansion losses. In normal star-forming galaxies only the first two mechanisms are significant for CR electrons associated with synchrotron emission at $\\sim$1~GHz \\citep{jc92}. Assuming an isotropic distribution of electron velocities,\\footnote[1]{For isotropically distributed electron velocities we assume \\(<\\sin^{2}\\theta> = \\twothirds\\), where $\\theta$ is the CR electron pitch angle.} the cooling timescale due to synchrotron and IC losses for CR electrons emitting at a frequency $\\nu$, in a magnetic field of strength $B$ is, \\begin{equation} \\left(\\frac{\\tau_{\\rm cool}}{\\rm yr}\\right) \\sim 2.8\\times10^{7} \\left(\\frac{\\nu}{\\rm GHz}\\right) ^{-1/2} \\left(\\frac{B}{\\rm \\mu G}\\right)^{1/2} \\left(\\frac{U_{\\rm B}+U_{\\rm rad}}{10^{-12}~{\\rm erg~cm^{-3}}}\\right)^{-1}, \\end{equation} where $U_{\\rm B}$ and $U_{\\rm rad}$ are the magnetic and radiation field energy densities, respectively.  In simple diffusion models, the propagation of CR electrons is usually characterized by an empirical energy-dependent diffusion coefficient, $D_{E}$, where \\(D_{E} = 10^{29}\\)~cm$^{2}$~s$^{-1}$ for $E < 1$~GeV and \\(10^{29}(E/{\\rm GeV})^{1/2}\\)~cm$^{2}$~s$^{-1}$ for $E \\geq 1$~GeV \\citep{ginz80}. Though derived for CR nuclei, the above values appear consistent with empirically measured diffusion coefficients for CR electrons \\citep[e.g.][]{dlg95}. According to synchrotron theory, CR electrons will emit at \\(\\nu \\propto B E^{2}\\) thereby allowing $D_{E}$ to be expressed as a function of $B$ for a fixed $\\nu$. Now, assuming that CR electrons diffuse a distance \\(l_{\\rm cool} = \\sqrt{D_{E}\\tau_{\\rm cool}}\\), we can estimate the relative importance of IC losses, ISM density ($n$), and $B$ on the distance CR electrons travel.  Fixing $B$ and scaling \\(U_{\\rm rad} \\propto \\Sigma_{\\rm SFR}\\), we plot the expected relation between $l_{\\rm cool}$ and $U_{\\rm rad}$ in Figure \\ref{SFRDleg} (dashed line). This relation has a slope $\\sim$9 times shallower than a line fit to the data (dotted line). It is more likely that $n$, $B$, and $U_{\\rm rad}$ will all scale together, with a $D_{E}$ dependence on $n$. Accordingly, we introduce the scaling relations, \\begin{equation} B \\propto n^{\\beta},~D_{E} \\propto n^{\\delta}B^{-1/4}, \\nonumber \\end{equation} where \\(\\onethird \\leq \\beta \\leq \\twothirds\\) and \\(-1 \\leq \\delta \\leq -\\onethird\\) are realistic index ranges \\citep{hb93}.  Setting \\(U_{\\rm B} = U_{\\rm rad}\\), which has been shown to be valid for a large sample of spirals \\citep{lvx96}, leads to \\(B \\propto \\Sigma_{\\rm SFR}^{1/2}\\), \\(\\tau_{\\rm cool} \\propto B^{-3/2}\\), and the dependence \\begin{equation} l_{\\rm cool} \\propto \\Sigma_{\\rm SFR}^{(\\delta/(4\\beta) - 7/16)}. \\end{equation} For the unlikely case that each galaxy takes the extreme values of \\(\\beta = \\onethird\\) and \\(\\delta = -1\\), we obtain the steepest possible index of $-19/16$ (solid line in Figure \\ref{SFRDleg}), which is still a factor of $\\sim$2 times shallower than the line fit to the data. Consequently, differences in the diffusion properties and/or cooling timescales of each galaxy's CR electron population alone (i.e. the case of steady-state star formation), cannot explain the observed trend in Figure \\ref{SFRDleg}.  On the other hand, by setting the best-fit scale-lengths equal to $l_{\\rm cool}$ and $B = 5~\\mu$G for each galaxy, we estimate mean ages ranging from $\\sim$1.1$\\times 10^{5}$ to $\\sim$8.3$\\times 10^{7}$~yr for these 1.4~GHz emitting populations of CR electrons. Even if we assume \\(B \\propto \\Sigma_{\\rm SFR}^{1/2}\\), this range of ages shrinks only slightly. While these ages are plausible, spanning nearly 3 orders of magnitude over a factor of 4 change in $\\Sigma_{\\rm SFR}$ seems unphysically large. This suggests that as $\\Sigma_{\\rm SFR}$ decreases the best-fit scale-length shifts from being dominated by young CR electrons associated with regions of star formation to being dominated by old CR electrons in the galaxy's underlying disk. Consequently, we conclude that galaxies with higher $\\Sigma_{\\rm SFR}$ are likely to have experienced a recent episode of enhanced star formation compared to galaxies with lower disk-averaged SFRs, leading to a larger fraction of relatively young CR electrons within their disks. We should also note that galaxies with higher star formation activity may also suffer from increased CR electron escape due to instability driven breaks in magnetic field lines \\citep{ep66}, though characterizing the importance of escape is beyond the scope of this {\\it Letter}.  The galaxies analyzed here exhibit a wide range of ages and spreading scale-lengths of their CR electrons, going from the youngest populations, recently injected within star-forming regions ($l \\sim$500~pc), to the oldest CR electrons, making up a galaxy's underlying synchrotron disk ($l \\sim$3~kpc). This exponential synchrotron disk results from the superposition of older CR electrons originating from all star-forming regions. In spite of an occasionally dominant mode, a given galaxy will contain a range of ages for CR electrons, so it is no surprise that a single smoothing kernel cannot describe perfectly the relation between infrared and radio disks, as evidenced by the recognizable structure in the residual maps. NGC~628, NGC~2403, NGC~3031, and NGC~3184 exhibit inflection points in their residual curves which is also suggestive of multiple CR electron populations, possibly from time-separated bursts of star formation. A more realistic description to account for the range of ages translates into a multiscale analysis of the galaxy images to separate the signatures of various CR electron populations. This will be the subject of a forthcoming paper."
    },
    "0609/nucl-th0609026_arXiv.txt": {
        "abstract": "The effects of the symmetry energy softening of the relativistic mean field (RMF) models on the properties of matter with neutrino trapping are investigated. It is found that the effects are less significant than those in the case without neutrino trapping. The weak dependence of the equation of state on the symmetry energy is shown as the main reason of this finding. Using different RMF models the dynamical instabilities of uniform matters, with and without neutrino trapping, have been also studied. The interplay between the dominant contribution of the variation of matter composition and the role of effective masses of mesons and nucleons leads to higher critical densities for matter with neutrino trapping. Furthermore, the predicted critical density is insensitive to the number of trapped neutrinos as well as to the RMF model used in the investigation. It is also found that additional nonlinear terms in the Horowitz-Piekarewicz and Furnstahl-Serot-Tang models prevent another kind of instability, which occurs at relatively high densities. The reason is that the effective $\\sigma$ meson mass in their models increases as a function of the matter density. ",
        "introduction": "\\label{sec_intro} Recently, the dynamical stability of the uniform ground state of multi components systems (i.e., electrons, protons, and neutrons) at low densities  has received considerable attentions~\\cite{Pethick,Douchin,Carr,Horowitz01,Provi1,Provi2}. This interest is motivated by the fact that the neutron star is expected to have a solid inner crust of nonuniform neutron-rich matter above its liquid mantle~\\cite{Carr} and the mass of its crust depends sensitively on the density of its inner edge and on its equation of state (EOS)~\\cite{Douchin}. Meanwhile, the critical density ($\\rho_c$), a density at which the uniform liquid becomes unstable to a small density fluctuation, can be used as a good approximation of the edge density of the crust~\\cite{Carr}. Using the Skyrme SLy effective interactions,  Ref.~\\cite{Douchin} found the inner edge of the crust density to be  $\\rho_{\\rm edge}$ = 0.08 $\\rm fm^{-3}$. Reference~\\cite{Carr} generalized the dynamical stability analysis of Ref.~\\cite{Horo1} in order to accommodate the various nonlinear terms in the relativistic mean field (RMF) model of Horowitz-Piekarewicz~\\cite{Horowitz01}. They found a strong correlation between $\\rho_c$ in neutron star and the symmetry energy ($a_{\\rm sym}$) which leads to a linear relation between   $\\rho_c$ and the skin thickness of finite nuclei.  These results suggested that a measurement of the neutron radius in $^{208}{\\rm Pb}$ will provide useful information on the $\\rho_c$ \\cite{Carr,Horowitz01}. Recently, a new version of the RMF model based on effective field theory  (ERMF) has been proposed by Furnstahl, Serot, and Tang~\\cite{Furn96}. The predictive power of the model in a wide range of densities is quite impressive (see the review articles~\\cite{sil,serot} for the details). On the other hand, an adjustment of the isovector-vector channel of the ERMF model in order to achieve a softer density dependence of $a_{\\rm sym}$ at high densities has been done in Ref.~\\cite{anto05}.  Similar to the previous calculation done in~\\cite{Horowitz01}, but using a different RMF model, it is found that by adjusting that channel a lower proton fraction ($Y_p$) in high density neutron star matter can be achieved without significantly changing the bulk properties~\\cite{anto05}. It is well known that $Y_p$ is related to the threshold of the direct URCA cooling process  and the trend of the density distribution of $Y_p$ is unique for each model which is sensitive to the different forms of nonlinear terms used. Therefore, an investigation of the EOS and the instability at low densities by using different RMF models could be quite interesting. Indeed, by comparing our results with the previous ones~\\cite{Carr,Horowitz01}, we can systematically study the influence of the form and the strength of isovector-vector nonlinear terms on the critical density.  The presence of electrons and the influence of the electromagnetic interaction on the unstable modes of asymmetric matter as well as the comparison between  dynamical instability region with the thermodynamical one have been investigated in Refs.~\\cite{Provi1,Provi2} by using the standard RMF model within the Vlasov formalism. They observed the important role of the Coulomb field in the large structure formation and also the role of the electron dynamics in restoring the large-wavelength instabilities \\cite{Provi2}.  In the actual situation, muons may also be present in the stellar matters. Their existence yields an additional electromagnetic contribution besides the electrons and protons ones. The significance of their contribution depends on the matter composition. Including muons in the dynamical instability analysis of uniform matter increases the electromagnetic effect on the $\\rho_c$ in medium. Moreover, in supernovae or protoneutron stars neutrinos can be trapped inside, if their mean free paths are smaller than the star radius. In general, the presence of trapped neutrinos in matter affects the stiffness of EOS and drastically changes the composition of neutral matter~\\cite{Prakash,Vidana,Guo,chiapparini,Bednarek}. It was also found that neutrino trapping causes the strange baryon~\\cite{Prakash,Vidana} and kaon condensate \\cite{Prakash, Guo} to appear at higher densities as compared to the case without neutrino trapping. In Ref.~\\cite{Bednarek}, the EOS of the strangeness rich protoneutron star and the EOS of neutron star matter with temperature different from  zero have been already studied by means of the ERMF model. Meanwhile, like neutron stars, protoneutron stars have dense cores and outer layers~\\cite{Bednarek} and in supernovae inhomogeneity can appear below the saturation density of nuclear matter ($\\rho_0$)~\\cite{shen,wata}. Therefore, an investigation of the boundary between the two phases, which leads to the determination of  $\\rho_c$, becomes crucial for a realistic and complete description of stellar matters.   In the present work, we shall extend the analysis of dynamical stability in uniform matter of Ref.~\\cite{Carr} in two ways, i.e., first, we consider the muon contribution and, second, we shall consider all possible mixings of vector and isovector contributions caused by the nonlinear terms in the ERMF model. The parameter set NL3 of the standard RMF~\\cite{lala} model as well as the parameter sets   Z271 and  Z271* of the Horowitz-Piekarewicz one~\\cite{Horowitz01} will be presented  for comparison. The results are used to study the effect of the neutrino presence on the $\\rho_c$ values. Here, we only use nucleons in the baryonic sector along with $\\sigma$, $\\omega$, and $\\rho$ mesons in the mesonic sector because $\\rho_c$ usually appears at a density lower than $\\rho_0$ and the strange particles appear at a density higher than $2\\rho_0$. Furthermore,  in the neutrino trapping case they appear at higher density than in the case  without neutrino trapping. Therefore, we can assume the strange particles to have only minor effects on the $\\rho_c$ and, as a consequence, their contributions may be neglected in this work.  On the other hand, although in the real situation the temperature of protoneutron stars is not equal to zero and the supernovae inner core can have a temperature around $T \\sim$ (10-50) MeV, the zero temperature approximation can be assumed here since the temperature effect on the EOS of supernovae matter and on the maximum mass of protoneutron star~\\cite{Guo,chiapparini} is smaller than that without neutrino trapping. In this approximation, the following constraints can be  used to calculate the fraction of every constituent in matter: \\begin{itemize} \\item balance equation for the chemical potentials \\be \\mu_{n}+\\mu_{\\nu_e} = \\mu_{p}+\\mu_{e}, \\ee \\item conservation of the charge neutrality \\be \\rho_e +\\rho_{\\mu}=\\rho_p, \\ee \\end{itemize} where the total density of baryon is given by \\be \\rho_B =\\rho_n+\\rho_p, \\ee while the fixed electronic-leptonic fraction $Y_{l_e}$ is defined as \\be Y_{l_e}=\\frac{\\rho_e+\\rho_{\\nu_e}}{\\rho_B}\\equiv Y_e+Y_{\\nu_e}, \\ee where $Y_{\\nu_e}$ and  $Y_e$ are neutrino electron and electron fractions, respectively. In addition, we will revisit the EOS of matter with neutrino trapping because we want to study the effect of different adjustments in the isovector-vector sector \\cite{anto05} on the EOS of matter with neutrino trapping. They were not fully explored in our previous works~\\cite{anto05,anto06}.  This paper is organized as follows. In Sec. II, a brief review of the model is given. Calculation of the dielectric function is presented in Sec. III, while numerical results along with the corresponding discussions are given in Sec. IV. Finally, we give the conclusion in Sec. V. ",
        "conclusions": "\\label{sec_conclu} The effects of the different treatments in the isovector-vector sector of RMF models on the properties of matter with neutrino trapping have been studied. The effects are less significant compared to those without neutrino trapping. Different dependences of the EOS and $B/A$ of both matters on $a_{\\rm sym}$ are the reason behind this.  The effects of the variation of the  neutrino fraction in matter on the EOS and $B/A$ have been also discussed.  The longitudinal dielectric function of the ERMF model for matter consisting of neutrons, protons, electrons, muons, and neutrinos has been derived. The result is used to study the dynamical instability of uniform matters at low densities. The behavior of the predicted $\\rho_c$ in matter with and without neutrino trapping has been investigated. It is found that different treatments in the isovector-vector sector of RMF models yield more substantial effects in matter without  neutrino trapping rather than in matter with neutrino trapping. Moreover, for  matter with neutrino trapping, the value of $\\rho_c$ does not significantly change with the variation of the models as well as  with the variation of the neutrino fraction in matter. In this case, the value of $\\rho_c$ is larger for matter with neutrino trapping. These are due to the interplay between the major role of matter composition and the role of the effective masses of mesons and nucleons. It is also found that the additional nonlinear terms of Horowitz-Piekarewicz and ERMF models prevent another instability at relatively high densities  to appear. This can be traced back to the effective $\\sigma$ mass which goes to zero when the density approaches $6.5~\\rho_0$ ."
    },
    "0609/astro-ph0609042_arXiv.txt": {
        "abstract": "We study a sample of approximately $16,500$ galaxies with $I_{ACS,AB}\\leq22.5$ in the central 38\\% of the COSMOS field (\\cite{scov06}, this volume), which are extracted from a catalog constructed from the Cycle 12 ACS F814W COSMOS data set. Structural information on the galaxies is derived by fitting single S\\'ersic (\\cite{sers68}) models to their two-dimensional surface brightness distributions. \\\\In this paper we focus on the disk galaxy population (as classified by the Zurich Estimator of Structural Types, ZEST; \\cite{scar06}, this volume), and investigate the evolution of the number density of disk galaxies larger than approximately 5 kpc between redshift $z\\sim1$ and the present epoch. Specifically, we use the measurements of the half-light radii derived from the S\\'ersic fits to construct, as a function of redshift, the size function $\\Phi(r_{1/2},z)$ of both the total disk galaxy population and of  disk galaxies split in four bins of bulge-to-disk ratio. In each redshift bin, the size function specifies the number of galaxies per unit comoving volume and per unit half-light radius $r_{1/2}$. Furthermore, we use a selected sample of roughly 1800 SDSS galaxies to calibrate our results with respect to the local universe. We find that:\\\\ $(i)$ The number density of disk galaxies with intermediate sizes ($r_{1/2}\\sim$ 5-7 kpc) remains nearly constant from $z\\sim1$ to today. Unless the growth and destruction of such systems exactly balanced in the last eight billion years, they must have neither grown nor been destroyed over this period.\\\\ $(ii)$ The number density of the largest disks ($r_{1/2}>7$ kpc) decreases by a factor of about two out to $z\\sim1$.\\\\ $(iii)$ There is  a constancy - or even slight increase - in the number density of large bulgeless disks out to $z\\sim1$; the deficit of large disks at early epochs seems to arise from a smaller number of {\\it bulged} disks. \\\\ Our results indicate that the bulk of the large disk galaxy population has completed its growth by $z\\sim1$, and support that secular evolution processes produce  -- or at least add stellar mass to -- the bulge components of disk galaxies. ",
        "introduction": "A detailed quantification of galaxy structure as a function of redshift is an essential tool to study the evolution of galaxies through cosmic time.  In particular, the evolution of galaxy sizes is a direct diagnostic of their star formation and mass assembly history. For disk galaxies, several theoretical predictions have been made, which start being testable with observations:  For example, under the assumption that the scale lengths of disks scale with their virial radii, their sizes are predicted to evolve as $r_{s}\\propto H(z)^{-2/3}$ at a fixed halo mass (\\citet{fallefstathiou80}, \\citet{mao98});  \\cite{bouwenssilk02}, instead, assume a dependence of the star formation rate on the local gas density within the disk, and infall of metal-free gas, and predict the sizes of disks to evolve as $R(z)/R(0) = 1 - 0.27\\cdot z$.  Pioneering observations, however, generally appear to support a substantially weaker size evolution of the disk galaxy population than predicted by theory. Data from the Canada-France Redshift Survey (CFRS) for a few tens of large disk galaxies up to redshift $z\\sim1$  indicate  about a magnitude of surface brightness  evolution for the disks, and a constant number density of large disks over this period (\\cite{schade95}; \\cite{schade96}; \\cite{lil98}). Similar results are reported by \\cite{roche98}. In contrast, \\cite{simard99} interpret their data as indicating no surface brightness evolution in disks over the same redshift range. The work of \\cite{trujillo04}, on the other hand, provides further evidence for surface brightness evolution in galaxy disks, although a possible decrease in the sizes of disks by approximately 30\\% at $z\\sim$ 0.7 is also proposed as an alternative to explain the data. Using 1500 disk galaxies at redshifts 0.25 $<z<$ 1.25 from the GOODS survey, \\cite{ravindranath04}  find at most a  very modest decrease of the number of disks with half-light radius $r_{1/2}>$ 4 kpc. The GEMS survey has been used to investigate the evolution with redshift of the  luminosity-size relation for disk galaxies (\\cite{barden05}; see also \\cite{trujillo05}), but  no analysis of the evolution of the number density of disk galaxies as a function of galaxy size has been presented based on the GEMS data.  Therefore, the observational evidence for a relatively constant number of (large) disks since $z\\sim 1$ relies on very small samples extracted from small areas in the sky.   In this article we use the COSMOS survey (\\cite{scov06}, this volume) to  investigate how the number density of disk galaxies with sizes above 5 kpc - the threshold above which our sample is complete - evolves with cosmic time from $z\\sim1$ to $z=0$. Thanks to the large statistics of COSMOS, we are able to extend the investigation so as to study this evolution as a function of the bulge-to-disk ratio of disk galaxies. Specifically, we use the ACS COSMOS database, and select a $I_{ACS,AB}\\leq22.5$ sample from the 38\\% of the COSMOS area that was imaged by the HST during Cycle 12. The complete $I_{ACS,AB}\\leq22.5$ sample contains more than 16,000 galaxies. We perform two-dimensional GIM2D (\\cite{marlsimard98}) fits to their surface brightness distributions, which we describe with a single S\\'ersic law (\\cite{sers68}). Free parameters in the fits are: the total flux $F_{tot}$, the S\\'ersic index $n$, the half-light radius $r_{1/2}$, the ellipticity $\\epsilon = 1-b/a$ (with $a$ and $b$ the semimajor- and minor-axis), and the position angle $\\phi$.  We compare the structural parameters obtained with the GIM2D  surface brightness fits with the structural measurements derived using  the Zurich Estimator of Structural Types (ZEST; \\cite{scar06}, this volume). ZEST uses individual measurements of four non-parametric diagnostics of galaxy structure, namely, the concentration $C$, the asymmetry $A$, the Gini coefficient $G$, and the second order moment of the brightest 20\\% of the galaxy pixels, $M_{20}$ (c.f. \\cite{scar06}, this volume, and references therein), plus the elongation of the distribution of the light, to classify galaxies in three morphological types: early-types, disk galaxies or irregular and peculiar galaxies. The disk galaxies are further split by ZEST into four classes of bulge-to-disk ratio ($B/D$). At any redshift, the multi-dimensional ZEST classification grid substantially reduces the intermixture of structurally different galaxy populations compared with simpler approaches. The comparison of our parametric structural quantities with the diagnostics supplied by ZEST shows the expected correlations (e.g., between S\\'ersic index $n$ and concentration $C$) and, by doing so, illustrates the robustness of the results of our GIM2D fits. It also demonstrates that studies of selected galaxy populations which are classified by a simple cut in S\\'ersic index $n$ - as customarily done, for instance, in studies of the SDSS galaxy population and in other surveys of the high-$z$ universe - are affected by  mixing of galaxies with different structural properties.  We use the ZEST classification to extract, from the total  $I_{ACS,AB}\\leq22.5$ sample, the subsample of approximately 12,000 disk galaxies (with different $B/D$ ratios), and   use the half-light radii provided by our S\\'ersic fits to compute  the number of disk galaxies per unit comoving volume and per unit half-light radius $r_{1/2}$ as a function of redshift (the {\\it size function} $\\Phi(r_{1/2},z)$).  In order to compare the results  at higher redshift with the local universe, we use a suitably selected sample of  about 1800 SDSS galaxies, which have been ``redshifted\" to $z=0.7$ by properly re-scaling for flux dimming, distance and image resolution (\\cite{kampczyk06}, in preparation).  We find that the number density of disk galaxies with intermediate sizes $r_{1/2}\\sim$ 5-7 kpc remains nearly constant from $z\\sim1$ to today, while that of the largest disks ($r_{1/2}>7$ kpc) shows a drop by a factor of approximately two. This deficit appears to stem mostly from a decline in the number density of {\\it bulged} disks at early epochs.\\footnote{In the rest of the paper, we consistently use AB magnitudes (\\cite{oke74}), and  thus drop the subscript ``AB\" from our notation. We adopt $\\Omega_{m}=0.25$, $\\Omega_{\\Lambda}+\\Omega_{m}=1$ and a present day Hubble parameter of $H_0=70~km~s^{-1}~Mpc^{-1}$.} ",
        "conclusions": "We have used measurements of the half-light radii obtained with analytic fits to the surface brightness distribution of roughly 12,000  disk galaxies in COSMOS to chart the evolution of the number density of intermediate-sized and large disks since redshift $z\\sim1$.  Our analysis has shown a general constancy of the size function for disk galaxies with half-light radii larger than about 5 kpc. This constancy, initially proposed by (\\cite{lil98}) based on much smaller statistics, and here clearly demonstrated from a much larger sample of disk galaxies of any bulge-to-disk ($B/D$) ratio, suggests that the massive disk galaxy population was largely in place with its current properties when the universe was only half as old as today.  Among the observed general constancy of the size function there is evidence for some subtler and possibly quite important trends: $(i)$ The number density of disk galaxies of all $B/D$-ratios in the intermediate size range $r_{1/2}\\in$ [5 kpc, 7 kpc[ remains constant to a good approximation up to $z\\sim1$. $(ii)$ The number density of the largest disks with half-light radii $r_{1/2}\\in$ [7 kpc, 17 kpc[ at $z\\sim1$ appears however to be only 60\\% of today's value in the disk galaxy population as a whole. $(iii)$ When this population of large-sized disk galaxies is split by the values of their $B/D$-ratio we find that the deficit at high redshift stems from disk galaxies with a bulge component.  It is conceivable that color differences between bulges and disks contribute to this result and we will study their effect in more detail.  Bearing this in mind, however, it is nevertheless likely that we might be  witnessing a secular transformation of bulgeless disks into disks with a bulge component.  \\clearpage  \\appendix"
    },
    "0609/astro-ph0609568_arXiv.txt": {
        "abstract": "We give a tracker solution for the quintessence scalar field for Ratra--Peebles or SUGRA potentials, holding before, during and after the equality epoch ($\\rho_m=\\rho_r$) and nicely fitting the numerical behavior. ",
        "introduction": "Dark Energy (DE) is one of the main puzzles of cosmology. It could be a scalar field $\\phi$, self--interacting through a potential $V(\\phi)$ (Wetterich 1988, Ratra \\& Peebles 1988), so that \\begin{equation} \\rho_{DE}=\\rho_{k,DE}+\\rho_{p,DE} \\equiv {{\\dot{\\phi }}^{2}/2a^{2}}+V(\\phi ),~ ~~~~ p_{DE}=\\rho_{k,DE}-\\rho_{p,DE} = w \\rho_{DE}~, \\end{equation} provided that $\\rho_{k,DE}/V\\ll 1/2$, so that $-1/3 \\gg w >-1$. Here \\begin{equation} ds^{2} = a^{2}(\\tau) (-d\\tau^{2}+ dx_i dx^i ) ~,~~~~~ (i=1,..,3) \\end{equation} is the metrics and dots indicate differentiation with respect to $\\tau $ (conformal time). This kind of DE is dubbed {\\it dynamical} DE (dDE) or {\\it quintessence}.  The most significant potentials are those allowing tracker solutions, so limiting the impact of initial conditions. Among them, we consider here the SUGRA (Brax \\& Martin 1999, 2001; Brax, Martin \\& Riazuelo 2000) and Ratra \\& Peebles (1988, 2003) potentials \\begin{equation} \\label{ptntl} SUGRA ~:~~~ V(\\phi) = (\\Lambda^{4+\\alpha}/\\phi^\\alpha) \\exp(4 \\pi \\phi^2/m_p^2) \\end{equation} \\begin{equation} RP ~:~~~ V(\\phi) = \\Lambda^{4+\\alpha}/\\phi^\\alpha \\end{equation} ($m_p=G^{-1/2}$: Planck mass); in the recent times, they yield $|w|$ close to unity and fastly variable or $|w|$ farther from unity and slowly variable, respectively. In a flat model, once the DE density parameters $\\Omega_{DE}$ is assigned, either $\\alpha$ or the energy $\\Lambda$ can still be freely chosen.  SUGRA and RP behave differently when $\\phi$ approaches $m_p$. Here, however, we shall not deal with late epochs and our expressions, worked out for RP, suit SUGRA as well.  Tracker solutions are then usually found by considering the equation of motion \\begin{equation} \\ddot \\phi + 2(\\dot a/a) \\dot \\phi = \\alpha a^2 \\Lambda^{4+\\alpha}/ \\phi^{1+\\alpha} \\label{phieq} \\end{equation} and seeking solutions of the form \\begin{equation} \\phi = \\phi_i (\\tau/\\tau_i)^\\beta \\end{equation} ($\\tau_i$ is a reference time and $\\phi_i$ is the field value at that time). They are fixed once the time dependence of the scale factor $a(\\tau)$ is set. In turn, we can however write \\begin{equation} \\dot a/a = u(\\tau)/\\tau~, \\end{equation} so that $u=1$ ($u=2$) in the radiation (matter) dominated eras and is actually $\\tau$ dependent around matter--radiation equality.  In this note we give a single expression for the tracker solution holding before, during and after equality, occurring at $\\tau_e$. This expression improves the tracker solution expression, known for $\\tau \\ll \\tau_{e}$, when $\\tau$ approaches $\\tau_e\\, $. It also neatly improves the tracker solution known for $\\tau \\gg \\tau_{e}$, for a large range of redshifts.  Furthermore, it fits the behavior of the $\\phi$ field across equality, with quite a small discrepancy from the numerical solution, keeping mostly well below $1\\, \\%\\, $. ",
        "conclusions": "Tracker solutions play an important role in the analysis of dDE cosmologies. A good tracker solution for the radiation dominated regime can be easily given. A tracker solution for the matter dominated regime also exists, but its usual analytical expression is rather far from a fair numerical behavior.  Here we gave a fair solution performing as well as the usual one in the radiation dominated era, fitting the numerical solution also about the equality (the residual discrepancy approaches $\\sim 1\\, \\% \\, $ just in a narrow interval) and fitting the numerical solution all through the matter dominated era. It reads \\begin{equation} \\phi = \\phi_e \\left( {\\tau \\over \\tau_e} {a \\over a_e} \\right)^b \\end{equation} with \\begin{equation} b={2 \\over 2+\\alpha}~,~~ \\phi_e = {\\alpha \\Lambda (a_e \\tau_e \\Lambda)^b \\over b^2 (u+1)^2 - b(u+1)(2u-1) }~. \\end{equation} It was found by using an exact analytical expression for $a(t) $ through the equality period, also to work out the $\\tau$ dependence of $u = \\tau \\dot a/a\\, $.  When initial conditions to any numerical problem are to be set, one usually needs to go back to the radiation dominated era, so to rely on a fair tracking. As an example of problems where initial conditions in the matter dominated era are useful, let us remind the study of the evolution of a spherical fluctuation, the prediction of cluster mass function and its redshift evolution (Mainini, Macci\\`o \\& Bonometto 2003, Mainini et al. 2003, Mainini 2005) using a Press \\& Schechter (1974) or similar (Sheth \\& Tormen 1999, 2002, Jenkins et al, 2001) approach, and performing N--body simulations (see, e.g., Klypin et al. 2003, Macci\\`o et al. 2004, Solevi et al. 2006).  Using the GT expression, given here, this can be safely avoided, and initial conditions can be given at any time, until DE itself begins to be important for the overall cosmic expansion.  \\vskip .3truecm  {\\bf Acknowledgments}  Loris Colombo and Roberto Mainini are gratefully thanked for discussions.  \\vskip .3truecm \\vskip .3truecm"
    },
    "0609/astro-ph0609797_arXiv.txt": {
        "abstract": "Turbulent Comptonization, a potentially important damping and radiation mechanism in relativistic accretion flows, is discussed. Particular emphasis is placed on the physical basis, relative importance, and thermodynamics of turbulent Comptonization.  The effects of metal-absorption opacity on the spectral component resulting from turbulent Comptonization is considered as well. ",
        "introduction": "\\label{s:1}  If a beam of photons interacts with a scattering layer consisting of relatively hot electrons such that $k_BT_e>>h\\nu$, the input seed photons gain energy while conserving their overall number, a process known as Compton up-scattering or ``Comptonization'' \\cite{ral79}.  In the non-relativistic limit, the shift in photon energy per scattering is small, so that $\\Delta E/E\\sim k_BT_e/ m_ec^2$.  The rare photons that remain in the scattering layer the longest are subsequently boosted to energies limited by the electron temperature  -- a process that is statistically identical to the classical theory of cosmic ray acceleration \\cite{fer49} -- and the resulting spectral energy distribution resembles a power law.  The fractional shift in energy per scattering event is proportional to the mean square electron velocity i.e., $\\Delta E/E\\sim \\left<v^2_e\\right>/c^2$ where $\\left<\\right>$ is taken over the electron distribution function.  Also, if the scattering layer moves in bulk with respect to the observer, then the radiation will experience a Doppler shift such that $\\Delta E/E\\sim v/c$, an effect often referred to in cosmological contexts as the ``kinetic Sunyaev- Zeldovich effect.''  However, if the scattering layer is turbulent and stationary with respect to the observer, then the Doppler shift term that is $\\propto v/c$ vanishes despite the presence of large scale bulk motions.  It follows that the contribution to Comptonization from turbulence leads to an amplification $\\Delta E/E\\sim v^2_{\\rm T}/c^2$, where $v_{\\rm T}$ is the turbulent velocity.  This effect, which we refer to as ``turbulent Comptonization,'' deforms the spectra of soft photons in a manner indistinguishable from that of thermal Comptonization \\cite{pal97},\\cite{tho94}.\\footnote{In the language of cosmology, turbulent Comptonization might be referred to as the ``second-order kinetic Sunyaev-Zeldovich effect.''}  In this note, we discuss certain aspects of turbulent Comptonization in the context of relativistic accretion disk theory.  Many of the results summarized here can be found in \\cite{sdb04}. ",
        "conclusions": "The mechanisms that result in the randomization of gravitational binding energy, transport of angular momentum, and the conversion of binding energy into heat and radiation must be simultaneously understood if a {\\it predictive} theory of accretion is desired.  It is widely believed that MRI turbulence is responsible for the randomization of gravitational binding energy and the transport of angular momentum\\cite{bah91}.  Also, sophisticated techniques used to model the spectra of relativistic accretion flows fit the data quite well.  Nevertheless, there is currently no accepted method of connecting the underlying MHD to the observable radiation spectrum, implying that accretion theory suffers from an overall lack of predictive power.  We argue that turbulent Comptonization alleviates some of these inadequacies even if it is not the dominant mechanism of radiative energy release, since it allows observers to peer into the inner workings of the turbulence which is thought to be the prime mover of the accretion process."
    },
    "0609/astro-ph0609332_arXiv.txt": {
        "abstract": "We have used the coronagraphic mode of the Advanced Camera for Surveys aboard the {\\it Hubble Space Telescope} to make the first polarization maps of the debris disk surrounding the nearby M star AU Microscopii (GJ 803). The linear polarization of the scattered light from the disk is unambiguously detected.  We find that the degree of polarization for the disk rises monotonically from about 0.05 to 0.35 between projected radii of 20 and 50 AU.  Polarized light is detectable out to about 80~AU, where the fractional polarization reaches a maximum observed value of $0.41\\pm 0.02$.  Polarization vectors are oriented perpendicular to the disk mid-plane, indicating that polarization originates from single scattering in an optically thin dust disk where the albedo is dominated by small ($x = 2\\pi a/\\lambda \\lesssim 1$) particles.  We use simple, optically thin disk models to infer the spatial structure of the disk and the scattering properties of the constituent grains by simultaneously fitting the surface brightness profile and the degree of linear polarization.  The best fit models require that the dust grains exhibit high maximum linear polarization and strong forward scattering. The inner disk ($<$40--50~AU) is depleted of micron-sized dust by a factor of more than 300, which means that the disk is collision dominated, i.e., grains that are dragged inward by corpuscular and Poynting-Robertson drag undergo a destructive collision. While the inferred optical properties are covariant with the radial distribution of dust, the only acceptable models have $p_{max} \\ge 0.50$ and $g \\ge 0.7$.  These constraints cannot be met by spherical grains composed of conventional materials.  A Mie scattering analysis implicates grains where the real part of the refractive index is about 1.03, which is a signature of highly porous (91--94\\%) media.  More reliable methods for calculating the scattering properties of aggregates confirm that the observed values of $p_{max}$ and $g$ can be matched by high porosity, micron-sized clusters of small particles. In the inner Solar System, porous particles form naturally in cometary dust, where the sublimation of ices leaves a ``bird's nest'' of refractory organic and silicate material.  In AU Mic, the grain porosity may be primordial, because the dust ``birth ring'' lies beyond the ice sublimation point.  The observed porosities span the range of values implied by recent laboratory studies of particle coagulation in the proto-solar nebula by ballistic cluster-cluster aggregation.  To avoid compactification, the upper size limit for the parent bodies is in the decimeter range, in agreement with theoretical predictions based on collisional lifetime arguments.  Consequently, AU Mic may exhibit the signature of the primordial agglomeration process whereby interstellar grains first assembled to form macroscopic objects. ",
        "introduction": "Approximately 15\\% of nearby main sequence stars manifest infrared excess due to the reprocessing of stellar radiation by dust grains in a circumstellar disk \\citep{1984ApJ...278L..23A, 1993prpl.conf.1253B}. These systems are known as ``debris disks'' because the lifetime of dust is orders of magnitude shorter than the stellar age, suggesting a continuous supply of fresh grains released from larger, undetected parent bodies.  In the solar system, interplanetary dust particles are resupplied by the collisional erosion of asteroids and the sublimation of comets.  Cometary dust grains retain the history of their interstellar origin, but are imprinted with structures that speak to their incorporation into larger bodies and subsequent return to interplanetary space \\citep{1990ApJ...361..260G}.  Therefore, cometary dust provides our closest link to the particle coagulation processes that occurred during the earliest phases of planet building in the Solar System.  Although debris-disk dust must be highly processed and modified during incorporation into, and subsequent attrition of, large bodies, these particles are our primary source of information regarding the growth of assembly of planetesimals in exoplanetary systems.  The laboratory study of low-velocity dust interactions thought to be characteristic of conditions in the early solar nebula suggests that particles grow under ballistic cluster-cluster aggregation into fractal assemblies \\citep{1998Icar..132..125W}. These bodies suffer restructuring when the aggregate diameters exceed a few centimeters. Smaller bodies are not expected to be subjected to impact compaction \\citep{2000Icar..143..138B}.  As these clusters have unique optical properties, the study of debris disks provides experimental validation of our ideas regarding growth of solid bodies from interstellar grains to macroscopic objects \\citep{2006A&A...449.1243K}.  At visible wavelengths debris-disk dust can be detected in scattered light, analogous to the Zodiacal light in our Solar System.  As in the Solar System, the polarization state of this scattered light is a key diagnostic of the grain properties.  In an optically thin disk, where single scattering dominates, a high degree of linear polarization is expected, with a characteristic orientation that is symmetric (concentric or radial) about the illuminating source \\citep[e.g.,][]{2003pid..book.....K}.  The observed degree of polarization is an important indicator of the size, shape, composition, physical structure and alignment of individual grains and their distribution along the line of sight.  Single particle scattering can be described by the matrix elements of the complex amplitude scattering function, $S$, which depend on the scattering angle \\citep{1981lssp.book.....V}.  The angular dependence of $S$ is a key clue to the nature of the particles. For example, the degree of forward scattering increases with particle size from the Rayleigh limit ($x=2\\pi a/\\lambda \\ll 1$), where the scattering asymmetry parameter, $g = \\langle \\cos \\theta \\rangle \\simeq 0$, to $g\\simeq 1$ as $x$ approaches unity.  Observations of spatially resolved disks can directly determine the azimuthal asymmetry due to asymmetric scattering \\citep[e.g.,][]{2005Natur.435.1067K}.  Because debris disks are optically faint relative to their illuminating star, there is a strong selection effect that favors the discovery and facilitates the study of high surface brightness, edge-on systems, e.g., $\\beta$ Pic, HD 32297, HD 139664, or AU Mic.  In an edge-on system, the azimuthal variation of surface brightness around the star is unobservable, and thus grain scattering properties versus phase angle cannot be measured directly.  Each line of sight through the disk includes a range of scattering angles.  Therefore, both the asymmetry parameter and the radial variation of dust opacity determine the disk surface brightness as a function of angle on the sky.  Unless there is prior information regarding $g$ or the dust distribution these properties cannot be uniquely disentangled from measurements of the radial surface brightness distribution.  Inclusion of further constraints, such as the spectral energy distribution \\citep[e.g.,][]{2005astro.ph.10527S}, can break this degeneracy. Polarization data can also play this role.  In general, the matrix elements of {\\sf S} have different angular dependences. Therefore, observation of the intensity, which depends on $|S_1|^2+|S_2|^2$, and the polarization state of scattered light, which depends on $|S_1|^2-|S_2|^2$, can be used to recover this otherwise lost information.  Here we present the first optical polarization study of the AU Mic debris disk.  AU Mic is a nearby (9.9~pc) dM1e star with Galactic space velocities that suggest a common origin with $\\beta$ Pic. \\citep{1999ApJ...520L.123B}.  The discovery of scattered visible light from a near-edge-on debris disk around AU Mic \\citep{2004Sci...303.1990K} supports the picture of $\\beta$ Pic, AU Mic and nearly 20 other stars as a coeval group with age 12$^{+8}_{-4}$~Myr \\citep{2001ApJ...562L..87Z}.  The $\\beta$ Pic debris disk has been studied extensively, including the finding of polarization along the mid-plane varying between 0.12 and 0.21 in optical data \\citep{ 1991MNRAS.252P..50G, 1995Ap&SS.224..395W}.  These measurements have been interpreted by several authors \\citep[e.g.,][]{1997AREPS..25..175A, 1999A&A...352..508V, 2000ApJ...539..424K}.  In recent ground based observations, polarization in the $K$ band has been detected \\citep{2006ApJ...641.1172T}.  Despite the large difference in stellar mass, \\citet{2005astro.ph.10527S} and \\citet{2006astro.ph..4313A} argue that the radiation pressure force that quickly expels small $\\beta$ Pic grains has a counterpart around AU Mic in the form of stellar wind. To first order this explains the initial finding that the disk mid-plane surface brightness distribution is nearly identical for both debris disks \\citep{2004Sci...303.1990K}.  However, the blue color gradient for AU Mic's mid-plane \\citep{2005ApJ...622..451M, 2005AJ....129.1008K, fitz06} contrasts against the red color of $\\beta$ Pic's mid-plane \\citep{2006AJ....131.3109G}.  This distinct color difference between the two disks points to a significant divergence in either the grain composition, grain size distribution, or minimum grain size.  Polarization observations of AU Mic are therefore valuable in constraining these grain properties and identifying the fundamental differences between the two disks.  Our purpose is to present Advanced Camera for Surveys (ACS)/High Resolution Camera (HRC) polarimetry of the AU Mic system and to shed some initial light on the disk structure and grain optical properties that are emphasized by the detection and measurement of the disk in polarized light.  The fidelity of any model increases with the observational challenges presented to it, thus it must fit the near-IR and optical emission \\citep{2004Sci...303.1990K, 2004Sci...305.1442L, 2005ApJ...622..451M, fitz06} and the IR to sub-mm spectral energy distribution \\citep{2004ApJ...608..526L, 2005ApJ...634.1372C}.  We defer this synthesis to a later study.  Section \\ref{observations} outlines the observations, point spread function (PSF) subtraction, calibration of the Stokes parameters, and correction for instrumental and interstellar polarization.  Section \\ref{results} reports the appearance of the AU Mic disk in polarized light.  Section \\ref{onedprofiles} provides a qualitative description of the extracted one-dimensional surface brightness profile and degree of linear polarization, compares AU Mic with $\\beta$ Pic, and draws some preliminary conclusions regarding the radial distribution of the dust.  We describe quantitative analysis using optically-thin edge-on disk models in \\S \\ref{diskmodels} and develop a method that simultaneously fits the observed surface brightness and fractional polarization to three dust models: a semi-empirical Henyey-Greenstein model, a Zodiacal dust model, and a Mie model.  The successes and failures of these models are addressed in \\S \\ref{discussion} and the evidence for porous grains is examined.  Our conclusions are summarized in \\S \\ref{summary}. ",
        "conclusions": "\\label{discussion}  Figure \\ref{hggzodi} and the corresponding fit parameters in Table \\ref{model_table} demonstrate that a variety of radial grain distributions can account for the observed surface brightness, although a large inner hole with radius of 40--50 AU is common to all models.  Taken together, the surface brightness and degree of linear polarization narrow down the range of acceptable grain optical properties.  Inspection of Figure \\ref{phase_function} shows that a combination of strong forward scattering and a high polarizing efficiency, with a Rayleigh-like $\\sin^2\\theta/(1+\\cos^2\\theta)$ angular variation is sufficient to describe these data.  Our analysis implies with high confidence that the constituent dust grains exhibit high maximum linear polarization ($p_{max} \\ge 0.50$) and strong forward scattering ($g \\ge 0.7$).  Plots of the phase function and polarization fraction versus phase angle shown in Figure \\ref{phase_function} help highlight the common features of the statistically acceptable Henyey-Greenstein and low-index Mie models, and how their scattering properties differ from conventional rocky or ice-grain models. Large rocky or icy grains typically have strong forward scattering, but polarization which oscillates with phase angle. Therefore, $Q$ tends to average to zero along a line of sight that integrates over a range of scattering angles (refer to \\S \\ref{linpol}).  As the conventional Mie fit shows, small grains are good polarizers, but scatter too isotropically to be consistent with the data.  The low index Mie model suggests a physical scenario that combines strong forward scattering and high polarization.  Application of the Clausius-Mossotti relation \\citep{1962clel.book.....J} implies that such a low value of $n$ must be associated with a very porous medium---the only terrestrial analog that comes to mind is silica aerogel.  Aerogels are transparent, highly porous materials of low density, ranging from 0.05 to 0.15 g~cm$^{-3}$, with a corresponding refractive index of 1.01 to 1.04, respectively.  Aerogel has porosity on a micron scale and is composed of individual silica grains with diameters of $\\simeq$ 10 nm, which are linked in a highly porous dendritic backbone.  Although aerogel is produced in a process that is unlikely to occur in an astrophysical setting---hydrolysis of methyl silicate in the presence of a solvent (ethanol) that is subsequently evaporated at high temperature and pressure---the comparison is not entirely frivolous.  The appearance of aerogel is often characterized as ``solid blue smoke'', because to a good approximation the scattering is Rayleigh scattering \\citep{1993JQWIISpectroscRadiarTransfer50}.  Thus, aerogel is an example of a bulk material that interacts with electromagnetic radiation in a way that is determined by its microscopic structure.  Allowing the real part of the refractive index to vary as described in \\S \\ref{mie_theory} while holding the imaginary part fixed violates the Kramers-Kronig relation.  An approach that has a better physical basis is to use an effective medium theory to compute the optical behavior of a porous particle described as vacuum matrix ($n=1$) with embedded inclusions \\citep{2003pid..book.....K}.  Using the Maxwell-Garnett rule we can choose a refractive index for the bulk material and use the grain porosity as a model parameter.  The best-fit porous grain model (see Fig. \\ref{hggzodi} and the first line of Table \\ref{model_table}) is practically identical to the variable index fit, yielding essentially the same structural parameters.  The grain porosity is 91--94\\%, depending on whether we assume that the matrix from which our grains are made is ice or rock.  \\clearpage  \\onecolumngrid  \\begin{figure}[t] \\plottwo{f7a.eps}{f7b.eps} \\caption{Normalized phase function (top) and polarization fraction (bottom) versus scattering angle. The heavy solid line is inferred from the best fit single component Henyey-Greenstein model with scattering asymmetry parameter $g=0.68$.  Here polarization denotes $-Q/I$, thus negative polarization indicates that the electric field is oriented parallel to the scattering plane.  On the left are the results of Mie calculations for spheres with $x$ = 1, 2 and 3 with conventional $m = 1.65-0.01i$ (``silicate''). The best fit single particle size Mie fit (Table 1), has $x = 1.63\\pm 0.01$.  The inability of a sphere to simultaneously match the phase function and the polarization explains why this model can only achieve a poor fit. Small dielectric spheres with $x \\simeq 3$ can account for the strong forward scattering, but they cannot simultaneously provide a high degree of polarized light perpendicular to the scattering plane.  On the right is shown the result when the real part of the refractive index is allowed to be a free parameter.  Grains with $m = 1.033-0.01i$ and $x = 3.25$ provide a satisfactory fit to the data---at least as good as the one component Henyey-Greenstein model. The low index means that the grains must be extremely porous ($\\ge 90$\\%). } \\label{phase_function} \\end{figure}  \\twocolumngrid      As our aerogel analogy reminds us, porous materials are likely highly anisotropic, and we may not be free to assume that we can neglect the effects of nearest neighbors within the matrix.  We therefore examine the results of numerical calculations (e.g., discrete-dipole approximation and transition-matrix) of light scattering by aggregates to understand whether or not our interpretation of the Mie results in terms of porous grains is credible.     A lucid exposition of the transition-matrix method applied to composite interstellar grains is given by \\citet{2004ApJ...615..286I}. Using this method \\citet {2000Icar..148..526P} show results for two instances of silicate ($m = 1.65 -0.01i$) grain clusters consisting of 31 particles or ``monomers'' each with $x_m = 1.5$.  A relatively compact aggregate with approximately 70\\% porosity has $g = 0.75$ and $p_{max} = 0.52$.  The more porous particle (81\\%) has $g = 0.75$ and $p_{max} = 0.65$.  In neither case does the degree of polarization oscillate with phase angle.  These clusters have optical properties which make them promising analogs of the material inferred to dominate the AU Mic disk.    \\citet{2006A&A...449.1243K} present additional results for larger, more porous aggregates.  Figure \\ref{aggregate_vs_mie} shows the phase function and polarization for a large ($x_c = 10.2$) porous (90\\%) silicate cluster ($m=1.6-0.01i$) composed of 128 small ($x_m=0.9$) monomers.  This particle has $g = 0.84$ and $p_{max} = 0.82$, and has optical properties which make it an excellent candidate material for the AU Mic disk.  Also shown is the corresponding Mie calculation with dielectric properties derived using the Clausius-Mossotti relation.  It is evident that the Mie calculation is only a rough approximation---the polarization curve is reasonably well reproduced, and $g$ is overestimated by about 15\\%. Although the qualitative conclusion that implicates porous grains is secure it seems unlikely that the accuracy of effective medium theories is sufficient, for example, to distinguish between different coagulation schemes that are characterized by different porosity.  It will be necessary to abandon Mie theory in favor of numerical modeling of aggregate scattering in the next stage of debris disk modeling.     If highly porous aggregate grains explain the polarization signature of AU Mic's debris disk, and if the dust beyond $r_1\\simeq 40$ AU originates from a ``birth ring'' of parent bodies $\\lesssim $10 cm in size, as envisioned by \\citet{2005astro.ph.10527S}, then this porosity may be a signature of the agglomeration process whereby interstellar grains first grew into macroscopic sized objects.  In the inner Solar System porous particles occur naturally in cometary dust, where the sublimation of ices leaves a ``bird's nest'' of refractory organic and silicate material \\citep{1990ApJ...361..260G}. Porous grains in the $\\beta$ Pic disk may originate from cometary activity \\citep{1997A&A...323..566L}. However, the birth ring in AU Mic lies safely outside the $\\sim$1 AU ice sublimation point.  Based on collisional lifetime arguments, the size of the parent-bodies that supply the observed dust in AU Mic is in the decimeter range \\citep{2005astro.ph.10527S}.  Though the existence of larger bodies that will suffer compaction and restructuring \\citep{2000Icar..143..138B} is not excluded, they are not the dominant reservoir for dust observed at optical and near-infrared wavelengths. Evidently, shock compression during attrition of the parent bodies in the birth ring is not significant. We envision these bodies as so weakly bound that even the most glancing collisions lead to their disruption.  Recent laboratory studies of particle coagulation in the proto-solar nebula by ballistic cluster-cluster aggregation \\citep{1998Icar..132..125W} leads to the formation of highly ($>$90\\%) porous aggregates.  Our evidence suggests that such a process may have mediated the initial growth of planetesimals.  Observations of scattered light at a single wavelength are primarily sensitive to grains with $x \\simeq 1$, and do not place strong constraints on the particle size distribution.  However, preliminary calculations show that the measured polarization is consistent with a Dohnanyi spectrum (\\citet{fitz06} and L. E. Strubbe 2006, private communication).  Polarization measurements in the UV through the near-infrared could be used to measure grain porosity as a function of grain size.  \\begin{figure} \\plotone{f8.eps} \\caption{The phase function and polarization (dots) for a large ($x = 10.2$) porous (90\\%) silicate particle ($m=1.6-0.01i$) composed of 128 small ($x=0.9$) monomers from \\citet{2006A&A...449.1243K}.  This particle has $g = 0.84$ and $p_{max} = 0.82$.  The dashed line is the phase function ($g=0.68$) and polarization ($p_{max} = 0.53$) for the Henyey-Greenstein model which best fits the AU Mic disk. The porous grain is more forward scattering and more polarizing than required by the AU Mic data, but nonetheless its optical properties show that highly porous aggregates constitute an excellent candidate for the AU Mic disk.  Also shown is the Mie approximation for the porous grain with dielectric properties derived using the Clausius-Mossotti relation ($m = 1.047-0.0007i$).  The Mie calculation is a useful first approximation to the polarization, but overestimates $g$ by about 15\\%.  } \\label{aggregate_vs_mie} \\end{figure}"
    },
    "0609/astro-ph0609618_arXiv.txt": {
        "abstract": "We present a model of the stellar populations in the fields observed by one of the SuperWASP-N cameras in the 2004 observing season. We use the Besan\\c con Galactic model to define the range of stellar types and metallicities present, and populate these objects with transiting extra-solar planets using the metallicity relation of \\cite{F&V05}.   We investigate the ability of SuperWASP to detect these planets in the presence of realistic levels of correlated systematic noise (`red noise'). We find that the number of planets that transit with a signal-to-noise ratio of 10 or more increases linearly with the number of nights of observations. Based on a simulation of detection rates across 20 fields observed by one camera, we predict that a total of $18.6 \\pm 8.0$ planets should be detectable from the SuperWASP-N 2004 data alone. The best way to limit the impact of co-variant noise and increase the number of detectable planets is to boost the signal-to-noise ratio, by increasing the number of observed transits for each candidate transiting planet. This requires the observing baseline to be increased, by spending a second observing season monitoring the same fields. ",
        "introduction": "Since the discovery of 51 Peg b by \\cite{M&Q95}, a total of 181 extra-solar planets around 155 stars (as at 2006 May) have been discovered\\footnote{See the {\\it Extrasolar Planets Encyclopaedia} at http://exoplanet.eu/catalog.php}, almost exclusively by radial velocity measurements. It was only with the discovery that HD209458b transits its host star \\citep{Charbonneau-etal00} that both the mass and radius of an extra-solar planet were first determined. A further eight transiting extra-solar planets have been discovered in the last few years. Most of these planets have been discovered with the OGLE transit survey (e.g. \\citealt{Konacki-etal03}), which is a deep-field survey. SuperWASP (Wide Angle Search for Planets) is a wide-field transit survey, which is capable of detecting planetary transits of stars that are bright enough for high precision radial velocity follow-up observations.  Transit surveys have the potential to find numerous hot Jupiter-like planets (HJs), that is Jupiter-sized planets which orbit close to their host star with periods of just a few days. The probability that a HJ system is aligned such that transits will occur is about 0.1; which is much more favourable than planets orbiting at greater distances (for an Earth-like orbit, the probability of transit alignment is 0.0046).  Crucially, the detection of transiting planets allows the determination of the average density of the planet. This is of interest for planet formation models, particularly since HD209458b appears to have an anomalously low density \\citep{Charbonneau-etal00}, which is not well understood.   Previous attempts have been made to estimate the expected detection rates of transiting HJs by shallow, wide-field transit surveys similar to SuperWASP (e.g. \\citealt{Brown03}). Based on an observing pattern consisting of 38 nights of observations spread over 91 days and a requirement that three or more transits are observed, Brown calculated the detection rate of HJs producing a transit of depth 1 per cent or greater to be 0.39 per $10^4$ stars. Brown also estimates that, for the same observing window function, 4.51 false alarms will be detected per $10^4$ stars -- indicating that only eight per cent of transit signals detected will be produced by planets. These planetary transit `impostors' are grazing eclipsing binaries and eclipsing binaries diluted by light from a third star (either a field star or the third member of a triple system).   The SuperWASP survey has the potential to define the population of extra-solar planets which transit nearby bright stars ($V < 13$). In order for the results of SuperWASP to be properly interpreted, it is essential that the selection effects that operate in the survey are well understood. In this work, we use the findings of \\cite{Pont06} and Pont, Zucker \\& Queloz (2006) to estimate SuperWASP's detection rate in the presence of realistic levels of systematic red noise. We find that in order to detect a significant number of transiting planets, the existence of red noise necessitates much longer observing baselines than previously thought. ",
        "conclusions": "We conclude that there is a significant component of systematic, red noise present in data from SuperWASP-N. The {\\sc SysRem} algorithm of \\cite{Tamuz-etal05} appears highly effective at reducing the level of red noise, but fails to eliminate it entirely. The remaining red noise is present in the data at a level of about 3~mmag on time-scales of 2.5 hours, roughly equivalent to a typical transit duration time. This remaining noise has a significant impact on the efficacy of planet detection, as demonstrated by a Monte Carlo simulation based on the Besan\\c con Galaxy model.  Our analysis reveals that if observations are conducted for only 51 nights, none of the simulated transiting planets produces a transit detection with a signal-to-noise ratio of 10 or more. A total of $3.72 \\pm 1.60$ planets which transit with $\\rmsub{S}{red}~\\ge~10$ are predicted for the fields observed in 2004 by one SuperWASP-N camera, which is representative of the other four cameras, so $18.6 \\pm 8.0$ planets are predicted in total. In order to improve the $\\rmsub{S}{red}$, and thus increase the number of detectable planets, a greater number of transits must be observed in the data set of a particular object. This requires observations  to be made over a longer time period.  On the basis of the transit detection rates predicted here, the SuperWASP consortium have decided to continue observing all the fields that were monitored during 2004. We expect that this will enhance greatly the number of planetary transit events detected at non-pathological periods."
    },
    "0609/astro-ph0609175.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory   \\authorrunning{Annibali et al.} \\titlerunning{Nearby early--type galaxies with ionized gas.III.}   \\abstract {} {The paper is devoted to the study  of the underlying stellar population of a sample of 65 nearby early-type galaxies, predominantly located in low density environments, a large fraction of which show emission lines. } { Ages, metallicities, and [$\\alpha$/Fe] ratios have been derived through the comparison of Lick indices measured at different galacto-centric distances (7 apertures and 4 gradients) with new simple stellar population (SSP) models that account for the presence of $\\alpha$/Fe--enhancement. The  SSPs cover a wide range of ages ($10^{9}-16 \\times 10^{9}$ yr), metallicities (0.0004 $\\le$ Z $\\le$0.05), and [$\\alpha$/Fe] ratios (0-0.8). To derive the stellar population parameters, we use an algorithm that provides,  together with the most likely solution in the (age, Z, [$\\alpha$/Fe]) space, the probability density function along the age-metallicity degeneracy.} {We derive a large spread in age, with SSP-equivalent ages ranging from a few to 15 Gyrs. Age does not show any significant trend with central velocity dispersion $\\sigma_c$, but E galaxies appear on average older than lenticulars. On the contrary, a significant increasing trend of metallicity and [$\\alpha$/Fe] with $\\sigma_c$ is observed, testifying that the chemical enrichment was more efficient and the duration of the star formation shorter in more massive galaxies. These latter two relations do not depend on galaxy morphological type.  We have also sought possible correlations with the local galaxy density $\\rho_{xyz}$, but neither metallicity nor $\\alpha$--enhancement show clear trends. However, we find that while low density environments (LDE) ($\\rho_{xyz} \\le 0.4$) contain very young objects (from 1 Gyr to 4 Gyr), none of the galaxies in the higher density environments (HDE) (40 \\% of galaxies with a measured density) is younger than 5 Gyrs. Considering the lack of environmental effect on the [$\\alpha$/Fe] ratio and the high value of [$\\alpha$/Fe] in some young massive objects, we argue that young galaxies in LDE are more likely due to recent rejuvenation episodes. By comparing the number of ``rejuvenated'' objects with the total number of galaxies in our sample, and by means of simple two-SSP component models, we estimate that, on average, the rejuvenation episodes do not involve more than 25 \\% of the total galaxy mass.  The good quality of the data also allow us to analyze the gradients of stellar populations. An average negative metallicity gradient $\\Delta \\log Z/\\Delta \\log (r/r_e) \\sim -0.21$ is firmly detected, while the age and $\\alpha$--enhancement spatial distributions within r$_e$/2 appear quite flat. These observations suggest that, in a given galaxy, the star formation proceeded on similar timescales all across the central r$_e$/2 region, but with an increasing efficiency toward the center.} {} ",
        "introduction": "This is the third paper of a series dedicated to the study of early-type galaxies with emission lines (Rampazzo et al. \\cite{Ram05}, hereafter Paper~I; Annibali et al. \\cite{Annibal05}, hereafter Paper~II). A recent work  by Falc\\' on-Barroso et al. (\\cite{Falcon06}) (see also Sarzi et al. \\cite{Sarzi06}) based on SAURON observations shows that the incidence of ionized-gas emission in early-type galaxies is about 75\\%. By morphological type, lenticular galaxies display a slightly higher content of ionized gas than elliptical galaxies (83\\% versus 66\\%), in good agreement with the pioneering survey performed by Macchetto et al. (\\cite{Mac96}), who found emission in 85\\% of the lenticulars and 68\\% of the ellipticals in their sample.  Falc\\' on-Barroso et al. (\\cite{Falcon06}) found similar percentages when their sample was divided according to environment (83\\% field, 66\\% cluster). Furthermore, the incidence of ionized gas shows no correlation with either luminosity or the presence of a bar in the galaxy. The above percentages suggest that the phenomenon of emission lines is quite common among early-type galaxies: early-type galaxies without emission lines are more exceptions rather than the rule in this class of galaxies. The first step to understanding the evolution of these galaxies and the nature of their ionized gas consists in studying the galaxy underlying stellar population. %In this context, with the aim at improving the %understanding of the evolution of these galaxies and the nature of the %ionized gas they contain, the first step of our study consists of the %analysis of the galaxy underlying stellar population.  For this purpose we make use of the of absorption-line indices that have been proven to be a powerful tool to disentangle age and metallicity  effects (e.g., Buzzoni et al.~\\cite{Buz92}; Worthey~\\cite{Wor92}; Gonz\\' alez \\cite{G93} (G93); Worthey~\\cite{Wo94}; Buzzoni et al.~\\cite{Buz94}; Worthey et al.~\\cite{W94}; Bressan et al. \\cite{Br96}; Leonardi \\& Rose~\\cite {LR96}; Worthey \\& Ottaviani~\\cite{WO97}; Trager et al.~\\cite{Tra98}; Longhetti et al.~\\cite{L98a}; Vazdekis~\\cite{Vaz99}; Longhetti et al.~\\cite{L99}; Longhetti et al.~\\cite{L00}; Trager et al.~\\cite{TFWG00}; Kuntschner ~\\cite{Ku00}; Beuing et al.~\\cite{Beu02}; Kuntschner et al.~\\cite{Ku02}; Mehlert et al.~\\cite{Mehl03}). More recently the interpretation has been faced with the more difficult task of also accounting for the presence of non--solar scaled abundance patterns in early-type galaxies, and of deriving age, metallicity, and [$\\alpha$/Fe] ratio at the same time (e.g., Tantalo et al. \\cite{Tan96}; Tantalo \\& Chiosi \\cite{TC04a}; Thomas et al. \\cite{Thom03}; Thomas et al. \\cite{Thom05}). Indeed, it has been shown that the $\\alpha$/Fe element ratio plays a fundamental role in deriving the formation epochs. This ratio, in fact, quantifying the relative importance of Type II and Type Ia supernovae in the enrichment of the ISM,  carries information about the timescale over which the star formation occurred (Greggio \\& Renzini \\cite{GR83}; Matteucci \\& Greggio \\cite{MG86}; Pagel \\& Tautvaisiene \\cite{PT95}; Thomas et al. \\cite{Thom98}).  In this paper we try to constrain the epoch of the baryon assembly of our sample of 65 nearby early-type galaxies predominantly located in low density environments, a large fraction of which show emission lines (Paper I and II), and seek possible correlations with their morphological type (elliptical and lenticular), velocity dispersion, and environmental density. A large part of this paper is dedicated to the development of our own $\\alpha$--enhanced models and to the comparison with the models already presented in the literature. The stellar population parameters are derived by means of a new algorithm based on the probability density function. The algorithm provides,  together with the most probable solution in the (age, Z, [$\\alpha$/Fe]) space, the solutions along the age-metallicity degeneracy that are within 1~$\\sigma$  error from the observed index values. Our analysis is based on the use of three different line--strength index diagnostics, including the Mgb index, the $<$Fe$>$ index, and a Balmer line (H$\\beta$, H$\\gamma$, or H$\\delta$). The use of higher order Balmer lines allows us to minimize spurious effects that may arise from uncertainties in the emission correction of the H$\\beta$ absorption line. The paper is organized as follows. Section~2 presents an overview of the sample, summarizing its main properties and focusing on the delicate step of the correction for emission infilling of the most commonly used age indicator, the H$\\beta$ absorption line. In Sect. 3 a complete description of the procedure adopted to compute the new $\\alpha$--enhanced SSPs is provided. Section~4 describes how the new SSP models have been implemented within an algorithm that allows us to derive the stellar population parameters, i.e., age, metallicity, and $\\alpha$/Fe ratio. Section~5 presents the results. In particular we analyze both the stellar population parameters derived for the apertures and gradients, and their relations with velocity dispersion and local galaxy density. In Sect. 6 we discuss our results and compare them with the literature. Finally, a summary and our conclusions are presented in Sect. 7. ",
        "conclusions": "We have analyzed the stellar populations of the sample of 65 early--type galaxies mainly located in the field, presented in Papers~I and II. Our sample is biased towards early-type galaxies selected on the basis of ISM traces, in particular optical emission lines, tracing the warm ISM component. On the other hand, considering the incidence of galaxies showing a warm ISM in a randomly selected sample of early-type galaxies quoted by Falcon-Barroso et al. (2006), i.e., 75 \\%, (83\\% in the field), we are confident that our sample is not very dissimilar from those with which we compare our results.  Stellar population parameters age, metallicity, and [$\\alpha$/Fe] ratio have been derived for all the apertures and gradients of the sample by the comparison of the Lick indices with new SSP models, presented in this work,  which account for the presence of non-solar element abundance patterns. New SSP models have been developed for a wide range of ages ($10^{9}-16 \\times 10^{9}$ yr), metallicities (Z=0.0004, 0.004, 0.008, 0.02, 0.05), and [$\\alpha$/Fe] ratios (0-0.8). The SSPs are based on the Padova stellar evolution tracks and accompanying isochrones (Bertelli et al. \\cite{Ber94}; Bressan et al. \\cite{Bre94}), which have been previously calibrated on color-magnitude diagrams of several star clusters in the Galaxy and in the Magellanic Clouds.  The standard composition line-strength indices are computed on the basis of the fitting functions of Worthey et al. (\\cite{W94}) and Worthey \\& Ottaviani (\\cite{WO97}). The enhanced SSPs are computed by applying an index correction, which we derive following the main guidelines provided by previous works in the literature (TMB03, Tantalo \\& Chiosi \\cite{TC04a}), although with revised index dependence on the element abundance, to the standard SSPs. More specifically, we find that a logarithmic dependence of the index with element abundance provides a more satisfactory description of the index behavior. This evidence is also supported by preliminary results obtained running new model atmospheres and synthetic spectra with the code ATLAS12, that consistently computes both interior and emergent spectra models with arbitrary abundances. The correction for the effects of $\\alpha$-enhancement is computed on the basis of the most recent specific index responses computed by Korn et al. (\\cite{K05}) for different metallicities.  The new $\\alpha$-enhanced SSPs have been implemented within an algorithm devised to derive stellar population parameters from the observed line--strength indices. The method, which rests on the probability density function, provides, together with the most probable solution in the (age, Z, [$\\alpha$/Fe]) space, the solutions along the age-metallicity degeneracy that are within 1~$\\sigma$ error from the observed index values.  Summarizing our results:  \\begin{itemize}  \\item The derived age distribution shows a wide spread, with SSP-equivalent ages from a few Gyr to 15 Gyr and with lenticular galaxies being generally younger than ellipticals (the average ages for the whole sample, E and S0 are 8, 8.7, and 6.3 Gyr, respectively); the metallicity distribution shows a broad peak at $0<[Z/H]<0.3$; finally, the [$\\alpha$/Fe] ratio definitely presents a peak at $\\sim$0.22, with a narrower distribution than for age and metallicity.  \\end{itemize}  We also seek possible correlation of the stellar population parameters with both central velocity dispersion and local galaxy density.  \\begin{itemize}  \\item Concerning the relation with the galaxy potential well, we do not find the clear signature of a global trend of age with velocity dispersion; on the other hand, significant positive trends are derived for metallicity and $\\alpha$/Fe enhancement.  \\item The robust metallicity-$\\sigma_c$ relation testifies that chemical enrichment is more efficient in more massive galaxies and,  within a universal initial mass function, the relation between the $\\alpha$/Fe enhancement and the velocity dispersion indicates that the overall duration of the star formation process is shorter within deeper potential wells. These two relations do not depend on galaxy morphological type, unlike the age-$\\sigma_c$ relation, indicating that the galaxy gravitational potential is the main driver of the chemical enrichment process of the galaxy.  \\end{itemize}  We analyzed the relations of the stellar population parameters with the local galaxy density $\\rho_{xyz}$ and we obtained the following results:  \\begin{itemize}  \\item Concerning the age--$\\rho_{xyz}$ relation, we derive a clear dichotomy: while very young objects (from 1 Gyr to 4 Gyr, some of which have high $\\sigma_c$) are found  in very low density environments ($\\rho_{xyz} \\le 0.4$), none of the galaxies in high density regions (40 \\% of the sample with measured density) is younger than 4-5 Gyr. No trend of the chemical enrichment path is found with environmental conditions. The lack of environmental effect on the ($\\alpha$-enhancement)--$\\sigma_c$  relation indicates that in very low environments rejuvenation episodes, rather than more prolonged star formation, are frequent. \\end{itemize}  We analyzed the stellar population gradients within the single galaxies:  \\begin{itemize}  \\item We definitely derive negative metallicity gradients (metallicity decreases from the central regions outwards) of the order of $\\Delta \\log Z/\\Delta \\log (r/r_e) \\sim -0.21$, consistently with literature values. On average the $\\alpha$-enhancement remains quite flat within r$_e$/2, but this is consistent with the observed metallicity gradient. The only strong correlation is that of the metallicity gradient with velocity dispersion: larger negative gradients are derived for less massive systems and become shallower at increasing galaxy mass.   \\end{itemize}  From the above results, we try to sketch a possible scenario for galaxy formation. The potential well of the galaxy seems to be the  main driver of the chemical path during its formation, with more massive galaxies exhibiting the more efficient chemical enrichment and shorter star formation timescale. Furthermore, galaxies in denser environments are on average older than galaxies in less dense environments. This is in agreement with recent models of galaxy formation that combine feedback from QSO and SN, and cooling of the gas (Granato et al. \\cite{G01}, \\cite{G04}).  Galaxies in less dense environments also show signatures of very recent rejuvenation episodes. By comparing the number of ``young'' objects with the total number of galaxies, and by means of simple two-SSP component models, we estimate that in these rejuvenation episodes (like major mergers), not more than 25\\% of the galaxy mass could be formed. Thus we estimate that, even if spheroidal galaxies are  affected by major mergers during their evolution, about 75\\% of the galaxy mass is formed during the initial epoch of formation."
    },
    "0609/astro-ph0609643_arXiv.txt": {
        "abstract": "We present the first comparison of profiles of H et He resonance lines observed by SUMER with theoretical profiles computed with our non-LTE radiative transfer code. We use the \\ion{H}{i} Lyman~$\\beta$, \\ion{H}{i} Lyman~$\\epsilon$, and \\ion{He}{i} $\\lambda$\\,584~\\AA\\ lines. Our code allows us to obtain the plasma parameters in prominences in conjunction with a multi-line, multi-element set of observations. The plasma temperature in the prominence core is $\\sim 8600$~K and the pressure is 0.03~dyn cm$^{-2}$. The Ly$\\beta$ line is formed in a higher temperature region (more than 11\\,000~K). ",
        "introduction": "The H and He resonance lines have a strong diagnostic potential. However, observations are difficult to interpret as the prominence plasma is optically thick at these wavelengths and out of LTE. It is thus necessary to use a non-LTE radiative transfer code to analyse the data. ",
        "conclusions": "The plasma temperature in the core of the observed prominence is found to be $\\sim 8600$~K and the pressure is 0.03~dyn cm$^{-2}$. The Ly$\\beta$ line is formed in a higher temperature region (more than 11\\,000~K). The \\ion{He}{i} line at 584 \\AA\\ has been succesfully used for the diagnostic of the prominence. The combination of the helium resonance lines with the Lyman lines of hydrogen provides new constraints for prominence models."
    },
    "0609/astro-ph0609369_arXiv.txt": {
        "abstract": "Using small automated telescopes in Arizona and Hawaii, the HATNet project has detected an object transiting one member of the double star system \\ads. This system is a pair of G0 main-sequence stars with age about 3 Gyr at a distance of $\\sim$139 pc and projected separation of $\\sim$1550 AU.  The transit signal has a period of 4.46529 days and depth of 0.015 mag. From follow-up photometry and spectroscopy, we find that the object is a ``hot Jupiter'' planet with mass about 0.53 \\mjup\\ and radius $\\sim$1.36 \\rjup\\ traveling in an orbit with semimajor axis 0.055 AU and inclination about 85\\fdg9, thus transiting the star at impact parameter 0.74 of the stellar radius. Based on a data set spanning three years, ephemerides for the transit center are: $T_C = 2453984.397 + N_{tr} * 4.46529$. The planet, designated \\hatp, appears to be at least as large in radius, and smaller in mean density, than any previously-known planet. ",
        "introduction": "\\label{sec:intro} The 10 currently known transiting extrasolar planets occupy a special place among the nearly 200 known exoplanets because their photometric transits yield unambiguous information on their masses and radii. Five of these transiting planets were found through spectroscopic and photometric follow-up of candidates discovered by the OGLE project \\citep[these are OGLE-TR-10b, 56b, 111b, 113b, and 132b; ] [respectively]{udalski02a,udalski02b,udalski02c,udalski03,udalski04}. For references to the follow-up confirmation papers for these objects, see the Extrasolar Planets Encyclopedia\\footnote{http://exoplanet.eu/}. The parent stars of these planets are all rather faint, making effective followup somewhat difficult. The other five\\footnote{Discovery of {\\mbox TrES-2} \\citep{odonovan06} was announced during the refereeing of this paper.} transiting exoplanets orbit nearby, bright stars: {\\mbox HD 209458b} \\--- \\citet{dc00,henry00}, {\\mbox TrES-1} \\--- \\citet{alonso04}, {\\mbox HD 149026b} \\--- \\citet{sato05}, {\\mbox HD 189733b} \\--- \\citet{bouchy05}, {\\mbox XO-1b} \\--- \\citet{mccullough06}. These planets are of special interest, as accurate parameter determination as well as other types of follow-up are facilitated.  For this reason, a number of wide-field surveys, using small telescopes, are underway \\citep[see, e.g. ][]{dc06b}. Indeed, two of the 5 known transiting planets around nearby stars, {\\mbox TrES-1} and {\\mbox XO-1b}, were detected by such surveys. (The other three were first detected by radial velocity measurements).  The HATNet project,\\footnote{www.hatnet.hu} initiated in 2003 by G.~\\'A.~B, is also a wide-field survey that aims for the discovery of transiting planets around bright stars. It currently comprises 6 small wide-field automated telescopes, each of which monitors $8\\arcdeg\\times8\\arcdeg$ of sky, typically containing 5000 stars bright enough to permit detection of planetary transits through the typically 1\\% photometric dips they induce on their parent stars. The instruments are spread in a two-station, longitude-distributed network, with four telescopes at the F.~L.~Whipple Observatory (FLWO) in Arizona, and two telescopes at the Submillimeter Array (SMA) at Hawaii. Technical aspects of HATNet will be described in a forthcoming paper, but for instrumentation, observations, and data flow, see \\citet{bakos02,bakos04}. Here we report on the first detection by HATNet of a transiting extrasolar planet. This planet, which we hereafter label as \\hatp, orbits a 10th magnitude star, and thus is well-suited for follow-up with large ground-based and space telescopes.  \\notetoeditor{This is the intended place of \\figr{stamps}}  \\ifthenelse{\\boolean{emulateapj}}{\\begin{figure*}[t]}{\\begin{figure}[t]} \\ifpdf \\plottwo{f1a.pdf}{f1b.pdf} \\else \\plottwo{f1a.eps}{f1b.eps} \\fi \\caption{ The left panel shows a $70\\arcmin\\times70\\arcmin$ (300 pixels wide) stamp of a typical image taken by HAT-5. The stamp width is about 1/7th of the original frame. \\ads\\ is in the middle (the two components are merged into a bright star). Peaks of saturated stars are masked, and shown in white.  The right panel is a stamp from an image taken with the FLWO \\tsize{1.2} telescope. It is $6\\farcm8\\times6\\farcm8$ (600 pixels wide), and corresponds to the small box on the left panel. \\adsb (marked with the cross) is well separated from \\adsa. \\label{fig:stamps}} \\ifthenelse{\\boolean{emulateapj}} {\\end{figure*}}{\\end{figure}}  The existence of \\hatp\\ was first inferred from a 0.6\\% (6 mmag) deep transit-like dip seen in the $I$-band \\lc\\ of the stellar system \\ads\\ \\citep[Aitken Double Star Catalogue; ][]{aitken32} obtained by the combination of the HAT-5 instrument at FLWO and HAT-8 at the SMA. The members of this binary stellar system are a pair of nearly identical G0V stars, \\adsa\\ ($I\\approx9$) and \\adsb\\ ($I\\approx9.6$) separated by 11\\farcs2.  The system appears as a single, elongated source in the HAT images because this separation is less than the 14\\pxs\\ image resolution.  Initial follow-up spectroscopy at the FLWO \\tsize{1.5} telescope showed the absence of radial velocity variations of either star at a level of 1 \\kms\\ or larger, demonstrating that the source of the photometric dips could not be due to a common false positive, namely an M dwarf star transiting in front of one of the stars. Subsequent spectroscopy at the Subaru telescope using N2K project observing time \\citep{fischer04,sato05}, and later at the W.~M.~Keck telescope revealed sinusoidal velocity variations of \\adsb\\ with an amplitude of 60\\ms, thus suggesting the existence of a 0.5\\mjup\\ planet transiting that star. After excluding blend scenarios, the combination of the photometric and spectroscopic data leads to the conclusion that this object, \\hatp, is indeed a planet, and has a radius comparable to that of {\\mbox HD 209458b}, and a somewhat lower mean density. In the remainder of this paper we discuss all of the above points in more detail. ",
        "conclusions": "\\label{sec:discus}  \\figr{exoplanet_m_r} shows the location of \\hatp\\ in a plot of planetary radius versus planetary mass for the now 11 known transiting exoplanets, plus Jupiter and Saturn.  It is immediately seen that \\hatp\\ is in an extreme position. It has a radius apparently as large as or larger than the largest known exoplanet, \\mbox{HD 209458b}, and at the same time it has a mass that is significantly smaller.  Thus, it has a lower mean density than any other known planet. Until now, \\mbox{HD 209458b} was the only significantly anomalous transiting planet. Thus, \\citet{guillot05} noted that all of the transiting planets known at that time, with the exception of \\mbox{HD 209458b}, fall within a mass-radius relation derived from modeling the interiors of strongly irradiated planets. All of the transiting planets detected subsequently, except \\hatp, also fit the same relation. Among them, OGLE-TR-10b had been thought to be similarly anomalous \\citep{konacki05}, but recent evidence \\citep[][M.~Holman, private communication]{holman05,santos06} points to a smaller planetary radius and normal density. Similarly, the XO-1b discovery paper \\citep{mccullough06} derives a larger radius than the value later refined by \\citep{holman06}.  \\notetoeditor{This is the intended place of \\figr{exoplanet_m_r}} \\begin{figure}[h] \\ifpdf \\plotone{f6.pdf} \\else \\plotone{f6.eps} \\fi  \\caption{ Mass-radius diagram depicting known transiting exoplanets, plus Saturn and Jupiter (for comparison). Sources for the mass and radius values (in this order) are listed after the name of the planet; \\mbox{\\hatp}: this work, \\mbox{HD 209458b}: \\citet{laughlin05}, \\citet{knutson06}, \\mbox{HD 189733b}: \\citet{bakos06}, \\mbox{OGLE-TR-111b}: \\citet{pont04}, \\mbox{OGLE-TR-10b}: \\citet{santos06}, \\mbox{OGLE-TR-132b}:  \\citet{moutou04}, \\mbox{OGLE-TR-56b}: \\citet{torres04a}, \\mbox{OGLE-TR-113b}: \\citet{bouchy04,konacki04}, \\mbox{TrES-1}: \\citet{sozetti04}, \\mbox{HD 149026}: \\citet{sato05}, \\citet{dc06a}, \\mbox{XO-1b}: \\citet{holman06}, \\mbox{TrES-2}: \\citet{odonovan06}. \\label{fig:exoplanet_m_r}} \\end{figure}  \\mbox{HD 209458b} has a radius some 10 to 20 percent larger than expected, given its age, mass, and expected temperature including effects of irradiation. Possible solutions include: (a) a fraction of a percent of the stellar flux irradiating \\mbox{HD 209458b} could be transported to deep layers of order 100 bar, and thereby keep the planet's radius puffed up \\citep{showman02}; (b) there is a small eccentricity of its orbit, forced by the presence of a so-far undetected planetary companion and leading to internal tidal heating \\citep{bodenheimer03}, and (c) HD 209458b might be in a Cassini State---a dynamical state where its rotation axis lies close to the plane of its orbit \\citep{winn05a}, leading to large obliquity tides that cause heating at significant depth leading to the planet's large radius. A natural question about the first of these mechanisms is why all close-in ``hot Jupiters'' are not similarly bloated by that mechanism. A problem with the last two mechanisms is that both are {\\em ad hoc}; they require unusual circumstances for which there is no independent evidence.  With the discovery of another planet, \\hatp, which has a radius similar to and probably even larger than \\mbox{HD 209458b}, and an even smaller density, mechanisms requiring unusual circumstances seem less likely.  Moreover, for \\mbox{HD 209458b}, explanation (b) appears ruled out because Doppler data \\citep{laughlin05b,winn05b} and transit timing variations data \\citep{miller05} have excluded companion planets capable of enacting it.  In the case of \\hatp, it remains possible that the system has a second planet capable of forcing its eccentricity and leading to internal tidal heating. Our current orbital solution (\\secr{rvorbit}) suggests $e=0.09 \\pm 0.02$. If confirmed, such eccentricity would be a strong indication for the presence of a third body (second planet), because otherwise tidal dissipation would have circularized \\hatp's orbit in $\\sim 225 (\\frac {Q}{10^6})$~Myr. The dissipation factor, $Q$, is poorly known for ``hot Jupiters'', but is expected to lie between the values for Jupiter and solar-like stars ($10^5$ to $10^6$). With forced orbital eccentricity, $e=0.09$, \\hatp\\ will be subject to internal tidal heating due to an energy dissipation rate \\begin{displaymath} \\frac{dE}{dt} = \\ \\frac{63}{4}\\frac{e^{2}n}{Q}\\left(\\frac{R_{p}}{a}\\right)^{5}\\frac{G\\mstar^{2}}{a} \\end{displaymath} where $G$ is the gravitational constant, $M_{\\star}$ is the mass of the star, $a$ is the semimajor axis, $n$ is the mean motion, and $Q$ is the specific dissipation factor. With $e = 0.09$, the rate of energy dissipation in \\hatp\\ would be $3.0\\times 10^{26} erg~s^{-1}$, corresponding to energy dissipation per unit Jupiter mass of $5.7 \\times 10^{26} erg~ s^{-1}\\mjup^{-1}$. This rate is almost identical to the dissipation rate per unit Jupiter mass of $5.8\\times 10^{26} erg~s^{-1} M_{J}^{-1}$ required to achieve a 20\\% increase in \\rpl\\ for \\mbox{HD 209458b} (no core model) and explain its anomalous low density \\citep{bodenheimer03}. Therefore we conclude that if the orbital eccentricity is confirmed, the anomalous \\rpl\\ and low density of \\hatp\\ could be easily explained by tidal heating. The perturbing second planet would be detectable by both transit timing variations and Doppler velocities.  For now, however, all that can be said is that bloated planets are not that unusual, and, since a second perturbing planet is lacking in HD 209458b's case, perhaps some other explanation should be sought that can explain such planets more generally."
    },
    "0609/astro-ph0609475_arXiv.txt": {
        "abstract": "{We report on XMM-Newton observations of \\sn\\ in the Large Magellanic Cloud.} {The large collecting area telescopes together with the European Photon Imaging Cameras (EPIC) provide X-ray spectra with unprecedented statistical quality and make it possible to investigate the spectral evolution during the brightening observed since the discovery in X-rays. High resolution spectra from the Reflection Grating Spectrometers yield a complementary view and allow us to perform more detailed investigations of prominent emission lines.} {The X-ray spectra were modeled with two-temperature emission components from a hot plasma in collisional ionization equilibrium and in non-equilibrium (NEI).} {We find a temperature for the equilibrium component of 0.24$\\pm$0.02 keV in January 2000 and April 2001 which increased to 0.30$\\pm$0.02 keV in May 2003 and also an indication for a temperature increase in the hot NEI component from $\\sim$2 keV to $\\sim$3 keV. Emission line ratios inferred from the RGS spectra suggest temperatures as low as 100 eV and an increase in the ionization state of oxygen and neon consistent with the observed temperature increases. The fast readout of the EPIC-pn instrument yields X-ray fluxes free of CCD pile-up effects which we used to normalize pile-up corrections for the published Chandra fluxes. The corrected X-ray light curve of \\sn\\ in the 0.5-2.0 keV energy band is best represented by a linear increase up to about day 4000 after the explosion and an exponential rise afterwards until the last published Chandra observation on day 6716. Modeling the light curve by emission from the inner ring which is approximated by a circular torus a central density n$\\sb{\\rm{H}}$~=~1.15$\\times$10$\\sp 4$~cm$\\sp{-3}$ is found. In this model the forward shock has just passed the center of the torus.} {\\sn\\ continues to brighten exponentially in soft X-rays. The X-ray spectra can be represented by pure thermal emission without significant contribution from a compact object yet.} ",
        "introduction": "The first supernova observed in 1987 (\\sn) went off in the Large Magellanic Cloud (LMC) on 23 February. This event with observations involving the entire electromagnetic spectrum and the first detection of cosmic neutrinos brought about a number of astrophysical surprises, some of which are puzzling still today. In this paper we concentrate on the observations of soft (0.1 - 3 keV) X-ray emission since the explosion, thought to arise from the interaction of shocks with the ambient circumstellar medium, and possibly a central compact object like a neutron star. We present and summarize for the first time the full set of previous data, which to some extent have been re-calibrated, and we add the XMM-Newton observations.  The first attempts to search for soft X-rays were made in August 1987 with sounding rockets carrying imaging telescopes. An upper limit of 1.5$\\times$10$\\sp{36}$ erg~s$\\sp{-1}$ could be established which excluded an ambient dense and slow wind typical for a red supergiant wind, but which is still compatible with a blue supergiant \\citep{1987Natur.330..232A}. X-rays from \\sn\\ became clearly visible in the ROSAT images not much before spring 1992 \\citep{1994A&A...281L..45B,1994ApJ...420L..25G}. The monitoring of \\sn\\ with ROSAT revealed a slow, but steady increase in flux over the first years \\citep{1996A&A...312L...9H}. The brightening continues since then but with increasing pace. \\citep{2000ApJ...543L.149B,2002ApJ...567..314P,2004ApJ...610..275P,2005ApJ...634L..73P}. Chandra has taken images of \\sn\\ since late 1999, which show a ring-like emission region steadily growing in diameter. Details of the ring structure became apparent in the meantime. These structural features are correlated with the optical hot spots, first discovered in earlier images taken with the Hubble Space Telescope (HST). These optically bright spots evolve quite rapidly and they now seem to completely engulf the so-called inner ring \\citep{2000ApJ...537L.123L,2005coex.conf...77M}.  The ring with its fairly high matter density existed before the explosion and was detected by re-emitted light of the prompt supernova EUV flash. A substantial heating, up to X-ray temperatures, of the ring was predicted for the period between 2002 and 2010 when the supernova shock wave would have reached and entered the ring. Both the optical and X-ray images and the X-ray light curve demonstrate that this state has been arrived at. Hope is that we can learn more about the details of the interaction process, the origin of the ring as well as of the progenitor star and its evolution by X-ray spectroscopy.  \\begin{table*} \\caption[]{XMM-Newton observations of \\sn\\ and data selection.} \\begin{center} \\begin{tabular}{llcclrrr} \\hline\\hline\\noalign{\\smallskip} \\multicolumn{1}{l}{Instrument} & \\multicolumn{1}{c}{Read-out} & \\multicolumn{1}{c}{Filter} & \\multicolumn{1}{c}{Sat. Revol.} & \\multicolumn{1}{c}{Date} & \\multicolumn{1}{c}{Time (UT)} & \\multicolumn{1}{c}{Net Exp.$^1$} & \\multicolumn{1}{c}{Counts$^2$} \\\\ \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{Mode} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{[s]} & \\multicolumn{1}{c}{} \\\\  \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} EPIC-pn      & FF, 73 ms & Medium &  21 & 2000 Jan. 19/20 & 16:19 -- 04:28 &  23250 &  2194 \\\\ EPIC-pn      & FF, 73 ms & Medium &  22 & 2000 Jan. 21    & 15:38 -- 19:34 &   5789 &   518 \\\\ EPIC-pn      & FF, 73 ms & Thin   &  22 & 2000 Jan. 21/22 & 20:32 -- 12:02 &  30990 &  3060 \\\\ \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} EPIC-pn      & FF, 73 ms & Medium & 244 & 2001 Apr. 8/9   & 20:35 -- 07:14 &  24534 &  3952 \\\\ \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} EPIC-pn      & SW, 6 ms  & Medium & 626 & 2003 May 10/11  & 16:41 -- 23:37 &  55855 & 25609 \\\\ EPIC-MOS1    & FF, 2.6 s & Medium &     &                 & 16:35 -- 23:37 & 101520 & 16020 \\\\ EPIC-MOS2    & FF, 2.6 s & Medium &     &                 & 16:35 -- 23:37 & 101510 & 16018 \\\\ RGS1         & Spectro   & --     &     &                 & 11:42 -- 23:38 & 112160 &  2509 \\\\ RGS2         & Spectro   & --     &     &                 & 10:47 -- 23:38 & 111670 &  3358 \\\\ \\noalign{\\smallskip}\\hline \\end{tabular} \\end{center}  $^1$Net exposures for the observations from revolutions 21/22 were derived from a comparison of the spectrum of N\\,157\\,B with that obtained from the ROSAT PSPC due to missing timing information in the early data \\citep{2001A&A...365L.208H}.\\\\ $^2$Net counts as used for the spectral analysis in the following energy bands: 0.2$-$9.0 keV for the EPIC instruments (except for EPIC-pn during revolutions 21/22 when we adopted 0.3$-$9.0 keV due to the higher low-energy threshold used at the beginning of the performance verification phase); 0.3$-$2.0 keV for RGS. \\label{xray-obs} \\end{table*}  The medium energy resolution X-ray spectrum measured by the Chandra ACIS detector is composed of two emission components with temperatures kT of $\\sim$2.4 keV and $\\sim$0.22 keV, which sometimes are interpreted as contributions from the blast wave shock and the decelerated slow shock, respectively \\citep{2002ApJ...574..166M,2004ApJ...610..275P}. The strong brightness increase made it also possible to obtain high resolution X-ray spectra with the Chandra LETG. A total exposure of 289 ks was accumulated between August 26 and September 5, 2004. Inferred temperatures from the emission line ratios range from 0.1 to 2 keV and expansion velocities are found between 300 and 1700 \\kms\\ \\citep{2005ApJ...628L.127Z}.  One other outstanding issue is the existence and the time of appearance of a central compact remnant like a neutron star. So far the results have been negative. The most stringent upper limits for an X-ray point source in the newly born supernova remnant are based on the analysis of Chandra images in the 2--10 keV band with \\lx\\ $<$ 1.5\\ergs{34} \\citep{2004ApJ...610..275P} and on INTEGRAL measurements in the 20--60 keV band \\citep[\\lx\\ $<$ 1.1\\ergs{36}, ][]{2005AstL...31..258S}. The HST derived optical limit is \\lopt\\ $<$ 5\\ergs{33} \\citep{2005ApJ...629..944G}.  XMM-Newton observed \\sn\\ on several occasions. Here, we report the results of the spectral analysis of the data with the highest statistical quality available to date, making use of the EPIC cameras and the RGS spectrometers. ",
        "conclusions": "We present the results from a spectral analysis of XMM-Newton data from \\sn\\ obtained in January 2000, April 2001 and May 2003. The spectra are consistent with thermal emission from the supernova remnant without a significant hard power-law component (with an absorption corrected 0.2-10.0 keV luminosity limit of 5.7\\ergs{34}) originating from a compact object. Modeling the spectra with two components involving emission from plasma in collisional ionization equilibrium and/or in non-equilibrium yields clear evidence for a temperature rise after the second epoch observation for both components, reaching $\\sim$0.3 keV and $\\sim$3 keV in the soft and hard component, respectively.  We have collected all available X-ray fluxes in the 0.5-2.0~keV band from ROSAT, XMM-Newton and Chandra for the light curve of \\sn\\ up to day 6716 after the explosion. This includes previously unpublished ROSAT data available in the archive and published Chandra fluxes which we corrected for photon pile-up effects present in the CCD spectra. The final light curve shows a linear increase up to about day 4000 and an exponential rise afterwards. We model the light curve with emission from a circular torus around the explosion site which represents the inner ring around \\sn. From the best fit we find that the forward shock just passed the center of the torus and that the mean shock velocity over 7000 days is 10400 \\kms. With a core density of the torus of 1.15$\\times$10$^4$ cm$^{-3}$ we find a total mass overrun by the forward shock of 0.45~M$_{\\odot}$. The model, assuming a constant blast wave velocity in its simplest form, predicts a flattening in the X-ray light curve for mid 2006. A new XMM-Newton observation is currently scheduled for the end of 2006 to further monitor the spectral evolution of \\sn."
    },
    "0609/astro-ph0609196_arXiv.txt": {
        "abstract": "We present medium-resolution VLT/FORS2 spectroscopy of six cataclysmic variables discovered by the Sloan Digital Sky Survey. We determine orbital periods for SDSS\\,J023322.61+005059.5 ($96.08\\pm0.09$\\,min), SDSS\\,J091127.36+084140.7 ($295.74\\pm0.22$\\,min), SDSS\\,J103533.02+055158.3 ($82.10\\pm0.09$\\,min), and SDSS\\,J121607.03+052013.9 (most likely $98.82\\pm0.16$\\,min, but the one-day aliases at 92\\,min and 107\\,min are also possible) using radial velocities measured from their H$\\alpha$ and H$\\beta$ emission lines. Three of the four orbital periods measured here are close to the observed 75--80\\,min minimum period for cataclysmic variables, indicating that the properties of the population of these objects discovered by the SDSS are substantially different to those of the cataclysmic variables found by other means. Additional photometry of SDSS\\,J023322.61+005059.5 reveals a periodicity of approximately 60\\,min which we interpret as the spin period of the white dwarf, suggesting that this system is an intermediate polar with a low accretion rate. SDSS\\,J103533.02+055158.3 has a period right at the observed minimum value, a spectrum dominated by the cool white dwarf primary star and exhibits deep eclipses, so is an excellent candidate for an accurate determination of the parameters of the system. The spectroscopic orbit of SDSS\\,J121607.03+052013.9 has a velocity amplitude of only $13.8 \\pm 1.6$\\kms, implying that this system has an extreme mass ratio. From several physical constraints we find that this object must contain either a high-mass white dwarf or a brown-dwarf-mass secondary component or both. ",
        "introduction": "\\label{sec:intro}  Cataclysmic variables (CVs) are interacting binary stars containing a white dwarf primary star and a low-mass secondary star in a close orbit. In the majority of systems the secondary star is unevolved, fills its Roche lobe and loses mass to the primary star via an accretion disc. Comprehensive reviews of the properties of CVs have been given by \\citet{Warner95book} and \\citet{Hellier01book}.  The evolution of CVs depends primarily on angular momentum loss (AML) from their orbit. CVs with long periods contain a secondary star which is only partially convective. It is thought that the presence of a magnetic field in this star causes AML by magnetic braking (\\citealt{VerbuntZwaan81aa}; \\citealt*{Rappaport++82apj}) until the orbital period decreases to about three hours. At this point the secondary star becomes fully convective and magnetic braking ceases. This change in the state of the star causes it to relax to its thermal equilibrium radius and so shrink inside its Roche lobe. Mass transfer ceases and weaker AML is now dominated by gravitational radiation \\citep{Paczynski67aca}. Once the period reaches about two hours, the shrinking Roche lobe re-establishes contact with the secondary star, restarting the mass transfer. The period continues to decrease to an observed minimum of approximately 80 minutes, at which point the secondary star becomes a degenerate brown-dwarf like object. As the mass donor is now degenerate, further mass transfer causes the period to increase \\citep{Patterson98pasp}.  Systems beyond this ''period bounce'' are theoretically predicted to represent up to 70\\% of the steady-state population of CVs; population synthesis models also predict that the vast majority (99\\%) of CVs have evolved to periods shorter than the 2--3 hour gap \\citep{Dekool92aa,DekoolRitter93aa,Kolb93aa,Politano96apj,KolbBaraffe99mn}. The models also predict that the low rate of change of period near the minimum value causes an increase in the population of CVs there, which should manifest itself as a spike in the period distribution. However, whilst several candidates exist, not one CV has yet been confirmed to have a secondary component with a mass appropriate for a period bounce CV \\citep*{Littlefair++03mn,Patterson++05pasp}. In addition, the observed period distribution of CVs does not have a spike at the minimum period and it does not contain far more short-period systems than those with periods beyond the well-known 2--3 hour gap \\citep[e.g.][]{Downes+01pasp,RitterKolb03aa}. Furthermore, the theoretical minimum period is generally found to be around 65 minutes, which is substantially shorter than the observed value. Many different AML mechanisms have been put forward as possible solutions to these problems (e.g.\\ \\citealt*{Andronov++03apj}; \\citealt{TaamSpruit01apj,Schenker+02mn}), without complete success. The AML rates given by different prescriptions can differ by several orders of magnitude \\citep{SchreiberGansicke03aa}.  Due to these difficulties, detailed observational population studies are of great importance to provide the constraints necessary to further develop the theory of compact binary evolution. The currently known population of CVs is afflicted by several strong biases which makes the characterisation of the true population -- and a comparison with models -- extremely difficult.  The Sloan Digital Sky Survey (SDSS\\footnote{\\tt http://www.sdss.org/}; \\citealt{York+00aj}) is a large project to obtain five-colour imaging and multi-object spectroscopy of a large area at high galactic latitudes in the Northern sky, with the main science goal being the measurement of redshifts for an extensive sample of galaxies and quasars. Objects for spectroscopic follow-up are chosen, on the basis of their photometric colours, to be blue and/or distant from the locus of the stellar main sequence. Approximately 200 CVs, with a wide range of photometric colours, have been detected in the SDSS spectroscopic database \\citep{Szkody+02aj, Szkody+03aj, Szkody+04aj, Szkody+05aj, Szkody+06aj}, and many more are expected to be found from future SDSS observations. The limiting magnitude of this population is $g \\sim 20$, much fainter than previous surveys (e.g.\\ \\citealt*{Green++86apjs}), so the SDSS CV population is both more homogeneous and on average fainter than the currently known sample of CVs. It is therefore expected to contain more short-period systems, and may include a number of intrinsically faint CVs which have evolved beyond the period bounce.  We are conducting a research program to measure the orbital periods of SDSS CVs (\\citealt{Gansicke05aspc, Gansicke+06mn}; Dillon et al., 2006, in preparation) in order to assess whether the study of a less biased sample of CVs can reconcile observations with theories of the evolution of interacting binary stars. In this work we present phase-resolved spectroscopy of six faint SDSS CVs (Table~\\ref{tab:iddata}) and measure orbital periods for SDSS\\,J023322.6+005059.5, SDSS\\,J091127.36+084140.7, SDSS\\,J103533.02+055158.3 and SDSS\\,J121607.03+052013.9. The spectra of SDSS\\,J103533.02+055158.3 show that it is an eclipsing system. In addition we observed SDSS\\,J101037.05+024914.9, finding that it was several magnitudes fainter than in the SDSS observations, and SDSS\\,J093238.2+010902.5, but did not detect a coherent radial velocity motion. Our selection policy was to pick faint CVs which were discovered by the SDSS and were observable from the VLT. We did not attempt to carefully select a representative sample of objects because the goal of our research project is to (eventually) measure the orbital periods of {\\it all} of the SDSS CVs with unknown periods. In this work we shall abbreviate the names of the targets to SDSS\\,J0233, SDSS\\,J0911, SDSS\\,J0935, SDSS\\,J1010, SDSS\\,J1035 and SDSS\\,J1216.  \\begin{table*} \\begin{center} \\caption{\\label{tab:iddata} Apparent magnitudes of our targets in the SDSS $ugriz$ passbands. $r_{\\rm spec}$ is an apparent magnitude calculated by the SDSS by convolving their flux-calibrated spectra with the $r$ passband function, and is obtained at a different epoch to the $ugriz$ magnitudes measured from the CCD imaging observations.} \\begin{tabular}{lllccccccc} \\hline SDSS name                & Short name & Reference     & $u$   & $g$   & $r$   & $i$   & $z$   & $r_{\\rm spec}$ \\\\ \\hline SDSS J023322.61+005059.5 & SDSS\\,J0233 & \\citet{Szkody+02aj} & 19.76 & 19.86 & 19.73 & 19.93 & 19.82 & 19.92 \\\\ SDSS J091127.36+084140.7 & SDSS\\,J0911 & \\citet{Szkody+05aj} & 19.63 & 19.70 & 19.16 & 18.66 & 18.28 & 19.51 \\\\ SDSS J093238.21+010902.5 & SDSS\\,J0932 & \\citet{Szkody+03aj} & 19.70 & 20.26 & 19.53 & 19.10 & 19.12 & 19.71 \\\\ SDSS J101037.05+024914.9 & SDSS\\,J1010 & \\citet{Szkody+03aj} & 20.39 & 20.75 & 20.35 & 20.71 & 20.46 & 20.63 \\\\ SDSS J103533.02+055158.3 & SDSS\\,J1035 & \\citet{Szkody+06aj} & 19.14 & 18.79 & 18.81 & 18.97 & 19.12 & 18.99 \\\\ SDSS J121607.03+052013.9 & SDSS\\,J1216 & \\citet{Szkody+04aj} & 20.32 & 20.13 & 19.98 & 20.26 & 20.27 & 20.10 \\\\ \\hline \\end{tabular} \\end{center} \\end{table*} ",
        "conclusions": "This work presents phase-resolved spectroscopy of six CVs recently discovered in SDSS spectroscopic survey observations. For four of these systems we have obtained a precise orbital period using radial velocities measured from the wings of H$\\beta$ or H$\\alpha$ emission lines. Three of the four systems have periods below the period gap: 96\\,min for SDSS\\,J0233, 82\\,min for SDSS\\,J1035 and 99\\,min for SDSS\\,J1216 (although we have not been able to completely rule out the possibility that the 92\\,min one-day alias is in fact the correct period). The fourth, SDSS\\,J0911, has a period of 296\\,min, which is substantially longwards of the period gap.  Whilst we have found a clear 96\\,min periodicity in the emission-line radial velocities of SDSS\\,J0233, a period of 60\\,min is present in the light curves of this system. We suggest that this system is a low-accretion-rate intermediate polar and the 60\\,min variation represents the spin period of the white dwarf.  The spectra of SDSS\\,J1035 show it to be an eclipsing system; subsequent ULTRACAM photometry has found the eclipses to be total (Littlefair et al., submitted). Its short period and spectrum dominated by a cool white dwarf are reminiscent of WZ\\,Sge systems (e.g.\\ \\citealt{Patterson+05pasp,Templeton06pasp}). From modelling the spectrum we find a white dwarf effective temperature of $11\\,700 \\pm 1500$\\,K and an upper limit of L0 to the spectral type of the secondary star, which is not detected in the spectrum. The high quality of the VLT/FORS2 observations has also allowed Doppler tomograms to be constructed for its accretion disc, and these show a clear bright spot and a disc with comparatively high-velocity motions.  We have found SDSS\\,J1216 to have a very low velocity amplitude of $13.8 \\pm 1.6$\\kms. To our knowledge this is the smallest amplitude so far measured for a CV from emission-line radial velocities (excluding lower limits set by non-detection of velocity variation). Through applying a series of physically reasonable constraints we find that the system must have either a high-mass white dwarf primary component and/or a secondary star with a mass in the brown dwarf regime. Systems with donor masses well below 0.1\\Msun\\ are expected to be very faint, because the secondary star is dim and the mass transfer rate is low, and so under-represented in the current observational population of CVs.  There are several candidates for CVs with brown-dwarf-mass secondary stars and their properties have been reviewed by \\citet{Littlefair++03mn}, to whose list should be added the recently-studied CVs HS\\,2331+3905 \\citep{Araujo+05aa} and RE\\,J1255+266 \\citep{Patterson++05pasp}. In general, inferences of the secondary mass have been drawn from the spectral energy distribution of the system, from calibrations between superhump period excess and mass ratio, or through detection of spectral lines of the secondary component in a few cases. \\citet{Littlefair++03mn} found that direct evidence for a brown dwarf secondary does not exist for any CV, but that some systems have substantial indirect evidence.  Three of the four CVs studied here have very short orbital periods, which is consistent with the expectation that studying successively fainter samples of CVs means that the targets become on average intrinsically fainter and of shorter period \\citep{Patterson++05pasp,Gansicke05aspc}. Our ongoing research program of CVs discovered by the SDSS will give orbital periods of a larger homogeneous sample of these systems, allowing us to accurately assess the biases affecting the observed sample of CVs and so accurately measure the properties of the intrinsic population of these objects."
    },
    "0609/astro-ph0609705_arXiv.txt": {
        "abstract": "We present a first analysis of a deep X-ray spectrum of the isolated neutron star RBS1223 obtained with XMM-Newton. Spectral data from four new monitoring observations in 2005/2006 were combined with archival observations obtained in 2003 and 2004 to form a spin-phase averaged spectrum containing 290000 EPIC-pn photons. This spectrum shows higher complexity than its predecessors, and can be parameterised with two Gaussian absorption lines superimposed on a blackbody. The line centers, $E_2 \\simeq 2E_1$, could be regarded as supporting the cyclotron interpretation of the absorption features in a field $B \\sim 4 \\times 10^{13}$\\,G. The flux ratio of those lines does not support this interpretation. Hence, either feature might be of truly atomic origin. ",
        "introduction": "\\label{intro} RBS1223 is one out of three XDINs (X-ray dim isolated neutron stars) found in the ROSAT Bright Survey (RBS), an optical identification program of the more than 2000 X-ray sources detected in the RASS with a count rate $>0.2$\\,\\crat\\ at high galactic latitude \\cite{r1}. The initial discovery \\cite{r2} was based on the soft X-ray spectrum and the steep SED, $f_{\\rm X}/f_{\\rm opt} > 10^4$. Follow-up Chandra observations revealed a periodically modulated signal \\cite{r3}, and HST-observations uncovered a candidate optical counterpart at $m_{\\rm 50CCD} = 28.56 \\pm0.13$\\,mag \\cite{r4}. Initial observations with XMM-Newton showed deviations from a Planckian energy distribution at energies below 500 eV which could be described with a Gaussian absorption line and interpreted as a cyclotron absorption line in a field of $(2-6) \\times 10^{13}$\\,G \\cite{fwh03}. The large number of photons collected with XMM-Newton also uncovered the true spin period of 10.31\\,s. In \\cite{asea05} the light curves in different energy bands were modeled in terms of a photospheric cap model. The temperature structure was found to be roughly compatible with the crustal field models of \\cite{urm04}. The implied field structure deviates from a simple centered dipole model. Spin-period changes appeared likely at that time thus stimulating  monitoring observations with Chandra and XMM-Newton. The spin history could be unequivocally fixed by a series of Chandra pointings \\cite{kap05} and revealed a spin-down of the star at a rate $\\dot{P} = 1.120 \\times 10^{-13}$\\,s\\,\\crat. Under the assumption that this is due to magnetic dipole torques, a characteristic age of 1.5 Myr and a magnetic field strength of $3.4 \\times 10^{13}$\\,G were inferred. Here we describe initial results of recent monitoring observations with XMM-Newton which, due to the large number of accumulated photons, besides accurate timing also allows spectroscopy with unprecedented accuracy.  \\begin{table}[t] \\caption{XMM-Newton observations of RBS1223 analysed in this contribution} \\centering \\label{t:obs} \\begin{tabular}{rclrcc} \\hline\\noalign{\\smallskip} Rev. & AO & Day & Exp. & Mode & Countrate \\\\ &     &     & [ks]     & &[\\crat]\\\\ \\tableheadseprule\\noalign{\\smallskip} 377 & 1& 2001 Dec 31 & 18.6 & SW & 2.49$^{\\rm a}$\\\\ 561 & 2& 2003 Jan 1 & 27.0 & FF& 2.53 \\\\ 743 & 3& 2003 Dec 30 & 30.2 & FF& 2.57 \\\\ 1015& 4& 2005 Jun 25 & 14.9 & FF& 2.56 \\\\ 1016& 4& 2005 Jun 26 & 12.9 & FF& 2.57 \\\\ 1025& 4& 2005 Jul 15 & 12.9 & FF& 2.56 \\\\ 1115& 4& 2006 Jan 10 & 14.9 & FF& 2.56 \\\\ \\noalign{\\smallskip}\\hline \\end{tabular}\\\\ \\begin{flushleft} -- Exposure time is given for EPIC-pn. \\newline -- Mean countrates are given for the whole XMM-Newton energy range (0.15 -- 10\\,keV)\\newline -- $^{\\rm a}$ countrate from \\cite{fwh03} \\end{flushleft} \\end{table} ",
        "conclusions": "The joint analysis of six observations of RBS1223 with XMM-Newton has led to a slight revision of the spectral parameters. The spectrum can be described with a blackbody and two Gaussian absorption lines superimposed. The blackbody temperature is significantly higher than previously found thanks to the detection of X-ray photons up to 2\\,keV. Also the parameters of the primary Gaussian absorption line are changed: it appears significantly narrower and more blueshifted than before. Whether we have seen the true continuum at any energy remains an open question.  Similarly unclear is the nature of the features which were parameterised with two Gaussian absorption lines. They have a flux ratio of about 5:1, their equivalent widths of EW$_1 \\sim 200$\\,eV (for $E_1 =0.23$\\,keV) and EW$_2 \\sim 180$\\,eV (for $E_2 = 0.46$\\,keV) are rather large. Their relative line energy and the rather large width of the lines are supportive of the proton cyclotron interpretation in a field of few times $10^{13}$\\,G (the line center of the primary and hence the derived magnetic field strength is not well contrained, in model 5 the best fit is achieved for $E_1 = 0.2$\\,keV, i.e.~it converges to the lower bound of the pre-set parameter range), consistent with the observed spin down rate of the star. The rather high flux ratio of the two lines, nevertheless, makes an interpretation of the two lines as being harmonics of each other unlikely, since the oscillator strength of the first harmonic is smaller than that of the fundamental by a factor $\\sim E/(mc^2)$, which becomes very small for proton masses \\cite{pav80}. An alternative could be to associate the two lines with different fields from e.g.~the two polar caps. Since the field was shown to be somehow non-dipolar or off-centered at least \\cite{asea05}, different field strength could be encountered at the two caps. It is difficult to test this hypothesis on the basis of the current data and it needs to be seen if the results of the phase-resolved spectroscopy reveal any helpful indication in this respect. However, if the two lines would be the fundamentals of two different cyclotron systems from two different regions, the infered field strengths would be $B_1 = 4.2 \\times 10^{13}$\\,G and $B_2 = 8.5 \\times 10^{13}$\\,G according to $B (10^{13}{\\mbox G}) = \\frac{1}{1.16} (1+z) \\frac{m_p}{m_e} E(\\mbox{keV})$ and assumed $z=0.18$ (see \\cite{asea05}), $E_1 = 0.23$\\,keV and $E_2 = 0.46$\\,keV. Hence, one field would be slightly below and one above the critical quantum field, where vacuum polarization becomes important in shaping the lines (see \\cite{laiho03} \\cite{laiho04}, and the discussion in \\cite{vkea04}). Since the second Gaussian line (higher field) is rather broader than the primary Gaussian, this scenario seems to be unlikely.  A possibility which needs to be explored in detail is a blend of magnetically shifted atomic transitions with proton cyclotron resonances. As the model computations of \\cite{laiho04} indicate, the amount of ionisation in the assumed hydrogen atmosphere will play a crucial role. An as accurate as possible temperature determination is thus very important for further progress. Due to its photometric stability on the one hand and the pronounced spin variability on the other hand, RBS1223 is a very promising target to gain further insight from even deeper X-ray spectroscopy."
    },
    "0609/astro-ph0609828_arXiv.txt": {
        "abstract": "We perform a self-consistent calculation of the thermal structure in the crust of a superbursting neutron star. In particular, we follow the nucleosynthetic evolution of an accreted fluid element from its deposition into the atmosphere down to a depth where the electron Fermi energy is 20~MeV. We include temperature-dependent continuum electron capture rates and realistic sources of heat loss by thermal neutrino emission from the crust and core. We show that, in contrast to previous calculations, electron captures to excited states and subsequent ${\\gamma}$-emission significantly reduce the local heat loss due to weak-interaction neutrinos. Depending on the initial composition these reactions release up to a factor of 10 times more heat at densities $<10^{11}\\nsp\\grampercc$ than obtained previously. This heating reduces the ignition depth of superbursts. In particular, it reduces the discrepancy noted by Cumming et al.\\ between the temperatures needed for unstable \\nuclei{12}{C} ignition on timescales consistent with  observations and the reduction in crust temperature from Cooper pair neutrino emission. ",
        "introduction": "\\label{s:introduction}  The ability to regularly monitor the X-ray sky with instruments such as \\rxte\\ and \\beppo\\ has produced many exciting discoveries in the last few years. An excellent example are superbursts \\citep[for a review see][]{Kuulkers2003The-observers-v}. Like normal type I X-ray bursts \\citep[for a review, see][]{strohmayer03:review}, superbursts are characterized by a rapid rise in the light curve followed by a quasi-exponential decay; when compared to type I X-ray bursts, superbursts are roughly $\\sim 10^{3}$ times more energetic, have cooling timescales of hours, and recur on timescales of years. For one source, \\source{4U}{1636}{-53}, three superbursts have been observed over a span of 4.7\\nsp\\yr\\ \\citep{wijnands:recurrent,kuulkers..ea:x-ray}. Although most superbursts are observed at mass accretion rates 0.1--0.3\\nsp\\MdotEdd, with $\\MdotEdd \\approx 10^{18}\\nsp\\grampersecond$ being the Eddington accretion rate, recently \\citet{.cornelisse.ea:superbursts} detected superbursts from the rapidly accreting source \\source{GX}{17}{+2}.  The currently favored scenario for these superbursts is the thermally unstable ignition of \\carbon\\ at densities $\\sim 10^{8}\\textrm{--}10^{9}\\nsp\\grampercc$ \\citep{cumming.bildsten:carbon,strohmayer.brown:remarkable}. \\citet{cumming.macbeth:thermal} demonstrated that the superburst lightcurve after the peak luminosity has been reached declines as a broken power law, with the early behavior dependent on the burst energy and the later behavior dependent on the ignition depth. This behavior has been observed \\citep{Cumming2005Long-Type-I-X-r}. The ability to constrain the burst energetics and ignition depth is crucial, not only for understanding the ignition of \\carbon\\ in a heavy-element bath, but also for probing the interior of the neutron star \\citep{brown:superburst,cooper.narayan:theoretical,Cumming2005Long-Type-I-X-r}. \\citet{brown:superburst} demonstrated that the temperature at densities $\\sim 10^{9}\\nsp\\grampercc$ is sensitive to the thermal conductivity of the inner crust and the neutrino emissivity of the core; this was confirmed by \\citet{cooper.narayan:theoretical} in a more extensive study. Recently, \\citet{Cumming2005Long-Type-I-X-r} showed that the neutrino emission from singlet-state Cooper pairing of neutrons in the inner crust limited the temperature at neutron drip to be $\\lesssim 5\\ee{8}\\nsp\\K$. Taken at face value, this would lead to superburst ignition column depths about an order of magnitude deeper than the $0.5-3\\ee{12}\\nsp\\GramPerSc$ deduced from observations, and correspondingly to recurrence times that are about an order of magnitude longer than observed. This led \\citet{Page2005Superbursts-fro} to speculate that the superbursters might in fact be strange stars.  The unstable ignition of \\carbon\\ depends critically on the thermal structure of the outer crust and is sensitive to the heat sources located there. The outer crust of the neutron star is composed of nuclei and degenerate electrons and exists where the electron chemical potential $\\mue \\lesssim 30\\nsp\\MeV$ \\citep[see][for a lucid discussion]{pethick98}. The ashes of hydrogen, helium, and carbon burning in the neutron star envelope are continuously compressed by the ongoing accretion of matter and incorporated into the crust. The increasing density induces nuclear reactions at various depths that can release energy. \\emph{We show that most of this energy is deposited in the crust, rather than being carried away by neutrinos. This leads to a hotter crust and decreases superburst ignition depths and recurrence times.}  Previous studies of the crust heat sources used a single representative isotope, \\iron\\ \\citep{sato79,haensel90a} or  \\palladium[106] \\citep{haensel.zdunik:nuclear} for the product of H/He burning. \\citet{haensel90a} found that electron captures would occur in two stages (see Fig.~\\ref{f:ec-schematic}). The first electron capture (labeled ``1'' in Fig.~\\ref{f:ec-schematic}) would occur when  $\\mue\\approx\\Egs^{(Z,A)}$ for capture onto an even-even nucleus: $(Z,A) + e^{-} \\to (Z-1,A)+\\bar{\\nu}_{e}$, where $Z$ and $A$ are even. Here the threshold is $\\Ethresh = \\Egs + \\Ex$, where the threshold for the ground-state-to-ground-state transition \\Egs\\ is computed from atomic mass differences and therefore includes the electron rest mass, and where \\Ex\\ is the energy of the lowest excited state into which the capture can proceed. Because of the odd-even staggering of the nuclear masses, $\\Egs^{(Z-1,A)}$ for electron capture onto the resulting odd-odd nucleus $(Z-1,A)$ is less than $\\Egs^{(Z,A)}$. Consequently, $\\mue > \\Egs^{(Z-1,A)}$ and a second electron capture $(Z-1,A)+e^{-}\\to(Z-2,A)+\\bar{\\nu}_{e}$ immediately follows.  Where experimental data were lacking, which is the case for the vast majority of nuclei, the heat deposited into the crust $Q$ was calculated assuming that the capture was into the ground state of the daughter nucleus (Fig.~\\ref{f:ec-schematic}, label ``4''). The heat deposited into the crust is then estimated as \\citep{haensel.zdunik:nuclear} $Q = (\\mue - \\Egs)/4$, where the factor 1/4 arises from the integral over the reaction phase space in the limit that $(\\mue-\\Ethresh)/\\mue\\ll 1$. This estimate assumes that final states are very low in excitation energy so that the contribution to the heat deposition from the radiative de-excitation of the daughter nucleus is always negligible compared to the heat release from the electron capture. In this picture, 75\\% of the nuclear energy release is emitted in form of neutrinos.  \\begin{figure}[tbh] \\centering{% \\includegraphics[width=3.1in]{electron-capture-schematic}} \\caption{Electron capture for which the first transition (labeled ``1'') proceeds to the ground state of nucleus $^{A}Z-1$, but the second transition instead goes to an excited state ``2'' followed by a radiative de-excitation ``3''. Previous works assumed that the second transition always went to the ground state ``4''. } \\label{f:ec-schematic} \\end{figure}  In this paper, we calculate heating from nuclear reactions in the outer crust (where the mass density is less than neutron drip, $\\sim 10^{11}\\nsp\\grampercc$) using realistic electron capture rates obtained by microscopic calculations of the transition matrix elements between the parent ground state and daughter excited states multiplied with the phase space factor corresponding to the $Q$-value of the decay. The rates depend on the excitation energy and the ambient temperature and density. We use a full reaction network having temperature- and density- dependences. Our thermal model of the crust, described in \\S~\\ref{s:thermal-structure}, is coupled to the reaction network (\\S~\\ref{s:thermal}), so that our solution is self-consistent. Instead of starting with a single nucleus, we follow the actual mix of nuclei initially produced by surface burning, both from the rp-process (\\S~\\ref{s:rp-ash}) and a mixture in nuclear statistical equilibrium formed during a superburst (\\S~\\ref{s:superburst-ash}). We find that electron captures deposit considerably more energy in the crust than in previous estimates \\citep{haensel.zdunik:nuclear}. This larger heat deposition is mainly due to the electron captures populating excited states with energy \\Ex\\ (Fig.~\\ref{f:ec-schematic}, label ``2''), rather than capturing into the ground state as assumed previously. The populated excited states then radiatively de-excite (Fig.~\\ref{f:ec-schematic}, label ``3''), depositing the entire excitation energy as heat into the crust. We also introduce an approximate model that gives a good estimate of the heating in the outer regions of the crust for an arbitrary mixture in \\S~\\ref{s:simple-model} and describe the implications for superburst ignition in \\S~\\ref{s:ignition}. ",
        "conclusions": "\\label{s:conclusions}  We show that allowing for electron captures into excited states drastically increases the heat deposited into the outer crust of accreting neutron stars. The fraction of the reaction energy carried away by neutrinos is sharply curtailed.  We also find that the amount of heating in the crust is very sensitive to the initial composition forged during the burning of light elements in the neutron star envelope. This sensitivity is a consequence of the pronounced shell and sub-shell structure of neutron-rich nuclei, which can lead to drastic changes in shape, single-particle level structure, and electron capture strength distributions, even between nuclei of similar $Z, A$.  We have explored how this increased deposition affects the ignition depth, and hence the recurrence time, of superbursts. For a realistic composition---steady-state rp-process ashes overlying superburst ashes---the inclusion of the excited states decreases the ignition depth by about 20\\%.  For the most optimistic case we consider, in which the neutrino emissivity from Cooper-pairing is suppressed, the ignition depth for  accretion at 0.3 Eddington is $\\yign=3.7\\ee{12}\\nsp\\GramPerSc$, closer to the value inferred from observations ($\\yign = 5\\ee{11}\\textrm{--}2.7\\ee{12}\\nsp\\GramPerSc$; \\citealt{Cumming2005Long-Type-I-X-r}), although a significant discrepancy remains unless the core neutrino emissivity is reduced below the modified Urca rate.  To explore the variation of superburst ignition depths with composition, we have constructed an approximate model of the heating in the crust. Using this model and including the neutrino emissivity from paired neutrons, we find that at 0.3 Eddington the superburst ignition depth can be as low as $1.3\\ee{12}\\nsp\\GramPerSc$ if the composition is dominated by nuclei with $A=40\\textrm{--}44$ or 82--86. The corresponding recurrence time for ignition at this depth, with the assumed accretion rate and neutron star parameters, is 2\\nsp\\yr. This is roughly consistent with the observations, without invoking any additional physics. We are not aware, however, of a mechanism that would, for example, avoid the production of nuclei around $A \\approx 60$ dominating the superburst ashes.  Our results also indicate that crustal heating is very sensitive to the underlying nuclear physics. Reliable estimates of the structure of very neutron-rich nuclei up to $A=106$ are needed to constrain the lowest lying electron capture strength, which determines the heating from the deexcitation of excited daughter states.  Nuclear masses are needed to determine reaction thresholds and neutron separation energies. Finally, masses, $\\beta^-$decay rates, and reaction rates for the proton- and helium-induced reactions on very neutron deficient nuclei during hydrogen and helium burning in the $\\alpha p$- and rp-processes are needed to accurately determine the initial composition. More work is also needed to extend our calculations into the inner crust to determine the heating there."
    },
    "0609/astro-ph0609533_arXiv.txt": {
        "abstract": "We present near-infrared narrow-band images of the supernova remnant W49\\,B, taken with the WIRC instrument on the Hale 200 inch telescope on Mt.~Palomar. The 1.64\\,$\\mu$m [Fe\\,II] image reveals a barrel-shaped structure with coaxial rings, which is suggestive of bipolar wind structures surrounding massive stars.  The 2.12\\,$\\mu$m shocked molecular hydrogen image extends 1.9\\,pc outside of the [Fe\\,II] emission to the southeast. We also present archival {\\em Chandra} data, which show an X-ray jet-like structure along the axis of the [Fe\\,II] barrel, and flaring at each end. Fitting single temperature X-ray emission models reveals:  an enhancement of heavy elements, with particularly high abundances of hot Fe and Ni, and relatively metal-rich core and jet regions\\@.  We interpret these findings as evidence that W49\\,B originated inside a wind-blown bubble ($R \\sim 5\\,{\\rm pc}$) interior to a dense molecular cloud. This suggests that W49\\,B's progenitor was a supermassive star, that could significantly shape its surrounding environment. We also suggest two interpretations for the jet morphology, abundance variations and molecular hydrogen emission: (1) the explosion may have been jet-driven and interacting with the molecular cavity (i.e.\\ a Gamma-ray burst); or (2) the explosion could have been a traditional supernova, with the jet structure being the result of interactions between the shock and an enriched interstellar cloud. ",
        "introduction": "\\label{intro.sec}  W49\\,B (G43.3-0.2) has the highest radio surface-brightness of all mixed-morphology  supernova remnants (SNRs) in the Galaxy \\citep{pye84,mof94}\\@. SNRs exhibiting centrally-filled X-rays inside an edge-brightened radio shell are referred to as mixed-morphology \\citep{rho98}\\@. A number of models have been proposed for mixed-morphology supernova remnants \\citep{whi91,cox99,she99,che99}; these models were designed to explain the larger (i.e. older) SNRs by invoking interactions with a denser-than-average interstellar medium. The very high radio brightness and the X-ray properties of W49\\,B make it a compelling object to study in detail, because it may be fundamentally different from other remnants of its class.  High resolution X-ray spectra of W49\\,B \\citep{hwa00} revealed elemental abundances enhanced in heavy elements, suggesting that W49\\,B was the product of a Type Ia explosion.  However, a subsequent study of \\ion{H}{1} absorption \\citep{bro01} shows that W49\\,B is at the same distance as the star forming region W49A \\citep[11.4\\,kpc,][]{gwi92}\\@.   Moreover, long wavelength radio observations \\citep{lac01} suggest that W49\\,B is absorbed by H$^+$ gas and is in a high pressure ($\\sim$$10^{6} \\,{\\rm cm^{-3} K}$) region of the Galaxy, again associating the remnant with the W49\\,A complex.  But at this distance, W49\\,B would have to have a massive progenitor, more akin to a Type II supernova.  Thus we have an observational inconsistency:  W49\\,B has enhanced iron abundances which are characteristic of a type Ia supernova, yet it is located in a star forming region and more likely a result of a core-collapse explosion.  Recently, \\citet{mic06} completed a study of W49B with XMM-Newton, where they concluded that the X-ray emission arises in a high-metallicity collisionally-ionized plasma, with a temperature gradient from west to east.  In this {\\em Letter} we present near-infrared narrow line imaging (\\S\\ref{ir.sec}), which we compare to our spectral analysis of {\\em Chandra} archival data (\\S\\ref{chandra.sec})\\@.  In \\S\\ref{discussion.sec} we interpret our results as evidence that W49\\,B was created inside a wind-blown bubble within a molecular cloud.  We also include two very different interpretations of our observations: (1) in \\S\\ref{grb.sec} we interpret our observations as consistent with a jet-driven explosion (i.e.~Gamma-Ray Burst); and in \\S\\ref{sn.sec} we interpret our observations as the result of a traditional supernova explosion inside a complex cavity. ",
        "conclusions": "We have presented evidence that W49\\,B resulted from the explosion of a supermassive star, inside a wind-blown bubble, which is in turn interior to a dense molecular cloud.  We have also given two interpretations for its morphological and chemical structure, each explanation require a set of extreme initial conditions.  One interpretation assumed that the axially-symmetric structure is caused by the explosion mechanism, resulting in the conclusion that W49\\,B was a jet-driven explosion akin to the current model of long/soft gamma-ray bursts. The other interpretation assumed that the same observations were rather dominated by complex cloud and magnetic field structures, and that the explosion could have been a standard isotropic supernova explosion.  Future work should consist of detailed modeling, with a realistic preexplosion medium, and with jet-driven and isotropic explosions, to explain all aspects of the W49\\,B observations.  Additional observations would also be useful, especially in order to understand  complete yields of nucleosynthesis with infrared spectroscopy by complementing the yields of the X-ray gas, and to independently map out the molecular cloud structure surrounding W49\\,B."
    },
    "0609/astro-ph0609019_arXiv.txt": {
        "abstract": "The large-scale dynamics of a two-fluid system with a time dependent interaction is studied analytically and numerically. We show how a rapid transition can significantly suppress the large-scale curvature perturbation and present approximative formulae for estimating the effect. By comparing to numerical results, we study the applicability of the approximation and find good agreement with exact calculations. ",
        "introduction": "There are many instances in cosmology, when the large-scale evolution of perturbations \\cite{Mukhanov:1990me} in a multi-component system needs to be solved \\cite{Wands:2000dp}. The use of curvature and entropy perturbations has become a \\textsl{de facto} standard in this formalism, pioneered by Kodama and Sasaki \\cite{Kodama:1985bj}. In a multi-component system, interactions between fluids are an important aspect in determining the evolution of the curvature perturbation \\cite{Malik:2004tf}. The importance of such systems is evident: interacting fluids are the cornerstone of many important cosmological scenarios, such as reheating at the end of inflation \\cite{Albrecht:1982mp,Den:1984tn,Kripfganz:1985mn,Bastero-Gil:2002xr,Dvali:2003em,Kofman:2003nx} or the curvaton scenario \\cite{Enqvist:2001zp,Lyth:2001nq,Moroi:2002rd,Moroi:2001ct,Lyth:2002my}.  The total curvature perturbation, $\\zeta$, generally evolves in multi-component fluid whenever the non-adiabatic pressure is non-zero. In particular, if the fluids are interacting the total curvature perturbation changes. Whether the evolution changes the amplitude of the large-scale curvature perturbation significantly or not, depends on the exact nature of the system. Efficient amplification, or damping, of $\\zeta$ on large scales by later evolution allows one to modify the previously produced perturbation spectrum, relaxing constraints on the scale at which the perturbations were produced. A natural framework for such mechanisms is \\eg within a traditional multiple inflationary scenario \\cite{minf} or a string landscape picture \\cite{cliff, multiverse}.  The transfer of energy between fluids can be described by a number of ways: for example in \\cite{Malik:2002jb} a constant interaction between different components is assumed whereas authors of \\cite{Lyth:2001nq} and \\cite{Lyth:2002my} use the so-called sudden decay approximation instead. In this approximation, the two fluids evolve independently until energy is transferred rapidly from one fluid to the other at a particular point in time. This method allows one to estimate the magnitude of the total curvature perturbation analytically. Here we relax these assumptions and allow for a time-dependent interaction while evolving the full large-scale perturbation equations with no further approximations. Research in this direction has recently been presented in \\cite{Dvali:2003em} and \\cite{Kofman:2003nx}, where locally different decay rates of the inflaton are generated by spatially varying reheating temperature and couplings. Our approach is to study how turning on the interaction affects the perturbations and the significance of the transition time scale. When the interaction is turned on rapidly so that it is well modeled by a step-function, we can proceed with analytical methods and derive a useful approximation to the numerical results.  In the fluid approach with constant interaction term, the microphysics cannot be properly accounted for as the fluid begins to decay simply when its decay width is of the order of the Hubble rate, $\\Gamma\\sim H$. Such a simple description does not always properly reflect the physical process, however, since there can be additional physical scales and processes involved that can play an important role. Cosmological processes involving different physical scales, such as phase transitions, do not need to conform to the simple fluid description but are better described by a time-dependent interaction.  In this paper we will consider time dependent interactions between different fluids. In section II we present the governing equations of perturbations in the Newtonian gauge. In the following section we compare the evolution of perturbations by means of numerical and analytical methods. We present analytical formulae which give an accurate approximation of the behavior of curvature perturbation variables and study the validity of these approximations. We end the paper with a short conclusion. Precise calculations of the evolution of the curvature perturbation of the decaying fluid during transition are presented in the Appendix. ",
        "conclusions": "In the present article we have discussed the evolution of perturbation variables in a two-component interacting fluid system. This has been studied previously by a number of authors, \\eg \\cite{Kodama:1985bj, Wands:2000dp, Malik:2004tf, Matarrese:2003tk}. Our treatment goes beyond the sudden decay approximation by including all relevant perturbations explicitly and including a time scale, $\\tau$, that describes how quickly the interaction is turned on. Such an interaction is well motivated by physical considerations and worth considering particularly in the fluid approach, where the microphysical properties of the fluids are otherwise hidden.  We have seen how the the curvature perturbation can be efficiently suppressed in a system with interacting fluids. A fluid that decays rapidly, $\\tau H_* \\ll 1$, $\\Gamma/H_* \\gg 1$, generally diminishes the curvature perturbation associated with that fluid which then subsequently affects the total curvature perturbation. The analysis shows that the decrease of perturbations depends crucially on the ratio $\\Gamma/H_*$ in addition to the initial conditions.  We have derived two expressions for different values of  $\\Gamma/H_*$ which give the magnitude of this jump and compared these analytical results to the numerical calculations, showing good agreement. We have also derived a formula to approximate the final value of the total curvature perturbation, $\\zeta$, after the transition. Again we find that the approximation is good when the decay is fast enough, \\ie $\\Gamma/H_*\\gsim 10^{-1}$.  Processes with efficient suppression of the large-scale curvature perturbation can in general be realized in a scenario where the three physical time scales have a hierarchy, $H_* \\ll \\Gamma,\\ 1/\\tau$. Such a hierarchy can be achieved naturally if the time scales are associated with different physical processes. If, on the other hand, the decay processes become effective only when $\\Gamma\\sim H_*$, no large net effect results. Interesting avenues to explore in this respect are processes including phase transitions and/or different scales, making  many cosmological phase transitions a potential source of interesting phenomena.  In addition, inflationary scenarios with multiple phases can provide a suitable framework for modifying the large-scale perturbations, \\eg  when the intermediate phases between inflationary periods are dominated by string networks \\cite{cliff}.   \\subsection*"
    },
    "0609/astro-ph0609255_arXiv.txt": {
        "abstract": "We have monitored the Neptune-mass exoplanet-hosting M-dwarf Gl 581 with the 1m Swope Telescope at Las Campanas Observatory over two predicted transit epochs. A neutral density filter centered at 550nm was used during the first epoch, yielding 6.33 hours of continuous light curve coverage with an average photometric precision of 1.6 mmags and a cadence of 2.85 min. The second epoch was monitored in B-band over 5.85 hours, with an average photometric precision of 1.2 mmags and 4.28 min cadence. No transits are apparent on either night, indicating that the orbital inclination is less than 88.1 deg for all planets with radius larger than 0.38 $R_{Nep}$ = 1.48 $R_{Earth}$. Because planets of most reasonable interior composition have radii larger than 1.55 $R_{Earth}$ we place an inclination limit for the system of 88.1 deg. The corresponding minimum mass of Gl 581b remains 0.97 $M_{Nep}$ = 16.6 $M_{Earth}$. ",
        "introduction": "\\label{sec:intro} Bonfils et al. (2005a) recently reported the detection of a Neptune-mass planet around the M-dwarf Gl 581. This planet, with an orbital period of 5.366 days and minimum mass $M_{2}sini$= 0.97$M_{Nep}$, is only the fifth one found around M-dwarfs. The two other M-dwarfs currently known to host planets are Gl 876, with two Jupiter-mass (Delfosse et al. 1998; Marcy et al. 1998) and one Neptune-mass (Rivera et al. 2005) planets, and Gl436 with a Neptune-mass planet (Butler et al. 2004). Therefore, GJ581b is only the third Neptune-mass class planet found around an M-dwarf.  Of the three Neptune-mass class planets, GJ876d and GJ436b have been photometrically searched for possible transits (Rivera et al. 2005; Butler et al. 2004), although no transits have been found. GJ581b, on the other hand, has not yet been searched for transits. The probability of transit for GJ581b is only 3.3\\%, assuming a stellar radius of 0.29$R_{\\sun}$ and an orbital separation of 0.041 AU (Bonfils et al. 2005a). However, if a transit does occur, the planet will cover a larger fraction of the M-dwarf than in the equivalent case of a planet orbiting a Sun-like star. Therefore, transits of GJ 581b would be in principle deeper and easier to detect. The detection of a transit would allow us to establish the absolute mass of the planet and its radius. From these parameters we could also determine the mean density of the planet and also estimate its chemical composition.  GJ 581 (V* HO Lib), an M3V star, has been catalogued as a BY Draconis variable. This type of variable is characterized by quasiperiodic photometric variations over time scales from less than a day to months, and amplitudes ranging from a few hundredths of a magnitude to 0.5 mags. This variability is generally attributed to surface spots and chromospheric activity, phenomena very common among M-type dwarfs. Weis (1994) monitored GJ 581 over eight years and found seasonal and long-term photometric variations of 6 -- 8 mmags in V, R, and I band. However, no short-term (in the scale of hours) variability measurements have been reported so far. Furthermore, no information is available about the photometric precision of the star at wavelengths bluer than V.  This paper presents the first short-term precision light curves of GJ581 over two predicted planetary transit epochs. The light curves were collected in optical and B bands and span over 6.33 and 5.85 consecutive hours, respectively. Section 2 describes the observations, with their analysis presented in section 3. The results of a geometrical search for transits are presented in section 4. ",
        "conclusions": ""
    },
    "0609/gr-qc0609024_arXiv.txt": {
        "abstract": "\\noindent We study the formation of black holes by spherical domain wall collapse as seen by an asymptotic observer, using the functional Schrodinger formalism. To explore what signals such observers will see, we study radiation of a scalar quantum field in the collapsing domain wall background. The total energy flux radiated diverges when backreaction of the radiation on the collapsing wall is ignored, and the domain wall is seen by the asymptotic observer to evaporate by non-thermal ``pre-Hawking radiation'' during the collapse process. Evaporation by pre-Hawking radiation implies that an asymptotic observer can never lose objects down a black hole. Together with the non-thermal nature of the radiation, this may resolve the black hole information loss problem. ",
        "introduction": "\\label{introduction}   Black holes embody the long-standing theoretical challenge of combining general relativity and quantum mechanics, with various proposals being advocated over the years to resolve paradoxes associated with black hole formation, evaporation and information loss. Resolution of these issues has become even more timely with the possible formation and evaporation of black holes in particle accelerators in the framework of higher dimensional models that have recently garnered much attention. The process of black hole formation is generally discussed from the viewpoint of an infalling observer. However, in all physical settings it is the viewpoint of the asymptotic observer that is relevant. More concretely, if a black hole is formed in the Large Hadron Collider, it has to be observed by physicists sitting on the CERN campus. The physicists have clocks in their offices and they watch the process of formation and evaporation in this coordinate frame. They must address questions such as: At what time did a black hole form? Is any information lost into the black hole? How long did it take for the black hole to evaporate? What is the spectrum of the decay products?  The process of gravitational collapse has been studied extensively over the last few decades, from many different viewpoints, including 1+1 dimensional models and modifications of general relativity ({\\it e.g.} see \\cite{BirrellandDavies}). Unlike a large subset of this work, our analysis is in 3+1 dimensions and within conventional general relativity. We model the general problem by choosing to study a collapsing spherical shell of matter, more specifically a vacuum domain wall. The physical setup of the problem and the functional Schrodinger formalism are described in Sec.~\\ref{setup}.  A crucial aspect of our analysis is that we address the question of black hole formation and evaporation {\\em  as seen by an asymptotic observer}. Initially, when the domain wall is large, the spacetime is described by the Schwarzschild metric, just as for a static star. From here on, the wall and the metric are evolved forward in time, always using the Schwarzschild time coordinate. We emphasize that all our discussion, unless explicitly stated, refers to the Schwarzschild time, $t$, and this defines the time slicing of the spacetime. As is well known, the Schwarzschild coordinate system breaks down at a black hole horizon, and there is danger that our analysis will also break down at some point during the gravitational collapse. However, we do not encounter any such difficulties, suggesting that our calculation is self-consistent. A second danger is that the coordinate system may provide an incomplete description of the gravitational collapse spacetime. This remains a possibility. However, we find that Schwarzschild coordinates are sufficient to answer the very specific set of questions we ask from the asymptotic observer's viewpoint. Namely, does the asymptotic observer see objects disappear into a black hole in the time that he sees the collapsing body evaporate? And, is the spectrum of the radiation received ever truly thermal (even in the semiclassical approximation)?  In Sec.~\\ref{classical} we verify the standard result that the formation of an event horizon takes an infinite (Schwarzschild) time if we consider classical collapse. This is not surprising and is often viewed as a limitation of the Schwarzschild coordinate system. To see if this result changes when quantum effects are taken into account, we address the problem of quantum collapse using a minisuperspace version of the functional Schrodinger equation \\cite{DeWitt} in Sec.~\\ref{qcollapse}. We find that even in this case the black hole takes an infinite time to form, contrary to some speculations in the literature \\cite{frolov}.  In Sec.~\\ref{semiclassradn} we consider the possible radiation associated with the collapsing shell by considering the interaction of a quantum scalar field and the classical background of a collapsing domain wall. We treat the problem using the functional Schrodinger picture, which we relate to the standard Bogolubov treatment carried out in Sec.~\\ref{appa}. Here we find that the shell, even as it collapses, radiates away its energy in a finite amount of time. With some assumptions about the metric close to the incipient horizon, we conclude that the evaporation time is shorter than what would be taken by objects to fall through a black hole horizon. This leads us to the conclusion that the asymptotic observer will see the evaporation of the collapsing shell before he can see any objects disappear.  We discuss our results from the point of view of an infalling observer in Sec.~\\ref{infallingobserver}, where we attempt to reconcile the fact that such an observer will not see substantial radiation with the observations made by an asymptotic observer. Our conclusions are summarized in Sec.~\\ref{discussion}, where we elucidate a possible resolution of the information loss problem suggested by our results, together with a discussion of possible loopholes and future directions. ",
        "conclusions": "\\label{discussion}  In this paper we have studied the collapse of a gravitating spherical domain wall using the functional Schrodinger equation. We would like to clearly delineate our analysis of the collapse and the emitted non-thermal quantum radiation from the interpretational issues about the formation of an event horizon.  First, we studied the collapse of a gravitating spherical domain in both classical and quantum theory, ignoring any evaporative processes. It has been suggested in the literature that quantum fluctuations can cause the collapse and formation of a black hole in a finite (Schwarzschild) time \\cite{frolov}. However, our results show that this is not the case and the horizon does not form in a finite time even in the full quantum treatment.  Then we studied radiation from the collapsing shell as seen by the asymptotic observer. In the process of gravitational collapse, there are two, perhaps related, sources of radiation: first is the radiation from particle creation in the changing gravitational field of the collapsing ball, and second may be Hawking-like  radiation due to a mismatch of vacua at early and late times. The functional Schrodinger analysis takes all such sources into account and therefore gives the total particle production. We have found a non-thermal distribution of particle occupation numbers, with departures from thermality as illustrated in Fig.~\\ref{Nvsomega} and discussed toward the end of Sec.~\\ref{semiclassradn}. In a limited range of frequencies, the spectrum is approximately thermal and the temperature fitted in a restricted range of frequencies is constant and roughly equal to the Hawking temperature $1/4\\pi R_S$. The radiation becomes thermal in the entire range of frequencies only in the limit $t \\rightarrow \\infty$, {\\it i.e.} when the horizon is formed. Further, the mode occupation number diverges in the infinite time limit, if the backreaction is neglected ({\\it i.e.} the background is held fixed). Since an outside  observer never sees formation of a horizon in a finite time, radiation observed by him is never quite thermal. (Non-thermal features also get greatly amplified once the background is also treated quantum mechanically \\cite{Vachaspati:2007hr}.) This non-thermal radiation has strong implications for the information loss paradox since it can carry information about the collapsing matter.  Without a rigorous calculation that includes backreaction, one can not give a definite answer to the final fate of a collapsing object. It may happen that the diagram in Fig.~\\ref{usualbh} is correct and some radical and elaborate solutions to the problems we mentioned in Sec.~\\ref{infallingobserver} are needed. However, one can imagine an alternative picture, different from the one in Fig.~\\ref{usualbh}, which seems to have fewer problems, and that is that an event horizon never forms. Since the mass of the shell is decreasing during the collapse, the shell will be chasing its own Schwarzschild radius, and the question is whether the shell will catch up to its own Schwarzschild radius or completely evaporate before that happens \\cite{Gerlach:1976ji}.  With backreaction included, the radiation should lead to a continual reduction of the Schwarzschild radius, $R_S$, occurring in the Ipser-Sikivie metric (see Sec.~\\ref{setup}). Then, as seen by the asymptotic observer, one of two possibilities occurs: either the collapsing domain wall evaporates and $R_S \\rightarrow 0$ in a finite time, or else backreaction causes the radiation rate to slow down and vanish in a finite time. This latter possibility is unlikely, as our estimates suggest that the rate of emission increases as $R_S$ decreases \\footnote{ Though our analysis only holds for objects of mass greater than Planck mass (see Sec.~\\ref{qcollapse}) and there could be qualitative changes in the collapse and radiation as the Planck mass is approached.}. We therefore conjecture that the backreaction due to particle production will cause the collapsing domain wall spacetime to completely evaporate in a finite time. In this case, the spacetime can either be as given in Fig.~\\ref{usualbh}, or have the same global spacetime structure as Minkowski space, as shown in Fig.~\\ref{minkowski}. If the latter picture is correct, it also means that the infalling observer will not encounter an event horizon, because this feature is simply absent from the spacetime. Another way to see this is to note that the causal relation between two events is the same for all observers. Hence if the asymptotic observer sees a signal from an infalling object after he sees the last radiation ray emitted by the evaporating wall, this will also be the sequence of signals seen by the infalling observer. As discussed in Sec.~\\ref{infallingobserver}, the infalling observer would expect to see an intense burst of radiation as the wall approaches the Schwarzchild radius, but can fail to do so because his detectors are too small to detect the emitted range of frequencies.    \\begin{figure} \\centerline{\\scalebox{0.80}{\\input{newminkowski.pstex_t}}} \\caption{The spacetime of a collapsing domain wall. During collapse the wall emits non-thermal (quasi-Hawking) radiation as depicted by the arrows. Our calculations indicate that the total energy flux between the point $i^0$ to some point indicated by F is equal to the energy of the initial domain wall. Hence we conjecture that the domain wall evaporates completely at point D. Between F and $i^+$, there is no radiation flux arriving at $\\scri^+$. The event horizon and singularity present in the customary treatment are not formed and the spacetime structure is the same as that of Minkowski spacetime. } \\label{minkowski} \\end{figure}  In the absence of an exact backreaction calculation, we also have to allow for the possibility that a value of the critical mass exists above which Fig.~\\ref{usualbh} applies and below which Fig.~\\ref{minkowski} holds. Also, as discussed by Hawking \\cite{Hawking:2005kf}, the question of ``whether a black hole forms'' is not sharp enough and may not make sense in the full quantum theory since all of the measurements are made by an asymptotic observer at infinity, while a collapsing object exists for a finite time and disappears by emitting radiation in the strong field region in the middle. An asymptotic observer can never be sure if a black hole formed because of underlying quantum uncertainty \\cite{Hawking:2005kf}.  The broad picture we have obtained is consistent with that proposed in Refs.~\\cite{Gerlach:1976ji,Boulware:1975fe}, though there are differences in the analysis and the conclusions. In particular, we find a non-thermal spectrum whereas Gerlach argues for thermality. Our picture also supports the interpretation of Hawking radiation given in Ref.~\\cite{Townsend} whereby particles are created during the process of gravitational collapse and are then radiated slowly to form what we call Hawking radiation. We have indeed found particle production during the collapse but the radiation is not quite thermal. It is only in the frequency range where the occupation number spectrum can be approximated by $T/\\omega$ (Eq.~(\\ref{smallfreqs})) that thermality holds at finite time. Also note that the non-thermality we find is in the mode occupation numbers. Propagation of the radiation in the background metric will cause further non-thermality due to greybody factors.  If we live in a world of low scale gravity, the collision of particles in high energy accelerators will lead to a situation where the particles are in a continual state of gravitational collapse from which non-thermal radiation is being emitted. The life-time of such a state can be estimated once we know the details of the radiation more precisely from an analysis which includes backreaction. However, on dimensional grounds, Hawking's estimate for the lifetime of a black hole ($\\sim R_S^3/G$) may well apply to the colliding particles as well.  In reality the collapse is further complicated by the fact that the collapsing object is not kept in isolation and there are external forces that can disrupt the collapse at any point in time. From the perspective of potential information loss, note that any infalling encyclopedias can be returned to the asymptotic observer if the collapse is disrupted at any time, as it could be, for example by a bomb set to go off at some late, but finite time. Most importantly, since we calculate that the radiation emitted during the gravitational collapse is never truly thermal, the classic information loss issue in black holes should, in this case, be a non-problem for the asymptotic observer \\footnote{To be sure, one must explicitly evaluate the information contained in the non-thermal radiation.}.  Our primary result, that no event horizon forms in gravitational collapse as seen by an asymptotic observer is suggestive of the possibility of using the number of local event horizons to classify and divide Hilbert space into superselection sectors, labeled by the number of local event horizons.  Our result suggests that no operator could increase the number of event horizons, but the possibility of reducing the number of pre-existing primordial event horizons is not so clear and would require that Hawking radiation not cause any primordial black hole event horizons to evaporate completely.  Our conclusions have been derived on the basis of a number of assumptions which we now discuss. The first is the truncation of superspace to minisuperspace. We have only included spherically symmetric field configurations. Even then, the metric is restricted to be the classical solution sourced by a spherical domain wall. A more general analysis would include more metric degrees of freedom, though it is hard to see how this would make a difference to our conclusion. Similarly, we have restricted ourselves to a zero thickness domain wall. A more general analysis would allow for a thick wall. Finally, the Wheeler-de Witt formalism as we have used it, does not allow for the creation and annihilation of domain walls (``third quantization''). Perhaps third quantization could allow for the spontaneous creation of a black hole and the annihilation of the wall, effectively leading to black hole formation. A fourth possibility is that our Lagrangian breaks down near the Schwarzschild horizon and ``quantum gravity'' effects become important. This is usually thought not  to be the case since the spacetime curvature near the horizon is small for large black holes.  Perhaps the most serious drawback of our analysis is that it does not include backreaction on the gravitational collapse due to radiation. While we do not expect such inclusion to alter our conclusions regarding the non-existence of event horizons for asymptotic observers, we  are currently exploring ways to extend our treatments to include backreaction.  No theoretical idea is complete without the possibility of experimental verification and so it is important to ask if the picture we have developed in this paper can also be tested experimentally. We have already mentioned the relevance of our conclusions to black hole production in particle accelerators provided low scale gravity is correct. However, there is an even more accessible experimental system where these theoretical ideas can be put to the test. These are condensed matter systems in which sonic black holes (dumbholes) may exist \\cite{Unruh:1980cg}. It is very hard to realize a dumbhole in the laboratory for various experimental reasons; the closest known realization seems to be the propagating He-3 AB interface in the experiment of Ref.~\\cite{Baretal00} as discussed in \\cite{Vac04}. Yet the crucial aspect of our work in this paper is that {\\em there is no need to produce a dumbhole in order to see acoustic ``pre-Hawking'' radiation}. The process of collapse toward a dumbhole will give off radiation. This is also the conclusion of Ref.~\\cite{Barceloetal} though the details of the analysis and conclusions are different -- for example, we find non-thermal emission whereas these authors claim thermal emission with a modified temperature that is lower than the Hawking temperature. In any case, it should be much easier to do experiments in the laboratory that do not go all the way to forming a dumbhole, and this could be an ideal arena to test pre-Hawking radiation.  Our conclusions are important not only for the general issue of the breakdown of unitarity via information loss, but also for more general studies of black hole formation, whether they be in the context of astrophysics ({\\it e.g.} galactic black holes) or in future accelerator experiments. In all these situations, we are asymptotic observers watching the gravitational collapse of matter, and we may never see effects associated with a black hole event horizon. Only effects occurring during the gravitational collapse itself appear to be visible.     In the range $t < 0$, $\\omega$ is a constant and the solution to Eq.~(\\ref{rhoeq}) is \\begin{equation} \\rho (\\eta ) = \\frac{1}{\\sqrt{\\omega_0}} \\end{equation}  In the range $0< t < t_f$, we do not have an analytic solution but we can derive certain useful properties. First note that in terms of $\\eta$ \\begin{equation} \\omega^2 = \\frac{\\omega_0^2}{1-\\eta/R_S} \\end{equation} Then the equation for $\\rho$ after rescalings can be written as: \\begin{equation} \\frac{d^2 f}{d{\\eta '} ^2} = - (\\omega_0 R_S )^2 \\left [ \\frac{f}{1-{\\eta '}} - \\frac{1}{f^3} \\right ] \\label{feq} \\end{equation} where ${\\eta '} = \\eta /R_S$, $f = \\sqrt{\\omega_0} \\rho$. The boundary conditions are \\begin{equation} f(0) = 1 \\ , \\ \\ \\ \\frac{df(0)}{d{\\eta '}} = 0 \\end{equation} The last term with the $1/f^3$ becomes singular as $f \\rightarrow 0$. Let us consider another equation with the $1/f^3$ replaced by something better behaved. For example, \\begin{equation} \\frac{d^2 g}{d{\\eta '}^2} = - (\\omega_0 R_S )^2 \\left [ \\frac{g}{1-{\\eta '}} - g \\right ] \\label{softereq} \\end{equation} with boundary conditions \\begin{equation} g(0) = 1 \\ , \\ \\ \\ \\frac{dg(0)}{d{\\eta '}} = 0 \\label{softbcs} \\end{equation} Eq.~(\\ref{softereq}) implies that $g(\\eta ')$ is monotonically decreasing as long as $g(\\eta ') > 0$. Furthermore, it is decreasing faster than the solution for $f$ as long as $f < 1$, since the $1/f^3$ term in Eq.~(\\ref{feq}) is a larger ``repulsive'' force than the $g$ term in the Eq.~(\\ref{softereq}). So \\begin{equation} g({\\eta '} ) \\leq f({\\eta '} ) \\end{equation} for all $\\eta '$ such that $g(\\eta ' ) > 0$.  Eqs.~(\\ref{softereq}) with initial conditions (\\ref{softbcs}) can be solved in terms of degenerate hypergeometric functions. For us, the important point is that the solution for $g$ is positive for all ${\\eta '}$ and, in particular, $g (1) > 0$ for all the values of $\\omega_0 R_S$ that we have checked. Therefore $f({\\eta '})$ is positive, at least for a wide range of $\\omega_0 R_S$.  Let $f_1 = f(1) \\ne 0$. Then the equation for $f$ can be expanded near ${\\eta '}=1$. \\begin{equation} \\frac{d^2f}{d{\\eta ' }^2} \\sim - (\\omega_0 R_S)^2 \\left [ \\frac{f_1}{1-{\\eta '}} - \\frac{1}{f_1^3} \\right ] \\end{equation} This shows that \\begin{equation} \\frac{df}{d{\\eta '}} \\sim  (\\omega_0 R_S)^2 f_1 \\ln (1-{\\eta '}) \\rightarrow - \\infty \\end{equation} as ${\\eta '} \\rightarrow 1$.  Hence $\\rho(\\eta = R_S)$ is strictly positive and finite while $\\rho_\\eta (\\eta = R_S) = -\\infty$ for finite and non-zero $\\omega_0$. Since $f = \\sqrt{\\omega_0} \\rho$, and $f \\rightarrow 1$ for $\\omega_0 \\rightarrow 0$, we also see that $\\rho \\rightarrow \\infty$ and $\\rho_\\eta \\rightarrow 0$ as $\\omega_0 \\rightarrow 0$.  In the range $t_f < t$, $\\omega$ is a constant. However, the solution for $\\rho$ is not a constant, unlike in the range $t < 0$, since the constant solution $1/\\sqrt{\\omega (t_f)}$ does not necessarily match up with $\\rho (t_f -)$ to ensure a continuous solution. Yet it is easy to check that in this region ${\\dot N} = 0$ and so there is no change in the occupation numbers. So we need only find $N(t_f-, {\\bar \\omega})$ to determine $N(t \\rightarrow \\infty , {\\bar \\omega})$."
    },
    "0609/astro-ph0609549_arXiv.txt": {
        "abstract": "Gravitationally unstable disks can fragment and form bound objects provided that their cooling time is short. In protoplanetary disks radiative cooling is likely to be too slow to permit formation of planets by fragmentation within several tens of AU from the star. Recently, convection has been suggested as a faster means of heat loss from the disk but here we demonstrate that it is only marginally more efficient than the radiative cooling. The crucial factor is the rate at which energy can be radiated from the disk photosphere, which is robustly limited from above in the convective case by the adiabatic temperature gradient (given a certain midplane temperature). Thus, although vigorous convection is definitely possible in disks, the inefficiency of radiative loss from the photosphere may create a bottleneck limiting the ability of the disk to form self-gravitating objects. Based on this argument we derive a set of analytical constraints which diagnose the susceptibility of an unstable disk to fragmentation and show that the formation of giant planets by fragmentation of protoplanetary disks is unlikely to occur at distances of tens of AU. At the same time these constraints do not preclude the possibility of fragmentation and star formation in accretion disks around supermassive black holes. ",
        "introduction": "\\label{sect:intro}   It has been known for a long time (Safronov 1960; Toomre 1964) that massive differentially rotating disks are prone to gravitational instability (hereafter GI) which sets in when the destabilizing effect of the self-gravity in the disk exceeds the combined restoring action of the pressure and Coriolis forces. Instability condition is conveniently expressed in terms of the so-called Toomre Q-parameter as \\ba Q\\equiv\\frac{\\Omega c_s}{\\pi G\\Sigma}\\lesssim Q_0, \\label{eq:Toomre} \\ea where $\\Omega,~c_s$, and $\\Sigma$ are the local angular frequency (disk is assumed to be in Keplerian rotation), sound speed, and surface density of the disk, and $Q_0\\approx 0.75-1.5$ (Kim, Ostriker, \\& Stone 2002; Boss 2002) is a threshold value of $Q$ at which GI sets in.  Nonlinear outcome of the GI depends on the thermodynamics of the disk, namely on its ability to cool rapidly (Gammie 2001; Johnson \\& Gammie 2003). Rapid cooling keeps pressure forces subdominant compared to the disk self-gravity in the non-linear stage of GI thus promoting the collapse of unstable parts of the disk and leading to the formation of bound objects. When the cooling is slow fragmentation is avoided and the disk settles to a quasi-steady regime of gravitoturbulence (Gammie 2001) characterized by dramatic perturbations of the disk surface density and fluid velocity (Durisen \\etal 2006). Gammie (2001), Rice \\etal (2003), and Rice, Lodato, \\& Armitage (2005) have demonstrated using simulations with fixed cooling time $t_{cool}$ that fragmentation requires \\ba \\Omega t_{cool}\\lesssim \\zeta, \\label{eq:frag_crit} \\ea where $\\zeta\\approx 3-12$ depending on the adiabatic index $\\gamma$ of the disk material, although in some situations $\\zeta$ can be much higher (Johnson \\& Gammie 2003).  Fragmentation resulting from GI is likely an important ingredient of the dynamics of accretion disks in the central regions of galaxies, around supermassive black holes (SMBHs; Paczynski 1978; Kolykhalov \\& Sunyaev 1980; Illarionov \\& Romanova 1988; Levin 2003; Goodman 2003;Goodman \\& Tan 2004; Thompson, Quataert, \\& Murray 2005), including the past disk around the black hole in the center of our Galaxy (Levin \\& Beloborodov 2003; Nayakshin 2005). On smaller scales, fragmentation of gravitationally unstable protoplanetary disks resulting in direct formation of gas giant planets has been put forward (Cameron 1978; Boss 1998) as an alternative to the conventional core instability model (Perri \\& Cameron 1974; Harris 1978; Mizuno 1980). Some simulations incorporating radiation transfer (Boss 2002; Mayer 2006) seem to demonstrate that planet formation by GI is possible. However, these results are in drastic conflict with the long radiative cooling times of optically thick protoplanetary disks (Rafikov 2005; hereafter R05) and the recent numerical results by another group (Cai \\etal 2006; Boley \\etal 2006).  To reconcile inefficient radiative cooling with fragmentation exhibited in their simulations Boss (2004) and Mayer \\etal (2006) have come forward with a suggestion that gravitationally unstable protoplanetary disks can be efficiently cooled by convection so that fragmentation becomes possible. The goal of this paper is to critically revise this possibility by setting in \\S \\ref{sect:up_limit} a lower limit on the cooling timescale of the disk in which energy is transported away from the midplane by convective motions. We then use this result in \\S \\ref{sect:fragmentation} to constrain the properties of the disks in which GI and convective cooling lead to fragmentation. ",
        "conclusions": "\\label{sect:concl}   We have investigated the possibility of rapid disk cooling by convection in the context of fragmentation of gravitationally unstable disks. We have shown that even the most extreme form of convective cooling (realized when $\\nabla=\\nabla_{ad}$) produces cooling rates which are only marginally higher than in the case of radiative energy transport. The major reason for the inefficient cooling of dense and cold convective disks is in the bottleneck caused by the need to eventually radiate from the photosphere the energy that has been transported there from the midplane, irrespective of how the latter has been done. Armed with this knowledge we demonstrate that inefficient cooling precludes direct formation of giant planets by GI in protoplanetary disks anywhere within several tens of AU from the parent star (analogous conclusion has been reached in R05). At the same time, disk cooling (convective or radiative) is fast enough to allow fragmentation of gravitationally unstable disks around black holes in the galactic centers, supporting the idea that the efficient star formation is possible in such disks."
    },
    "0609/astro-ph0609063_arXiv.txt": {
        "abstract": "The Pierre Auger Observatory in Malarg\\\"ue, Argentina, is designed to study the origin of ultrahigh energy cosmic rays with energies above $10^{18}$\\,eV.  The energy calibration of the detector is based on a system of four air fluorescence detectors. To obtain reliable calorimetric information from the fluorescence stations, the atmospheric conditions at the experiment's site need to be monitored continuously during operation.  One of the components of the observatory's atmospheric monitoring system is a set of four elastic backscatter lidar stations, one station at each of the fluorescence detector sites.  This paper describes the design, current status, standard operation procedure, and performance of the lidar system of the Pierre Auger Observatory. ",
        "introduction": "\\label{sec:intro}   The Pierre Auger Observatory in Malarg\\\"ue, Argentina, the world's largest facility to detect ultrahigh energy cosmic rays with energies above $10^{18}$\\,eV, is currently nearing completion.  Data collection started in January 2004, and the observatory has already accumulated an ultrahigh energy cosmic ray data set of unprecedented size.  Due to their extremely low flux, primary cosmic rays at ultrahigh energies cannot be observed directly.  Rather, they enter the atmosphere and interact with air molecules, inducing a cascade of secondary particles, called an extensive air shower.  The properties of the primary particle (energy, arrival direction, and chemical composition) have to be deduced from indirect measurements of the extensive air shower.  Two detector types have traditionally been used to measure air showers: surface detectors, which detect the particles of the air shower cascade that reach the ground; and air fluorescence detectors, which make use of the fact that the particles in the air shower excite air molecules, causing UV fluorescence.  Observing this UV light from air showers with photomultiplier cameras allows us to image the air shower development and obtain a nearly calorimetric energy estimate of the shower.  For both detector types, the atmosphere acts as the calorimeter and provides the large detector volume needed to measure the small cosmic ray flux at the highest energies.  The Pierre Auger Observatory is a hybrid detector which combines a large surface detector (SD) array of area $3000\\,\\mathrm{km}^{2}$ and a system of four fluorescence telescopes (FD) at the same site.  While the SD provides the well-defined aperture and the $100\\,\\%$ duty cycle needed to achieve high statistics, the FD provides, for a subset of showers, the calorimetric information needed to calibrate the SD.  For the calibration to be meaningful, the properties of the calorimeter, {\\it i.e.} the atmosphere, must be well-known.  Particulates in the atmosphere, Cherenkov radiation and the weather need to be factored into this calibration.  Fluorescence light from air showers is affected by the absorption and scattering on molecules and aerosols, and, in addition, some fraction of the Cherenkov radiation from the relativistic particles in the air shower is scattered into the detector.  This Cherenkov light cannot be separated from the fluorescence light, and so the contamination by Cherenkov light needs to be modeled and subtracted to give an accurate energy estimate for the cosmic ray shower. Moreover, the light transmission properties of the atmosphere are not constant over the large detector volume of the Pierre Auger Observatory, and equally important, they are not constant in time.  Changing weather conditions inevitably mean changing transmission properties.  Consequently, the properties of the atmosphere need to be monitored continuously.  The Pierre Auger Observatory has an extensive program to monitor the atmosphere within the FD aperture and measure atmospheric attenuation and scattering properties in the 300 to 400\\,nm sensitivity range of the FDs.  A summary of the atmospheric monitoring system can be found in~\\cite{Cester:2005}.  Within this system, a central role is played by a system of four elastic\\footnote{Here, the term {\\it elastic} refers to the light scattering process.  In elastic lidar applications, the return signal is measured at the same wavelength as the original laser signal.} backscatter lidar ({\\it li}ght {\\it d}etection {\\it a}nd {\\it r}anging) stations, one at each fluorescence site. system is currently under construction, with three out of four stations operating.  At each lidar station, a high-repetition UV laser sends short laser light pulses into the atmosphere in the direction of interest. The backscattered signal is detected as a function of time by photomultiplier tubes at the foci of parabolic mirrors.  Both the laser and the mirrors are mounted on a steering frame that allows the lidar to cover the full azimuth and elevation of the sky.  During each hour of FD data taking, the four lidars perform a routine scan of the sky over each FD.  The data provide information about the height and coverage of clouds as well as their depth and opacity, and the local aerosol scattering and absorption properties of the atmosphere.  In addition to this routine operation, the lidar system is used for real time monitoring of the atmospheric homogeneity between the FDs and selected cosmic ray events. For example, if a high energy ``hybrid'' event is observed with the SD and one or more FDs, the routine scan is interrupted and, within 2 to 4 minutes of the event detection, the lidar scans the atmosphere in the vicinity of the air shower reported by the FD.  This procedure is called \"shoot-the-shower\" (StS), and allows for a rejection of events where the light profile from the track is distorted by clouds or other aerosol non-uniformities that are not characterized well by the average hourly aerosol measurements. Both light reflection and opacity can distort the light profile.  This paper describes the design, standard operation procedure, and performance of the lidar system of the Pierre Auger Observatory.  It is organized as follows. After a brief description of the SD and FD detectors (Section\\,\\ref{sec:auger}), we give a summary of the relevant atmospheric parameters (Section\\,\\ref{sec:atmo}). Section\\,\\ref{sec:hardware} describes the current lidar hardware.  In Section\\,\\ref{sec:ops}, the daily operating procedure is summarized, including a description of the routine scan and the shoot-the-shower operation.  First results of the analysis of atmospheric parameters like cloud height and extinction are summarized in Section\\,\\ref{sec:analysis}.  A section on inelastic Raman lidar measurements and a summary conclude the paper.  \\begin{figure}[t] \\centering \\includegraphics[width=.7\\textwidth]{fig1.eps} \\caption{\\it Schematic layout of the Pierre Auger Observatory.} \\label{f:observatory} \\vskip1cm \\end{figure} ",
        "conclusions": "\\label{sec:conclusion}  Since March 2006, three of the four lidar stations are routinely operated each night of FD data taking by a designated lidar shift taker at the central campus. The fourth site is currently under construction and will start data taking in 2007.  The three operating lidars perform continuous and discrete scans of the local atmosphere, and the data are analyzed for cloud coverage, cloud height, and the optical depth of cloud layers.  An analysis of the aerosol scattering and absorption parameters will be added in the near future.  As the part of the lidar program that has the largest potential to interfere with the FD operation, the shoot-the-shower program is still in a testing phase.  Algorithms to maximize the efficiency of StS are currently under development.  This includes the implementation of software to filter background events that can erroneously trigger the StS, such as lightning.  As a possible future extension of the elastic backscatter lidar program, an inelastic Raman lidar is operated at the beginning of each data taking night.  Another future upgrade currently under consideration is the installation of new mirrors with a larger focal length to reduce the speed of the lidar optics.  The large field of view of the current system implies that the background rate is high and the laser beam enters the field of view of the mirrors very early.  To lower the amount of background light and overcome saturation at short range, a segmented mirror with a 1\\,m diameter and 1.1\\,m focal length is currently under development.   \\ack We are grateful to the following agencies and organizations for financial support: Comisi\\'on Nacional de Energ\\'{\\i}a At\\'omica and Gobierno de la Provincia de Mendoza, Argentina; INFN (Istituto Nazionale di Fisica Nucleare) (Italy); Slovenian Research Agency; National Science Foundation (USA) (contract numbers NSF-PHY-0500492, NSF-PHY-0134007); Department of Energy (DOE) Office of Science (USA) (DE-FG03-92ER40732)."
    },
    "0609/astro-ph0609580_arXiv.txt": {
        "abstract": "{} {We present a study of the neutron star X-Ray Transient Cen X-4. Our aim is to look for any evidence of irradiation of the companion with a detailed analysis of its radial velocity curve, relative contribution of the donor star and Doppler tomography of the main emission lines.} {To improve our study all our data are compared with a set of simulations that consider different physical parameters of the system, like the disc aperture angle and the mass ratio.} {We conclude that neither the radial velocity curve nor the orbital variation of the relative donor's contribution to the total flux are affected by irradiation. On the other hand, we do see emission from the donor star at H${\\alpha}$ and HeI $\\lambda$5876 which we tentatively attribute to irradiation effects. In particular, the H${\\alpha}$ emission from the companion is clearly asymmetric and we suggest is produced by irradiation from the hot-spot. Finally, from the velocity of the HeI $\\lambda$5876 spot we constrain the disc opening angle to $\\alpha=7^{\\circ}-14^{\\circ}$.} {} ",
        "introduction": "Soft X-Ray Transients (SXRTs), a subclass of Low Mass X-Ray Binaries, are interacting binary systems which alternate short periods (weeks to months) of high X-ray luminosity with long (several years) intervals of quiescence. During quiescence, such 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 companion stars (see Campana et al. 1998 for a review).  Cen X-4 is a well-known neutron star X-ray transient, from which two bright X-ray outburst were detected in 1969 (\\cite{Co69}) and in 1979 (\\cite{Ka80}). From then on it remained in the quiescent state. With an evolved 0.3 $\\msole$, K3-K7 star (\\cite{SNC93}, \\cite{To02}, D'Avanzo et al. 2005), Cen X-4 is one of the brightest SXRTs in quiescence, with $V = 18.7$ and low interstellar absorption ($A_V = 0.3$ mag). A 15.1 hr orbital period was determined from the sinusoidal variation of the optical light curve (Cowley et al. 1988; Chevalier et al. 1989; McClintock \\& Remillard 1990). Both the optical spectrum and emission line maps obtained with the Doppler tomography technique (\\cite{MH88}) show clear evidence of disc emission (\\cite{To02}, \\cite{PDA05}).  In general, the effect of heating of the secondary star leads to a change in absorption lines strength, with the consequent effect to move the effective centre of the secondary away from the centre of mass of the star. This causes phase-dependent distortions of the radial velocity from a sinusoidal fit, leading to incorrect determinations of the radial velocity amplitude.  In our previous work on Cen X-4 evidence of irradiation was proposed in light of a residual emission of H${\\alpha}$ and HeI from the companion star. We present here a search for evidence of irradiation in Cen X-4, comparing our observations presented in D'Avanzo et al. (2005) with the simulations obtained using the code presented in ~Mu\\~noz-Darias, Casares \\& Mart\\'inez-Pais (2005). ",
        "conclusions": "We have examined the H$\\alpha$ and HeI $\\lambda$5876 bright spots associated to the companion star in Cen X-4. The HeI spot is well centered on the line joining the two stars and we use its velocity to constrain the disc flaring angle to be 7-14$^{\\circ}$. On the other hand, the H$_{\\alpha}$ spot concentrates around the L1 point and is shifted towards the leading side of the Roche lobe. We tentatively explain this asymmetry through irradiation from the hot-spot and shielding by the gas stream.  We have also searched for evidence of irradiation effects in the radial velocity curve of the photospheric absorptions, by comparing our data with model simulations, but found none. This lead us to conclude that the donor's metallic lines are not affected by irradiation and hence the orbital parameters presented in previous papers (Torres et al. 2002, D'Avanzo et al. 2005) are safe from bias."
    },
    "0609/astro-ph0609313_arXiv.txt": {
        "abstract": "We apply the principle of energy conservation to the motion of the test particle in gravitational field by requiring that its energy, gained by gravitation, has to be balanced by decrease of its rest mass. Due to the change of mass in gravitational field Newton's force law between gravitating bodies is modified, too. With this modified force law we build up the the classical field theory of gravitation in which all relevant field quantities are in the definition domain $r\\in [0,\\infty)$ finite and positive. We show that under such circumstances, the energy release at any gravitational collapse is finite. On the other side, the energy conservation leads to an equation which relates the mass change of the test particle due to  gravitation and the metric of the corresponding gravitational field. The mass change in Newton's gravitational field lead to a remarkable simple metric which shifts, in contrast to the Schwarzschild metric, the horizon of events to the gravity center of the gravitational collapse.\\\\  KEYWORDS: energy conservation, naked singularity, black hole, dark matter ",
        "introduction": "It is generally accepted that if a star-like object with a sufficient high mass is collapsing over a certain limit radius, it becomes unstable and due to stronger and stronger gravitation forces it shrinks to a mathematical point. In Newtonian physics the final state of such gravitational collapse - the naked singularity - is a state of infinite mass density not separated from the outside word by a horizon of events and its formation is accompanied by radiating an infinite amount of heat. The general relativity offers the possibility of stable end state of collapsing mass body called black hole representing a mass object from which elm radiation can not escape. Beside the general feeling of physicists who have a psychological resistance against the existence of any singular state of matter in both cases one also faces difficulties with the physical interpretation of the final product of collapsing star. The classical naked singularity represents a highly unphysical state of matter and the existence of black hole in its general form is conditioned with the cosmic censorship hypothesis which appears not generally prove so far. Therefore, the endeavor to find alternatives to the classical naked singularity as well as to the black holes seems still to be justified. In what follows, we will investigate the possibility of creation of the naked singularity and the classical black hole by strict application principle of energy conservation to the gravitational interaction.  As is well-known, the {\\it energy} is one of the most fundamental concepts of physics and its conservation is evident in all its subdisciplines. This is why we apply the principle of energy conservation to the motion of the test particle in gravitational field by requiring that its energy, gained by gravitation, has to be balanced by the decrease of its mass. Consider, for simplicity, the static spherically symmetrical gravitational field of a large mass $M$ in which a test particle $m$ occurs at distance $r$ from $M$. We demand that the sum of the inertial {\\it as well as} the gravitational mass of the test particle at a spatial point $r$ and the work being done by its adiabatic translation ($v\\approx 0$) from infinity to $r$ must be equal to its rest mass $m_{\\infty}$ in infinity, i.e. $$m_{\\infty}=m(r)+1/c^2\\int_{\\infty}^r F(x)dx,\\quad\\eqno(a)$$ where $F(r)$ is the force acting on the test particle in gravitational field. Inserting Newton's law into Eq.(a) one obtains $$m(r)=m_{\\infty}\\gamma(r) =m_{\\infty}\\exp{(-\\lambda/r)};\\quad \\lambda=\\frac{GM}{c^2}\\quad\\eqno(b)$$ and the new modified Newton's force law reads $$F_N(r)=-\\frac{GMm_{\\infty}\\exp{(-\\lambda/r)}}{r^2}.\\quad\\eqno(c)$$ The factor $\\exp{(-\\lambda/r)}$, appearing in the Newton-like force law, has a remarkable consequence, namely when the test particle is approaching sufficiently near to a (point-like) gravitating mass, the force acting on it becomes weaker and weaker till it completely ceases at the gravity center. On the other side, $F_N(r)$ coincides asymptotically with common Newton's law in the regions where the field is weak ($\\lambda\\ll 1$).  In our approach, the particle mass in gravitational field apparently depends on its position. On the other side, it is generally assumed that particle rest mass is, in classical gravity as well as in general relativity, constant. This assumption, as stressed by Dicke \\cite{DIC}, is very partially supported by experiment. Especially, this constancy has never been tested in strong gravitational fields. The concept of variable rest mass, motivated by different reasons, is not new. It appears, e.g. in Dicke reformulation \\cite{D} of Brans-Dicke theory \\cite{BD}, in Hoyle and Narlikar's theory of gravitation \\cite{HO}, in Malin \\cite{MA} cosmological theory of variable rest masses, in Beckenstein theory of rest mass field \\cite{BE} and in Vondr\\'a\\v cek \\cite{V} Mach's theory of gravitation, to mention only a few.  Since at each spatial point of gravitational field the force acting on the test particle is given by Eq.(c) one can consider the neighborhood of $M$ as a classical force field, i.e. one can determine its intensity, its potential, its field energy density, and its total field energy. For the Newton-like field, all these field quantities assume everywhere finite and positive values. As a consequence, the energy conservation, the total energy released by any gravitational collapse can not be larger than $nm_{\\infty}c^2$, $n$ being the number of particles participating on the gravitational collapse. Since the energy released by a gravitational collapse is restricted no naked singularity, in classical sense, can be formed in nature. Moreover, a star-like objects that is larger than approximately $2M_{sun}$-Chandrasekhar limit- might exist in equilibrium due to weakening of gravitation toward the gravity center (see \\cite{GM}).  There is a link between the Newton-like gravitation theory and the geometrical structure of spacetime. To show this we consider a spherically symmetrical, asymptotically flat gravitation field with local Lorentz reference frame co-moving with the freely falling elevator. The velocity of this system is changing  and, accordingly, the corresponding relativistic length contraction and time dilatation. Therefore, one can ascribe to each spatial point of gravitational field the special-relativistic factor $\\gamma=\\sqrt{1-v^2/c^2}$. Knowing this factor as a function of position $r$ the line element in the reference frame of rest observer turns out to be $$ds^2=\\gamma^{-2} dr^2+r^2(d\\theta^2 +\\sin^2 \\theta d\\phi^2)-\\gamma^{2}c^2dt^2,\\qquad \\gamma=\\sqrt{1-\\frac{v^2}{c^2}} \\quad\\eqno(d)$$ The special-relativistic energy conservation for a freely falling particle leads to the formula $$\\frac{m_{\\infty}c^2\\gamma(r)}{\\sqrt{1-v^2/c^2}}=m_{\\infty}c^2,\\quad\\eqno(e)$$ where $m_{\\infty}$ is the mass of particle in infinity. This formula, which tell us that the total energy at each spatial point of gravitational field is constant equal to $m_{\\infty}c^2$, represents a bridge between mass decrease function (MDF, for short) $\\gamma(r)$ of a particle in gravitational field and its geometrical structure. Inserting MDF of the Newton-like field $\\gamma(r)=\\exp{(-\\lambda/r)}$ into Eq.(e) then the line element (d) becomes the form $$ds^2=\\exp{(\\frac{2\\lambda}{r})} dr^2+r^2(d\\theta^2 +\\sin^2 \\theta d\\phi^2)-\\exp{(-\\frac{2\\lambda}{r})}c^2dt^2.\\quad\\eqno(f)$$ Apparently, this line element coincides for weak field ($\\lambda/r\\ll 1)$ asymptotically with the Schwarzschild one.  The metric, associated with the Newton-like field, represents a non-vacuum solution of Einstein's equations. In order to determine the source terms assigned to this metric we solved the inverse problem of general relativity, i.e. to find the source terms given the metric. The (tt)-component of energy-momentum, assigned to the metric (f) represents a regular function of $r$ assuming at $r=0$ a finite value.  As is common knowledge, gravitationally collapsing object of sufficient mass are doomed to form black holes defined by an event horizon within which resides the singularity of general relativistic equations. Characteristic feature of the metric associated with the Newton-like field is that, in the classical black hole model, the surface of events and the infinite red shift are shifted into center of a collapsing mass, which implies the non-existence of common black-hole. Hence, strict application of energy conservation has at least to important consequences: the impossibility of forming the naked singularity and the non-existence of black holes.  It appears that the 'field' and 'geometrical' aspects of gravitation are mutual related. Given MDF one can determine the geometry of spacetime where the gravitational field occurs, and given the geometry one can determine the corresponding MDF and so its field quantities.  The organization of this article is as follows. In Section 2, we derive MDF and the modified force law from the requirement of the constancy of the energy content of the test particle during its motion in gravitational field. In Section 3, we describe properties of the Newton-like field which arises if one inserts into Eq.(a) the common Newton force law. In Section 4, we determine the metric associated with the Newton-like field and calculate the source terms of Einstein's equations by inserting components of metric tensor into left-hand sides of these equations. In Section 5, we study the actual gravitational fields and metrics associated with them. In section 6, we  put forward the hypothesis concerning gravitation as a force field embedded into metric determined by the given mass decrease function. ",
        "conclusions": "Taking into account what has been said so far we put forward the following hypothesis concerning gravitation. We assume that the neighborhood of a gravitating object is the space where the rest mass of a particle is transformed into its kinetic energy, according to Einstein's equivalence principle between mass and energy, whereby its total energy remains conserved. Due to transformation of mass into energy the velocity of particles in gravitational field are changed and with this change also the local Lorentz factor and the metric. Spacetimes in regions, where gravitation occurs, appears to rest observer generally as curved. Their metrics are determined by the density of the ordinary matter as well as by the energy density of the force fields. Gravitational field appears as a force field embedded into metric determined by its mass decrease function.  The presented picture of gravitation has at least two important consequences:\\\\ (i) {\\it In classical gravitation theory}. When a star run out of nuclear fuel, the only force left to sustain it against gravity is the pressure associated with the zero-point oscillation of its constituent fermions. This is valid if the gravitational force obeys Newton's law. If one take, instead of Newton's, the Newton-like force law then the force is weakening near the gravity center as a consequence of the exhausting of particle rest mass. Due to principle of the energy conservation, therefore the total energy released by the gravitational collapse can not be larger than the sum of all masses participating at it. The strict application of energy conservation prevent so the forming of the naked singularity. The energy balance for an isolated collapse (nothing is put into or taken out of it), i.e. energy release plus the internal energy of masses has to be during the whole process constant. On the other side, the weakening of force  near to the gravity center in the Newton-like field  makes it possible the existence of star-like objects in equilibrium which are more collapsed than neutron stars (see \\cite{GM}. The work done by translation of a test particle from infinity to the center of gravity is, in contrast to the Schwarzschild field, for the Newton-like field just equal to its rest mass in infinity.\\\\ (ii) {\\it In the general relativity}. The strict application of energy conservation yields an important relation between a gravitational force field and the associated spacetime geometry. The metric associated with the Newton-like field is remarkable simple $g_{tt}=\\exp{(-2\\lambda/r)}$ and $g_{rr}=\\exp{(2\\lambda/r)}$. It coincides for weak field ($\\lambda/r\\ll 1)$ asymptotically with the Schwarzschild metric and represents a non-vacuum solution of Einstein's equations. The (tt)-component of energy-momentum represents assigned to this metric represents a regular function of $r$ assuming at $r=0$ a finite value. Its integral over the whole range of radius $r$ is likewise a finite number. In the common black hole model, the metric associated with the Newton-like field shifts the surface of an infinite red shift and the horizon of events into the center of gravity, which implies the non-existence of the 'classical' black holes.\\\\  It is noteworthy that the dependence of mass on position satisfies one requirement of Mach' principle, namely that that mass of particle should depend on neighboring mass distribution (see, e.g. \\cite{N},\\cite{E}).  The application of the principle of energy conservation to the motion of particle in gravitational field leads to a modified force law and implies a remarkable simple metric with the important consequences. It prevents the naked singularity and leads to a metric preventing also the forming of classical black holes. The energy of gravitational field can be interpreted as one component of dark matter.\\\\   The aim of this article was only to outline the basic ideas concerning the strict application of energy conservation in the gravitation theory, therefore everything is simplified and many important issues remained open which will be subject of subsequent works."
    },
    "0609/astro-ph0609125_arXiv.txt": {
        "abstract": "We develop a new method to predict the density associated with weak lensing maps of (un)relaxed clusters in a range of theories interpolating between GR and MOND (General Relativity and Modified Newtonian Dynamics). We apply it to fit the lensing map of the bullet merging cluster 1E0657-56, in order to constrain more robustly the nature and amount of collisionless matter in clusters {\\it beyond} the usual assumption of spherical equilibrium (Pointecouteau \\& Silk 2005) and the validity of GR on cluster scales (Clowe et al. 2006).  Strengthening the proposal of previous authors we show that the bullet cluster is dominated by a collisionless -- most probably non-baryonic -- component in GR as well as in MOND, a result consistent with the dynamics of many X-ray clusters.  Our findings add to the number of known pathologies for a purely baryonic MOND, including its inability to fit the latest data from the Wilkinson Microwave Anisotropy Probe.  A plausible resolution of all these issues and standard issues of Cold Dark Matter with galaxy rotation curves is the ``marriage\" of MOND with ordinary hot neutrinos of 2eV. This prediction is just within the GR-independent maximum of neutrino mass from current $\\beta$-decay experiments, and is falsifiable by the Karlsruhe Tritium Neutrino (KATRIN) experiment by 2009. Issues of consistency with strong lensing arcs and the large relative velocity of the two clusters comprising the bullet cluster are also addressed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609131_arXiv.txt": {
        "abstract": "We carry out three-dimensional, high resolution (up to $1024^2\\times 256$) hydrodynamic simulations of the evolution of vortices in vertically unstratified Keplerian disks using the shearing sheet approximation.  The transient amplification of incompressible, linear amplitude leading waves (which has been proposed as a possible route to nonlinear hydrodynamical turbulence in disks) is used as one test of our algorithms; our methods accurately capture the predicted amplification, converges at second-order, and is free from aliasing. Waves expected to reach nonlinear amplitude at peak amplification become unstable to Kelvin-Helmholtz modes when $\\mid W_{\\rm max}\\mid\\gtrsim \\Omega$ (where $W_{\\rm max}$ is the local maximum of vorticity and $\\Omega$ the angular velocity). We study the evolution of a power-law distribution of vorticity consistent with Kolmogorov turbulence; in two-dimensions long-lived vortices emerge and decay slowly, similar to previous studies. In three-dimensions, however, vortices are unstable to bending modes, leading to rapid decay. Only vortices with a length to width ratio smaller than one survive; in three-dimensions the residual kinetic energy and shear stress is at least one order of magnitude smaller than in two-dimensions.  No evidence for sustained hydrodynamical turbulence and transport is observed in three-dimensions. Instead, at late times the residual transport is determined by the amplitude of slowly decaying, large-scale vortices (with horizontal extent comparable to the scale height of the disk), with additional contributions from nearly incompressible inertial waves possible. Evaluating the role that large-scale vortices play in astrophysical accretion disks will require understanding the mechanisms that generate and destroy them. ",
        "introduction": "Many aspects of the structure and evolution of astrophysical accretion disks are controlled by the mechanism for angular momentum transport. When the disk is dynamically coupled to a weak magnetic field, the magnetorotational instability (MRI, e.g., Balbus \\& Hawley 1998) generates MHD turbulence which is a vigorous source of angular momentum transport (for a recent review, see Balbus 2003). However, in neutral or very weakly ionized disks (e.g., proto-planetary disks, and dwarf nova systems in quiescence) the magnetic field may be decoupled from the gas, and the MRI may not operate (Salmeron \\& Wardle 2005; Fromang, Terquem, \\& Balbus 2002; Gammie \\& Menou 1998).  It these cases, it is important to investigate purely hydrodynamical mechanisms for the transport of angular momentum.  Although Keplerian shear flows are linearly stable according to the Rayleigh criterion, it has long been supposed that there exists some mechanism that produces turbulence and transport via nonlinear instabilities at high Reynolds number. To date, however, no mechanism for such instabilities has been explicitly demonstrated, and no evidence for self-sustained hydrodynamical turbulence has been found in direct numerical simulations of Keplerian disks at the same Reynolds number needed to capture known nonlinear instabilities in Cartesian shear flows (Balbus, Hawley, \\& Stone 1996; Hawley, Balbus \\& Winters 1999).  Even if nonlinear instabilities exist in Keplerian shear flows, Lesur \\& Longaretti (2005) argue that current simulations already show the associated transport would be negligable.  Recently, new interest has focused on the role that transient growth of incompressible, leading waves might play in generating hydrodynamical turbulence (Chagelishvili et al. 2003; Umurhan \\& Regev 2004; Yecko 2004). The amplification factor of incompressible (vortical) waves is of order $(k_{x,0}/k_y)^2$, where $k_{x,0}$ and $k_y$ are the initial wavenumbers in the radial and azimuthal directions respectively (Johnson \\& Gammie 2005a, hereafter JGa). For leading waves that are initially tightly wound $\\mid k_{x,0}\\mid/k_y \\gg 1$, large amplification factors are possible, potentially leading to nonlinear amplitudes. This has been suggested as one step in the route to self-sustained hydrodynamical turbulence in disks (Afshordi et al 2005; Mukhopadhyay et al. 2005; but see Balbus \\& Hawley 2006). One purpose of this paper is to investigate the effect of transient growth of leading waves in fully three-dimensional (3D) compressible hydrodynamical disks.  On the other hand, there is continued interest in the role that vortices play in disks.  There is particular interest in the context of proto-planetary disks, where it has been suggested that vortices play a role in trapping dust and aiding in grain-growth (e.g., Barge \\& Sommeria 1995; Adams \\& Watkins 1995; Johansen, Anderson, \\& Brandenburg 2004). In addition, if vortices produce outward angular momentum transport they may play an important role in the structure and evolution of proto-planetary disks as well as other weakly ionized accretion disks (e.g., Li et al. 2001; Barranco \\& Marcus 2005, hereafter BM; Johnson \\& Gammie 2005b, hereafter JGb).  A variety of authors have studied the evolution of random vorticity in both 2D and 3D, and in both incompressible and compressible gases.  Long-lived, anticyclonic (with rotation opposite to the background shear) vortices have been found in 2D Keplerian disks (e.g., Godon \\& Livio 1999, 2000; Umurhan \\& Regev 2004; JGb).  The work of JGb is particularly relevant to this study in that they considered the evolution of a spectrum of vorticity consistent with Kolmogorov turbulence in fully compressible 2D simulations. In every case, they found the kinetic energy and transport decays from the initial conditions (sustained turbulence was not evident), although the rate of decay was affected by factors such as the horizontal extent of the domain.  These authors argued that compressibility significantly increases the outward angular momentum transport during the decay. Fully 3D simulations of the evolution of vorticity in stratified disks in the anelastic approximation have been presented by BM; they found vortices survive (and in fact, are produced) in the upper regions of the disk above a few scale heights.  In the midplane, three-dimensional instabilities destroy vortices and reduce transport.  Here we extend the work of JGb and BM by studying vortices in fully {\\em compressible} unstratified disks in 3D. We will report on 3D studies of stratified disks elsewhere.  This paper is organized as follows. In Section 2 we describe the numerical methods and shearing sheet model used in our simulations. As a test, we study the evolution of both incompressible and compressible planar shearing waves in two-dimensions in section 3; we point out that at large amplitude vortical waves become Kelvin-Helmholtz unstable.  In section 4 we discuss the evolution of random vortical perturbations in both 2D and 3D. We summarize and conclude in section 5. ",
        "conclusions": "Using the 3D version of the Athena code (Gardiner \\& Stone 2005; 2006), we have carried out hydrodynamical simulations of the evolution of both planar waves and random vortical perturbations in unstratified Keplerian disks using the shearing sheet approximation. Our results can be summarized by the following four points.  (1) Our numerical methods reproduce the evolution of the amplitude of both compressible and incompressible (vortical) plane waves predicted by linear theory with negligible dissipation as long as the numerical resolution is 16 grid points per wavelength or larger.  At lower resolution, the wave amplitude is smoothly damped.  Significantly, there is no evidence of aliasing of trailing into leading waves as they are sheared down to the grid resolution.  We have demonstrated the amplitude error in plane waves in the shearing sheet converges at better than second-order with our methods.  (2)  We have shown that incompressible plane waves become KH unstable and are destroyed when $\\mid W_{\\rm max}\\mid\\gtrsim\\Omega$, the condition that the growth rate of the KH instability exceeds the angular velocity. Whether the wave becomes KH unstable only depends on the initial vorticity amplitude in the initial conditions. In this case the planar shearing wave amplitude is damped well before it gets amplified by the transient growth mechanism.  (3) In 2D, the evolution of large amplitude random vorticity perturbations follows the results reported by others (Godon \\& Livio 1999, 2000; Umurhan \\& Regev 2004; JGb).  In particular, coherent, large-scale (horizontal extent larger than the scale height), anticyclonic vortices survive and decay slowly.  At late times the dimensionless angular momentum flux is of the order $10^{-3}$ averaged over $100<\\Omega t<200$, consistent with JGb.  This flux is dominated by the residual motions in the large-scale anticyclonic vortices.  (4) In 3D, the evolution of large amplitude random vorticity perturbations is similar to the 2D case, except the elliptic instabilities destroy vortex columns whose vertical extent exceeds their horizontal extent. This greatly increases the rate of decay of kinetic energy and stress. The resulting volume averaged shear stress $\\alpha$ at late-times is at least one order of magnitude smaller than that for the 2D case, and is probably associated with linear amplitude, low-$m$ inertial waves that remain in the box, and which decay slowly.  The destruction of planar vortical waves by the KH instability has implications for transient amplification as a means to drive strong turbulence. Although this might appear attractive as a feedback mechanism to drive turbulence (Chagelishvili et al. 2003; Yecko 2003; Afshordi, Mukhopadhyay, \\& Narayan 2005), we find instability destroys the wave and results in rapid decay of kinetic energy, rather than re-seeding new leading waves. We speculate that most of the energy in the vortical wave is converted into compressible modes by the KH instability, which are not amplified by shear. Moreover, Balbus \\& Hawley (2006) found that these planar vortical wave solutions are actually exact to all orders and cannot serve as a route to self-sustained turbulence.   Our results show no evidence for sustained hydrodynamical turbulence in the shearing sheet.  While there is non-zero transport, it is associated with large scale vortices that are introduced in the initial conditions, or which emerge from mode interactions during the nonlinear decay phase.  At late times the Reynolds stress is oscillatory, which may indicate long-lived linear amplitude inertial waves also contribute to the time-averaged shear.  To understand the relevance of large-scale vortices to the dynamics of astrophysical disks, it will be important to understand how they are generated and destroyed. BM have shown that such vortices can be generated in stratified disks, but given the importance of compressible waves on the decay of vortices (JGb), it is important to investigate vortex production and evolution in fully compressible stratified disks."
    },
    "0609/astro-ph0609307_arXiv.txt": {
        "abstract": "The temporal variability of the telluric sodium layer is investigated by analyzing 28 nights of data obtained with the Colorado State University LIDAR experiment. The mean height power spectrum of the sodium layer was found to be well fit by a power law over the observed range of frequencies,  $10\\ \\mu$Hz to $4$ mHz. The best fitting power law was found be be $10^\\beta\\nu^\\alpha$, with $\\alpha = -1.79 \\pm 0.02$ and $\\beta = 1.12 \\pm 0.40$.  Applications to wavefront sensing require knowledge of the behavior of the sodium layer at kHz frequencies. Direct measurements at these frequencies do not exist. Extrapolation from low-frequency behavior to high frequencies suggests that this variability may be a significant source of error for laser-guide-star adaptive optics on large-aperture telescopes. ",
        "introduction": "The advances made in adaptive optics (AO) over the last decade have enabled many ground-based optical astronomical projects that were never before possible\\cite{beck}. Most of these studies have been performed with natural guide star AO systems (NGSAO), which require a bright star near to the desired target. While there are many examples of strikingly successful studies done with NGSAO, the fraction of the sky accessible to them is rather small. In order to access the whole sky, one requires the ability to sample the atmospheric distortion anywhere, not just toward the direction of a bright star. The telluric sodium layer is a layer of sodium atoms that is typically present at an altitude of $\\sim 90$ km over an extent of $\\sim 15$ km. By illuminating the layer with a laser tuned to the wavelength of the $D_2$ transition ($589$ nm), one may excite emission from these sodium atoms. The resulting bright spot acts as an artificial guide star, and enables the operation of laser guide star AO (LGSAO) systems.  The development of a LGSAO system is of particular importance to the Thirty-Meter telescope (TMT). In order to take full advantage of the potential spatial resolution of the large aperture, it is imperative for many of the TMT's instruments to be fed by the proposed LGSAO system, the Narrow Field InfraRed Adaptive Optic System (NFIRAOS). The development of NFIRAOS presents many technical challenges. As discussed in O'Sullivan et al. (2000)\\cite{osull}, one of these challenges---the refocusing of the telescope on the moving laser guide star spot---is due to the variability of the sodium layer. Though the physical mechanism is not entirely understood, atmospheric gravity wave propagating upward in the mesopause region is thought to be the most likely explanation. A typical Shack-Hartmant wavefront sensor will determine the centroid of a illuminated spot on a CCD, and will therefore respond to the instantaneous mean height of the sodium layer. A change in the mean height, $\\Delta h$, of the sodium layer produces a RMS wavefront error, averaged over the aperture, \\begin{equation} \\sigma=\\frac{1}{8\\sqrt{3}}\\left(\\frac{D}{h}\\right)^2\\Delta h, \\end{equation} where $D$ is the diameter of the telescope aperture, and $h$ is the instantaneous density-weighted mean height of the sodium layer \\cite{herr}. Note that $\\sigma$ is proportional to the square of the diameter of the telescope. Thus, for the TMT, we expect the effect to be an order of magnitude larger than for the largest optical telescopes presently in operation. For the TMT, $\\sigma = 8$ nm for every meter change of the mean height.   While an AO system can use laser and natural guide stars to correct for focus error caused by the atmosphere and variations in the height of the sodium layer, it can not do so instantaneously. Therefore, the correction applied by the AO system will be imperfect, and a residual wavefront error (RWFE) will remain. While the mean height of the sodium layer is known to be variable on the order of several kilometers over timescales of minutes, LGSAO systems sample the atmosphere many times per second. Therefore, it is necessary to know how the mean height of the sodium layer will vary on millisecond timescales in order to estimate the RWFE. In order to assess the variability of the sodium layer on short time scales we examined data taken by the Colorado State University (CSU) Light Detection and Ranging (lidar) experiment over 89 nights in 2003. We determined the typical RWFE expected by multiplying the mean height power spectrum of the sodium layer  by a model rejection transfer function determined specifically for the TMT. ",
        "conclusions": "The data for each night with at least one hour of acceptable quality (76 of 89) were analyzed. It was found that nights with sections of usable data that were less than $7$ hours (48 nights) had a shallower slope with respect to those nights with longer sections of usable data (28 nights). Using just the longer nights we find the best-fit power law has the parameters $\\alpha=-1.79\\pm0.02$ and $\\beta=1.12\\pm0.40$. We found that RWFE for those nights has a median value of $20$ nm and a $90^{th}$ percentile value of $34$ nm. This effect is an important contribution to the overall error budget of NFIRAOS, and more generally, will likely be for most, if not all, of the LGSAO systems designed for the next generation of large ground-based telescopes.  This analysis was limited by the relatively long integration times of the CSU lidar experiment, which in turn set a low Nyquist frequency. The variability of the sodium layer height in the frequency range observed is mainly due to the upward passing of atmospheric waves. Atmospheric turbulence at the sodium layer height can perturb the mean height of the sodium layer on time scales of seconds or shorter. However, changes caused by atmospheric turbulence are likely to be small (cm or less), unlike the kilometer-scale changes reported here. The assumption that the mean height power spectrum can be fit by a single unbroken power law from $\\mu$Hz to kHz is perhaps na\\\"ive, but reasonable. The actual high-frequency behavior of the mean sodium height power spectrum, of course, can only be determined by further experimentation. The temporal resolution could be improved by designing a lidar experiment with either a larger collecting telescope, or by sacrificing spatial resolution."
    },
    "0609/astro-ph0609594_arXiv.txt": {
        "abstract": "We present optical spectroscopy of a sample of 77 luminous AGN and quasars selected on the basis of their mid-infrared colors. Our objects are selected from the {\\em Spitzer} Extragalactic First Look Survey and SWIRE XMM-LSS fields, with a typical 24$\\mu$m flux density of 5mJy. The median redshift is 0.6, with a range of $\\sim 0.05-4$. Only 33\\% (25/77) of these objects are normal type-1 quasars, with no obscuration. 44\\% (34/77) are type-2 objects, with high-ionization, narrow emission lines, 14\\% (11/77) are dust-reddened type-1 quasars, showing broad lines but a dust-reddened or unusually weak quasar continuum. 9\\% (7/77) show no sign of an AGN in the optical spectrum, having either starburst spectra or spectra which could be of either starburst or LINER types. These latter objects are analogous to the X-ray detected population of AGN with weak or non-existent optical AGN emission (the ``XBONGs''). 21 of our objects from the SWIRE field fall within moderately-deep XMM exposures. All the unobscured quasars, and about half the obscured quasars are detected in these exposures. This sample, when taken together with other samples of {\\em Spitzer} selected AGN and quasars, and results from X-ray studies, confirms that obscured AGN dominate the AGN and quasar number counts of all rapidly-accreting supermassive black hole systems, at least for $z\\stackrel{<}{_{\\sim}}4$. This implies a high radiative efficiency for the black hole accretion process. ",
        "introduction": "Quasars whose optical emission is hidden by dust in the optical and gas in the soft X-ray have always been hard to find. However, a large type-2 (hidden) quasar population has long been predicted, both from the statistics of Seyfert-1 and 2 galaxies, and from those of radio galaxies and radio-loud quasars. A large population of hidden AGN is also predicted from models of the X-ray background (e.g.\\ Comastri et al.\\ 1995; Worsley et al.\\ 2005), though most models predict these AGN to be of Seyfert rather than quasar luminosity. Some success in finding dust obscured quasars has been obtained using selection based on the Two Micron All Sky Survey (2MASS) to find lightly dust-reddened (rest frame extinctions $A_V \\sim 2$) quasars, based either on 2MASS colors alone (Cutri et al.\\ 2001), or on 2MASS combined with a radio detection and very red optical-IR colors (Glikman et al.\\ 2005). Only recently, however, have large samples of type-2 quasars, with rest-frame $A_V\\stackrel{>}{_{\\sim}}5$ been found. Zakamsaka et al.\\ (2003) used the Sloan Digital Sky Survey (SDSS) to find objects with the high-ionization narrow-line spectra characteristic of type-2 AGN. Combining {\\em ISO} and 2MASS data allowed Haas et al.\\ (2004a) and Leipski et al.\\ (2005) to find significant numbers of obscured AGN, including several type-2 quasars. {\\em Spitzer} data has been used to find highly-obscured AGN, both through mid-infrared color selection (Lacy et al.\\ 2004 [hereafter Paper I], 2005a; Stern et al.\\ 2005; Alonso-Herrero et al.\\ 2005; Polletta et al.\\ 2006) and joint radio-infrared selection (Mart\\'\\i nez-Sansigre et al.\\ 2005, 2006; Weedman et al.\\ 2006; Donley et al.\\ 2005). Type-1 quasars selected through joint Spitzer and optical colors and/or morphologies have been targeted by several groups (Brown et al.\\ 2005; Hatziminaoglou et al.\\ 2005; Siana et al.\\ 2006). Also, the mid-infrared properties of X-ray selected AGN have been discussed by Treister et al.\\ (2005) and Barmby et al.\\ (2005).  Mid-infrared selection of AGN is effective because dust heated by the AGN forms a ``calorimeter'', whose luminosity is not as strongly affected by obscuration as that in most other wavelength bands (except the radio) (e.g.\\ Spinoglio et al.\\ 1995; Meisenheimer et al.\\ 2001; Ogle et al.\\ 2006). Richards et al.\\ (2006) show that the bolometric correction in the mid-infrared varies little for different subsamples of type-1 quasars (e.g.\\ optically or infrared luminous, optically-red or optically blue), or luminosity, and that about 15\\% of the bolometric luminosity of a quasar is emitted at mid-infrared wavelengths. Integration over the Richards et al.\\ (2006) mean quasar SED as a whole suggests that about 40\\% of the emission from a typical quasar is reprocessed by dust in the infrared. This means that mid-infrared selection of AGN is probably a good proxy for selection on the basis of bolometric luminosity, at least for those objects which are not so highly obscured that their mid-infrared luminosity is significantly supressed. Much of the dust emission is thought to arise in the ``obscuring torus'' required by Unified Schemes (e.g.\\ Antonucci 1993), though it is becoming increasingly clear that host dust in the narrow line region of AGN also contributes significantly (e.g.\\ Mason et al.\\ 2006). AGN are an important global contributor to the mid-infrared luminosity density. Brand et al.\\ (2006) show that the luminosity density of mid-IR emission in samples of MIPS 24$\\mu$m sources selected at mJy levels is dominated by AGN, only at submJy levels do star forming galaxies begin to dominate.  In this paper, we focus on optical spectroscopic follow-up of AGN candidates selected according to the mid-infrared color criteria presented in Paper I. By selecting relatively bright objects from wide-area {\\em Spitzer} surveys we are able to find many obscured AGN with mid-infrared luminosities typical of quasars at moderate redshifts. These objects provide a good complement to the mostly lower-luminosity or higher redshift AGN found in surveys such as the Great Observatories Origins Deep Survey (GOODS) and the Chandra Deep fields. Preliminary results from the spectroscopic follow-up of the sample presented in Paper 1 were given in Lacy et al.\\ (2005a). In this paper, we describe the selection of samples of mid-infrared selected AGN candidates flux-limited at 24$\\mu$m from both the Extragalactic First Look Survey (XFLS) and the {\\em Spitzer} Wide-area Infrared Extragalactic Survey (SWIRE) XMM-LSS field (Lonsdale et al.\\ 2005), and their spectrosopic follow-up observations. We classify the optical spectra according to the type of AGN, and use them to place a lower bound on the fraction of obscured quasars and luminous Seyfert galaxies. We also investigate the X-ray fluxes of those objects which fall on archival observations made with the X-ray Multiple Mirror telescope ({\\em XMM-Newton}). A more detailed analysis of the Spectral Energy Distributions (SEDs) and X-ray properties of the AGN will be carried out in a later paper (Petric et al.\\ 2006).  We assume a cosmology with $\\Omega_{\\rm M}=0.3$, $\\Omega_{\\Lambda}=0.7$ and $H_0=70 {\\rm kms^{-1}Mpc^{-1}}$. ",
        "conclusions": "The results presented in this paper, taken together with studies of higher redshift objects selected from mid-infrared deep fields and objects found using joint mid-infrared -- radio selection show that a large population of luminous obscured AGN exist at $z\\stackrel{>}{_{\\sim}}0.3$, including a significant population with bolometric luminosities above the Seyfert-quasar divide. We find that the obscured population dominates in number density over the unobscured one over a wide range in AGN luminosity, from Seyferts to luminous quasars. Comparison with Alonso-Herrero et al.\\ (2006) suggests little variation in the obscured to unobscured ratio with redshift for AGN with strong mid-infrared emission.  Our results contrast with X-ray studies, which typically find that obscured AGN are confined to low redshifts ($z\\stackrel{<}{_{\\sim}} 1$) and luminosities below the Seyfert-quasar divide $L_X \\stackrel{<}{_{\\sim}} 10^{44} {\\rm ergs^{-1}}$ (e.g.\\ Trump et al.\\ 2006), suggesting that high luminosity and high redshift obscured quasars are missing from current X-ray based surveys, presumably due to limitations of areal coverage and depth. Although objects with columns $N_H \\sim 10^{23}{\\rm cm^2}$ can be found in deep X-ray surveys, locally $\\sim 50$\\% of AGN are Compton thick, with columns $> N_H \\sim 10^{24}{\\rm cm^{-2}}$, and the equivalents of these objects at high redshifts would be very hard to find in the X-ray, even allowing for the positive k-correction. This has important implications for models of the X-ray background (e.g.\\ Worseley et al.\\ 2006). However, it is  clear that all the techniques used to find obscured AGN are incomplete. Mid-infrared selection presumably fails for weak AGN, for highly obscured AGN, and at high redshifts. Our highest redshift objects are selected not on the basis of their hot dust emission, but on the basis of dust-reddened accretion disk emission, so we will inevitably miss the most obscured objects at $z\\stackrel{>}{_{\\sim}}3$.  For most objects, classification of galaxy type (AGN or starburst) based on optical emission line ratios and classification based on mid-infrared IRAC color appear to agree fairly well. The mid-infrared criterion is seen to fail for $z\\stackrel{<}{_{\\sim}} 0.3$ AGN with strong PAH emission in the IRAC 8.0$\\mu$m band from an accompanying starburst. There are also some objects with AGN-like IRAC colors which lack the high-ionization emission line spectra of AGN. For most of these, their nature (AGN or starburst) remains to be resolved.  The discovery of a large population of obscured AGN has important implications for accretion efficiency and black hole demographics. The mass density of black holes in the Universe today may be accurately estimated using the bulge-luminosity black hole mass relation, and the amount of accretion luminosity arising from unobscured accretion may be estimated from the optical quasar luminosity function (Soltan 1982). Recent estimates using this technique place the radiative efficency of accretion close to that expected from a non-spinning (Schwartzschild) black hole (e.g.\\ Yu \\& Tremaine 2002). However, if most accretion is actually hidden, to maintain the same mass density of black holes today in the face of more accretion activity requires a {\\em higher} radiative efficiency, which can only be obtained from rapidly spinning Kerr black holes."
    },
    "0609/astro-ph0609077_arXiv.txt": {
        "abstract": "We present a spectroscopic survey of the newly-discovered Canes Venatici dwarf galaxy using the Keck/DEIMOS spectrograph.  Two stellar populations of distinct kinematics are found to be present in this galaxy: an extended, metal-poor component, of half-light radius $7\\mcnd8^{+2.4}_{-2.1}$, which has a velocity dispersion of $13.9^{+3.2}_{-2.5} \\kms$, and a more concentrated (half-light radius $3\\mcnd6^{+1.1}_{-0.8}$) metal-rich component of extremely low velocity dispersion.  At 99\\% confidence, the upper limit to the central velocity dispersion of the metal-rich population is $1.9\\kms$. This is the lowest velocity dispersion ever measured in a galaxy. We perform a Jeans analysis on the two components, and find that the dynamics of the structures can only be consistent if we adopt extreme (and unlikely) values for the scale length and velocity dispersion of the metal-poor population.  With a larger radial velocity sample and improved measurements of the density profile of the two populations, we anticipate that it will be possible to place strong constraints on the central distribution of the dark matter in this galaxy. ",
        "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.}  Kinematical analyses of the dwarf satellite galaxies show them to be the astrophysical structures with the highest measured fraction of dark matter, with some claims of the fraction of total to luminous matter exceeding $M/L \\sim 500$ \\citep{kleyna05}.  Until recently, only 11 such dwarf galaxies were known to orbit the Milky Way, however analysis of the Sloan Digital Sky Survey (SDSS) has revealed the presence of several lower luminosity dwarf galaxies that had remained undetected due to their extremely low surface brightness.  Among the newly-discovered dwarf spheroidal galaxies (dSph) uncovered by the SDSS is Canis Venatici (CVn), a small structure located at a distance of approximately $220\\kpc$ \\citep{zucker06}.  The CVn dSph has an absolute magnitude of $M_V = -7.9 \\pm 0.5$, similar to the Draco dSph ($M_V=-8.8$, \\citealt{mateo}), a Milky Way satellite galaxy $71\\kpc$ distant \\citep{odenkirchen} that provides us with one of the strongest cases for containing an extremely large fraction of dark matter. CVn is somewhat more extended than Draco however, possessing a half-light radius of $550\\pc$ (or $8\\mcnd5$, \\citealt{zucker06}), compared to $200\\pc$ ($9\\mcnd7$, \\citealt{wilkinson}) for Draco.  Although Draco is one of the most convincing examples that dwarf galaxies possess huge mass-to-light ratios, the tidal effects of the Milky Way on this galaxy are a significant concern for analyses of its dark matter content, and the dark matter properties that have been deduced from this galaxy are contingent on the system not being significantly tidally perturbed.  For this reason the CVn dwarf is an interesting candidate for a follow-up study to understand the dark matter. It is similar to Draco, but more distant, and can be expected therefore to suffer less from the Galactic tides, a reasonable assertion supported by the larger half mass radius.  Here we present a first spectroscopic study aimed to reveal the dark matter properties of this intriguing galaxy. ",
        "conclusions": "We have uncovered two stellar components in the Canes Venatici dwarf spheroidal galaxy with distinct structure and kinematics.  Below a metallicity of ${\\rm [Fe/H] = -2}$ this galaxy has a velocity dispersion of $\\sim 10\\kms$, typical of other dSph galaxies. At metallicities richer than this limit, an intriguing additional component is detected, possessing a velocity dispersion that is barely resolved with our measurement accuracy, and has an upper limit of $1.9\\kms$. The nature of the metal-rich component is unclear, though the wide metallicity spread it displays suggests that it was formed in situ from gas reprocessed from the previous generation of metal-poor stars. The possibility that the metal-rich population is younger is supported by the overlap of the red-giant branches of both populations in the CMD.  Thus the CVn dwarf presents us with a very interesting puzzle: why does this younger population possess such a small velocity dispersion?  Naively, one would expect this to attest to the presence of a cored potential in the central regions of the galaxy: where the potential gradients are flat one can hide much dark mass without structures that are significantly smaller in scale than the core radius being substantially affected.  However, the Jeans analysis we have performed does not support this scenario, due to the relatively small scale length of the metal-poor population.  Instead, we find that the model parameters must be forced to highly unlikely extremes to achieve a reasonable consistency between the components. We feel, however, that the paucity of well-measured stars in the kinematic sample, especially at large radius, introduces a large modelling uncertainty, and we therefore avoid stating categorically at the present time that the two populations are inconsistent.  If future studies confirm the values for the velocity dispersion and component scale lengths presented here, it will be a clear indication that one or more of the assumptions made in the Jeans analysis is invalid. The system may not be in dynamical equilibrium: despite the previous discussion, it is possible that the metal-rich population is a recent accretion that has not had time to relax. Given typical dynamical times, this seems very unlikely, though we note that in other dSph fossil substructure has been detected, and its longevity can be interpreted as consistent with the presence of a cored dark matter distribution \\citep{kleyna03}. Alternatively, the adopted assumption of spherical symmetry may be poor. The metal-rich component may also be a disk-like component seen face-on, though this is hard to envisage if it was formed from the existing apparently non-rotating metal-poor population, given the similarity in their scale lengths.  It may also turn out to be that the velocity distribution is not isotropic.  Although the present uncertainties in the model parameters do not allow us to provide strong constraints on the dark matter distribution, we anticipate that if the predictions of the two components can be made consistent with improved data, this will yield a reliable dark matter profile, from which we may confidently determine whether the dark matter conforms to the predictions of CDM theory or not.     \\newcommand{\\mnras}{MNRAS} \\newcommand{\\nat}{Nature} \\newcommand{\\araa}{ARAA} \\newcommand{\\aj}{AJ} \\newcommand{\\apj}{ApJ} \\newcommand{\\apjl}{ApJ} \\newcommand{\\apjs}{ApJSupp} \\newcommand{\\aap}{A\\&A} \\newcommand{\\aaps}{A\\&ASupp} \\newcommand{\\pasp}{PASP}"
    },
    "0609/astro-ph0609288_arXiv.txt": {
        "abstract": "We study the contamination of the $B$--mode of the Cosmic Microwave Background Polarization (CMBP) by Galactic synchrotron in the lowest emission regions of the sky. The 22.8-GHz polarization map of the 3-years WMAP data release is used to identify and analyse such regions. Two areas are selected with signal-to-noise ratio $S/N<2$ and $S/N<3$, covering $\\sim 16$\\% and $\\sim 26$\\% fraction of the sky, respectively. The polarization power spectra of these two areas are dominated by the sky signal on large angular scales (multipoles $\\ell < 15$), while the noise prevails on degree scales. Angular extrapolations show that the synchrotron emission competes with the CMBP $B$--mode signal for tensor-to-scalar perturbation power ratio $T/S = 10^{-3}$--$10^{-2}$ at 70-GHz in the 16\\% lowest emission sky ($S/N<2$ area). These values worsen by a factor $\\sim 5$ in the $S/N<3$ region. The novelty is that our estimates regard the whole lowest emission regions and outline a contamination better than that of the whole high Galactic latitude sky found by the WMAP team ($T/S>0.3$). Such regions allow $T/S \\sim 10^{-3}$ to be measured directly which approximately corresponds to the limit imposed by using a sky coverage of 15\\%. This opens interesting perspectives to investigate the inflationary model space in lowest emission regions. ",
        "introduction": "The Cosmic Microwave Background Polarization (CMBP) allows the study of the first stages of the Universe and is one of the hot topics of cosmology. One of its components, the $B$--mode, is sensitive to the primordial Gravitational Wave (GW) background left by inflation giving a way to investigate the physics of the very early Universe (e.g. \\citealt{ka98}). In fact, its power spectrum on degree scales has a linear dependence on the tensor-to-scalar perturbation power ratio $T/S$, which measures the amount of primordial GW (e.g. \\citealt{boyle06,kinney06}). In combination with parameters measured by the CMB Temperature spectrum (the scalar perturbation spectral index $n$ and its running $dn / d\\ln k$), it allows us to distinguish among the numerous models of inflation.  The $B$-mode signal is faint and only upper limits of $T/S$ have been set so far. \\citet{seljak06} and \\citet{martin06}, using different analyses, find consistent results of $T/S < 0.22$ and $T/S < 0.21$ (95\\% C.L.), respectively. On the other hand, the plethora of inflation models does not help constrain $T/S$. This can vary by orders of magnitude and can be even smaller than $10^{-4}$, although \\citet{boyle06} show that only models with high fine tuning degree can feature $T/S < 10^{-3}$. Correspondingly, the peak $B$-mode signal can vary from $\\sim200$~nK of the present upper limits down to values even smaller than $3$~nK for $T/S < 10^{-4}$.  \\begin{figure} \\centering \\includegraphics[angle=90, width=1.0\\hsize]{fig1a.ps} \\includegraphics[angle=90, width=1.0\\hsize]{fig1b.ps} \\includegraphics[angle=90, width=1.0\\hsize]{fig1c.ps} \\caption{{\\it Top.} Polarized intensity map at 22.8-GHz of the WMAP experiment \\citep{page06}. Data are averaged on $\\sim 4^\\circ$ scale using the HEALPix \\citep{gorski05} pixelation with parameter $N\\!side = 16$ (resolution r4). The map is in Galactic coordinates with the Galactic centre in the middle. The coordinate grid is $15^\\circ$ spaced both in longitude and latitude. Lowest emission regions are the blue spots at mid-high Galactic latitudes out of the large SNR loops. {\\it Mid and bottom.} Lowest emission regions in the 22.8-GHz WMAP polarization map: the $4^\\circ$-pixels of the two regions with $S/N < 2$ (mid) and $S/N < 3$ (bottom) are blanked (white). \\label{WMAPFig} } \\end{figure}  The weakness of the CMB $B$-mode makes it easily contaminated by foreground emissions from both the Galaxy and extragalactic sources. The study of astrophysical foregrounds is thus crucial to set the capability of CMBP experiments to investigate the early Universe and help set which part of the inflationary model space is accessible.  At frequencies lower than 60--70~GHz, the most relevant contaminant is expected to be the Galactic synchrotron emission. Recent results based on 3-years Wilkinson Microwave Anisotropy Probe (WMAP) data at 22.8-GHz show that even the high Galactic latitudes are normally strongly contaminated \\citep{page06}. These authors use about 75~percent of the sky (all high latitudes but strong large local structures like the Northern Galactic Spur) and find that at 60--70~GHz the synchrotron emission competes with the cosmic $B$-mode signal even for models with $T/S = 0.3$--0.5, which are already disfavoured by the present upper limits. Similar results are obtained by \\citet{laporta06} through the analysis of the 1.4~GHz sky at latitude $|b| > 20^\\circ$ using the DRAO survey data \\citep{wolleben06}.  \\citet{laporta06} also analyse areas\\footnote{ Three areas defined as follows: A) $180^\\circ<l<276^\\circ$, $b>45^\\circ$; B) $193^\\circ<l<228^\\circ$, $b<-45^\\circ$; C) $65^\\circ<l<180^\\circ$, $b>45^\\circ$; $l$ and $b$ are Galactic longitude and latitude. } with smaller extension (for a total of about 10\\% of the sky) and lower emission, finding that the Galactic signal at 70-GHz is comparable to the CMB $B$-mode in case of $T/S \\sim 0.1$. Even though lower, this signal would largely contaminate the CMB $B$--mode for most of the $T/S$ values presently allowed. However, these areas are not in the lowest emission regions visible in the 22.8-GHz WMAP polarization map.  Information about the lowest emission regions are available, instead, from three independent small areas observed at 1.4 and 2.3-GHz \\citep{carretti05b,carretti06,bernardi06}. The synchrotron emission here actually looks weaker, competing with the CMB $B$-mode only for cosmological models with $T/S = 10^{-3}$--$10^{-2}$ at 70-GHz \\citep{carretti06}.  Such values would give more chances to investigate a large part of the inflationary model space, opening interesting perspective for the study of the inflation physics. However, the areas observed so far cover just few square degrees ($\\sim 20$-deg$^2$ in total) and represent few sparse samples of the lowest emission regions. A real knowledge of the typical Galactic synchrotron contribution of such regions is still an open question, and requires wider areas to be analysed.  This paper reports the first study carried out to identify and analyse the areas with the lowest synchrotron emission at microwave wavelengths, which allow negligible Faraday rotation effects (still appreciable at $\\sim 1$-GHz; e.g. see \\citealt{carretti05a}). We use the 22.8-GHz polarization map of the 3-years WMAP data release to select them (Section~\\ref{WMAPSec}), and compute their polarized angular power spectra (Section~\\ref{specSec}). Finally, we estimate the contamination of the CMBP by Galactic synchrotron in the CMBP frequency window (70--90~GHz) and discuss implications for the detection of the cosmic $B$--mode in lowest emission regions (Section~\\ref{CMBSec}). ",
        "conclusions": "\\label{CMBSec}  The contamination of the CMB $B$-mode signal can be estimated by extrapolating the spectra of Sect.~\\ref{specSec} to 70-GHz, a frequency in the range where the combined contribution of synchrotron and dust is believed to be minimum \\citep{page06,carretti06}. It is worth noting that typical values of rotation measure at high Galactic latitudes ($RM \\sim 10$--20~rad~m$^{-2}$) generate polarization angle rotations of 6--12~arcmin at 22.8-GHz. Thus, Faraday rotation effects can be considered negligible, allowing safe frequency extrapolations.  Our spectra are noise dominated at the multipole $\\ell$ the CMB $B$-mode peaks at ($\\ell\\sim 90$) and a direct extrapolation of them would provide just upper limits. Then, we adopt another approach: we use the spectra measured in the $\\ell$--range where they are signal--dominated and perform angular extrapolations to the scale of interest. The mean emission is that provided by the last column of Table~\\ref{fitTab}. As for the extrapolation, the analysis of the 1.4-GHz DRAO survey data shows that the angular behaviour of Galactic synchrotron spectra is well represented by a power law \\begin{equation} C_\\ell \\propto \\ell^\\beta \\label{plEq} \\end{equation} with slope $\\beta$ varying in the range [-2.5, -3.0] for $\\ell = 2$--300 \\citep{laporta06,burigana06}. This is consistent with the value of $\\beta=-2.6$ quoted by the WMAP team for the whole 75\\% of the sky they consider \\citep{page06}. We use as slope both the two edge values of the $\\beta$ range (-2.5 and -3.0) and start extrapolations from the middle of the $\\ell$-range used to evaluate the mean spectrum, namely $\\ell_0 = 8$,~14 for $S/N < 2$,~$<3$, respectively.  Frequency extrapolations are performed assuming the brightness Temperature of the Galactic synchrotron follows a power law $T_b^{\\rm synch} \\propto \\nu^\\alpha$, with $\\alpha = -3.1$ \\citep{bernardi04}.  Results are reported in Figure~\\ref{spec70Fig}. \\begin{figure} \\centering \\includegraphics[angle=0, width=1.0\\hsize]{fig3.ps} \\caption{$B$--mode power spectra of the Galactic synchrotron emission estimated at 70-GHz in the regions with $S/N < 2$ (green shaded area) and $S/N < 3$ (gray shaded area). The two angular extrapolations with $\\beta = -2.5$ and $-3.0$ are reported (upper and lower dashed line of shaded areas). A slope $\\alpha=-3.1$ is assumed for the frequency extrapolation. The case of a slope $\\alpha=-3.5$ is also reported for the $S/N < 2$ region (yellow shaded). For comparison, the plot shows CMB spectra for three different values of $r=T/S$, the synchrotron emission detected in the two low emission areas in the target fields of the BOOMERanG ($C^B_{\\rm BOOM}$, \\citealt{carretti05b}) and BaR-SPOrt experiments ($C^B_{\\rm BaR}$, \\citealt{carretti06}), and the upper limit found in the DASI fields ($C^B_{\\rm DASI}$, \\citealt{bernardi06}). Finally, the general contamination at high Galactic latitude, as estimated by the WMAP team using 74.3\\% of the sky is also reported: the spectrum at 22.8-GHz is considered \\citep{page06} and extrapolated to 70-GHz. \\label{spec70Fig} } \\end{figure} For the $S/N < 2$ case, which concerns $\\sim 16$\\% of the sky, the contamination range covered by the two angular extrapolations at $\\ell \\sim 90$ competes with the CMBP $B$-mode for models with $T/S = [3 \\times 10^{-3}$,~$1 \\times 10^{-2}]$. The comparison with previous measurements shows these values are significantly lower than the mean contamination estimated by the WMAP team for most of high Galactic latitudes. On the other hand, our results are similar to the contamination estimated in the small sky patches observed in low emission regions (\\citealt{carretti06} and references therein) giving a valuable consistency of results in regions selected with similar criteria (areas in lowest emission regions). Although with the uncertainties of an angular scale extrapolation, our results extend to the whole lowest emission regions the promising indications obtained in those few small sample areas covering just few square degrees, and strongly support the possibilities to access small $T/S$ values in the lowest emission part of the sky.  The use of frequency slope $\\alpha = -3.1$ can be considered as a conservative approach. In fact, \\citet{hinshaw06} fit the behaviour of the synchrotron total intensity, finding out a steepening of the slope at the highest frequencies of the WMAP range. In particular, they fit $\\alpha=-3.5$ in the range 23--60~GHz. With such a slope the synchrotron contamination at 70~GHz further reduces by a factor 2.5 with respect to $\\alpha=-3.1$ and even $T/S \\sim [1 \\times 10^{-3}$,~$3 \\times 10^{-3}]$ becomes accessible (yellow shaded area of Figure~\\ref{spec70Fig}).  The situation is slightly worse for $S/N < 3$, as expected, since higher emission pixels are included in the analysis. However, $T/S = [2\\times10^{-2}$,~$4\\times10^{-2}]$ looks still accessible even in this larger region, assuming $\\alpha=-3.1$. These are values better than that for all the high Galactic latitudes (WMAP results on $\\sim 75\\%$ of the sky), or in the areas analysed by \\citet{laporta06}.  For the lowest emission areas, the minimum combined synchrotron-dust contamination has been found to be somewhere between 70 and 90-GHz \\citep{carretti06}. It is thus worthy estimating the synchrotron contamination of our best case ($S/N < 2$) even at 90-GHz. A further factor $\\sim 4$ is gained with respect to 70-GHz (in thermodynamic temperature), and the Galactic signal competes with models with $T/S \\sim [1 \\times 10^{-3}$,~$3 \\times 10^{-3}]$.    Despite the significant general contamination of high Galactic latitudes,  the results obtained by our analysis depict a situation with better and interesting conditions in the lowest emission regions. The possibility to reach $T/S$ values as low as $10^{-3}$--$10^{-2}$ looks now extended from few small good sky patches to an area representing $\\sim 16$\\% of the sky, that could be the right place for deep CMBP observations looking at the $B$-mode.  To limit observations in small regions, however, imposes intrinsic constraints on the minimum detectable $T/S$, mainly because of the leakage from $E$- into the weaker $B$-mode. In fact, \\citet{amarie05} find that an all-sky survey would allow a detection of $T/S$ with a theoretical sensitivity limit of $\\Delta(T/S) = 1.5 \\times 10^{-5}$, that becomes $\\Delta(T/S) = 3.2 \\times 10^{-5}$ when 70\\% of the sky is available, and $\\Delta(T/S) = 10^{-3}$ for 15\\% (3-$\\sigma$ C.L.). It is worth noting that the class of inflationary models with minimal fine-tuning have $T/S$ values ranging within $10^{-3}$ and $10^{-1}$ \\citep{boyle06}, for which a 15\\% sky portion would be large enough for the first detection of the tensor CMBP component.  In spite of these limitation, the minimum $T/S$ value detectable in 15\\% of the sky almost matches the limit imposed by the foregrounds we have found in lowest emission regions. Although the uncertainties because of the angular extrapolations we need to apply, the search for the $B$-mode in a sky portion of such a size could thus be a good trade-off between intrinsic and foregrounds limits.  In addition, the weakness of the $B$-mode signal makes already a challenge to detect the signal for $T/S=0.1$ with the present technology (e.g. \\citealt{cortiglioni06}). It is likely that significant technological improvements will be necessary before cosmologists can face an all-sky mapping mission with a sensitivity able to match the intrinsic limit of $\\Delta(T/S)\\sim 10^{-5}$. Therefore, an experiment aiming at detecting the $B$-mode in a smaller region (10-15\\% of the sky) can be a valuable intermediate step that would allow us to probe inflation models down to $T/S = 10^{-3}$."
    },
    "0609/astro-ph0609527_arXiv.txt": {
        "abstract": "{ {\\it Aims.} Non-local thermodynamical equilibrium (NLTE) line formation for neutral and singly-ionized calcium  is considered through a range of spectral types when the Ca abundance varies from the solar value down to [Ca/H] = $-$5. We evaluate the influence of departures from LTE on Ca abundance determinations and inspect the possibility of using \\ion{Ca}{i} / \\ion{Ca}{ii} line-strength ratios as indicators of surface gravity for extremely metal-poor stars.  \\noindent {\\it Methods.} A comprehensive model atom for \\ion{Ca}{i} and \\ion{Ca}{ii} is presented. Accurate radiative and electron collisional atomic data are incorporated. The role of inelastic collisions with hydrogen atoms in the statistical equilibrium of \\ion{Ca}{i/ii} is estimated empirically from inspection of their different influence on the \\ion{Ca}{i} and \\ion{Ca}{ii} lines in selected stars with well determined stellar parameters and high-quality observed spectra.  \\noindent {\\it Results.} The dependence of NLTE effects on the atmospheric parameters is discussed. Departures from LTE significantly affect the profiles of \\ion{Ca}{i} lines over the whole range of stellar parameters considered. However, at [Ca/H] $\\ge -2$, NLTE abundance correction of individual lines may be small in absolute value due to the different influence of NLTE effects on line wings and the line core. At lower Ca abundances, NLTE leads to systematically depleted total absorption in the line and positive abundance corrections, exceeding +0.5\\,dex for \\ion{Ca}{i} $\\lambda\\,4226$ at [Ca/H] = $-$4.9. In contrast, NLTE effects strengthen the \\ion{Ca}{ii} lines and lead to negative abundance corrections. NLTE corrections are small, $\\le 0.02$\\,dex, for the \\ion{Ca}{ii} resonance lines over the whole range of stellar parameters considered. For the IR lines of multiplet $3d$ - $4p$, they  grow in absolute value with decreasing Ca abundance exceeding 0.4\\,dex in metal-poor stars with [Fe/H] $\\le-3$. As a test and first application of the \\ion{Ca}{i/ii} model atom, Ca abundances are determined on the basis of plane-parallel LTE model atmospheres for the Sun, Procyon (F IV-V), and seven metal-poor stars, using high S/N and high-resolution spectra at visual and near-IR wavelengths. Lines of \\ion{Ca}{i} and \\ion{Ca}{ii} give consistent abundances for all objects (except Procyon) when collisions with hydrogen atoms are taken into account. The derived absolute solar Ca abundance (from \\ion{Ca}{i} and \\ion{Ca}{ii} lines) is $\\eps{Ca,\\odot}$ = 6.38\\,$\\pm$\\,0.06. For Procyon, the mean Ca abundance from \\ion{Ca}{i} lines is markedly subsolar, [Ca/H] = --0.14$\\pm$0.03. All metal-poor stars within our sample show an overabundance of calcium relative to iron with [Ca/Fe] abundance ratios of 0.26 to 0.46 typical for the halo population. The $W$(\\ion{Ca}{i}4226) / $W$(\\ion{Ca}{ii}8498) equivalent width ratio is predicted to be sensitive to surface gravity for extremely metal-poor stars, while this is not the case for the ratio involving the \\ion{Ca}{ii} resonance line(s). ",
        "introduction": "Calcium is one of the best observable chemical elements in late-type stars. The resonance lines of \\ion{Ca}{i} at $\\lambda\\,4226\\,\\AA$ and \\ion{Ca}{ii} at $\\lambda\\,3933$ and 3968~\\AA\\, lie in the visual spectral range and can be measured even in extremely metal-poor stars with metallicity [Fe/H] $< -5$ (Christlieb et al. \\cite{chri02}; Frebel et al. \\cite{frebel05}). In such stars, Ca is the only chemical element which is visible in two ionization stages, and the \\ion{Ca}{i} and \\ion{Ca}{ii} lines can be potent tools in the derivation of accurate values for fundamental stellar parameters and for the Ca abundance itself. Calcium is an important chemical element to study the history of $\\alpha-$process nucleosynthesis in the Galaxy. The subordinate lines of neutral Ca located in the relatively uncrowded yellow-to-red spectral regions are suitable for spectroscopic analysis over a wide range of Ca abundance from super-solar values down to [Ca/H] = $-$4 (Cayrel et al. \\cite{cayrel04}). The subordinate lines of ionized Ca at $\\lambda\\,8498$, 8542, and 8662~\\AA\\, are among the strongest features in the near-infrared spectra of late-type stars with metallicity down to [Fe/H] = --3 (Mallik \\cite{mallik97}). The \\ion{Ca}{ii} triplet lies at the focus of cool star research with the advent of large spectroscopic surveys of the Galaxy like the Radial Velocity Experiment (RAVE, Steinmetz et al.~\\cite{steinmetzetal06}) and the upcoming ESA Gaia satellite mission (Perryman et al.~2001). These broad lines are also powerful abundance and metallicity indicators out to large distances via medium-resolution spectroscopy, favorably located near the flux maximum of red giants. Applications range from quantitative analyses of individual stars in globular clusters (e.g. Armandroff \\& da Costa~\\cite{ardaco91}; Rutledge et al.~\\cite{rutledgeetal97}) to stars and stellar populations in nearby galaxies (e.g. Tolstoy et al.~\\cite{tolstoyetal01}; Ibata et al.~\\cite{ibataetal05}). Moreover, the \\ion{Ca}{ii} triplet is useful for the analysis of unresolved stellar systems like early-type galaxies (e.g. Saglia et al.~2002). However, as an increasing number of high-quality (echelle) spectra of stars and unresolved stellar systems becomes available, an equally high level in the theoretical modelling is required.  Previous investigations of the statistical equilibrium (SE) of neutral Ca in the Sun and Procyon (Watanabe \\& Steenbock \\cite{ca1}) and in the models with effective temperatures $\\Teff$ = 4500\\,K -- 6200\\,K, surface gravities $\\log g$ = 4.5 -- 1.0 and metal abundance, [Fe/H] = 0 and $-$1 (Drake \\cite{Drake}) found rather small departures from LTE. Drake (\\cite{Drake}) concluded that ``overionization effects (for \\ion{Ca}{i}) are not expected to increase in severity in metal-poor stars\". NLTE analysis of neutral Ca in the sample of 252 dwarf and subgiant stars was performed by Idiart \\& Thevenin (\\cite{th-mg}). Previous NLTE investigations for \\ion{Ca}{ii} were concerned only with the resonance lines in the Sun (we cite only the first paper, Shine \\& Linsky \\cite{ca74}) and the IR lines of multiplet $3d - 4p$ ($\\lambda\\,8498$, 8542 and 8662\\,\\AA) in moderately metal-poor stars by J\\o rgensen et al. (\\cite{ca92}) for [Fe/H] $\\ge -1$ and by Andretta et al. (\\cite{andretta05}) for [Fe/H] $\\ge -2$.  The need for a new analysis is motivated by two points. Firstly, a fairly extensive set of accurate atomic data on photoionization cross-sections and oscillator strengths was recently calculated in the Opacity Project (OP; see Seaton et al. (\\cite{OP}) for a general review). We include in the model atom 66 terms of \\ion{Ca}{i}, contrary to 16 terms in the works of Watanabe \\& Steenbock (\\cite{ca1}) and Drake (\\cite{Drake}), and 36 terms of \\ion{Ca}{ii}, contrary to 3 terms considered by J\\o rgensen et al. (\\cite{ca92}) and Andretta et al. (\\cite{andretta05}). This allows to investigate the detailed line formation of the high excitation lines of \\ion{Ca}{i} and \\ion{Ca}{ii}. All our results are based on line-profile analysis and take advantage of the advanced theory of collisional broadening by atomic hydrogen treated recently by Anstee \\& O'Mara (\\cite{omara_sp}), Barklem \\& O'Mara (\\cite{omara_pd}, \\cite{omara_ion}) and Barklem et al. (\\cite{omara_df}). Hereafter, these four important papers are referred to collectively as $A\\&O'M$. The second point is connected with an extension of the metallicity range observed in stars down to [Fe/H] = $-$5.45 (Aoki et al. \\cite{frebel06}). It is highly desirable to have an understanding of the likely influence of any departures from LTE over this whole range. In stars differing significantly in their metal abundance, different subsets of Ca lines are used to calculate the \\ion{Ca}{ii}/\\ion{Ca}{i} ionization equilibrium and Ca abundance. We investigate here NLTE line formation of our extended list of \\ion{Ca}{i} and \\ion{Ca}{ii} lines in the metallicity range between [Fe/H] = 0 and $-$4.34. An exception is the resonance lines of \\ion{Ca}{i} and \\ion{Ca}{ii} and the \\ion{Ca}{ii} lines of multiplet $3d - 4p$ in solar metallicity stars: their cores and inner wings are most probably influenced by the chromospheric temperature rise and a non-thermal and depth-dependent chromospheric velocity field neither of which are part of the homogeneous photospheric models used in this study.  For the statistical equilibrium of atoms in cool stars, an important issue debated for decades, from Gehren (\\cite{gehren75}) to Belyaev \\& Barklem (\\cite{belyaev03}), is the role of inelastic collisions with hydrogen atoms. One of our aims is to empirically constrain the efficiency of this type of collisions in the SE of \\ion{Ca}{i/ii} from inspection of their different influence on the \\ion{Ca}{i} and \\ion{Ca}{ii} lines in the Sun, Procyon and seven metal-poor ($-1.35 \\ge$ [Fe/H] $\\ge -2.43$) stars with well determined stellar parameters and high quality observed spectra.  The paper is organized as follows. The model atom of \\ion{Ca}{i/ii} and atomic data are described in Sect.\\,\\ref{NLTE}. Departures from LTE for \\ion{Ca}{i} and \\ion{Ca}{ii}, NLTE trends with metallicity, effective temperature and surface gravity, and errors of NLTE results caused by the uncertainties of atomic data are discussed in Sect.\\,\\ref{effect}. In Sect.\\,\\ref{sun}, the solar Ca spectrum is studied to provide the basis for further differential analysis of stellar spectra. We determine the absolute solar Ca abundance using four different sets of oscillator strengths based on laboratory measurements, OP calculations, and collected from the NIST (http://physics.nist.gov/PhysRefData) and VALD (Kupka et al. \\cite{vald}) databases and, thus, estimate the accuracy of the available atomic data. Observations and stellar parameters of our sample of stars are described in Sect.\\,\\ref{obs}. For the selected stars, Sect.\\,\\ref{stars1} investigates whether or not Ca abundances derived from the \\ion{Ca}{i} and \\ion{Ca}{ii} lines agree. NLTE formation of the \\ion{Ca}{i} resonance line in the atmospheres of metal-poor stars is tested in Sect.\\,\\ref{stars2}. The predictions are made in Sect.\\,\\ref{extrim} for the \\ion{Ca}{i}/\\ion{Ca}{ii} line-strength ratios as indicators of surface gravity for extremely metal-poor stars. Our recommendations and conclusions are given in Sect.\\,\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  In this study, NLTE line formation of an extended list of \\ion{Ca}{i} and \\ion{Ca}{ii} lines was considered for the temperatures ranging between $\\Teff = 5000 K$ and $\\Teff = 6000 K$, surface gravities $\\log g =$ 3.0 and 4.0, and metallicities from [Fe/H] = 0 down to [Fe/H] = $-$4.34. For every \\ion{Ca}{i} line, departures from LTE affect significantly its profile over the whole range of stellar parameters. If a line is strong (multiplets 2, 3, 4, 18, 21, and 47 at [Fe/H] $\\ge -2$ and the remaining multiplets at [Fe/H] $\\ge -1$), its wings are weakened but the core is strengthened compared to the LTE case. The value and sign of NLTE abundance correction are defined by a relative contribution of the core and the wings to the overall line strength. When the line becomes weak due to decreasing Ca abundance, NLTE leads to depleted total absorption in the line and positive abundance correction. NLTE effects are very large at extremely low Ca abundance. For example, at [Ca/H] = $-$4.9 NLTE abundance correction for \\ion{Ca}{i} $\\lambda\\,4226$ can exceed 0.5\\,dex. Thus, for a given line, $\\Delta_{\\rm NLTE}$ depends on $\\Teff$, $\\log g$, [Ca/H], and microturbulence value, and, for a given model, $\\Delta_{\\rm NLTE}$ is different in value and sign for a variety of \\ion{Ca}{i} lines. Any interpolation of NLTE results for \\ion{Ca}{i} has to be performed with caution. A different situation is found for \\ion{Ca}{ii}. For every \\ion{Ca}{ii} line, NLTE leads to enhanced absorption in the line core and negative abundance correction over the whole range of stellar parameters. The absolute value of $\\Delta_{\\rm NLTE}$ is defined by the relative contribution of the core and the wings to equivalent width. For example, NLTE corrections remain very small, $\\le 0.02$\\,dex, for the resonance lines and grow in absolute value with decreasing Ca abundance for the IR lines of multiplet $3d$ - $4p$ exceeding 0.4\\,dex in a metal-poor model with $T_{\\rm eff}$\\,=\\,6000\\,K, $\\log g$\\,=\\,3.0 and [Fe/H]\\,=\\,$-$3. Such corrections are important to be considered for the correct interpretation of the near-IR spectroscopic data collected in the RAVE survey and the future ESA Gaia satellite mission.  Empirical evidence is found from the analysis of stellar spectra that inelastic collisions with hydrogen atoms serve as an additional source of thermalization in cool (metal-poor) stars. Our best choice for a scaling factor to the formula of Steenbock \\& Holweger (\\cite{hyd}) for calculation of hydrogenic collisions is \\kH\\ = 0.1. Disregarding hydrogenic collisions completely will lead to overestimated NLTE effect at low metallicity.  Taking advantage of our SE approach and accurate atomic data for the investigated lines, we obtain good agreement, within 0.04\\,dex, between absolute solar \\ion{Ca}{i} (from 23 lines) and \\ion{Ca}{ii} (from 8 lines) abundances, and the average of the two is $\\eps{Ca}^\\odot$ = 6.38\\,$\\pm$\\,0.06.  Calcium abundances from two ionization stages are examined for the first time for five metal-poor stars. We show that NLTE largely removes obvious discrepancies between \\ion{Ca}{i} and \\ion{Ca}{ii} obtained under the LTE assumption.  Solving the restricted NLTE problem for calcium in ``classical'' one-dimensional LTE model atmospheres, we qualitatively understand the formation of the \\ion{Ca}{i} resonance line in metal-poor stars, where it is of purely photospheric origin. In agreement with observations, NLTE predicts weakening of the line wings and strengthening of the line core compared to the LTE case. Residual discrepancies may be related to the use of classical model atmospheres with adjustable parameters like microturbulence."
    },
    "0609/astro-ph0609711_arXiv.txt": {
        "abstract": "The Baikal Neutrino Telescope  NT200 takes data since April 1998. On April 9th, 2005, the 10 Mton scale detector NT200$+$ was put into operation in Lake Baikal. Selected results obtained during 1998-2002 with the neutrino telescope NT200 are presented. ",
        "introduction": "The Baikal Neutrino Telescope  NT200 takes data since April 1998. On April 9th, 2005, the 10-Mton scale detector NT200$+$ was put into operation in Lake Baikal. Description of site properties, detector configuration and performance have been described elsewhere \\cite{NE05,HEB,NANP05,RW}.  In this paper we present main results obtained from the analysis of data taken with the Baikal neutrino telescope NT200 between April 1998 and February 2003. ",
        "conclusions": "The Baikal neutrino telescope NT200 is taking data since April 1998. The upper limit obtained for a diffuse ($\\nu_e+\\nu_{\\mu}+\\nu_{\\tau}$) flux with $E^{-2}$ shape is $E^2 \\Phi = 8.1 \\times 10^{-7}$cm$^{-2}$s$^{-1}$sr$^{-1}$GeV. The limits on fast magnetic monopoles and on a $\\bar{\\nu_e}$ flux at the resonant  energy 6.3$\\times$10$^6$GeV are presently the most stringent. To extend the search for diffuse extraterrestrial neutrinos with higher sensitivity, NT200 was significantly upgraded to NT200$+$, a detector with about 5 Mton enclosed volume, which takes data since April 2005 \\cite{RW,NT+2}. The three-year sensitivity of NT200$+$ to the all-flavor neutrino flux is approximately $2\\times10^{-7}$cm$^{-2}$s$^{-1}$sr$^{-1}$GeV for $E>$10$^2$ TeV (shown in Fig. \\ref{fig7}). In parallel with exploiting NT200$+$ we started research \\& development activities towards a Gigaton Volume Detector in Lake Baikal.  \\vspace{-0.5cm} \\ack This work was supported by the Russian Ministry of Education and Science, the German Ministry of Education and Research and the Russian Fund of Basic Research (grants 05-02-17476, 04-02-17289, 04-02-16171, 05-02-31021, 05-02-16593), and by the Grant of  the President of Russia NSh-1828.2003.2."
    },
    "0609/astro-ph0609182_arXiv.txt": {
        "abstract": "{} % {We study  jet-driving protostars and T Tau stars to characterize their X-ray emission. We seek soft spectral components that may be due to shock emission, and search for soft, shock-induced emission displaced from the stellar position. } {We study two stellar samples, the first consisting of lightly absorbed T Tau stars with strong jets, the other  containing protostars with disks seen nearly edge-on. The former sample was observed in the {{\\it XMM-Newton Extended Survey of the Taurus Molecular Cloud}} (XEST), while {{\\it Chandra}} archival data provided observations of the latter.} {We confirm the previously identified peculiar spectrum of DG Tau A and find similar phenomenology in GV Tau and DP Tau, suggesting a new class of X-ray  spectra.   These  consist of a lightly absorbed, very soft component and a strongly absorbed, very hard component. The latter is flaring while little variability is detected in the former. The absorption of the harder component is about an order  of magnitude higher than expected from the optical extinction assuming a standard gas-to-dust mass ratio.  For the absorbed protostars, only the hard, stellar X-ray component is found. } {The flaring hard component represents active coronal emission. Its strong absorption is attributed to mass inflow from the accretion disk. The optical extinction is small because the dust has sublimated at larger distances. The weakly absorbed soft component cannot originate from the same location. Because the stars drive strong jets, we propose that the X-rays are generated in shocks in the jets. We find that for the three peculiar X-ray sources, the luminosity of the soft component  roughly scales with the equivalent width of the [O~I] $\\lambda$~6300 line formed in the jets, and with the mass outflow rate. In the more strongly obscured protostars, the soft component is entirely absorbed, and only the hard, coronal component penetrates the envelope or the near-edge-on disk.} ",
        "introduction": "Pre-main sequence stars show various signs of accretion and outflow,  such as circumstellar disks detected at radio and millimeter wavelengths (e.g., \\citealt{simon00}), molecular outflows observed in molecular lines (e.g., \\citealt{bachiller96} for a review), and accompanying optical (e.g., \\citealt{hirth97, eisloeffel98}) and radio jets (e.g., \\citealt{anglada95}).  The most visible manifestation of outflows are the optically visible jets that  may be excited close to the star, in internal shocks along the mass stream, or at distances up to several parsecs as the fast stream encounters the interstellar medium where it shock-ionizes the gas in Herbig-Haro objects (for a review of Herbig-Haro flows, see \\citealt{reipurth01}). Under ideal circumstances (low extinction, strong ionization), jets can be identified at distances as close as 0.1\\arcsec\\ to the star \\citep{bacciotti02}. The same compact jets are also routinely detected at radio wavelengths, where the emission mechanism  is thought to be bremsstrahlung from the shock-heated gas \\citep{rodriguez95, anglada95}. Radio brightness temperatures suggest overall gas temperatures of order  $10^4$~K. This picture is ambiguous, however, as a number of non-thermal jets have been suggested from radio polarization or synchrotron-like spectral shapes \\citep{yusef90, curiel93, hughes97, ray97, smith03, loinard05}. Magnetic fields may thus play a role not only in the launching of jets, but in their propagation as well.   Two principal families of models for the jet-launching mechanism have been proposed: The magnetocentrifugal mechanism \\citep{blandford82, koenigl00} posits that outflows are launched on the disk along outward-bent, open magnetic fields owing to the increasing centrifugal forces along the field lines that are co-rotating with their disk footpoints. In contrast, the X-wind model \\citep{shu00} assumes that outflows are launched from magnetized regions between the star and the disk. Magnetohydrodynamic simulations show that magnetic reconnection and subsequent plasmoid ejection may lead to field-aligned outflows \\citep{hayashi96}.  Accretion and outflow processes are prone to producing X-rays, given that shocks with shock jump velocities of order several  hundred km~s$^{-1}$ are possible and  sometimes observed. Accretion onto the star essentially extracts gravitational energy and transforms it into thermal energy through shocks. The relevant theory has been presented by \\citet{ulrich76}, \\citet{calvet98} and \\citet{lamzin99}. Material falling along magnetic fields toward the stellar surface with nearly free-fall velocity, $v_{\\rm ff}$, develops strong shocks near or within the photosphere, heating plasma to temperatures of $T \\approx 3v_{\\rm ff}^2\\mu m_p/16k \\approx 5.2\\times 10^6 M/R~$~[K] where $M$ and $R$ are the stellar mass and the stellar radius, respectively, in solar units, $\\mu$ is the mean molecular weight, $m_p$ is the proton mass, and $k$ is the Boltzmann constant. For T Tau stars, heating to a few MK is thus entirely plausible (as, e.g., proposed for TW Hya, \\citealt{kastner02}), but  the temperatures in excess of 10~MK   that often dominate X-ray emitting plasmas in T Tau stars  \\citep{skinner03, preibisch05} cannot be due to accretion shocks. Rather, high electron temperatures in pre-main sequence stars are conventionally attributed to magnetized plasma trapped in  corona-like or large-scale ``magnetospheric'' magnetic fields around the stars.  Shocks also form in jets and outflows, in particular in Herbig-Haro (HH) objects. The relevant theory and a simple model have been discussed  by \\citet{raga02}. The strong-shock temperature can be expressed as $T\\approx 1.5\\times 10^5v_{\\rm 100}^2$~K (for fully ionized gas) where $v_{100}$ is the shock  speed relative to a target in units of 100~km~s$^{-1}$. Jet speeds are typically of order $v = 300-500$~km~s$^{-1}$ \\citep{eisloeffel98, anglada95, bally03}, allowing for shock speeds of similar magnitude. If a flow shocks a standing medium at 400~km~s$^{-1}$, then $T \\approx 2.4$~MK.  Observations have explicitly demonstrated that faint, soft X-ray emission originates from some  HH objects \\citep{pravdo01, pravdo04, favata02, tsujimoto04, pravdo05, grosso06}. \\citet{bally03} used a {\\it Chandra} observation to show that  X-rays form within an arcsecond of the protostar L1551 IRS-5 while the star itself is too  heavily obscured to be detected. As this example illustrates, the jet-launching region of powerful protostellar jets is usually inaccessible to optical, near-infrared or X-ray studies of protostellar outflow sources due to excessive absorption. However, there are a number of optically-revealed, classical T Tau stars that drive appreciable jets and outflows \\citep{hirth97}. Their jets  were  initially  dubbed ``micro-jets'' although recent studies have shown them to reach parsec scales and the mass loss rates  compete with those of more embedded sources \\citep{mcgroarty04}.  One of the most prominent examples, the classical T Tau star DG Tau A, has recently obtained scrutiny with high-resolution {\\it Chandra} X-ray observations. \\citet{guedel05} discovered a new type of X-ray spectrum in which a lightly absorbed spectral component from a rather cool (2--3~MK) plasma is complemented by a strongly absorbed component from a very hot (20--30~MK)  plasma. The latter was interpreted as being related to the usual coronal/magnetospheric source often seen in T Tau stars, subject to some excess absorption. In the light of this interpretation, the cooler plasma is unlikely to be located close to the magnetospheric component because the source of absorption would likely affect it as well.  The strong jets and high mass-loss rates of DG Tau A motivated the suggestion that the soft X-rays are produced in shocks in the acceleration region of the jet (the ``base'' of the jet), which is unresolved in {\\it Chandra} or {\\it XMM-Newton} images. Complementary  evidence was seen in faint and very soft X-ray emission along the optical jet, outside the stellar point-spread function and co-spatial with a strong bow-shock feature seen in the optical. The spectral distribution of the counts in the jet source was reminiscent of the count distribution in the soft stellar component. A spectrum requiring two absorbers was also reported for a proplyd in the the Orion Nebula Cluster, with a variable hard component but a constant soft component \\citep{kastner05}. \\citet{skinner06} have detected a similar two-component spectrum in the strongly accreting FU Ori, the prototype of FU Ori-type stars. The cooler, weakly absorbed component, however, reached temperatures more common to T Tau stars ($\\approx 8$~MK) and was therefore identified with a coronal component, while the very hot component was again strongly absorbed.  The recently conducted {\\it XMM-Newton Extended Survey of the Taurus Molecular Cloud} (XEST henceforth; \\citealt{guedel06})  offers unique access to numerous stellar objects in the Taurus Molecular Cloud (TMC). It covers  most stellar concentrations in the TMC molecular filaments and offers high-quality spectra and light curves for any ``typical'' T Tau star in the field of view. XEST also covers a number of well-studied jet-driving T Tau stars and protostars. The observations confirm the initial discovery of the anomalous DG Tau A and add substantial evidence for a jet-related origin. We find further clear examples of these {\\it two-absorber X-ray} sources (TAX sources henceforth). The present paper describes these new observations and presents a detailed discussion on the possible origin of the X-ray emission. For a comparison, we also present X-ray spectra of jet-driving sources that are deeply absorbed by their thick circumstellar disks seen nearly edge-on. The soft component, if present, would thus be absorbed {\\it unless} it is displaced sufficiently from the star to make it separate from the stellar X-ray image. These complementary data are based on high-resolution {\\it Chandra} observations.     \\begin{table*}[t!] \\centering \\caption{Principal targets and observations} \\label{obslog} \\begin{tabular}{lllll} \\hline \\noalign{\\smallskip} Parameter                  &        DG Tau A                &    GV Tau                        & DP Tau                           &    CW Tau        \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} XEST number                &  02-022                        &\t13-004                     &  10-045                          &   20-046 \\\\ ObsID  {\\it (XMM-Newton})  &  0203540201                    &\t0203541301                 &  0203542201                      &  0203542001  \\\\ Start time  (UT)           &  2004-08-17\\ 06:08:10          &  2004-08-25\\ 11:11:23            &  2005-03-05\\ 05:56:38            &  2004-09-12\\ 07:04:43      \\\\ End time (UT)              &  2004-08-17\\ 17:32:46          &  2004-08-25\\ 19:59:19            &  2005-03-05\\ 14:44:30            &  2004-09-12\\ 15:52:37    \\\\ Exposure time  (s)         &  41076\t\t                &  31676                           &  31672                           &  31674  \\\\ Boresight RA~(J2000.0)     &  04h 27m 19.6s                 &  04h 29m 52.0s                   &  04h 42m 20.9s                   &  04 14 12.9     \\\\ Boresight $\\delta$~(J2000.0)&  $26\\deg\\ 09\\arcmin\\ 25\\arcsec$&  $24\\deg\\ 36\\arcmin\\ 47\\arcsec$  &  $25\\deg\\ 20\\arcmin\\ 35\\arcsec$  &  28 12 12  \\\\ Star: RA~(J2000.0)$^a$     &  04h 27m 04.70s                &  04h 29m 23.73s\t           &  04h 42m 37.70s                  &  04h 14m 17.00s  \\\\ Star: $\\delta$~(J2000.0)$^a$&  $26\\deg\\ 06\\arcmin\\ 16\\farcs 3$&  $24\\deg\\ 33\\arcmin\\ 00\\farcs 3$ &  $25\\deg\\ 15\\arcmin\\ 37\\farcs 5$& $28\\deg\\ 10\\arcmin\\ 57\\farcs 8$   \\\\ \\hline \\end{tabular} \\vskip 0.5truecm \\begin{tabular}{lllll} \\hline \\noalign{\\smallskip} Parameter                  &   DG Tau B                     &    FS Tau B                      &  HH 34 IRS                       &    HH 111 IRS              \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} XEST number                &  C3-2                          &\t11-054 = C2-1              &  ...                             &   ...        \\\\ ObsID {\\it (Chandra})      &  4487                          &  4488                            &  4489                            & 4490         \\\\ Start time  (UT)           &  2004-01-11\\ 02:58:51          &  2003-11-08\\ 12:57:58            &  2003-12-05\\ 08:59:23            & 2004-11-21\\ 17:05:50    \\\\ End time (UT)              &  2004-01-11\\ 11:52:21          &  2003-11-08\\ 21:56:34            &  2003-12-05\\ 17:38:31             & 2004-11-22\\ 01:36:42     \\\\ Exposure time  (s)         &   29717                        &  29674                           &  28590                           & 28795       \\\\ Boresight RA~(J2000.0)     &   04h 27m 02.3s                &  04h 22m 00.2s                   &  05h 35m 32.07s                  & 05h 51m 54.97s       \\\\ Boresight $\\delta$~(J2000.0)& $26\\deg\\ 04\\arcmin\\ 56\\arcsec$ &  $26\\deg\\ 58\\arcmin\\ 07\\arcsec$  & $-06\\deg\\ 26\\arcmin\\ 50\\farcs 0$ & $02\\deg\\ 48\\arcmin\\ 55\\farcs 3$ \\\\ Star: RA~(J2000.0)$^b$      &  04h 27m 02.55s\t        &  04h 22m 00.70s                  & 05h 35m 29.84s                   & 05h 51m 46.31s              \\\\ Star: $\\delta$~(J2000.0)$^b$& $26\\deg\\ 05\\arcmin\\ 30\\farcs 9$&  $26\\deg\\ 57\\arcmin 32\\farcs 50$ & $-06\\deg\\ 26\\arcmin\\ 58\\farcs 40$& $02\\deg\\ 48\\arcmin 29\\farcs 72$ \\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[$^{\\mathrm{a}}$] For origin of coordinates, see \\citet{guedel06} \\item[$^{\\mathrm{b}}$] For origin of coordinates, see  Tables~\\ref{proto} and \\ref{GVcoords} \\end{list} \\end{table*}  The plan of the paper is as follows. In Sect.~\\ref{sect:targets}, we briefly describe our data analysis and  introduce the targets of this study. Sec.~\\ref{sect:results} presents the results that are further discussed in Sect.~\\ref{sect:discussion}. We conclude in Sect.~\\ref{sect:summary}. ",
        "conclusions": "\\label{sect:summary}  Jets are assumed to be launched within a few AU from the central star and have been resolved down to 0.1\\arcsec\\ (14~AU) from the latter at the distance of the Taurus Molecular Cloud  \\citep{bacciotti00}. These regions remain unresolved from the central star in present-day X-ray observations. Nevertheless, several jet-driving T Tau stars reveal peculiar X-ray spectra that have been taken as evidence for shock formation at the base of the jets (\\citealt{guedel05}, and our discussion above). In all cases, a very weakly attenuated ($N_{\\rm H} \\approx 10^{21}$~cm$^{-2}$), luminous but very soft spectral component indicates electron temperatures of  no more than a few MK. While this component produces photons mostly between 0.5--1~keV, a strongly absorbed ($N_{\\rm H} > 10^{22}$~cm$^{-2}$), hard spectral component dominates above 2~keV, resulting in a shallow spectral shape with two peaks. The hard component  originates from hot plasma of several tens of MK. In each case, it reveals flares while no significant  variability is detected from the soft component  whose flux is approximately  constant, even on time scales of 6--8 months (Sect.~\\ref{longterm}).  We attribute the hard component to a magnetospheric corona, as observed on most T Tauri stars. Although these T Tauri stars are unusually strong accretors (with $\\dot{M}_{\\rm acc}$ up to $\\approx 10^{-6}~M_{\\odot}$~yr$^{-1}$) we do not attribute the soft component to an accretion shock close to the surface of the star as proposed for the so far unique CTTS TW Hya \\citep{kastner02}. Such emission would be subject to the same high absorption experienced by the hard component, while the observed soft component is little absorbed. Instead, we propose that the soft component originates further away from the star where the first shocks form at the base of the jets. Indeed, the X-ray luminosity of the soft component in these three stars appears to scale with the mass outflow rate and with the equivalent width of [O~I] from the jet.   It is interesting to note that the visual extinction of the T Tau stars showing Two-Absorber X-ray (TAX) spectra is modest and does not agree with the photoelectric absorption of the hard component if standard gas-to-dust mass ratios compatible with the interstellar medium are assumed. We suggest that the hard emission is absorbed by accreting material an that this accreting material is dust-depleted. This is supported by the dust sublimation radius of our sources being several stellar radii. Our observations thus provide both evidence for dust sublimation in the innermost part of the accretion disk and for massive gas accretion streams from the disk to the star enshrouding the hot coronal plasma. Because the line of sight toward DG Tau A is inclined against the axis at a rather small angle of $37.7\\pm 3$~deg \\citep{eisloeffel98}, the absorbing mass accretion streams must reach high stellar latitudes rather than fall toward the equatorial regions. This fully supports the common picture of magnetic accretion (e.g., \\citealt{calvet98}).  Detection of the soft jet component requires low attenuation by gas, i.e., it is inaccessible in embedded protostars although such X-rays may locally exist in all jet-driving sources. We have indeed found only hard components in our strongly obscured targets, and they coincide very closely with the position of the stars. No displaced soft components were identified although we recall that a soft jet source was reported previously for DG Tau A  (\\citealt{guedel05}; see also Fig.~\\ref{images})."
    },
    "0609/astro-ph0609461_arXiv.txt": {
        "abstract": "$JHK$ photometry in the Mauna Kea Observatory (MKO) near-infrared system is presented for 115 stars.  Of these stars, 79 are UK Infrared Telescope (UKIRT) standards from Hawarden et al., and 42 are Las Campanas Observatory (LCO, or NICMOS) standards from Persson et al.  The average brightness of the sample in all three bandpasses is 11.5 magnitudes, with a range between 10 and 15.  The average number of nights each star was observed is 4, and the average of the  internal error of the final results is 0$\\fm$011.  These $JHK$ data agree with those reported by other groups to 0$\\fm$02, for stars in common, which is consistent with the uncertainties. The measurements are used to derive colour transformations between the MKO $JHK$ photometric system and the UKIRT, LCO and Two Micron All-Sky Survey (2MASS) systems. The 2MASS$-$MKO data scatter by 0$\\fm$05 for redder stars, which is consistent with a dependence on stellar luminosity: the 2MASS $J$ bandpass includes H$_2$O features in dwarfs and the MKO $K$ bandpass includes CO features in giants. We stress that colour transformations derived for stars whose spectra contain only weak features cannot give accurate transformations for objects with strong absorption features within one, but not both, of the filter bandpasses. We find evidence of systematic  effects at the 0$\\fm$02 level in the photometry of stars with $J<11$ and $H$,$K<10.5$ presented here and in Hawarden et al.. This is due to an underestimate of the linearity correction for stars observed with the shortest exposure times; very accurate photometry of stars approaching the saturation limits of infrared detectors which are operated in double-read mode is difficult to obtain. There are indications that four stars in the sample, GSPC S705-D, FS 116 ( B216-b7), FS 144 (Ser-EC84) and FS 32 (Feige 108), may be variable.  There are 84 stars in the sample presented here that have $11 < J < 15$ and $10.5 < H$,$K < 15$, are not suspected to be variable, and have magnitudes with an estimated error $\\leq 0\\fm$027; 79 of these have an error of  $\\leq 0\\fm$020. These represent the first published high-accuracy $JHK$ stellar photometry in the MKO near-infrared photometric system; we recommend these objects be employed as primary standards for that system. ",
        "introduction": "In the 1990's infrared astronomy changed radically when array cameras started to replace single-channel photometers. Not only did the field of view and spatial resolution increase markedly, but much fainter limits could be reached. One repurcussion was that the existing photometric standards were too bright for the new detectors, and many groups published lists of fainter standards. These included: \\citet{Bo, Ca, Hu, Ps, Ha, Gt}.   Towards the end of the 1990's and extending into this decade, modern infrared sky surveys have come on line.  To maximise the output from the surveys requires a well-understood photometric system and a large grid of standard stars. Calibration of the Two Micron All Sky Survey (2MASS) is described in \\citet{2M} and of the Deep Near-Infrared Survey (DENIS) in \\citet{Fo}. Even deeper surveys have now started at the United Kingdom Infrared Telescope (UKIRT, the UKIRT Infrared Deep Sky Survey or UKIDSS, see \\citet{He}), and at the Canada France Hawaii Telescope (CFHT, their WIRCAM instrument, \\citet{Pu}).   Another area of recent improvement is photometry at longer infrared wavelengths.  The UKIRT group presented 4-5$\\mu$m $L^{\\prime}M^{\\prime}$ standards in \\citet{Le}. Also, Cohen and collaborators have produced a series of papers extending absolute flux calibrations from the near-infrared to the mid-infrared, with particular application to mid-infrared satellites. The first paper and the most recent paper in this series, respectively, are \\citet{iras} and \\citet{Pr}. Calibration of the recently launched $Spitzer$ infrared space telescope is described by \\citet{Re}.   Despite this progress, a fundamental problem remained for infrared astronomers -- there was no single filter set or photometric system.  Not only did this mean that it was difficult to compare data, but existing filters tended to be too broad and include poor regions of the atmosphere, leading to additional noise and thus lower-accuracy data.  This situation was resolved when A. Tokunaga formed a consortium to purchase a set of well-defined filters that better matched the atmospheric windows \\citep{Si, To, To05}.  These filters have now been widely adopted; as well as UKIRT (and UKIDSS) the following facilities have purchased Mauna Kea Observatory (MKO) near-infrared filter sets: Arizona, Boston, Cornell, Florida, Hawaii, Kyoto, Montreal, Ohio State, Tohoku, Tokyo, Virginia and Wyoming Universities; the Anglo-Australian Observatory, CalTech, Center for Astrophysics, CFHT, European Southern Observatory, Gemini, Grenoble Observatory, the NASA Infrared Telescope Facility, the Isaac Newton Group, Keck, Korea Observatory, National Astronomical Observatory of Japan, National Optical Astronomy Observatory, Nordic Optical Telescope, Osservatorio Astrofisico de Arcetri, Rome Observatory, and Telescopio Nazionale Galileo.   In this paper we present a set of standard stars observed at UKIRT using the UKIRT Fast Track Imager (UFTI) and $JHK$ filters in the MKO system. The sample is described in \\S 2, the photometric system in \\S 3, the observing and analysis techniques in \\S 4 and the results are given in \\S 5. Comparison to published observations of the sample are given in \\S 6. Our conclusions are given in \\S 7.  \\begin{figure} \\includegraphics[scale=.45]{Leggett_fig1.ps} \\caption{$JHK$ filter transmission profiles, where the dotted line includes typical site-dependent atmospheric effects. Central wavelengths and half-power bandwidths are indicated for filter$+$atmosphere profiles.  See \\citet{To05} for a discussion of isophotal and effective wavelengths for the MKO filter set.} \\label{filt-fig} \\end{figure}  \\section[]{The Sample}  The sample consists primarily of UKIRT Faint Standards \\citep{Ha} and Las Campanas (LCO, or NICMOS) standards \\citep{Ps}. Two stars from Hawarden et al., FS 18 and 118, have been removed from this sample due to the presence of a nearby star that makes aperture photometry less accurate. Three additional unreddened A0 stars, or stars with zero colours, were taken from \\citet{La}.  Such stars are useful for extrapolating magnitudes to other bandpasses. The sample is listed in Table 1; coordinates are taken from the Simbad astronomical database, and proper motions and types from sources provided by the database. References are given in Table 1.  For stars without spectral types in the literature we have estimated an approximate type from colour, where possible.  \\begin{table*} \\centering \\begin{minipage}{300mm} \\caption{Standard Stars.} \\begin{tabular}{@{}rlrrccccc@{}} \\hline UKIRT &  Name  & \\multicolumn{2}{c}{Coordinates (J2000)} & \\multicolumn{2}{c}{Proper Motion (mas/yr)} & P. M. &Spectral & Type\\\\ FS No. &       & RA & Dec &   RA       &         Dec               & Ref. & Type    & Ref.   \\\\ \\hline 101 & CMC 400101       &  0:13:43.6 & $+$30:38:00 & $-$5.3 & $-$8.9 & 1 & F0 & 2  \\\\ \\nodata & BRI B0021$-$0214 &  0:24:24.6 &  $-$1:58:20 & $-$136  & $+$162 & 3 & M9.5V & 4   \\\\ 102 & GSPC P525-E      &  0:24:28.4  &  $+$7:49:02  & \\nodata & \\nodata & \\nodata & G3  & 5 \\\\ 1 & G158-100         &  0:33:54.6 & $-$12:07:59 & $+$154  & $-$183 & 6 & DK-G & 7  \\\\ 103 & GSPC P241-G      &  0:36:29.6 & $+$37:42:54 & \\nodata & \\nodata & \\nodata & K2 & 5  \\\\ 2 & SA92-342         &  0:55:09.9 & $+$0:43:13  & $-$2.2  & $-$0.2 & 1 & F5 & 8  \\\\ \\nodata & GSPC S754-C      &  1:03:15.9  & $-$4:20:44    & \\nodata & \\nodata & \\nodata & F/G$^\\dag$  & \\nodata \\\\ 3 & Feige 11         &  1:04:21.7 & $+$4:13:37  & $+$12.3 & $-$28.6 & 9 & sdB & 10  \\\\ 104 & GSPC P194-R      &  1:04:59.6 & $+$41:06:31 & $+$0.4 & $-$4.3 & 1 & A7 & 5   \\\\ 105 & GSPC P527-F      &  1:19:08.2 & $+$7:34:12  & \\nodata & \\nodata & \\nodata & K1 & 5  \\\\ 106 & GSPC P152-F      &  1:49:46.9 & $+$48:37:53 & \\nodata & \\nodata &  \\nodata & K4 & 5  \\\\ 107 & CMC 600954       &  1:54:10.0 & $+$45:50:38 & $-$25.1 & $-$4.0 & 11 & G0 & 2  \\\\ 5 & Feige 16         &  1:54:34.7 & $-$6:46:00  & \\nodata & \\nodata & \\nodata & A0 & 7  \\\\ 4 & SA 93-317        &  1:54:37.7 & $+$0:43:00  & $-$9.3  & $-$18.5 & 1 & F5 & 8   \\\\ 6 & Feige 22         &  2:30:16.6 & $+$5:15:51  & $+$71.2 & $-$24.6 & 9 & DA3 & 12  \\\\ \\nodata & GSPC P530-D      &  2:33:32.2  & $+$6:25:38    & \\nodata & \\nodata & \\nodata  & G$^\\dag$ & \\nodata \\\\ 7 & SA 94-242        &  2:57:21.2 & $+$0:18:39  & \\nodata & \\nodata & \\nodata & A2  & 8 \\\\ 108 & CMC 502032       &  3:01:09.8 & $+$46:58:48 & $+$1.1 & $-$0.7 & 1 & F8 & 2   \\\\ \\nodata & TVLM 832-38078   &  3:04:01.8  & $+$00:45:50   & \\nodata & \\nodata & \\nodata & M-V & 13  \\\\ 109 & LHS 169          &  3:13:24.2 & $+$18:49:38 & $+$1283 & $-$1061 & 3 & esdK7 &  14  \\\\ \\nodata & GSPC P247-U      &  3:32:03.0  & $+$37:20:41   & \\nodata & \\nodata & \\nodata & G/K$^\\dag$ & \\nodata \\\\ 110 & GSPC P533-d      &  3:41:02.4  & $+$6:56:13    & \\nodata & \\nodata & \\nodata & G5  & 5 \\\\ 111 & CMC 601790       &  3:41:08.6 & $+$33:09:36 & $+$2.5  & $+$3.0 & 11 & G5 & 2 \\\\ 112 & GSPC S618-D      &  3:47:40.7 & $-$15:13:14 & \\nodata & \\nodata & \\nodata & G0 & 5  \\\\ 10 & GD 50            &  3:48:50.2 & $-$0:58:31  & \\nodata & \\nodata & \\nodata & DA1 & 12 \\\\ 113 & GSPC P117-F      &  4:00:14.1 & $+$53:10:39 & \\nodata & \\nodata & \\nodata & K0 & 5   \\\\ 114 & Melotte 25 LH 214 & 4:19:41.6  & $+$16:45:22   & \\nodata & \\nodata & \\nodata & M7V & 15  \\\\ 115 & B216-b5          &  4:23:18.2  & $+$26:41:15   & \\nodata & \\nodata & \\nodata &  \\nodata & \\nodata \\\\ 116 & B216-b7          &  4:23:50.2  & $+$26:40:07   & \\nodata & \\nodata & \\nodata &  \\nodata & \\nodata \\\\ 117 & B216-b9          &  4:23:56.5  & $+$26:36:38   & \\nodata & \\nodata & \\nodata & \\nodata & \\nodata  \\\\ 11 & SA 96-83         &  4:52:58.9 & $-$0:14:41  & $+$0.3  & $-$3.2 & 1 & A3 & 8   \\\\ 119 & HD 289907        &  5:02:57.5 & $-$1:46:43  & $+$1.3 & $-$5.7 & 1 & A2 & 2  \\\\ \\nodata & GSPC S840-F      &  5:42:32.2  & $+$0:09:04    & \\nodata & \\nodata & \\nodata  & G$^\\dag$ & \\nodata  \\\\ 12 & GD 71            &  5:52:27.6 & $+$15:53:13 & $+$92   & $-$189 & 12 & DA1 & 12  \\\\ 13 & SA 97-249        &  5:57:07.6 & $+$0:01:12  & $+$19.3 & $+$2.4 & 1 & G5V  & 8 \\\\ 120 & LHS 216          &  6:14:01.2 & $+$15:09:53 & $+$628  & $-$1249 & 3 & sdM2 & 14    \\\\ \\nodata & GSPC S842-E      &  6:22:43.7  & $-$0:36:30    & \\nodata & \\nodata &\\nodata & G/K$^\\dag$ &\\nodata  \\\\ \\nodata & SA 98-653        &  6:52:05.0 & $-$0:18:18  &  $-$0.3 & $-$3.5  & 1 & B9  & 8 \\\\ 121 & GSPC S772-G      &  6:59:46.8 & $-$4:54:33  & \\nodata & \\nodata & \\nodata & K3 & 5  \\\\ 122 & GSPC P161-D      &  7:00:52.0  & $+$48:29:24   & \\nodata & \\nodata & \\nodata & G0 & 5  \\\\ 14 & Rubin 149A       &  7:24:14.0    & $-$0:32:42    & \\nodata & \\nodata & \\nodata & O9-B2p & 7  \\\\ \\nodata & Rubin 149D       &  7:24:15.4 & $-$0:32:48  & $+$0.70 & $-$8.30 & 1 & A0 & 2   \\\\ \\nodata & GSPC P309-U      &  7:30:34.6  & $+$29:51:12   & \\nodata & \\nodata & \\nodata  & G$^\\dag$ & \\nodata \\\\ \\nodata & GSPC S495-E      &  8:27:12.5  & $-$25:08:01   & \\nodata & \\nodata & \\nodata & G/K$^\\dag$ & \\nodata  \\\\ \\nodata & GSPC P545-C      &  8:29:25.2  & $+$5:56:08    & \\nodata & \\nodata & \\nodata & F$^\\dag$ & \\nodata  \\\\ \\nodata & LHS 2026         &  8:32:30.5 & $-$1:34:39  & $+$155  & $-$473 & 3 & M6Ve & 16    \\\\ \\nodata & GSPC S705-D      &  8:36:12.5  & $-$10:13:39   & \\nodata & \\nodata & \\nodata & F/G$^\\dag$ & \\nodata  \\\\ 15 & M67-I-48         &  8:51:05.7 & $+$11:43:46 & \\nodata & \\nodata & \\nodata & G5IV-V & 5   \\\\ 123 & GSPC P486-R      &  8:51:11.8 & $+$11:45:22   & $-$7.7 & $-$5.6 & 9 & B8V & 17   \\\\ 16 & M67-IV-8         &  8:51:15.1 & $+$11:49:21 & \\nodata & \\nodata & \\nodata & G1V & 5   \\\\ 17 & M67-IV-2         &  8:51:19.7 & $+$11:52:11 & \\nodata & \\nodata & \\nodata & G4V & 5   \\\\ 124 & LHS 254          &  8:54:12.3 & $-$8:05:00  & $+$919  & $-$789 & 3 & M6.5V  & 14 \\\\ 125 & GSPC P259-C      &  9:03:20.6 & $+$34:21:04 & \\nodata & \\nodata & \\nodata & G8 & 5  \\\\ 126 & GSPC P487-F      &  9:19:18.7 & $+$10:55:54 & \\nodata & \\nodata & \\nodata & K3 & 5  \\\\ \\nodata & GSPC S852-C      &  9:41:35.9  & $+$0:33:12    & \\nodata & \\nodata & \\nodata & G$^\\dag$ & \\nodata  \\\\ \\nodata & GSPC S708-D      &  9:48:56.5  & $-$10:30:32   & \\nodata & \\nodata & \\nodata & K$^\\dag$ & \\nodata  \\\\ 127 & GSPC P212-C      & 10:06:29.0  & $+$41:01:27 & \\nodata & \\nodata & \\nodata & F9  & 5 \\\\ 19 & G 162-66         & 10:33:42.8 & $-$11:41:38 & $-$342  & $-$27.2 & 18  & DA2 & 12  \\\\ \\nodata & GSPC P550-C      & 10:33:51.9  & $+$4:49:05    & \\nodata & \\nodata & \\nodata & F$^\\dag$ & \\nodata  \\\\ 128 & LHS 2347         & 11:05:10.5 & $+$7:06:50  & $-$466  & $-$151 & 19 & M5V  & 16 \\\\ 20 & G 163-50         & 11:08:00.0 & $-$5:09:26  & $-$63   & $-$440 & 3 & DA3 & 12  \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*} \\begin{table*} \\begin{minipage}{300mm} \\contcaption{} \\begin{tabular}{@{}rlrrccccc@{}} \\hline UKIRT &  Name  & \\multicolumn{2}{c}{Coordinates (J2000)}   & \\multicolumn{2}{c}{Proper Motion (mas/yr)} & P. M. & Spectral & Type \\\\ FS No. &       & RA & Dec &   RA       &         Dec               & Ref. & Type   & Ref.     \\\\ \\hline 129 & LHS 2397aAB      & 11:21:49.3 & $-$13:13:08 & $-$508  & $-$80 & 3  & M8V+L7.5 & 20  \\\\ 130 & GSPC P264-F      & 11:24:55.9 & $+$34:44:39 & \\nodata & \\nodata & \\nodata & K4 & 5  \\\\ 21 & GD 140           & 11:37:05.1 & $+$29:47:58 & $-$146.9 & $-$5.9 & 9 & DA3  & 12 \\\\ 131 & GSPC P266-C      & 12:14:25.5  & $+$35:35:55   & \\nodata & \\nodata & \\nodata & F8 & 5  \\\\ 132 & GSPC S860-D      & 12:21:39.4  & $-$0:07:13    & \\nodata & \\nodata & \\nodata & G1 & 5  \\\\ 33 & GD 153           & 12:57:02.3 & $+$22:01:53 & $-$33   & $-$206 & 3 & DA1 & 12  \\\\ 133 & GSPC P172-E      & 13:15:52.8 & $+$46:06:37 & \\nodata & \\nodata & \\nodata & G9 & 5  \\\\ \\nodata & GSPC S791-C      & 13:17:29.6  & $-$5:32:37    & \\nodata & \\nodata & \\nodata  & F/G$^\\dag$ & \\nodata \\\\ 23 & M3-193           & 13:41:43.7 & $+$28:29:51 & $+$3.4  & $+$2.0 & 21 & G8III  & 5 \\\\ \\nodata & GSPC P133-C      & 13:58:40.3  & $+$52:06:24   & \\nodata & \\nodata & \\nodata  & F$^\\dag$ & \\nodata \\\\ \\nodata & GSPC P499-E      & 14:07:34.0  & $+$12:23:51   & \\nodata & \\nodata & \\nodata  & G$^\\dag$ & \\nodata \\\\ 134 & LHS 2924         & 14:28:43.3 & $+$33:10:38 & $-$337  & $-$747 & 19 & M9Ve  & 16  \\\\ 135 & GSPC S867-V      & 14:40:58.0  & $-$0:27:48    & \\nodata & \\nodata & \\nodata & G5  & 5 \\\\ \\nodata  & GSPC P272-D      & 14:58:33.1  & $+$37:08:33   & \\nodata & \\nodata & \\nodata & G$^\\dag$ & \\nodata  \\\\ 136 & GSPC S868-G      & 14:59:32.1 & $-$0:06:17  & \\nodata & \\nodata & \\nodata & K2 & 5  \\\\ \\nodata & TVLM 868-53850   & 15:00:26.4  & $-$0:39:28    & \\nodata & \\nodata & \\nodata & M5Ve  & 22 \\\\ \\nodata & TVLM 868-110639  & 15:10:17.2  & $-$2:41:07    & \\nodata & \\nodata & \\nodata & M9V & 4  \\\\ \\nodata & GSPC S870-T      & 15:39:03.6  & $+$0:14:54    & \\nodata & \\nodata & \\nodata  & G$^\\dag$ & \\nodata  \\\\ \\nodata & GSPC P177-D      & 15:59:14.0 & $+$47:36:42 & \\nodata & \\nodata & \\nodata & G2V & 23  \\\\ 137 & GSPC P565-C      & 16:26:42.8  & $+$5:52:20    & \\nodata & \\nodata & \\nodata  & G1 & 5  \\\\ 138 & GSPC P275-A      & 16:28:06.7 & $+$34:58:48 & $-$11.7 & $+$5.3 & 1 & A1 & 5   \\\\ \\nodata & GSPC P330-E      & 16:31:33.9 & $+$30:08:47 & \\nodata & \\nodata & \\nodata  & G2V & 23 \\\\ 139 & GSPC P137-F      & 16:33:53.0 & $+$54:28:22 & \\nodata & \\nodata  & \\nodata  & K1 & 5  \\\\ 27 & M13-A14          & 16:40:41.3 & $+$36:21:13 & $+$3.9 & $-$1.8 & 24  &  G8IV/V & 5  \\\\ 140 & GSPC S587-T      & 17:13:22.7 & $-$18:53:34 & \\nodata & \\nodata & \\nodata  & G9  & 5 \\\\ \\nodata & GSPC P138-C      & 17:13:44.6  & $+$54:33:21   & \\nodata & \\nodata & \\nodata & G$^\\dag$ & \\nodata  \\\\ 141 & P489-D           & 17:48:58.9 & $+$23:17:44 & \\nodata & \\nodata & \\nodata & G2 & 5  \\\\ 35 & GSC 00441-01200  & 18:27:13.5 & $+$4:03:09  & $-$9.0  & $-$4.4 & 18 & K0 & 5   \\\\ 143 & Ser-EC68         & 18:29:53.9  & $+$1:13:31    & \\nodata & \\nodata & \\nodata & \\nodata & \\nodata  \\\\ 144 & Ser-EC84         & 18:29:57.0 & $+$1:12:47  & \\nodata & \\nodata & \\nodata & \\nodata & \\nodata  \\\\ \\nodata & GSPC P182-E      & 18:39:33.7  & $+$49:05:38   & \\nodata & \\nodata & \\nodata  & G$^\\dag$ & \\nodata \\\\ \\nodata & LDN 547          & 18:51:15.6  & $-$4:16:02    & \\nodata & \\nodata & \\nodata & \\nodata & \\nodata  \\\\ 146 & GSPC P280-U      & 18:54:04.0 & $+$37:07:19 & \\nodata & \\nodata & \\nodata & K1 & 5  \\\\ 147 & GSPC P230-A      & 19:01:55.3 & $+$42:29:19 & $-$1.4 & $-$0.1 & 1 & A0  & 2 \\\\ \\nodata & GSPC S808-C      & 19:01:55.5  & $-$4:29:12    & \\nodata & \\nodata & \\nodata & G/K$^\\dag$  & \\nodata  \\\\ 148 & GSPC S810-A      & 19:41:23.4 & $-$3:50:57  & $-$1.5 & $-$4.5 & 9 & A0 & 25   \\\\ 149 & GSPC P338-C      & 20:00:39.2 & $+$29:58:38 & $+$4.5 & $-$4.3 & 1 & B7.5V  & 26 \\\\ 150 & CMC 513807       & 20:36:08.4 & $+$49:38:24 & $+$8.2 & $+$9.3 & 1 & G0 & 2   \\\\ \\nodata & GSPC S813-D      & 20:41:05.2  & $-$5:03:42    & \\nodata & \\nodata & \\nodata  & G$^\\dag$ & \\nodata \\\\ 34 & EG 141           & 20:42:34.8 & $-$20:04:35 & $+$354.6 & $-$97.9 & 9 & DA2.5  & 12  \\\\ \\nodata & GSPC P576-F      & 20:52:47.4  & $+$6:40:05    & \\nodata & \\nodata & \\nodata & G$^\\dag$  & \\nodata \\\\ 151 & GSPC P340-H      & 21:04:14.8 & $+$30:30:21 & \\nodata & \\nodata & \\nodata & G2 & 5  \\\\ 29 & G 93-48          & 21:52:25.4 & $+$2:23:20  & $+$23.0 & $-$303.0 & 9 & DA3 & 12   \\\\ \\nodata & BRI B2202-1119   & 22:05:35.7 & $-$11:04:29 & $-$271  & $-$170 & 3 & M5.5V & 4   \\\\ 152 & GSPC P460-E      & 22:27:16.1 & $+$19:16:59 & \\nodata & \\nodata & \\nodata  & K1  & 5 \\\\ 30 & SA 114-750       & 22:41:44.7 & $+$1:12:36  & \\nodata & \\nodata  & \\nodata  & B9  & 8 \\\\ 153 & S820-E           & 23:02:32.1 & $-$3:58:53  & \\nodata & \\nodata & \\nodata  & K2 & 5  \\\\ 31 & GD 246           & 23:12:23.1 & $+$10:47:04 & $+$127 & $-$11 & 12 & DA1 & 12  \\\\ 32 & Feige 108        & 23:16:12.4 & $-$1:50:35  & \\nodata & \\nodata & \\nodata & sdB & 27   \\\\ 154 & GSPC S893-D      & 23:18:10.1  & $+$0:32:57    & \\nodata & \\nodata & \\nodata & G0 & 5  \\\\ \\nodata & GSPC S677-D      & 23:23:34.5  & $-$15:21:06   & \\nodata & \\nodata & \\nodata  & F$^\\dag$  & \\nodata \\\\ \\nodata & GSPC P290-D      & 23:30:33.5  & $+$38:18:57   & \\nodata & \\nodata & \\nodata & G$^\\dag$ & \\nodata\\\\ \\nodata & PG 2331$+$055A   & 23:33:44.5  & $+$5:46:41  & \\nodata & \\nodata & \\nodata  & sd & 10 \\\\ 155 & CMC 516589       & 23:49:47.8 & $+$34:13:05 & \\nodata & \\nodata & \\nodata & K5V & 28   \\\\  \\hline \\end{tabular} \\end{minipage}\\\\ $^\\dag$ Spectral type estimated from $B-V$ and $V-K$ from Lasker et al. 1988 and this work.\\\\ References: (1) The Tycho Reference Catalogue, Hog et al. 1998 (2) AGK3 Catalogue, Heckmann \\& Dieckvoss  1975 (3) Salim \\& Gould 2003 (4) Kirkpatrick, Henry \\& Simons 1995 (5) Hawarden et al. 2001 (estimated from colour or colour-magnitude diagrams) (6)  Harrington \\& Dahn 1980 (7) Turnshek et al. 1990 (8) Drilling \\& Landolt 1979 (9) The Hipparcos Catalogue, Perryman et al. 1997 (10) The Palomar-Green Survey, Green, Schmidt \\& Liebert 1986 (11) PPM North Star Catalogue, Roeser \\& Bastian 1988 (12) McCook \\& Sion 1999 (13) Tinney 1993 (14) Gizis 1997 (15) Reid \\& Hawley 1999 (16) Leggett 1992 (17) Pesch 1967 (18) The Second U.S. Naval Observatory CCD Astrograph Catalog (UCAC2), Zacharias et al. 2004 (19) Bakos, Sahu \\& N\\'{e}meth 2002 (20) Freed, Close \\& Siegler 2003 (21) Tucholke, Scholz \\& Brosche 1994 (22) Gizis 2002 (23)  Colina \\& Bohlin 1997 (24) Kadla 1966 (25) Henry Draper Catalogue, Cannon \\& Pickering 1989 (26) Straizys \\& Kalytis 1981 (27) Greenstein \\& Sargent 1974 (28) Stephenson 1986. \\\\ \\end{table*} ",
        "conclusions": "We have presented $JHK$ magnitudes for 115 stars in the brightness range of 10--15th magnitude. The photometry is in the widely adopted MKO photometric system, which uses well-defined filters matched to the atmospheric windows. The internal accuracy is on average 0$\\fm$011, with the stars being observed typically over 4 nights.   The results presented here should be used in preference to those of \\citet{Ha} if observations are being made in the MKO system, as the filters, and hence the photometric systems, are significantly different from the UKIRT system.  We present colour transformations between the MKO and the 2MASS, LCO and UKIRT photometric systems.  However these transformations can be luminosity-class dependent at the 5\\% level for redder stars with strong absorption bands.  For very accurate transformations between systems, the exact spectral type needs to be known and transformations calculated and applied as a function of type (or the stars should be observed in the required system).  We find that the difference between the $JHK$ magnitudes presented here and those presented previously by \\citet{Ha} shows a dependence on brightness.  We explain this in terms of an under-estimate of the linearity correction for stars with $J<11$ and $H$,$K<10.5$, such that the photometry presented here and by Hawarden et al. is affected by $\\lesssim$0$\\fm$02. The effect is likely to be more significant in the results of Hawarden et al. than in the present results: this is supported by comparisons with 2MASS data.  There are 84 stars in the sample presented here that have $11 < J < 15$ and $10.5 < H$,$K < 15$, are not suspected to be variable, and have magnitudes with an estimated error $\\leq 0\\fm$027; 79 of these have an error of  $\\leq 0\\fm$020. We recommend that these be adopted as primary standards for the MKO near-infrared ($JHK$) photometric system."
    },
    "0609/astro-ph0609377_arXiv.txt": {
        "abstract": " ",
        "introduction": "An ensemble of gravitational waves (GWs) from a number of inspiraling binaries of compact objects at different redshifts can be observed as stochastic gravitational wave background radiation (GWBR). Coalescing supermassive black hole (SMBH) binaries with masses in the rage $10^6$ -- $10^9~M_{\\odot}$ emit GWs, as a results of the merging of their host galaxies. Recent studies\\cite{jaffe03,enoki04} predict that these SMBH binaries  produce GWBR of frequencies $\\sim 1~{\\rm n}$ -- $1~\\mu{\\rm Hz}$. In the frequency range  $0.1~{\\rm m}$ -- $10~{\\rm mHz}$, it is believed that extragalactic close binaries of white dwarfs are dominant sources of GWBR.\\cite{farmer03}   Gravitational waves with frequencies in the range $1~{\\rm n}$ -- $100~{\\rm nHz}$ can be detected by pulsar timing measurements.\\cite{detweiler79} \\ Thus we should be able to detect the GWBR from SMBH binaries directly. Recently, Jenet et al.\\cite{jenet05} developed a new method for detecting GWBR using multiple pulsar timing data. In the case of the Parkes Pulsar Timing Array (PPTA) project,\\cite{hobbs05}\\footnote{See http://www.atnf.csiro.au/research/pulsar/array/} their  results showed that with regular timing observations of 20 pulsars with a timing accuracy of $100~{\\rm ns}$, we should be able to directly detect the predicted levels of the GWBR from SMBH binaries within five years. \\cite{jenet05} \\  In previous studies \\cite{jaffe03,enoki04} of the GWBR from SMBH binaries, it was assumed that all binaries are in circular orbits. However, binary orbits are generally eccentric. Before entering into the GW emitting  regime, the evolution of a SMBH binary is determined by its dynamical interaction with field stars in the center of its host galaxy. Mikkola and Valtonen \\cite{mikkola92} constructed approximate expressions relating the energy and angular momentum transfer rates for a heavy binary and the surrounding field of light stars, and they showed that the length of the semi-major axis decreases and the eccentricity increases owing to dynamical friction. Fukushige et al.\\cite{fukushige92} investigated the evolution of a SMBH binary in a uniformly distributed background of field stars and found that the dynamical friction on the eccentric binary is most effective at the apocenter, where the orbital velocity is minimum, and thus the eccentricity increases. Iwasawa et al.\\cite{iwasawa05} performed $N$-body simulations of the dynamical evolution of triple SMBH (a binary of SMBHs and a single SMBH) systems in galactic nuclei and showed that the eccentricity of the SMBH binary reaches $e >0.9$ in the GW emission regime. They found that both the Kozai mechanism and the thermalization of eccentricity due to the strong binary-single SMBH interaction drive the increase of the orbital eccentricity. Matsubayashi et al. \\cite{matsubayashi05} investigated the orbital evolution of intermediate mass black hole (IMBH)-SMBH systems in galactic centers, using $N$-body simulations. They found that the eccentricity approaches unity ($e > 0.8$) and the IMBH can cause the SMBH to rapidly coalesce as a result of GW emission. Armitage and Natarajan \\cite{armitage05} studied the evolution of the SMBH binary eccentricity which is excited by the interaction between a binary and circumbinary gas disk. They estimated a typical eccentricity at one week prior to coalescence to be $e \\sim 0.01$. In the case of extreme-mass ratio binaries, higher eccentricities ($e \\sim 0.1$) are possible.  An eccentric binary emits GWs at all integer harmonics of the orbital frequency.\\cite{peters63,fitchett87} \\  Thus, the spectral energy distribution (SED) is different from that of a binary in a circular orbit, even if the masses and semi-major axes of the two binaries are the same. For larger eccentricities, radiation from higher harmonics are greater and thus the peak of the SED shifts toward higher frequency. Moreover, the orbital evolution of a binary due to GW radiation depends strongly on its eccentricity. As a result, the power of GW radiation of an eccentric binary is larger than that of a circular binary, and the timescale of GW radiation of an eccentric binary is shorter than that of a circular binary. Therefore, in order to predict the power spectrum of GWBR from compact binaries, it is necessary to take account of orbital eccentricities of binaries.  For the cases of GWBRs from galactic and extragalactic neutron star binaries and black hole MACHO binaries, some studies have taken account of the effect of eccentricity.\\cite{farmer03,ioka99} \\ However, for the case of the GWBR from SMBH binaries, none of models include the effect of eccentricity.  In this study, we investigate the effect of orbital eccentricity on the expected GWBR from SMBH binaries. First, we formulate the power spectrum of GWBR from cosmological eccentric binaries, taking into account eccentricity evolution. In order to formulate the power spectrum, we adopt a simple relationship between the power spectrum of the GWBR produced by cosmological GW sources, the total time-integrated energy spectrum of individual sources, and the comoving number density of GW sources found by Phinney.\\cite{phinney01} Next, we apply this formulation to the case of the GWBR from coalescing SMBH binaries. In order to estimate the number density of SMBH binaries,  we use a semi-analytic (SA) model\\cite{enoki03} \\ in which SMBH formation is incorporated into galaxy formation.\\cite{nagashima01}   This paper is organized as follows. In \\S\\ref{sec.binary} we briefly review the GW emission of a binary in an eccentric orbit. In \\S\\ref{sec.gwbr} we formulate the power spectrum of GWBR from cosmological compact binaries in eccentric orbits. In \\S\\ref{sec.smbh} we present the power spectrum of the GWBR from eccentric SMBH binaries. In \\S\\ref{sec.summary} we present a summary and conclusions. ",
        "conclusions": "\\label{sec.summary}  In this study, we have investigated how orbital eccentricities of binaries affect power spectra of GWBR from coalescing SMBH binaries. A compact binary in an eccentric orbit radiates GWs at all integer harmonics of its orbital frequency. Owing to this harmonic radiation, the SED, the power and the timescale of the GW emission of a binary in an eccentric orbit are different from those of a binary in a circular orbit. Therefore, first, we formulated the power spectrum of GWBR from cosmological compact binaries in eccentric orbits. Then using this formulation  and our SA model for galaxy and SMBH formation, we calculated the power spectra of GWBR from coalescing SMBH binaries in eccentric orbits.   We found that the calculated power spectra of the GWBR from SMBH binaries in eccentric orbits are suppressed owing to the harmonic radiation for lower frequencies ($f \\lesssim 1~{\\rm nHz}$) if the initial eccentricity satisfies $e_0 > 0.2$ at $a = 300~R_{\\rm S}$. The degree of this suppression depends strongly on the initial eccentricity. In this paper, for simplicity and in order to clarify the effect of the eccentricity, we assumed that all binaries have the same initial eccentricity. The results of our study suggest that elucidating the initial eccentricity distribution is essential for the investigation of the power spectrum of GWBR from compact binaries.  Because the  number density of coalescing SMBH binaries is determined by the processes of galaxy formation, the power spectrum of the GWBR from SMBH binaries depends strongly on these processes, especially processes related to star formation, which regulate the amount of cold gas accreted onto SMBHs. The prediction of our model shows that the overall shape and amplitude of the power spectrum of the GWBR from coalescing SMBH binaries depend on galaxy formation processes. However, for low frequencies ($f \\lesssim 1~{\\rm nHz}$), the shape of the power spectrum of the GWBR also depends on the initial eccentricity. Because the pulsar timing measurements can provide limits on the level of GWBR with frequencies $f \\sim 1~{\\rm n}- 100~{\\rm nHz}$, pulsar timing observations, such as PPTA project, should provide constraints not only on the number density of coalescing SMBH binaries  but also on the effect of orbital eccentricity on the GWBR from SMBH binaries."
    },
    "0609/astro-ph0609694_arXiv.txt": {
        "abstract": "Time delays between lensed multiple images have been known to provide an interesting probe of the Hubble constant, but such application is often limited by degeneracies with the shape of lens potentials. We propose a new statistical approach to examine the dependence of time delays on the complexity of lens potentials, such as higher-order perturbations, non-isothermality, and substructures. Specifically, we introduce a reduced time delay of the dimensionless form, and explore its behavior analytically and numerically as a function of the image configuration that is characterized  by the asymmetry and opening angle of the image pair. In particular we derive a realistic conditional probability distribution for a given image configuration from Monte-Carlo simulations. We find that the probability distribution is sensitive to the image configuration such that more symmetric and/or smaller opening angle image pairs are more easily affected by perturbations on the primary lens potential. On average time delays of double lenses are less scattered than those of quadruple lenses. Furthermore, the realistic conditional distribution allows a new statistical method to constrain the Hubble constant from observed time delays. We find that 16 published time delay quasars constrain the Hubble constant to be $H_0=70\\pm 6 \\,{\\rm km\\,s^{-1}Mpc^{-1}}$, where the value and its error are estimated using jackknife resampling. Systematic errors coming from the heterogeneous nature of the quasar sample and the uncertainty of the input distribution of lens potentials can be larger than the statistical error. After including rough estimates of the sizes of important systematic errors, we find $H_0=68\\pm 6({\\rm stat.})\\pm8({\\rm syst.}) \\,{\\rm km\\,s^{-1}Mpc^{-1}}$. The reasonable agreement of the value of the Hubble constant with other estimates indicates the usefulness of our new approach as a cosmological and astrophysical probe, particularly in the era of large-scale synoptic surveys. ",
        "introduction": "\\label{sec:intro}  It has been known that time delays between multiple images of strong gravitational lens systems offer an interesting method to measure the Hubble constant $H_0$, the most fundamental cosmological parameter that governs the length and time scale of our universe \\citep{refsdal64}. A huge advantage of this method is that it does not rely on so-called distance ladder and can measure the global Hubble constant independently of any local measurements. Motivated by this time delays have been measured in more than 10 lensed quasar systems \\citep[see, e.g.,][]{kochanek06a}. The situation as it presents is, however, somewhat confusing and controversial. \\citet{kochanek02,kochanek03} claimed from the analysis of several lens systems that the Hubble constant should be relatively low, $H_0\\sim 50 \\,{\\rm km\\,s^{-1}Mpc^{-1}}$. The time delay of SDSS J1650+4251 also prefers the low Hubble constant \\citep{vuissoz07}. On the other hand, \\citet{koopmans03} performed systematic mass modeling of B1608+656 using all available data from radio to optical and found constrained the value of the Hubble constant to be $H_0=75^{+7}_{-6}\\,{\\rm km\\,s^{-1}Mpc^{-1}}$. The analysis of the smallest separation lens B0218+357 yields $H_0=78\\pm 3\\,{\\rm km\\,s^{-1}Mpc^{-1}}$ \\citep{wucknitz04}. By combining time delays in 10 lensed quasar systems \\citet{saha06a} obtained $H_0=72^{+8}_{-11}\\,{\\rm km\\,s^{-1}Mpc^{-1}}$.  The large variation of derived values of the Hubble constant from time delays reflects the fact that time delays are quite sensitive to mass distributions of lens objects, which leads to strong degeneracies between lens mass profiles and the Hubble constant. The most fundamental, mathematically rigorous degeneracy is the mass-sheet degeneracy \\citep{falco85}; inserting an uniform sheet instead of decreasing the mass normalization of lenses changes estimated values of the Hubble constant while leaving unchanged the image observables. This degeneracy implies that the derived values of the Hubble constant also degenerates with radial density profiles of of lens galaxies \\citep{refsdal94,witt95,keeton97b,witt00,saha00,tada00, wucknitz02,treu02,kochanek02,kochanek03,oguri03b,rusin03,schechetr05, mortsell05,kawano06}. In addition, it may degenerate with the angular structure of lenses as well \\citep{zhao03,saha06b}. The sensitivity of time delays on mass profiles suggests that by assuming the Hubble constant, which can be determined by other methods \\citep{freedman01,tegmark06,spergel07}, we can put constraints on mass distributions of lenses, in particular radial density profiles \\citep[e.g.,][]{oguri04a,kochanek06b,dobke06}.  One way to overcome the degeneracies is to adopt lens systems whose lens potentials can well be constrained by many observational constraints.  An example of such constraints comes from lensed images of quasar host galaxies. The lensed host galaxies often form Einstein rings, which accurately and independently determines the structure of the lens potential as well as the shape of the lensed host galaxy \\citep[e.g.,][]{keeton00,kochanek01,koopmans03}. The small-scale structure of lensed quasars, such as radio jets and sub-components, determine lens potentials in a similar way as Einstein rings \\citep[e.g.,][]{cohn01,wucknitz04}.  Furthermore, the measurement of velocity dispersions of lens galaxies serves as useful constraints on mass distributions, helping to break degeneracies between mass models \\citep[e.g.,][]{treu02}. Lens systems for which these strong additional constraints are available, sometime referred as ``golden lenses'', have been thought to be an effective probe of the Hubble constant.  Another potentially powerful, but less studied, method to measure the Hubble constant is a statistical approach. Even if each individual lens lacks strong constraints that allow detailed investigations of the mass distribution, by combining many lens systems one can put tight constraints on the Hubble constant. As mentioned above, this approach was in some sense demonstrated recently by \\citet{saha06a} who combined 10 lensed quasar systems to constrain the Hubble constant. A caveat is that lensed quasars sometimes suffer from selection effects. For instance, brighter lensed quasars  with larger image separations are more likely to lie in dense environments such as groups and clusters \\citep{oguri05b,oguri06a}, thus the Hubble constant from those lensed quasars may be systematically higher than the true value without any correction of the effect of dark matter along the line of sight. This indicates the importance of well-defined statistical sample of lensed quasars for which we can quantitatively estimate and correct the selection effect. While the statistical sample has been very limited so far, containing $\\sim 20$ lenses even for the largest lens sample \\citep{myers03,browne03}, larger lens surveys will be soon available by ongoing/future lens surveys such as done by the Sloan Digital Sky Survey \\citep{oguri06b}, the Large Synoptic Survey Telescope (LSST)\\footnote{http://www.lsst.org/}, and the Supernova/Acceleration Probe (SNAP)\\footnote{http://snap.lbl.gov/}. Future lens surveys will also find strong lensing of distant supernovae \\citep[e.g.,][]{oguri03a} for which time delays can easily be measured \\citep[but see][]{dobler06}. Therefore the statistical approach is growing its importance.  In this paper, we study how time delays depend on various properties of the lens potential and image configurations, which is essential for the determination of the Hubble constant from time delays. Using both analytic and numerical methods, we show how time delays are affected by several complexity of the lens potential, such as radial mass profiles, external perturbations, higher order multipoles, and substructures. We then derive the expected distributions of time delays by adopting realistic lens potentials. The distribution, in turn, can be used to derive statistically the value of the Hubble constant from observed time delays. This approach differs from the statistical argument by \\citet{saha06a} in the sense that they fit image positions of individual lens systems to constrain the Hubble constant and then combined results of all the lens systems: Our new approach does not even require modeling of each lens system. This has an advantage that we can include lens systems that have too few constraints to determine the lens potential. In this sense the approach is extension of study by \\citet{oguri02} in which only spherical  halos are considered to compute time delay distributions. We note that the methodology is similar to that adopted by \\citet{keeton03a,keeton05} who derived distributions of flux ratios of image pairs to identify small-scale structure in lens galaxies.  In addition to the measurement of the Hubble constant, the sensitivity of time delays on mass models, which we explore in this paper, offers guidance on the usefulness of each time delay measurement as a cosmological or astrophysical probe: If time delays at some image configuration is quite sensitive to detailed structure of the lens such as higher-order multipole terms with small amplitudes or substructures, which are difficult to be constrained even for best-studied lens systems,  it is almost hopeless to use these time delays to extract either radial mass profiles or the Hubble constant. Our result can be used to assess quantitatively which lens systems are more suitable for detailed studies, i.e., less sensitive to the complexity of the lens potential. There has been several insightful analytic work \\citep[e.g.,][]{witt00,kochanek02}, but here we perform more systematic and comprehensive survey of model dependences of time delays by parameterizing image configurations of lensed quasar systems using dimensionless quantities.  This paper is organized as follows. In \\S \\ref{sec:model} we introduce several dimensionless quantities that are used to explore the model dependence of time delays. We study the sensitivity of time delays on the various lens potentials analytically in \\S \\ref{sec:ana}. Section \\ref{sec:pot} is devoted to construct the conditional distribution of time delays from realistic Monte-Carlo simulations. We compare it with observed time delays in \\S \\ref{sec:obs}, and constrain the Hubble constant in \\S \\ref{sec:hub}. Finally discussion of our results and conclusion are given in \\S \\ref{sec:dis}. Throughout the paper we adopt a flat universe with the matter density $\\Omega_M=0.24$ and the cosmological constant $\\Omega_\\Lambda=0.76$ \\citep{tegmark06}, although our results are only weakly dependent on the specific choice of the cosmological parameters.  The Hubble constant is sometimes described by the dimensionless form $h\\equiv H_0/(100\\,{\\rm km\\,s^{-1}Mpc^{-1}})$. ",
        "conclusions": "\\label{sec:dis}  In this paper, we have studied time delays between multiply imaged quasars. Adopting the reduced time delay, which is a measure of how the lens potential is complicated compared with the simple isothermal form, we have explored the dependence of time delays on various complex structure of lens potentials such as external perturbations, non-isothermality, and substructures. The distribution of time delays has been studied as a function of image configuration which we characterize using two dimensionless quantities, the asymmetry and opening angle of an image pair. We have pointed out that the sensitivity on lens potentials is quite dependent on the image configuration. For instance, more symmetric image pairs are more affected by a small change of the lens potential. Image pairs with smaller opening angles are also more sensitive to lens potentials. In particular time delays of close image pairs are very sensitive to higher-order external perturbations and substructures that are very hard to be constrained from mass modeling even for best studied lens systems. In addition, those pairs usually have the largest relative uncertainties of time delay measurements. Therefore we conclude that it is quite difficult to extract any information on the Hubble constant or the mass distribution from close image pairs. It is interesting to note that perturbations on lens potential not only introduce scatter around the mean but also can systematically shift the distribution of time delays. One such example is smaller time delays for smaller opening angle image pairs, which is caused by external shear.  We have performed Monte-Carlo simulations to derive a probability distribution of reduced time delays for each image configuration. Input distributions are determined from observational and theoretical constraints. The distribution is weighted by the lensing cross section and magnification bias, allowing a realistic estimate of time delay distributions. The probability distribution was then compared with observed time delays. We have shown that the distribution of time delays computed from our simulations is in good agreement with observed time delays. In particular, distributions of observed time delays also exhibit strong dependence of image configuration in a consistent manner with our theoretical expectations. The probability distribution can be used to constrain the Hubble constant. We have found that 16 published time delay quasars constrain it to be $h=0.70^{+0.03}_{-0.02}$ at 68\\% confidence using the effective chi-square or $h=0.70\\pm0.06$ estimated using jackknife resampling, consistent with other estimates.  An important caveat is that our lensed quasar sample is quite heterogeneous. In particular, it should be noted that current time delay quasars (see Table \\ref{tab:xi}) have significantly larger image separations on average compared with the other quasar lenses: The median image separation of time delay quasars listed in Table \\ref{tab:xi} is $1\\farcs7$ (image separations before and after the median are $1\\farcs5$ and $2\\farcs1$), whereas that of all lensed quasars is $\\sim 1\\farcs4$. Quasar lenses with larger image separations are more likely to lie in dense environments because both the image separation and biased cross section are boosted by surrounding dark matter \\citep{oguri05b,oguri06a}, thus the Hubble constant inferred from those lenses are more affected by the environmental convergence. Indeed, the association of group/cluster has been reported for more than half of the time delay quasars \\citep[e.g.,][]{morgan05,fassnacht06,momcheva06,williams06,auger07}. In addition, our input distributions of external perturbations may be underestimated for these large separation lenses. To estimate this systematic effect, we exclude four lensed quasars with image separation larger than $3''$ and repeat the analysis done in \\S \\ref{sec:hub}. The resulting Hubble constant $h=0.67^{+0.04}_{-0.03}$ at 68\\% confidence from the effective chi-square is consistent with our full result, thus we conclude that the effect of lens galaxy environments is not so drastic here. However, to minimize the systematic effect, in the future we should apply our statistical method to well-defined samples of lensed quasars such as the CLASS \\citep{myers03,browne03}, SQLS \\citep{oguri06b}, and those obtained in future lens surveys.  Another source of the systematic effect is the uncertainty of our input distribution of lens potentials. Among others, the most important systematic error comes from the uncertainty of the mean value of the slope of the radial profile, $\\alpha$. While $\\alpha=1$ (isothermal) for the mean appears to be a reasonable choice, direct studies of lens galaxies \\citep[e.g.,][]{treu04,rusin05,hamana07} indicate that the error on the mean $\\alpha$ could be as large as 0.1. The derived Hubble constant depends on the slope as $h\\propto 2-\\alpha$, therefore the change of the mean $\\alpha$ systematically shifts the best-fit value. The scaling relation suggests that the 0.1 error of the mean results in 10\\% error on $h$, indicating that the systematic error may be even larger than the statistical error. Another important systematic effect is caused by external convergence from lens galaxy environments. Since its effect on the Hubble constant is straightforward \\citep[e.g.,][]{keeton04}, one can estimate the effect of external convergence rather easily even without including the distribution in the simulation. From the result of \\citet{oguri05b}, we expect that the posterior distribution of external convergence with lensing bias taken into account is roughly $\\kappa_{\\rm ext}\\sim 0.03\\pm0.03$ unless the image separation is too large. This, combined with the fact that the Hubble constant scales as $h\\propto 1-\\kappa_{\\rm ext}$, suggest that the effect of external convergence is not so dominant here (note that external convergence was already taken into account for two extreme lenses, Q0957+561 and RXJ0911+0551). By including rough estimates of these two systematic errors, we obtain $h=0.68\\pm0.06({\\rm stat.})\\pm0.08({\\rm syst.})$, indicating the importance of reducing the systematic error.  Since the Hubble constant is now determined fairly well by other methods, time delays are sometimes used to study mass distributions of lens objects. Our statistical technique offers a new method to study the lens mass distribution. By comparing the probability distributions for different input distributions of lens potentials (e.g., different median slopes of the primary lens galaxy), one can infer which input model is most plausible. Unlike previous statistical studies \\citep[e.g.,][]{oguri02}, this new method allows us to include various complexity of lens potentials relatively easily, particularly the non-spherically symmetric nature of lens potentials.  We note that our input distribution of lens potentials was designed for galaxies, but it is straightforward to modify it to that of lensing by other populations, e.g., lensing by a cluster of galaxies.  In summary, our new statistical approach is invaluable for the study of both cosmological parameters and structure of lens potentials.  We believe its importance grows more and more in the era of large-scale synoptic surveys such as LSST and SNAP: Quasar lens candidates are easily recognized in these synoptic surveys by making use of strong time variability of quasars \\citep{pindor05,kochanek06c}. Strong lensing of distant supernovae offers additional interesting opportunity to apply our statistical technique. The statistical analysis is essential to make efficient use of the large homogeneous samples of strong lenses provided by these surveys."
    },
    "0609/astro-ph0609141_arXiv.txt": {
        "abstract": "We present 24\\,$\\micron$ {\\it Spitzer}/MIPS photometric observations of the $\\sim$\\,50\\,Myr open cluster IC\\,2391. Thirty-four cluster members ranging in spectral type from B3-M5 were observed in the central square degree of the cluster. Excesses indicative of debris disks were discovered around 1 A star, 6 FGK stars, and possibly 1 M dwarf. For the cluster members observed to their photospheric limit, we find a debris disk frequency of 10$_{-3}^{+17}\\%$ for B-A stars and 31$_{-9}^{+13}\\%$ for FGK stars using a 15$\\%$ relative excess threshold. Relative to a model of decaying excess frequency, the frequency of debris disks around A-type stars appears marginally low for the cluster's age while that of FGK stars appears consistent. Scenarios that may qualitatively explain this result are examined. We conclude that planetesimal activity in the terrestrial region of FGK stars is common in the first $\\sim$\\,50\\,Myr and decays on timescales of $\\sim$100\\,Myr. Despite luminosity differences, debris disk evolution does not appear to depend strongly on stellar mass. ",
        "introduction": "Nearly all stars are believed to form with primordial accretion disks, but it is not clear whether the formation of planets is also a nearly-universal process of stellar evolution. Answering this question would help us understand the incidence of planetary systems in the Galaxy. Current planet detection techniques, while continuously improving, all suffer from some instrument sensitivity limitation. Many of the planets may very well be insufficiently massive to be detected through gravitational recoil, too faint against the glare of the central star to be imaged directly, too small to significantly reduce the star's measured brightness, or positioned unfavorably along the line of sight to produce a lensing event.  Planetary debris disks, however, provide an additional approach. Debris disks contain micron-sized dust grains predominantly produced in collisions between larger-sized bodies (such as rocks). These dust grains are heated by the parent star and re-radiate at longer wavelengths. A key facet of debris disks is that the dust grains must be short-lived compared to the age of the system given the efficiency of typical loss mechanisms like Poynting-Robertson drag and radiation pressure (with timescales of 10$^{6}$ to 10$^{7}$ years). The dust must therefore be regenerated either through a continuous collisional cascade or through stochastic collisions. Therefore, the presence of dust implies the existence of larger bodies that can collide and produce dusty debris \\citep{bac93,lag00,zuc01}. The largest of these bodies could be meter-sized up to planet-sized, and we may well refer to them generally as planetesimals. {\\em Therefore, any system with excess thermal emission implies planet formation at least to the extent of forming planetesimals.} The ability to measure thermal emission in the mid-infrared is therefore a powerful technique in identifying systems in which planetary system formation has occurred or is occurring. However, since debris disks are cool, optically and geometrically thin, and gas-poor, they are generally harder to detect than the primordial, optically-thick accretion disks found around very young stars ($\\lesssim$\\,10\\,Myr).  The {\\em Spitzer} Space Telescope's unprecedented sensitivity in the mid-infrared allows for the first time a statistical study of debris disks and their evolution across a wide spectral range. Excesses detected at 24\\,$\\micron$ generally imply temperatures on the order of 100\\,K. This equilibrium temperature is achieved in the vicinity of 1-5\\,AU for spectral types FGK and 5-30\\,AU for the more luminous B and A stars. By probing these distances in the mid-infrared, we are therefore studying potential planet-forming and planet-bearing regions around other stars. Building on earlier work from the {\\em Infrared Astronomy Satellite} (IRAS) and {\\em Infrared Space Observatory} (ISO) which showed that the amount of dust in debris disks steadily declines over time \\citep[e.g.][]{spa01,hab01}, \\cite{rie05} showed that more than half of A-type stars younger than $\\sim$\\,30\\,Myr have mid-infrared excess. This result implies that planetary system formation occurs frequently around stars a few times more massive than the Sun. However, the same result also shows that up to $\\sim$\\,50$\\%$ of the youngest stars have small or nonexistent excesses in the mid-infrared, pointing to a possible range of planetesimal formation and clearing timescales.  Can we expect similar behavior for lower-mass, longer-living, solar-like stars? Both {\\em IRAS} and {\\em ISO} were in general not sufficiently sensitive to detect the photospheric emission from lower-mass stars. Only with {\\em Spitzer} have mid-infrared surveys of lower-mass stars begun \\citep{gor04,you04,mey04,sta05,kim05,bei5b,che05,bry06,sil06,bei06,gor06,che06}. These studies conducted at 24 and/or 70\\,$\\micron$ have shown that debris disks exist around solar-like stars at a wide range of possible distances ($\\sim$1-50\\,AU) and temperatures ($\\sim$10-650\\,K) with an age-dependent frequency. It is one of the goals of this investigation to constrain this age-dependence. Continued surveys of stars with known ages at mid-infrared wavelengths will bring us nearer to understanding how debris disks evolve and ultimately will provide constraints on planet formation time scales.  In this investigation we use the 24\\,$\\micron$ channel on {\\em Spitzer} to study the incidence of debris disks in the open cluster IC\\,2391. IC\\,2391 is estimated to be 50$\\pm$5\\,Myr old \\citep{byn04}, an age consistent with both theoretical \\citep{cha01} and observational \\citep{kle02} timescales of terrestrial planet formation. It is believed to be of intermediate size with $\\sim$100-200 members. At a distance of 154\\,pc \\citep{for01}, the cluster is amongst the closest and best studied. Furthermore, its proximity allows for the detection of photospheric emission at mid-infrared wavelengths from low-mass stars. With little observed 24\\,$\\micron$ cirrus structure and visible extinction \\citep[{\\it E(B-V)}=0.006\\,$\\pm$\\,0.005;][]{pat96}, IC\\,2391 offers an attractive combination of age, distance, and background in which to study debris disks.  The aim of this investigation is to measure the incidence of debris disks found around $\\sim$\\,50\\,Myr stars across a broad range of spectral types. We discuss the ensemble properties of excesses in IC\\,2391 by placing our data in context with other relevant samples. In the process we begin characterizing the evolution of debris disks around FGK stars, and compare this result to that previously established for more massive A stars. ",
        "conclusions": "\\label{bozomath}  We have conducted a photometric survey for dusty debris disks in the $\\sim$\\,50\\,Myr open cluster IC\\,2391 with the MIPS 24\\,$\\micron$ channel on {\\it Spitzer}. This wavelength probes regions $\\sim$\\,5 to 30\\,AU around A-type stars and regions $\\sim$\\,1 to 5\\,AU around FGK stars. Due to the cluster's proximity, fluxes of stars as late as spectral type $\\sim$\\,K4 can be measured down to the photospheric level. Of the 34 cluster members detected, only 10$_{-3}^{+17}\\%$ (1/10) of the A-type stars had 24\\,$\\micron$ flux densities $\\geq$\\,15$\\%$ that of the photosphere. This is lower than the 31$_{-9}^{+13}\\%$ (5/16) frequency measured for FGK stars in the cluster as well as marginally lower than A-type stars located in other young clusters. However, it is possible that this difference simply reflects random statistical variations.  In comparison, 31$_{-9}^{+13}\\%$ of the FGK stars in IC\\,2391 have excesses. From their behavior, we find the following: \\begin{enumerate} \\item A high level of planetesimal activity (collisions) is still occurring in terrestrial planet regions ($\\sim$\\,1-5\\,AU) at $\\sim$\\,50\\,Myr. \\item The fraction of FGK stars with 24\\,$\\micron$ excesses decreases significantly on timescales of $\\sim$100\\,Myr. This decay over time corresponds to the observed decline of the frequency of collisions within the inner parts of these systems analogous to the asteroid belt of our own Solar System. \\item The decay and variation of 24\\,$\\micron$ excess ratios around FGK stars is very similar to that measured around A-type stars. Despite an overall decaying excess ratio evolution, there are large fractions of young stars with no excess at the youngest ages and rare large excesses at older ages indicative of episodic and stochastic events. \\item Despite differences in luminosity and in the annuli probed at 24\\,$\\micron$ between A-type and FGK stars, debris disk evolution does not appear to be strongly influenced by stellar mass (for this range of spectral type).   \\end{enumerate}"
    },
    "0609/astro-ph0609188_arXiv.txt": {
        "abstract": "We address the effect of an extended local foreground on the low--$\\ell$ anomalies found in the CMB. Recent X--ray catalogues point us to the existence of very massive superstructures at the 100~$h^{-1}$Mpc scale that contribute significantly to the dipole velocity profile. Being highly non--linear, these structures provide us a natural candidate to leave an imprint on the CMB sky via a local Rees--Sciama effect. We show that the Rees--Sciama effect of local foregrounds can induce CMB anisotropy of $\\Delta T/T \\sim 10^{-5}$ and we analyse its impact on multipole power as well as the induced phase pattern on largest angular scales. ",
        "introduction": "At largest angular scales which correspond to small multipole moments $\\ell$ there exist puzzling features in the Cosmic Microwave Background (CMB). The near vanishing of the two--point angular correlation function in all wavebands for angular scales between $60^\\circ$ and $170^\\circ$ is one of the longest known anomalies, already detected in the data of the Cosmic Background Explorer's Differential Radiometer (COBE--DMR). It has been confirmed and persists in the three--year Wilkinson Microwave Anisotropy Probe data [WMAP(3yr)] \\cite{wmap3,lambda}. Among the two--point angular correlation functions it has been shown that none of the almost vanishing cut--sky wavebands matches the full sky and again neither one of these is in accordance with the best fit $\\Lambda$ cold dark matter ($\\Lambda$CDM) model \\cite{Copi2}. The disagreement turned out to be even more distinctive in the WMAP(3yr) data than in WMAP(1yr) and is unexpected at $99.97\\%$ C.L. for the updated Internal Linear Combination map [ILC(3yr)]\\cite{Copi2}.  Besides the lack of power, there are a number of remarkable anomalies regarding the phase relationships of the quadrupole and octopole within the WMAP data\\cite{schwarz,eriksen,tegmark}. In order to be able to make distinct statements with respect to a phase analysis of multipoles we make use of the \\emph{multipole vectors formalism}\\cite{mvec_form}. Looking at quadrupole plus octopole vectors from WMAP(3yr) the alignment with the equinox (EQX) and with the ecliptic is found to be unlikely at $99.8\\%$ C.L. and $96\\%$ C.L. respectively \\cite{Copi2}. The correlation with the dipole direction and with the galactic plane is found to be odd at $99.7\\%$ C.L. and $99\\%$ C.L. respectively. Moreover from the combined full sky map of $\\ell = 2+3$ one infers that the octopole is quite planar and that the ecliptic strongly follows a zero line of the map, leaving the two strongest extrema in the southern hemisphere and the two weakest in the northern hemisphere. Some of these effects are statistically dependent, e.g. given the observed quadrupole--octopole alignment, the significance of alignment with the galactic plane is reduced to unremarkable $88\\%$ C.L.  These findings support the conclusion that either the Universe as seen by WMAP is not statistically isotropic on largest scales, or that the observed features are due to unexpected foregrounds, hidden systematics or new physics challenging the standard cosmological model. Diverse attempts for explanation can be found in the literature: considering anisotropic or inhomogeneous models [Bianchi family, Lema\\^itre--Tolman--Bondi (LTB) models]\\cite{bianchi,jaffe,alnes,moffat,rakic,tomita}, Solar system foreground\\cite{frisch}, lensing of the CMB\\cite{vale} and moving foregrounds\\cite{cooray}, Sunyaev--Zel'dovich (SZ) effect\\cite{abramo,hansen} and Rees-Sciama (RS) effect\\cite{silk,rakic}, considering a non--trivial topology of the Universe\\cite{luminet,cornish}, considering modifications and refinements of the standard simplest scenario of inflation\\cite{boyanovsky,campanelli,contaldi1,ferrer,gordon,contaldi2,wolunglee}, considering possible phenomenology of loop quantum gravity\\cite{hofmann,tsujikawa}. In this talk we update and expand our previous work \\cite{rakic} in the light of the WMAP(3yr) data release. ",
        "conclusions": "Recent astrophysical data cataloguing our neighborhood in the X--ray band point us to the existence of massive non--linear structures like the SSC at distances $\\sim$~100~$h^{-1}$Mpc. Besides its significant contribution to the dipole velocity profile such a structure is able to induce anisotropies $\\sim 10^{-5}$ via its RS effect. Regarding CMB modes, the spherical symmetry (LTB) which we use to approximate the local superstructure reduces to an axial symmetry along the line connecting our position and the centre of the superstructure where we locate the SSC. We produced statistically isotropic and gaussian MC maps of the CMB and computed their $S$--statistics (\\ref{eq_snx}) for alignment with generic astrophysical directions like the NEP, EQX and NGP. The additional zonal harmonics have been added with increasing strength (see ref.\\cite{web} for full--sky maps). When gauging the preferred axis to the direction of local motion (WMAP dipole) the consistency of the data with theory becomes even worse, albeit with less significance w.r.t. WMAP(3yr). On the other hand an orthogonally directed (Solar system) effect would be more consistent with the three--year data.  \\begin{figure*}[b!] \\epsfig{figure=figSnNEP_wmap_col.eps,width=1.15in,angle=270} \\epsfig{figure=figSnEQN_wmap_col.eps,width=1.15in,angle=270} \\epsfig{figure=figSnNGP_wmap_col.eps,width=1.15in,angle=270}\\\\ \\epsfig{figure=figSnNEP_NEP_col.eps,width=1.15in,angle=270} \\epsfig{figure=figSnEQN_NEP_col.eps,width=1.15in,angle=270} \\epsfig{figure=figSnNGP_NEP_col.eps,width=1.15in,angle=270} \\caption{WMAP one- and three--year ILC maps compared to the alignment (\\ref{eq_snx}) of quadrupole and octopole normals with physical directions (NEP, EQX, NGP in columns) for two orthogonal realisations of the preferred direction $\\hat{\\boldsymbol z}$ (WMAP dipole, NEP in rows). The bold histograms represent statistically isotropic and gaussian skies. Increasing the axial contribution makes the anomalies worse for $\\hat{\\boldsymbol z}=\\rm{WMAP \\, dipole}$, but with the exclusions being less significant for the ILC(3yr) than for the ILC(1yr). At the same time a Solar system effect is preferred by the data.} \\label{fig2} \\end{figure*}  Considering extended local foregrounds Abramo et al.\\cite{abramo} recently proposed that a cold spot in the direction of the local Supercluster could account for the cross alignments of quadrupole and octopole. The cold spot would be realised by the SZ effect of CMB photons scattering of the hot intracluster gas. On the other hand Inoue and Silk\\cite{silk} suggest a certain geometrical pattern of two identical voids to account for the cross alignment as well as for the octopole planarity via the RS effect. Each of the latter approaches alone is not fully satisfactory. Nevertheless a combined approach enfolding the RS effect as well as the SZ effect from extended foregrounds seems promising for the future. Moreover, since the local RS effect can contribute up to $10^{-5}$ to the temperature anisotropies on large angular scales, a detailed study is important for cross--correlating CMB data (including upcoming Planck data) with astrophysical observations on the local large--scale structure."
    },
    "0609/astro-ph0609231_arXiv.txt": {
        "abstract": "The profile of the silicate 10~$\\mu$m IR band contains important information about the evolutional stage of dust in circumstellar environments and the possible ongoing process of planetesimal formation. In order to extract this information, the observed band profiles are compared with calculated or laboratory-measured absorption cross sections of amorphous and crystalline grains with different sizes and compositions. We present in this study the first laboratory measurements of the 10~$\\mu$m band profiles of nonembedded, i.e. free-flying, particles of amorphous and crystalline Mg$_2$SiO$_4$ (with two different particle shapes), amorphous and crystalline MgSiO$_3$, and crystalline olivine. We compare the spectra with those measured on embedded grains and discuss the potential of the new experimental method for comparison with observed spectra, as well as for future studies of agglomeration and surface manipulation of the grains. ",
        "introduction": "Dust particles are important players in astrophysical processes such as star and planet formation. Changes in size, composition, and crystallinity of the dust grains indicate dynamical processes (e.g.~in accretion disks) and hold a key to understanding specifically the early formation of terrestrial (Earth-like) planets in the environments of young stars. The most important technique, which makes it possible to obtain information about the dust populations in different objects (and nowadays even at different regions within one object), is infrared spectroscopy. The most important role here is played by the 10~$\\mu$m Si-O stretching band of silicate dust because silicates are the major dust component in most dusty media in space, and this band can be observed from the ground.  In fact, solid particles in the interstellar medium consist mainly of amorphous silicate dust grains \\citep[see the review by][]{whi92}, whereas presence of crystalline silicates has been confirmed via mid-infrared observations in circumstellar disks around the relatively bright young Herbig Ae/Be stars \\citep[e.g.][]{mal98,mee01}, for evolved stars \\citep{mol02}, the T Tauri stars \\citep[TTSs;][]{mee03}, the TTS Hen 3-600 \\citep{hon03}, and the Vega-type stars HD 145263 and $\\beta$-Pic \\citep[e.g.][]{hon04,oka04}. Analysis of the infrared emission bands and especially the 10~$\\mu$m band profile of such silicate grains has allowed one to find indications of grain growth and crystallization that trace the early evolution of young circumstellar disks toward planetary systems \\citep{bo05}. However, the analysis of observational infrared spectra is very complex for an accurate interpretation, since not only mineralogical properties such as crystallinity and chemical composition determine the shape and intensity of the distinctive spectral feature, but also size, shape, and agglomeration of the grains exert an influence. Thus, exhaustive laboratory studies of these influences on the silicate dust bands are essential for an appropriate interpretation of the observed spectra.  Currently, band profiles that are used for comparison with observed spectra are mostly calculated, assuming simple geometrical models such as spherical or ellipsoidal grain shapes. The predictions of these calculations are uncertain because in reality the grain shape might be irregular. So far, laboratory measurements on real particle ensembles were not able to provide exact band profiles for a direct comparison due to the use of an embedding medium (KBr) in the sample preparation, which changes the band profile substantially by the influence of its electromagnetic polarization \\citep{fab01}. In this Letter, we present a new method for laboratory measurements of extinction spectra avoiding the influence of electromagnetic interaction with solid embedding media. Applying the aerosol technique \\citep{hin99}, we obtain for the first time extinction spectra in the mid-infrared region of both amorphous and crystalline silicate grains \\citep[for amorphous SiO$_2$ particles, see][]{tam06}, which are directly comparable to observed spectra. We investigate the differences to KBr-measured band profiles by deriving peak positions from the band profiles measured with both methods. Furthermore, we use electron-microscopic investigation of the grains to associate morphological (size, shape, agglomeration) influences on the measured profile. \\section {Experimental procedures and samples} Regarding the aerosol measurements, we used a White-type long-path infrared cell (MARS-8L/20V, Gemini Scientific Instr.) of about 20~m path length, attached to a Fourier transformation infrared spectrometer (FTIR; Bruker ver. 113), to measure the extinction of the infrared radiation by the silicate dust particles suspended in an aerosol cloud filling the cell (diameter about 15~cm, length 0.6~m). By multiple reflection between the two gold mirrors at both ends of the cell, the IR radiation passed 32 times through the particle cloud, thus enhancing the sensitivity of the detection to optical depths around unity in the 10~$\\mu$m band. An available dust flow generator (Palas RBG 1000), which utilized a rotating brush to disperse a pressed powder from a dust storage hole into a nitrogen gas stream, produced the dense aerosol that was led through stainless steel tubes toward the cell. In order to concentrate small-sized particles in the aerosol, we placed a two-stage impactor between the dust flow generator and the cell. The first stage retained clumps of particles above a certain size limit whereas the second stage densified the aerosol up to a concentration of 10$^6$ particles per cubic centimeter. A fraction of the aerosol was sampled through polyester membrane filters for scanning electron microscope (SEM) inspection of the particle morphology. After filling the aerosol into the cell, we stopped the aerosol flow from the generator and trapped the suspended particles in the cell during the spectroscopic analysis. Particles of 1~$\\mu$m size and smaller are sufficiently coupled to the gas at atmospheric pressure to form a stable cloud on the timescale of hours. Larger particles and clumps of particles settled to the bottom of the cell. A small gas flow (50 liters hr$_{-1}$) from both ends of the cell toward the outlet valve was maintained during the measurement in order to protect the mirrors. This caused a continuous loss of aerosol and therefore a reduction of the particle column density during the experiment. Measurements shown here were usually taken after 10 minutes in order to allow for some sedimentation of remaining clumps. Apart from the decreasing absorption strength, there were no changes of the band profiles observed after this time.  In order to make a comparative study, we have applied the classical KBr pellet technique where a granular sample (e.g. silicate particles) was mixed with potassium bromide powder (KBr), which has high transmission throughout the mid-infrared range, in a mass ratio of 1:500.  After very careful mixing in an agate mortar for segregating the particles as much as possible, 0.2~g of each mixture were pressed with 10~ton of force load to make about 0.5~mm thick clear pellets with a diameter of 1.3~mm, the transmission of which was measured in the FTIR spectrometer.  We investigated six powdered silicate samples of five different sorts with respect to composition and crystal structure. The crystalline Mg$_2$SiO$_4$ (forsterite) samples were commercial products whereas for the amorphous silicates we made use of original samples that were produced in our laboratory by the sol-gel technique \\citep{jag03} and glass melting \\citep{dor95}. Table~1 gives a list of the samples we used, together with their properties. We subjected the samples to a sedimentation process in a solvent (acetone), which allowed us to obtain a size fraction $leq$ 1.0~$\\mu$m in diameter except CF2 and AF samples, which were originally small enough for the experiments. \\begin{table*} \\caption{Properties of the silicate samples.} \\label{table:1} \\centering \\begin{tabular}{c c c c c c} \\hline Sample & Chemical Formula & Product Information & Size ($\\mu$m) & Shape \\\\ \\hline CF1 & crystalline Mg$_2$SiO$_4$ &  Alfa Aesar Johnson & d$<$1.0       & irregular \\\\ CF2 & crystalline Mg$_2$SiO$_4$ &  Marusu Yuuyaku     & d$\\approx$0.2 & ellipsoid\\\\ AF & amorphous Mg$_{2.3}$SiO$_4$ &  sol-gel process    & d$<$1.0        & irregular \\\\ CE & crystalline MgSiO$_3$     &  melting     & d$<$1.0        & irregular \\\\ AE & amorphous MgSiO$_3$     &  melting/quenching     & d$<$1.0        & irregular \\\\ CO & crystalline Mg$_{1.9}$Fe$_{0.1}$SiO$_4$ &  natural (San Carlos)   & d$<$1.0       & irregular \\\\ \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "We have applied the aerosol technique to IR spectroscopic measurements of the 10~$\\mu$m band profile of crystalline and amorphous silicates and demonstrated its application to obtaining realistic band profiles that can be compared directly to observed spectra for the first time.  Compared to the traditional KBr technique, the extinction measurement in aerosol has two disadvantages: it requires about 100~mg of material and it is, unfortunately, not quantitative. The column density of the dust grains along the spectrometer beam can hardly be estimated or measured to an accuracy higher than a factor of 2. A better accuracy could be achieved only by accompanying quantitative KBr measurements or by a determination of the material's optical constants, for example, by reflection measurements. In such a case and given that a theoretical model can be adjusted to represent correctly the morphological properties of the particulate, the aerosol spectra can be renormalized.  Conversely, there are distinct advantages of the aerosol method. First, it lacks the problem of the embedding medium, which is inherent to the KBr method. A second major advantage is that the structures of the grains can be analyzed by either in situ or ex situ microscopic imaging. This allows investigation in detail of the influence of the grain morphology on the spectra and will hopefully help to calibrate theoretical approaches.  Concerning practical aspects, the preparation of the aerosol is easier than that of a KBr pellet with an ideal homogeneous grain distribution. Large clumps of particles can be receded or settled down within a short time. Even separation of different grain sizes is uncomplicated with the impactor technique. An important point is that collision timescales in the aerosol can probably be made smaller than the duration of an experiment, so that fractal aggregate formation could be directly studied by IR spectroscopy. So far, the particle densities were too small in our experiment to observe grain growth by aggregation. With the KBr technique, this is hardly possible, because agglomerates are modified or destroyed by the preparation process. Finally, with the aerosol technique, it should be possible to manipulate the dust in situ by surface condensation or irradiation, so that its influences on the IR spectra can be measured well."
    },
    "0609/gr-qc0609089_arXiv.txt": {
        "abstract": "We study the deformation of a long cosmic string by a nearby rotating black hole. We examine whether the deformation of a cosmic string, induced by the gravitational field of a Kerr black hole, may lead to the formation of a loop of cosmic string.  The segment of the string which enters the ergosphere of a rotating black hole gets deformed and, if it is sufficiently twisted, it can self-intersect chopping off a loop of cosmic string. We find that the formation of a loop, via this mechanism, is a rare event. It will only arise in a small region of the collision phase space, which depends on the string velocity, the impact parameter and the black hole angular momentum. We conclude that generically, the cosmic string is simply scattered or captured by the rotating black hole. ",
        "introduction": "Cosmic strings are one-dimensional topological defects~\\cite{td1,td2,td3} which could have been formed at the end of phase transitions associated with spontaneously broken symmetries in the early universe, via the Kibble mechanism~\\cite{kibble}. Even though cosmic strings do not play a dominant r\\^ole in structure formation, these objects are expected to be generically formed in the framework of supersymmetric grand unified theories~\\cite{jrs}.  Cosmic microwave background measurements allow a small, but non negligible, contribution of strings~\\cite{cmbstr}, which can be realized in supersymmetric inflationary models~\\cite{johnm}. Recently, cosmic strings gained a lot of interest since it was pointed out that fundamental strings, D1-branes and D$(p-2)$-branes wrapping around $(p-3)$-cycles in the compactified dimensions, could also play a cosmic string-like r\\^ole in cosmology~\\cite{dstr}.  Cosmic string interactions can lead to a network of string loops. Intercommutations of two long strings in two points, or self-intercommutations of a long string in one point, may chop off a loop of cosmic string. In what follows, we explore another mechanism which may lead to the formation of cosmic string loops.  More precisely, we examine the possibility that the scattering of a long cosmic string by a rotating black hole can cause a deformation of the long string, which may lead to the formation of a loop of cosmic string. As a segment of a long cosmic string enters the ergosphere of a rotating black hole, with rotation axis not parallel to the axis of the cosmic string, this segment may be deformed and, in some cases, it may get twisted. If the long string self-intersects, as a result of the deformation, a loop of cosmic string may be formed.  The formation of loops of cosmic strings is important for the astrophysical implications of a string network. In particular, cosmic strings can be detected through the gravitational wave background produced by oscillating string loops~\\cite{dv}~\\footnote{We note that long cosmic strings also emit gravitational waves, since these are wiggly due to string intercommutations~\\cite{s}.}.  The scattering of a long cosmic string by a rotating black hole has been previously studied in the case of a black hole whose axis of rotation is parallel to the axis of the long string, concluding that the long string will be either scattered, or captured by the black hole. Which of the two regimes will take place depends on the impact parameter. The critical value of the impact parameter which separates the two regimes is given in Ref.~\\cite{snajdr2}. The effect of the scattering of a long cosmic string by a rotating black hole was studied in Ref.~\\cite{snajdr1}, where it was shown that the string is displaced in the direction perpendicular to its velocity by an amount which depends on the impact parameter~\\cite{snajdr1}. In this article we generalize the study of Ref.~\\cite{snajdr1} for different configurations of a long cosmic string approaching a rotating black hole.  We conclude that the formation of a loop from the scattering of a long string by a rotating black hole is a rare event. It will only arise in a small region of the collision phase space, which depends on the string velocity, the impact parameter and the black hole angular momentum. More frequently, the cosmic string is either simply scattered or captured by the rotating black hole.  In what follows we assume that the string thickness is negligible as compared to its length. Moreover, the string linear mass density is very small.  Then, assuming that the string size is much larger than the Schwarzschild radius of the black hole, while the mass of the string is much smaller than that of the black hole, we can treat the cosmic string as a test object, and we can neglect its gravitational back reaction.  The string is initially sufficiently far from the black hole, so it can be taken as a straight string. As it approaches the black hole, it will move in a Kerr geometry, however since the string is taken to be sufficiently long, the end regions of the infinitely long string will not be influenced by the presence of the gravitational field of the Kerr black hole.  Let us consider the axis of rotation of the Kerr black hole to be along the direction ${\\cal X}$, in a system of coordinates $({\\cal X, Y, Z})$. A cosmic string being initially far away from the black hole is moving towards it.  One can investigate the following setups: \\begin{itemize} \\item The string is at first parallel to the axis of rotation of the Kerr black hole, and it moves towards the black hole with a velocity perpendicular to the axis of rotation. Thus, the cosmic string is originally along the ${\\cal X}$-axis and its velocity is along the ${\\cal Z}$-axis. [We call this case {\\bf (i)}.] \\item The  string is at first perpendicular to the axis of rotation of the Kerr black hole and it moves towards the black hole with velocity parallel to axis of rotation. Thus, the cosmic string is originally along the ${\\cal Z}$-axis and its velocity is along the ${\\cal X}$-axis. [We call this case {\\bf (ii)}.] \\item The string is at first perpendicular to the axis of rotation of the Kerr black hole and it moves towards the black hole with velocity which is also perpendicular to the axis of rotation. Thus, the cosmic string is originally along the ${\\cal Z}$-axis and it moves towards the Kerr black hole with velocity along the ${\\cal Y}$-axis. [We call this case {\\bf (iii)}.] \\end{itemize} Case {\\bf (i)} has been studied in Ref.~\\cite{snajdr1}; in what follows we will focus on cases {\\bf (ii)} and {\\bf (iii)}. Case {\\bf (iii)} is the most promising configuration to lead to the formation of a cosmic string loop.  We organize the rest of the paper as follows: In Section II we briefly discuss the string equations of motion in a curved background. We give the string equations of motion for a string in a general background geometry and then we derive them for a straight string far away from a black hole. We proceed with first the Newtonian and then the Lense-Thirring scattering of a straight string in the linearized approximation. In Section III we study the scattering of a long straight string by a Kerr black hole. Our numerical approach is presented in Section IV. We round up with our conclusions in Section V.  We will use units where $G=c=1$. ",
        "conclusions": "We have studied the deformation of a long cosmic string entering the ergosphere of a Kerr black hole. The aim of this study is to investigate how often such a deformation can be accompanied by a sufficient twist, such that there is a string self-intersection leading to a loop formation.  Among the configurations we have studied, the case of a long cosmic string being at first perpendicular to the axis of rotation of the Kerr black hole and moving towards it with velocity perpendicular to the rotation axis, is the most prominent to lead to loop formation. [It is our case {\\bf (iii)}.]  A numerical investigation has shown us that loop formation is a rather rare event, while in general the cosmic string will be just scattered from the black hole or captured by it."
    },
    "0609/astro-ph0609007_arXiv.txt": {
        "abstract": "Stability  properties  of ``magnetic  tower''  jets  propagating in  the gravitationally stratified  background have been  examined by performing three-dimensional       magnetohydrodynamic       simulations.       The current-carrying,  Poynting  flux-dominated  magnetic tower  jet,  which possesses  a highly  wound helical  magnetic  field, is  subject to  the current-driven instability  (CDI). We find that,  under general physical conditions  including  small  perturbations  in the  initial  background profiles,   the   propagating    magnetic   tower   jets   develop   the non-axisymmetric, $m=1$  kink mode of the  CDI.  The kink  mode grows on the local Alfv\\'en  crossing time scale. In addition,  two types of kink modes  appear in  the  system.   At the  central  region where  external thermal pressure confinement  is strong, only the internal  kink mode is excited and  will grow.  A large  distance away from  the central region where the external thermal pressure  becomes low, the external kink mode is observed.  As  a result, the exterior of magnetic  tower jets will be deformed  into  a  large-scale   wiggled  structure.   We  also  discuss extensively  the different  physical  processes that  contribute to  the overall stability  properties of the magnetic  tower jets. Specifically, when the jet propagates in  an initially unperturbed background, we find that they  can survive the kink  mode beyond the point  predicted by the well-known Kruskal-Shafranov (K-S) criterion.  The stabilization in this case  comes mainly  from  the dynamical  relaxation  of magnetic  twists during the propagation of magnetic towers; the magnetic pitch is reduced and  the corresponding  K-S critical  wavelength becomes  longer  as the tower jet  proceeds.  Furthermore, we show that  the pressure-driven and Kelvin-Helmholtz instabilities do not  occur in the magnetic tower jets. This  strongly suggests  that  the CDI  is  the primary  reason for  the wiggling structures in jets. ",
        "introduction": "\\label{sec:INT} Magnetohydrodynamic  (MHD) mechanisms are  often invoked to  explain the  launching, acceleration  and collimation  of  jets from Young Stellar  Objects, X-ray  binaries, Active Galactic  Nuclei (AGNs), Microquasars,  and  Quasars   \\citep[see,  {\\it  e.g.},][and  references therein]{DLM01}.   Strongly magnetized jets,  particularly those  with a strong toroidal field encircling the collimated flow, are often referred to as ``current-carrying'' or ``Poynting flux-dominated'' (PFD) jets.  A large  current flowing  parallel  to  the jet  flow  is responsible  for generating a  strong, tightly wound helical magnetic  field.  The global picture of a  current-carrying jet with a closed  current system linking the magnetosphere of the central engine and the hot spots was introduced by \\citet[]{B78,  B06} and  applied to AGN  double radio  sources.  This closed  current  system  includes  a  pair  of  current  circuits,  each containing both  a forward electric  current path (the jet  flow itself, with  its  toroidal magnetic  field,  toward  the  lobe), and  a  return electric current path (along some path back to the AGN core).   Recent high resolution observations,  in X-ray, optical and radio bands, show that many  AGN and quasar jets display  wiggles/kinks, which may be associated with  a helical motion of  the underlying bulk  flow, in wide spatial ranges from sub-parsec  to kpc scales \\citep[{\\it e.g.},][]{R89, H92, B97, F99, LZ01, M01, EO02, S06}.  Many ideas have been proposed for these   helical    distortions:   MHD   instabilities    such   as   the Kelvin-Helmholtz  instability   (KHI),  the  current-driven  instability (CDI), and  the pressure-driven instability (PDI), or  the precession of the  jet ejection  axis due  to  the existence  of a  binary black  hole \\citep[]{B80}, or an encounter with another galactic core.   It  is  well known  that  a cylindrical  plasma  column  with a  helical magnetic configuration  is subject to  the $m=0$ (sausage) mode  of PDI, the $m=1$  (kink) mode of CDI,  and the other higher  order modes, where $m$ is  the azimuthal ($\\phi$) mode number  \\citep[see, {\\it e.g.},][for details]{K66,  BA78, F82}.  CDI  is driven  by parallel  (axial) current flows and can  exist even in a zero pressure,  force-free plasma. On the other  hand, PDI is  powered by  perpendicular current  flows and  it is often known as  the interchange (internal) modes.  KHI  is considered to be very important  at the shearing boundary between  the flowing jet and the external  medium, particularly  in the kinetic  flux-dominated (KFD) jets.   In principle,  CDI  and  KHI have  both  external (surface)  and internal (body) modes.  PFD jets should be especially susceptible to CDI rather  than PDI and  KHI.  This  is because  the strong  axial electric current is  present inside  the PFD jets,  which is responsible  for the highly wound  helical field configurations, and the  Lorentz force plays an  important  role in  the  jet  dynamics.   The investigation  of  the destructive influence of CDI on PFD  jet flows is an important avenue of research in the study of astrophysical jets.   Much  more  attention  on  studies  of stabilities  (disruption  of  the interior/exterior of astrophysical jets) have  given to KHI than to CDI. This is because jets were  believed to be super-Alfv\\'enic or super-fast magnetosonic  KFD flows  ({\\it  i.e.}, kinetic  energy  was expected  to greatly  exceed the magnetic  energy).  However,  it is  not necessarily always the case; a  super-Alfv\\'enic or super-fast magnetosonic PFD flow can exist  in both nonrelativistic \\citep[]{KS97,  V00} and relativistic \\citep[]{VK01} MHD outflow solutions.  Analytic  studies on the CDI as a possible  explanation for  jet disruption  have been  performed  as well \\citep[]{E93,  S97,  B98,  L99}.   Linear  analyses  of  nonrelativistic force-free  jets,   which  are   thermally  confined  by   the  external nonmagnetized medium,  were performed  in super-fast \\citep  []{AC92} or trans-fast  \\citep[]{A96}  magnetosonic regimes.   They  found that  the growth rate of CDI is always  substantially smaller that those of KHI in their  specific  treatment.   Only  for quite  small  fast  magnetosonic numbers,  $M_{\\rm  F} \\lesssim  0.5$,  both  become comparable.   Linear analyses   of   the   relativistic   force-free   jets,   in   case   of constant-longitudinal field  $B_{z}$ along the jet  axis, were performed by \\citet{IP94, IP96}.  They concluded that such a jet is stable against CDI, while \\citet[]{L99} analyzed  the relativistic force-free jets with $B_{z}$ decreasing outwards and found that they are unstable.   Nonlinear development  of the CDI for cold  super-fast magnetosonic jets has  been   studied  by  \\citet{L00},  based  on   the  linear  analysis \\citep[]{A00}.  It  was found that the current  density is redistributed within the inner  part of the jet radius on  a characteristic time scale of the order of the Alfv\\'en crossing time in the jet frame.  Nothing in their numerical  results indicated a  possible disruption of the  jet by the CDI  sausage ($m=0$) or kink ($m=1$)  modes \\citep[see also][]{LF00, F00}.  In  general, these linear/nonlinear  considerations conclude that ``current-driven  instabilities are  unlikely  to be  dangerous for  the integrity of astrophysical jets'' \\citep[]{A96} and this has been widely accepted  in theoretical/observational  communities.   In the  nonlinear regime,  an interplay between  KHI and  CDI can  occur in  the helically magnetized flow even with slightly super-fast magnetosonic speed $M_{\\rm F} \\gtrsim 1$ \\citep[]{BK02}.  This nonlinear interaction can contribute to jet  survival, and  the large-scale magnetic  deformations associated with  the CDI  mode  development can  effectively  saturate KHI  surface vortices and prevent jet disruption.   By contrast, the first attempt  to investigate the nonlinear behavior of visible distortion in  KFD jets along the large-scale  magnetic field by the CDI  kink mode was performed  by \\citet[]{T93} to  model Herbig Haro objects.  In  a similar study,  \\citet[]{N01} examined the  formation of wiggled structures by the $m=1$ mode  in PFD jets to apply for AGN jets. Their numerical result is  applicable for the systematic distribution of the observed  rotation measure in 3C 449  \\citep[]{F99}, associated with the  jet deformation \\citep[]{K04}.   The rotating  PFD/KFD jets  can be potentially  stabilized  against  the  $m=1$ mode  beyond  the  original Kruskal-Shafranov criterion, however, sudden destabilization of rotating jets  due to  the $m=1$  mode can  occur via  the angular  momentum loss triggered by MHD shock waves \\citep[]{NM04}.   The purpose of the present  paper is to discuss the stability properties of  ``magnetic tower'' jets,  which have  not been  studied in  the past work.  Of  particular interest here is  the nonlinear growth  of the CDI kink  mode.   The   CDI  kink  mode  in  magnetic   tower  jets  may  be distinguished from that in the  classical MHD jets along the large-scale magnetic field  \\citep[]{T93, N01, NM04}, while it  should be comparable directly to ``non-disrupting'' models  \\citep[]{A96, A00, L00}, in which the  MHD  jets  are  assumed  to  be in  pressure  equilibrium  with  an unmagnetized ambient medium.   This  is the  third of  our series  of papers  to examine  the nonlinear magnetic tower jets.  The  first of our series, \\citet[][hereafter Paper I]{L06},  described the  basic  assumptions and  the  approaches in  the numerical modeling of  magnetic tower jets.  The evolution  of the tower jets in a constant  density/pressure background has been examined there. The second of our series, \\citet[][hereafter Paper II]{N06} investigated the  global structure  of  magnetic tower  jets  in the  gravitationally stratified  atmosphere (a  more realistic  astrophysical  situation), in terms of the MHD waves  structure, the radial force equilibrium, and the collimation.   In  the  present  paper,  the third  of  our  series,  we investigate the stability properties of magnetic tower jets.  We examine our numerical  results in two  cases; one without  initial perturbations and  another  with  initial  perturbations.   The  former  exhibits  the quasi-axisymmetric  evolution without any  growth of  MHD instabilities, while  the latter  shows non-axisymmetric  evolution by  the  CDI modes. This paper is  organized as follows.  In \\S 2,  we outline our numerical methods.  In  \\S 3, we  describe our numerical results.   Discussion and conclusion are given in \\S 4 and \\S 5. ",
        "conclusions": "Stability  properties  of magnetic  tower  jets  have  been examined  by performing  3-D MHD  simulations. Our  numerical results  show  that the magnetic tower  jets in the gravitationally  stratified atmosphere could survive the current-driven kink instability beyond the Kruskal-Shafranov criterion  even  in  the  nonlinear  regime.   One  of  the  potentially stabilizing effects  on a traditionally  unstable magnetic field  may be due to  the dynamical relaxation  of magnetic twists in  the propagating magnetic tower jets.  This causes an increased threshold of the unstable critical wavelength.   However,  the  magnetic  tower  jets  are eventually  distorted  by  the non-axisymmetric  perturbations with a  few percent  of the  local sound speed in  the system.  The current-driven kink  modes grow predominantly on  time  scales  of the  order  of  the  local Alfv\\'en  crossing  time $\\tau_{\\rm A}$.  In the gravitationally stratified atmosphere, two types of  kink modes appears:  the internal  and external  modes.  At  a large distance away  from the  central region, the  external kink  mode grows, while  only the  internal kink  mode exhibits  near the  central region. Therefore, the exterior  of magnetic tower jets will  be deformed into a large-scale wiggled structure by the external kink mode.   Possible mechanisms  of stabilizing/destabilizing in  the magnetic tower jets    under   the    gravitationally   stratified    atmospheres   are discussed. None  of the  growing surface/body modes  of Kelvin-Helmholtz instability are identified due to the small velocity shear and slow bulk speed  (compared with  the  local Alfv\\'en  speed). The  pressure-driven instabilities  are also  inhibited  due  to the  low  plasma $\\beta$  ($ \\lesssim 0.1$) inside the  magnetic towers. The ``reversed field pinch'' field configuration in the towers  and the external unmagnetized gas may provide some stabilization effects on the current-driven external modes.   Non-axisymmetric current-driven  instabilities ($m \\ge  1$) are absolute instabilities, i.e., they grow but do not propagate in the jet co-moving (rest)  frame   \\citep[]{A00,  NM04}.   If   the  dissipation  timescale $\\tau_{\\rm d}$ of  the current associated with jets  is much longer than the  jet  dynamical  timescale  $\\tau_{j}  \\lesssim  \\tau_{\\rm  A}  \\sim \\tau_{\\rm CDI} \\ll \\tau_{\\rm d}$,  the patterns created by the CDI modes will persist  for some  time as  the flow advances.   And the  bulk flow itself will appear to travel in a true 3-D helical pattern as it follows the  magnetic backbone  of  the helix,  as  observationally expected  in several radio sources as 3C 345 \\citep[]{Z95}, 3C 120 \\citep[]{G01}, and 3C 449 \\citep[]{F99}."
    },
    "0609/hep-th0609095_arXiv.txt": {
        "abstract": "The Hamiltonian structure of general relativity provides a natural canonical measure on the space of all classical universes, {\\it i.e.}, the multiverse. We review this construction and show how one can visualize the measure in terms of a ``magnetic flux\" of solutions through phase space. Previous studies identified a divergence in the measure, which we observe to be due to the dilatation invariance of flat FRW universes. We show that the divergence is removed if we identify universes which are so flat they cannot be observationally distinguished. The resulting measure is independent of time and of the choice of coordinates on the space of fields. We further show that, for some quantities of interest, the measure is very insensitive to the details of how the identification is made. One such quantity is the probability of inflation in  simple scalar field models. We find that, according to our implementation of the canonical measure, the probability for $N$ e-folds of inflation in single-field, slow-roll models is suppressed by of order ${\\rm exp} (-3N)$ and we discuss the implications of this result. ",
        "introduction": "The problem of comparing different possible histories of the universe, and assigning a probability to each, is central to theoretical cosmology.  We cannot expect a fundamental theory to predict precisely what we see today: at best, it should predict an ensemble or ``multiverse\" of possible universes, with the universe we observe being a typical member.  Of course there are many attitudes that one may take to probability not only  in the context of cosmology, but in science more generally. The stance adopted in this paper  is that in practising science we often adopt Bayesian methods which make essential use of a priori probabilities (even if one does not accept this, or has doubts about what these probabilities mean, this is certainly a fair characterisation of much of the current observational literature).  In fact, following Laplace's Principle of Indifference~\\cite{Laplace}, we often start with the least informative probability distribution, {\\it i.e.}, a flat distribution, and then sharpen it in the light of additional information. This is,  as we see it, the essence of the Bayesian approach. It requires a well defined a priori measure, {\\it i.e.}, one with finite total measure.  At the risk of falling into pedantry let us spell this out more formally. We define the {\\it a priori probability} $P(U)$ to be the probability that the universe is $U$, according to some fundamental theory. Likewise we define $P(O)$ to be the probability of making an observation $O$ in {\\it any} universe. The joint probability that the universe is $U$ and that we make an observation $O$ is $P(U\\cap O)$. The conditional probability of making an observation $O$ in a universe $U$ is $P(O|U)$ (called the {\\it likelihood}), and the conditional probability that we are in the universe $U$ given that we have made an observation $O$ is $P(U|O)$ (called the {\\it a posteriori probability}). It follows from elementary considerations that \\ben P(U|O) P(O)= P(U \\cap O) = P(O|U)P(U) \\een whence the  Cosmic Bayes's Theorem~\\cite{Bayes} tells us that \\ben P(U|O)={P(O|U)P(U) \\over \\int_M P(O|U)P(U) dU } \\label{aa1} \\een where the integral is over what we call the {\\sl multiverse}, $M$, and $dU$ is a measure on the multiverse. Equation (\\ref{aa1}) describes how the a priori probability $P(U)$ is updated in the light of observations.  In this paper, we define the multiverse $M$ to be the set of all possible universes, or, better, the set of all model universes. These are by definition disconnected from one another. {\\sl  This differs from other popular interpretations} in which the multiverse is a connected spacetime containing many causally disconnected regions and, roughly speaking, probabilities are taken to be proportional to the numbers of such regions. This latter concept has been termed the {\\it metauniverse}~\\cite{Moss,Vilenkin}. It is not obvious to us that this second meaning can be made quantitively precise in such a way that the probabilities are well defined. The basic problem is that the natural measure on spacetime, $\\int d^4 x \\sqrt{-g}$, is usually infinite in the scenarios being considered, and the infinity is not removed by restricting attention to so-called reheating surfaces. Much of the current literature on inflation appears to reflect this difficulty.  By contrast to Laplaces's  Principle, one may regard a Proposal for the State of The Universe as some choice of probability distribution function $P(U)$ relative to the a priori measure $dU$. It might possibly constitute an ``explanation\" if derived from some underlying fundamental theory which, again, must be well defined to be meaningful. {\\sl Note that in this paper we are not setting out to make a specific proposal for the universe, i.e.,} we are not advocating a specific formula for $P(U)$, but rather, as we shall argue in detail later, we are adopting Laplace's Principle of Indifference and finding that indeed this is inadequate to favour inflation. A more specific, sharper probability distribution is required to explain inflation. One may quantify the sharpness of a probability distribution by defining (following von Neumann, Shannon and others) the {\\it information} \\ben I=\\int_M P(U) \\ln \\left( P(U)\\right) dU, \\label{aa2} \\een which, after imposing the constraint $\\int P(U) dU =1$, is minimised by the uniform distribution $P(U)=$constant, as one can easily check by computing the first and second variational derivatives and using $P(U)\\geq 0$. The quantity $I$ measures the information (or lack of randomness) carried in a typical realisation of $P(U)$: the least information is obtained if all possible outcomes are equally likely, and the greatest if a single outcome is certain. The {\\it information entropy} $S$ is then defined to be the negative of $I$, which is of course maximised in the uniform distribution. As is well-known, maximising $S$ under various additional constraints including the conservation of energy yields the usual formulae of statistical mechanics.  In this paper, we show that, with the canonical definition of $dU$ given by Hamiltonian dynamics, Laplace's principle does not typically predict inflation. Hence, even if a fundamental theory allows inflation, a sharper $P(U)$is required in order to explain why inflation actually occurred. It remains as a challenge to fundamental theory to explain how such a $P(U)$ might arise. The predictiveness of any particular proposal can be assessed from (\\ref{aa2}); what we would be most happy with is a highly predictive theory, whose predictions were consistent with observation. In fact this idea can be taken further, as we have been  reminded by Don Page. If we regard a proposal $P(U)$ as a {\\it hypothesis}  then we may, given an appropriate measure  on the infinite dimensional set of hypotheses  $\\{P(U)\\}$, adopt Bayesian methods to evaluate their relative probabilities.  We shall not attempt this formidable task in this paper.  The ideas sketched above are not new, and many people have pursued them before in some, largely qualititive, way. In this paper we attempt , following \\cite{GibbonsStewart}, to make them quantitative.  In fact our own viewpoint was influenced in part by trying to make some of Penrose's arguments \\cite{Penrose1,Penrose2,Penrose3} mathematically precise.  There are several additional problems peculiar to cosmology when attempting to construct a statistical theory. First, there there is the problem of general covariance. There is no absolute notion of space or time in general relativity: these are properties of each particular classical solution of the field equations. Second, the solutions generically possess singularities in the past or the future, where the field equations break down in finite time.  Even if one makes the drastic simplification of restricting attention to Friedmann-Robertson-Walker (FRW) universes, {\\it i.e.}, so-called mini-superspace, it is not obvious how to compare different classical spacetimes which are solutions to the same laws of physics. Finally, in theories like string or M-theory, where the low energy effective description involves many additional fields, there may not even be a preferred space-time metric: one can change coordinates on the space of fields (for example via Weyl transformations) in an arbitrary way and physical results should not depend on these choices.  In spite of these difficulties, the problem of constructing sensible measures on the space of solutions is of undeniable importance to the evaluation of various cosmological scenarios. Most of these scenarios focus on some subset of classical solutions, for example arguing that they are ``generic\" because they exhibit dynamical attractor behavior. But to understand how predictive the proposals really are, in the absence of any further information, we need to quantify the extent to which the flat probability measure, without any additional input, succeeds in narrowing down the range of allowed classical solutions.  In recent times, the problem of the measure in cosmology has grown hugely in importance. The claimed vast ``landscape\" of possible compactifications and moduli in string theory seems to offer a truly bewildering range of classical cosmologies\\cite{DouglasKachru}. When combined with recent observations of dark energy, the landscape picture has encouraged many physicists to pursue anthropic explanations. However, the lack of an {\\it a priori} measure on the space of model universes is a serious flaw in such attempts, as even the most ardent proponents will admit. Nevertheless, remarkably little attention has been devoted to constructing measures whereby the different possible classical spacetimes within the ``multiverse\" can be compared \\cite{Coule,Kirchner}. In a similar vein to the anthropic arguments, a \"top down\" approach to cosmology has been proposed by Hawking and Hertog~\\cite{HertogHawking}, in which one imposes constraints on the state of the universe today and attempts to determine its most probable history, for example by constraining the initial state to satisfy Hartle and Hawking's no-boundary proposal~\\cite{HartleHawking}. However, it is hard to judge the success or otherwise of this approach without being able to quantify the extent to which it narrows down the multiverse.   Nearly two decades ago, Gibbons, Hawking and Stewart~\\cite{GibbonsStewart} identified a natural measure on the set of classical cosmological solutions. They argued that any such measure should satisfy the following requirements: (i) it should be positive, (ii) it should depend only on the intrinsic dynamics and neither on any choice of time slicing nor on the choice of dependent variables, and (iii) it should respect all the symmetries of the space of solutions without introducing any additional ad hoc structures (e.g. ``Planck mass cutoffs\") not arising from the field equations themselves. They showed that a measure satisfying all these requirements arises naturally from the Hamiltonian structure of general relativity and that it can be used to count the number of different solutions of the classical field equations in a well-defined way. Note that the set of all Cauchy data is unsuitable for this purpose because different Cauchy data can yield the same classical solution, evaluated on a different time-slice. The proposed measure avoids this overcounting problem by counting each distinct classical solution, in its entirety, only once.   One obvious application of such a measure on the multiverse is in determining the probability of inflation, {\\it i.e.}, how likely is it that the universe started out in an inflationary state, within different models of the laws of fundamental physics. This question is potentially an Achilles' heel for inflation: if inflation is itself highly improbable, then it cannot be claimed to solve the classic cosmological fine tuning puzzles, of large-scale homogeneity, isotropy and flatness. Initially, it seemed that the canonical measure did indeed favor inflation~\\cite{GibbonsStewart}, in simple inflationary models, but more detailed investigations, in particular by Hawking and Page~\\cite{HawkingPage}, identified a serious ambiguity. They found that the canonical measure is infinite, and that both inflationary and non-inflationary solutions typically have infinite measure. Unfortunately, this left the problem of whether inflation is likely, or unlikely, unresolved. Since then, the problem has not received much attention, although various intuitive arguments, generally not satisfying the three above-mentioned conditions, have been given for or against inflation.  In this paper, we shall revisit the canonical measure, showing that when treated with sufficient care, a finite measure, still satisfying the Gibbons-Hawking-Stewart conditions, can be obtained. In agreement with Hawking and Page, we find the canonical measure is infinite. However, we identify the divergence as being due to the dilatation symmetry of flat FRW universes, and we show it is removed if one identifies solutions which cannot be observationally distinguished. Concretely, we impose a limit on the parameter known in the observational literature as $\\Omega_k$, measuring the ratio of the space curvature term in the Friedmann equation to the square of the Hubble parameter. If all universes with $|\\Omega_k|$ smaller than some observational limit $\\Delta \\Omega_k \\ll 1$ are identified, then the canonical measure becomes well-defined. Furthermore, provided the limit on $|\\Omega_k|$ is imposed once the matter fields have entered a phase of evolution in which the expansion of the universe acts adiabatically, {\\it i.e.}, as a slow variation of parameters in the matter Lagrangian, our measure reduces to the canonical measure on the matter fields alone, and is an adiabatic invariant. In this regime the measure becomes independent of the conjugate ``angle\" variable. This statement is independent of the time, the values of the cosmological parameters, or indeed the value of $\\Delta \\Omega_k$ when the cutoff is imposed. These good properties motivate us to reconsider the probability of inflation with various numbers of e-foldings, using the proposed canonical measure. The probability of $N$ e-folds of inflation depends critically on the ``angle\" variable, and we find that the probability of obtaining $N$ e-folds of inflation in simple inflationary models is suppressed by a factor of ${\\rm exp} (-3 N)$. Since $N > 50-60 $ is typically required in realistic models (depending on the efficiency of reheating), it follows that inflationary solutions are, in fact, tremendously rare in the space of classical solutions. We emphasize these statements are independent of the details of the cutoff imposed on the canonical measure.   Our conclusion is, of course, at variance with much of the inflationary literature over the last two decades.  It has some resonance with the well-known arguments (see e.g. Refs.~\\cite{Stewart,Penrose1,Penrose2,Penrose3}) that inflation cannot possibly solve the classical cosmological puzzles for arbitrary initial conditions because {\\it any} physically permissible current state of the universe would, if run back in time, correspond to {\\it some} initial conditions, and also with Holland and Wald's arguments~\\cite{HollandsWald} that any canonical measure would naturally assign the same probability to deflation as to inflation, since Einstein's equations are time-reversal invariant. Because the canonical measure is a measure on the set of universes, independent of their time-orientation, we do indeed find this result: both inflation and deflation are predicted to be exponentially rare phenomena among the set of all classical solutions of the field equations.  The structure of this paper is as follows. In Section 2 we review the construction of the canonical measure, clarifying the geometrical structure in elementary terms and in particular demonstrating that the measure is independent of the choice of initial slice in the space of dynamical variables (contrary to the apparent assertion of Hollands and Wald on this matter~\\cite{HollandsWald}). In Section 3, we discuss the classical dynamics of simple scalar field models and we compute the canonical measure for scalar field matter in FRW spacetimes with arbitrary spatial curvature. At face value, the measure exhibits a divergence at large scale factor, {\\it i.e.}, in the flat space limit, but we argue that since the scale factor is unobservable in this limit, one must factor out dilatations from the result. We show that this results in a sensible measure on the space of classical solutions, {\\it i.e.}, one of finite total measure. In Section 4 we solve the classical field equations, identifying the solutions undergoing $N$ e-folds of inflation and computing the associated measure.  Section 5 compares our results with those obtained in previous discussions, and explains why our conclusions are so different from those reached on the basis of intuitive reasoning from ``chaos at the Planck density\". Section 6 concludes. ",
        "conclusions": "We have developed the canonical measure for homogeneous, isotropic universes and, we hope, clarified its geometrical structure. We applied this measure to simple inflationary models and found it to be divergent, in agreement with Hawking and Page. However, we identified the source of the divergence as being due to the geometrical and physical degeneracy of solutions which are spatially very flat, {\\it i.e.}, for which $|\\Omega_k| \\ll 1$. Our proposal for removing the divergence is to identify universes which cannot be observationally distinguished, {\\it i.e.}, for which $|\\Omega_k|$ is smaller than a critical value. For this purpose, it was convenient to work in coordinates $\\phi, H$ and $a$ in which the limit is easily imposed. The resulting integral over the scale factor $a$ may then be performed, leading to an integral over the matter variables alone.  We have shown that the resulting measure, when evaluated for low Hubble constant $H_S$ on the measure surface, reduces to an adiabatic invariant for the matter fields. In particular, it is independent of the conjugate ``angle\" variable. It is this ``angle\" variable which becomes exponentially focussed during inflation. Hence we find that the probability of inflation, evaluated with the canonical measure, is exponentially small, and this result is very insensitive to the cutoff.  Had we instead chosen to implement our prescription at a very large value of $H_S$, one for which many e-folds of inflation are possible, we would have found, on the contrary, a strong dependence of the results on the cutoff parameter $\\Delta \\Omega_k$. One could choose to live with this conclusion, and indeed if $H_S$ is chosen large enough, and $\\Delta \\Omega_k$ small enough, one would find that most classical trajectories inflate. But there are two reasons for being suspicious about the conclusion. First, the inferred probability measure will be strongly cutoff-dependent, allowing no firm conclusions to be drawn. Second, we find it hard to justify why, if one adopts a time-dependent measure, the decision as to whether universes are or are not geometrically distinct should be made at the Planck time, rather than today. These arguments suggest that some new ingredient or dynamical principle is needed in the theory, in order to  explain why inflation began.   On the other hand, in an anthropic, or ``top down\" approach to cosmology, the idea is to select universes on the basis of what they would be like at low $H$, not at their beginning. In that case, it seems clear that the appropriate measure should be a late-time measure like the one we have used, not an early one. The same statement would apply to any approach based on computing the probability of asymptotic ``out\" states.  Even though we have argued that the measure we have developed is reasonable, and fulfils the conditions originally proposed by Gibbons, Hawking and Stewart, we have no argument that it is unique or even physically relevant.  It seems to be a perfectly satisfactory {\\it a priori} measure, that is, an unbiased estimator of our ignorance. However, its status as such is rather different from probability distributions which have arisen as a result of some equilibration process. In our case, there is no obvious physical mechanism which allows different members of the ensemble to interact, and without such interactions, the notion of equilibration is not relevant. What the canonical measure does is allow one to discuss in a quantitative way how different proposals for the big bang, or the beginning of the universe, cut down the space of classical trajectories and hence make predictions about the state of the universe today.  The fact that the canonical measure, with what seems a sensible resolution of its divergence, strongly disfavors many e-folds of inflation, poses a serious puzzle for inflationary theory. It is important, we believe, for inflationary theory to explain why the kinetic-dominated trajectories which we have found to overwhelmingly dominate the canonical measure are somehow excluded. Indeed, the measure we have calculated is generous to inflation in that we have assumed spatial homogeneity and isotropy, and the canonical measure gives equal weight to every distinct classical solution, even those with very high potential energy density early on. The main conclusion we draw from this work is that the question of why or how inflation started remains a deep mystery, and a challenge for fundamental theory. Until that question is answered, we should remain cautious about claiming that cosmology's classic puzzles are ``solved\".   \\medskip  {\\bf Acknowledgements} We thank Stephen Hawking, Alexei Starobinsky and other participants in the first Cambridge-Mitchell Texas Conference for stimulating discussions, and Bob Wald for correspondence on this topic. This work was funded in part by PPARC (UK)."
    },
    "0609/astro-ph0609715.txt": {
        "abstract": "% context heading (optional) {} % aims heading (mandatory) { We present the first NIR \\textit{spectro-interferometry} of the LBV %Luminous Blue Variable \\object{$\\eta$~Carinae}. The observations were performed with the AMBER instrument of the ESO \\textit{Very Large Telescope Interferometer} (VLTI) using baselines from 42 to 89~m. The aim of this work is to study the wavelength dependence of $\\eta$~Car's optically thick wind region with a high spatial resolution of 5~mas (11~AU) and \\textit{high spectral resolution.} } % methods heading (mandatory) { The observations were carried out with three 8.2 m Unit Telescopes in the $K$-band. The raw data are spectrally dispersed interferograms obtained with spectral resolutions of 1,500 (MR-K mode) and 12,000 (HR-K mode). The MR-K observations were performed in the wavelength range around both the \\ion{He}{I}\\,2.059\\,$\\mu$m and the Br$\\gamma$\\,2.166\\,$\\mu$m emission lines, the HR-K observations only in the Br$\\gamma$ line region. } % results heading (mandatory) { The spectrally dispersed AMBER interferograms allow the investigation of the \\textit{wavelength dependence} of the visibility, differential phase, and closure phase of $\\eta$~Car. In the $K$-band continuum, a diameter of $4.0\\pm0.2$~mas (Gaussian FWHM, fit range 28--89~m baseline length) was measured for $\\eta$~Car's optically thick wind region. If we fit Hillier et al.\\ (2001) model visibilities to the observed AMBER visibilities, we obtain 50\\% encircled-energy diameters of 4.2, 6.5 and 9.6~mas in the 2.17$\\,\\mu$m continuum, the \\ion{He}{I}, and the Br$\\gamma$ emission lines, respectively. In the continuum near the Br$\\gamma$ line, an elongation along a position angle of $120\\degr\\pm15\\degr$ was found, consistent with previous VLTI/VINCI measurements by van Boekel et al.\\ (2003). We compare the measured visibilities with predictions of the radiative transfer model of Hillier et al.\\ (2001), finding good agreement. Furthermore, we discuss the detectability of the hypothetical hot binary companion. For the interpretation of the non-zero differential and closure phases measured within the Br$\\gamma$ line, we present a simple geometric model of an inclined, latitude-dependent wind zone. Our observations support theoretical models of anisotropic winds from fast-rotating, luminous hot stars with enhanced high-velocity mass loss near the polar regions. } % conclusions heading (optional), leave it empty if necessary %{{\\bf tbd}. %} {} ",
        "introduction": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  The enigmatic object $\\eta$ Car is one of the most luminous and most massive ($M\\sim100\\,M_\\odot$) unstable Luminous Blue Variables suffering from an extremly high mass loss \\citep{dh97}. Its distance is approximately 2300$\\pm$100~pc \\citep{dh97,davidsonetal01,smi06}. $\\eta$~Car, which has been subject to a variety of studies over the last few decades, is surrounded by the expanding bipolar Homunculus nebula ejected during the Great Eruption in 1843. The inclination of the polar axis of the Homunculus nebular with the line-of-sight is $\\sim 41\\degr$ with the southern pole pointing towards us \\citep{davidsonetal01,smi06}. The first measurements of structures in the innermost sub-arcsecond region of the Homunculus were obtained by speckle-interferometric observations \\citep{we86,hw88}. These observations revealed a central object (component A) plus three compact and surprisingly bright objects (components B, C, and D) at distances ranging from approximately 0.1$^{''}$ to 0.2$^{''}$. HST observations of the inner $1^{''}$ region \\citep{wei95} provided estimates of the proper motion of the speckle objects B, C, and D (velocity $\\sim50$ km/s; the low velocity suggests that the speckle objects are located within the equatorial plane), and follow-up HST spectroscopy unveiled their unusual spectrum \\citep{dav95}. The central object (speckle object A) showed broad emission lines, while the narrow emission lines came from the speckle objects B, C, and D. Therefore, A is certainly the central object while B,C, and D are ejecta. Recent observations of $\\eta$ Car by \\cite{chesneau05} using NACO and VLTI/MIDI revealed a butterfly-shaped dust environment at $3.74\\,$ and $4.05\\,\\mu$m and resolved the dusty emission from the individual speckle objects with unprecedented angular resolution in the NIR. Chesneau et al.\\  also found a large amount of corundum dust peaked $\\sim1^{''}$ south-east of the central object.  Spectroscopic studies of the Homunculus nebula showed that the stellar wind of $\\eta$~Car is aspherical and latitude-dependent, and the polar axes of the wind and the Homunculus appear to be aligned \\citep[bipolar wind model;][]{smi03}. Using Balmer line observations obtained with HST/STIS, \\citet{smi03} found a considerable increase of the wind velocity from the equator to the pole and that the wind density is higher in polar direction \\citep[parallel to the Homunculus; PA  of the axis $\\sim$132\\degr; ][]{davidsonetal01} than in equatorial direction by a factor of $\\sim$2. \\citet{van03} resolved the optically thick, aspheric wind region with NIR interferometry using the VLTI/VINCI instrument. They measured a size of 5 mas (50\\% encircled-energy diameter), an axis ratio of $1.25\\pm0.05$, and a position angle (PA) of the major axis of $134\\degr\\pm7\\degr$, and derived a mass-loss rate of $1.6\\times10^{-3}\\,M_\\odot/{\\rm yr}^{-1}$. The aspheric wind can be explained by models for line-driven winds from luminous hot stars rotating near their critical speed \\citep[e.g.,][]{owo96,owo98}. The models predict a higher wind speed and density along the polar axis than in the equatorial plane. In addition, van Boekel et al.\\ showed that the $K$ broad-band observations obtained with VINCI are in agreement with the predictions from the detailed spectroscopic model by \\citet{hillier01}.  The \\citet{hillier01,hillier06} model was developed to explain STIS HST spectra. The luminosity of the primary ($5 \\times 10^6\\,L_\\odot$) was set by observed IR fluxes \\citep[see discussion by ][]{dh97} and the known distance of 2.3~kpc to $\\eta$~Car. Any contribution to the IR fluxes by a binary companion was neglected.  Modeling of the spectra was undertaken using CMFGEN, a non-LTE line blanketed radiative transfer developed to model stars with extended outflowing atmospheres \\citep{himi98}. For the modeling of $\\eta$~Carinae, ions of H,  He, C, N, O, Na, Mg, Al, Si, S, Ca, Ti, Cr, Mn, Fe, Ni, and Co were included. The mass loss was derived from the strength of the hydrogen lines and their associated electron scattering wings. Due to a degeneracy between the mass-loss rate and the He abundance, the H/He helium abundance ratio could not be derived, but was set at 5:1 (by number), which is similar to that found by \\citet{davidson86} from nebula studies. CNO abundances were found to be consistent with those expected for full CNO processing. With the exception of Na (which was found to be enhanced by at least a factor of 2), the adoption of solar abundances for other metal species was found to yield satisfactory fits to the STIS spectra. A more recent discussion of the basic model, with particular reference to the UV and outer wind, is given by \\citet{hillier06}.  Because the wind is optically thick, the models are fairly insensitive to the radius adopted for the hydrostatic core (i.e., the radius at which the velocity becomes subsonic). One exception was the \\ion{He}{I} lines, which decreased in strength as the radius increased and, in general, were very sensitive to model details. Additional HST STIS observations show that the \\ion{He}{I} lines are strongly variable and blue-shifted throughout most of the 5.54-year variability period. These observations cannot be explained in the context of a spherical wind model. It now appears likely that a large fraction of the \\ion{He}{I} line emission originates in the bow shock and an ionization zone, associated with the wind-wind interaction zone in a binary system \\citep{davidsonetal99,davidson01, hillier06,nielsen06}. Consequently, the hydrostatic radius derived by \\citet{hillier01} is likely to be a factor of 2 to 4 too small. Because the wind is so thick, a change in radius will not affect the Br$\\gamma$ formation region, and it will only have a minor influence on the Br$\\gamma$ continuum emitting region. If this model is correct, the \\ion{He}{I} emission will be strongly asymmetrical and offset from the primary star.   A variety of observations suggest that the central source of $\\eta$ Car is a binary. \\citet{dam96} first noticed the 5.5-year periodicity in the spectroscopic changes of this object \\citep[see][]{dam97,dam00, duncan99,ishibashi99,davidsonetal99,davidsonetal00,vangend03,sd04,whitelock04,corcoran05,weis05}. On the other hand, to date, the binary nature of the central object in $\\eta$~Car and its orbital parameters are still a matter of debate \\citep[see, e.g.,][]{ZWS84,davidson99,davidson01,davidsonetal99, davidsonetal00,davidsonetal05,ishibashi99,smith00,feast01,ishibashi01,pico02,smi03,martin06}.  The 1997.9 X-ray peak with the subsequent rapid drop to a few-month-long minimum was detected by RXTE \\citep[see][]{corcoran05}. Then the first spectra with HST/STIS were obtained at 1998.0, demonstrating changes in both the central star and the aforementioned speckle objects \\citep{davidsonetal99,gull99}. \\citet{pico02} demonstrated that the CHANDRA X-ray spectrum can be explained by the wind-wind collisions of the primary star ($\\dot{M}=2\\times10^{-4}\\,M_\\odot/{\\rm yr}^{-1}$ at 500 km/s) and a hot companion ($\\dot{M}=10^{-5}\\,M_\\odot/{\\rm yr}^{-1}$ at 3,000 km/s). \\citet{verner05b} used models calculated with the CLOUDY code to demonstrate that during the spectroscopic minimum, the excitation of the speckle objects is supported by the primary stellar flux, but that the UV flux of a hot companion consistent with an O7.5V, O9I, or early WN star was probably necessary to excite the speckle objects during the broad spectroscopic maximum.  In this paper we present the first spectro-interferometric $K$-band observations of $\\eta$ Car obtained with the VLTI beam-combiner instrument AMBER with medium and high spectral resolution and in the projected baseline range from 28 to 89~m.  %%% %%% Tab. 1: Overview of AMBER observations %%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{table*}[t] \\caption{Summary of the AMBER $\\eta$~Car observations using the UT2, UT3, and UT4 telescopes.} \\label{tab:observations} \\centering \\begin{tabular}{lll cccccccc} \\hline\\hline Date [UT]    &\\multicolumn{2}{c}{Time [UT]} &Orbital  &spectral &line within                   &DIT$^d$  &$N_{\\eta~Car}^f$&Calibrator &$N_{calib.}^g$   &Calibrator\\\\ &Start  &End                                &phase$^e$&mode     &spectral range                &         &                &           &               &uniform disk \\\\ &       &                                   &         &         &                              &[ms]     &                &           &               &diameter [mas]\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 2004 Dec. 26 & 07:52 & 08:16                & 0.267  &MR-K     & Br$\\gamma$                &40                &7,500         &HD\\,93030   &5,000 & 0.39$^a$\\\\ & 08:19 & 08:32                & 0.267  &MR-K     & \\ion{He}{I}~2.059\\,$\\mu$m &40                &5,000         &HD\\,93030   &5,000 & 0.39$^a$\\\\ 2005 Feb. 25 & 04:33 & 04:43                & 0.298  &MR-K     & Br$\\gamma$                &50                &5,000         &HD\\,89682   &2,500 & 3.08$^b$\\\\ & 04:55 & 05:05                & 0.298  &MR-K     & \\ion{He}{I}~2.059\\,$\\mu$m &50                &5,000         &HD\\,89682   &2,500 & 3.08$^b$\\\\ 2005 Feb. 26 & 08:16 & 08:57                & 0.299  &HR-K     & Br$\\gamma$                &82                &7,500         &L\\,Car      &2,500 & 2.70$^c$\\\\ \\noalign{\\smallskip} \\hline \\end{tabular}  \\begin{flushleft} \\hspace{0mm}Notes~--~$^{a}$ Uniform disk (UD) diameter estimated using the method described by~\\citet{dyc96}.\\\\ \\hspace{0mm}$^{b}$ UD diameter taken from the CHARM2 catalog~\\citep{ric05}.\\\\ \\hspace{0mm}$^{c}$ UD diameter of L~Car at the time of the AMBER high-resolution observations derived from the limb-darkened diameter $d_{\\rm LD}=2.80\\,$ mas at the L~Car pulsation phase $\\phi$=0.0 \\citep{ker04a,ker06} and $d_{\\rm UD}/d_{\\rm LD}=0.966$ \\citep{ker04b}.\\\\ \\hspace{0mm}$^{d}$ Detector integration time per interferogram.\\\\ \\hspace{0mm}$^{e}$ The orbital phase was computed assuming a zero point at JD 2\\,450\\,800.0 and a period of 2024~days \\citep{corcoran05}.\\\\ \\hspace{0mm}$^{f}$ Number of $\\eta$ Car interferograms.\\\\ \\hspace{0mm}$^{g}$ Number of calibrator interferograms.\\\\ \\end{flushleft} \\vspace*{-3mm}  \\end{table*} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  The paper is organized as follows: In Sect.\\ 2 we give an overview of the AMBER observations of $\\eta$ Car and describe the data reduction procedure in detail, and in Sect.\\ 3, the analyses of the continuum data and the measurements within the Br$\\gamma$ and \\ion{He}{I} lines are discussed individually.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": " \\begin{itemize} \\item In the $K$-band continuum, we resolved $\\eta$ Car's optically thick wind. From a Gaussian fit of the $K$-band continuum visibilities in the projected baseline range from 28--89~m, we obtained a FWHM diameter of $4.0\\pm0.2$ mas. Taking the different fields-of-view into account, we found good agreement between the AMBER measurements and previous VLTI/VINCI observations of $\\eta$ Car presented by \\citet{van03}. \\item When comparing the AMBER {\\it continuum} visibilities with the NLTE radiative transfer model from \\citet{hillier01}, we find very good agreement between the model and observations. The best fit was obtained with a slightly rescaled version of the original Hillier et al.\\ model (rescaling by 1--2\\%), corresponding to FWHM diameters of 2.27 mas at $\\lambda=2.040\\,\\mu$m and 2.33~mas at $\\lambda=2.174\\,\\mu$m. \\item If we fit \\citet{hillier01} model visibilities to the observed AMBER visibilities, we obtain, for example, 50\\% encircled-energy diameters of 4.2, 6.5, and 9.6~mas in the 2.17$\\,\\mu$m continuum, the \\ion{He}{I}, and the Br$\\gamma$ emission lines, respectively. \\item In the continuum around the Br$\\gamma$ line, we found an asymmetry towards position angle PA=$120\\degr\\pm15\\degr$ with a projected axis ratio of $\\xi=1.18\\pm0.10$. This result confirms the earlier finding of \\citet{van03} using VLTI/VINCI and supports theoretical studies which predict an enhanced mass loss in polar direction for massive stars rotating close to their critical rotation rate \\citep[e.g.\\ ][]{owo96,owo98}. \\item For both the Br$\\gamma$ and the \\ion{He}{I} emission lines, we measured non-zero differential phases and non-zero closure phases within the emission lines, indicating a complex, asymmetric object structure. \\item We presented a physically motivated model which shows that the asymmetries (DPs and CPs) measured within the wings of the {\\it Br$\\gamma$ line} are consistent with the geometry expected for an aspherical, latitude-dependent stellar wind.  Additional VLTI/AMBER measurements and radiative transfer modeling will be required to determine the precise parameters of such an inclined aspherical wind. \\item Using a simple binary model, we finally looked for a possible binary signature in the AMBER closures phases. For separations in the range from 4 to 14~mas and arbitrary PAs, our simple model reveals a minimum $K$-band flux ratio of $\\sim$50 with a 90\\% likelihood. \\end{itemize}  Our observations demonstrate the potential of VLTI/AMBER observations to unveil new structures of $\\eta$~Car on the scales of milliarcseconds. Repeated observations will allow us to trace changes in observed morphology over $\\eta$~Car's spectroscopic 5.5~yr period, possibly revealing the motion of the wind-wind collision zone as predicted by the $\\eta$~Car binary model. Furthermore, future AMBER observations with higher accuracy might be sensitive enough to directly detect the hypothetical hot companion.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0609/astro-ph0609761_arXiv.txt": {
        "abstract": "{}  {We have performed optical  polarimetric observations of the SN~2006aj  associated to the  X-ray flash (XRF) of  February 18, 2006, \\object{XRF~060218} that  provide information on its expansion geometry.}  {The  data were acquired  in the $R$-band with  the 0.7m telescope of Crimea,  2.5m Nordic Optical Telescope and  the 2.2m of Calar  Alto.}   {We  report  the detection  of  linear  polarization between 3  and 39  days after the  gamma-ray event  ($t-t_0$).  This represents the  first polarization detection of a  Ic supernova (SN) associated  to  an  XRF.   Our  data  exhibit  a  degree  of  linear polarization  ($P$)  around $P\\sim4\\%$  at  $t-t_0  \\sim 3-5$  days, followed by  a constant  polarization phase with  $P \\sim  1.4\\%$ at $13.7 \\lesssim t-t_0 \\lesssim 39$ days.  Our data suggest a decay in $P$,  and  more  interestingly,  show a  position  angle  ($\\theta$) rotation of $\\sim 100^{\\circ}$ comparing data taken before and after the   $R$-band  lightcurve   peak.}    {The  reported   polarization measurements can be  explained by the evolution of  an asymmetric SN expansion.  We discuss on several ingredients that could account for the observed $\\theta$ rotation. } ",
        "introduction": "The {\\it  Swift} satellite detected a long-soft  gamma-ray burst (GRB) at $t_0$ = Feb.  18.149 2006 UT, (Cusumano et al.  2006), showing both a low redshift ($z=0.033$, Mirabal  \\& Halpern 2006) and a peak energy (Campana et al.  2006),  typical of X-ray flashes (\\object{XRF~060218} hereafter).    Further  observations  associated   it  to   a  Ic-type supernova,   \\object{SN~2006aj}   (Masetti   et  al.    \\cite{Mase06}; Soderberg et  al.  \\cite{Sode06};  Pian et al.   \\cite{Pian06}).  This finding  strengthened  even   more  the  already  solid  long-duration GRB/Supernova   link   (Hjorth   et   al  \\cite{Hjor03};   Stanek   et al. \\cite{Stan03}).  Since 1999 (Covino et  al.  \\cite{Covi99}) several GRB afterglows have shown optical polarization at a level of $\\sim$1-3\\% (see Gorosabel et al.  \\cite{Goro04} and references therein).  Optical supernovae (SNe), including typical  Ic-types not related to  GRBs, do not  tend to show optical polarization  above $\\sim1$\\%.  However,  some counterexamples exist (Leonard  et al.  \\cite{Leon06}).  The  very few polarimetric studies of  Ic SNe related to GRBs  do not allow yet  to infer clear global polarimetric differences with  normal Ic SNe.  However, as we report in  this Letter, some  normal Ic SNe  like \\object{SN~2002ap} (Mazzali  et  al.    \\cite{Mazz02})  share  some  similarities  with \\object{SN~2006aj};  polarization levels  above $\\sim1$\\%  and, more interestingly,  long-term   rotations  in  the   polarization  angle (Kawabata et al.  \\cite{Kawa02}).  \\vspace{-0.4cm}  Most of  the SN intrinsic  polarization detections are  attributed to the break  of the symmetry  around line of  sight, due to  a Thompson photon scattering  through an aspherical  SN expansion (Kasen  et al. \\cite{Kase03}).   This polarization  is  expected to  decay as  $\\sim t^{-2}$, as the  atmosphere expands and the optical  depth (and hence the  polarization) falls  (Leonard  et al.   \\cite{Leon06}). In  most cases the afterglow is so  strong that overshines the SN polarization component   (for    instance   in   GRB~030329;    Greiner   et   al. \\cite{Grei03}).   In this  Letter  we report  the first  polarization detection to date  of a SN associated to an XRF,  not coming from the GRB ejecta forward shock emission (Fan et al. \\cite{Fan06}; Soderberg et   al.   \\cite{Sode06}).   Sect.~\\ref{Observations}   reports  the observations      carried      out     for      \\object{SN~2006aj}. Sect.~\\ref{results} presents our results and Sect.~\\ref{Discussion} discusses     the     implications     of     our     measurements. Sect.~\\ref{Conclusions} summarises the final conclusions. ",
        "conclusions": "\\label{Conclusions}  We  report  the  detection   of  $R$-band  linear  polarization  for XRF~060218/SN~2006aj in observations done since $t-t_0 \\sim 3$ to 39 days. These  observations represent  the first detection  of optical polarization in  an XRF.  The  measured high degree  of polarization gives  further credence to  the XRF/GRB  off-axis scenario,  since a face-on SN would not  likely produce any intrinsic polarization.  As Patat et al (\\cite{Pata01}) noted  there is a degeneracy between the viewing  angle and  the  asymmetry level,  so  a detailed  geometric picture based  on only our  data is not  possible.  Our data  show a $P\\sim4\\%$  value  in the  first  $\\sim  3-5$  days after  the  GRB, followed by a  constant $P \\sim 1.4\\%$ polarization  phase at $t-t_0 \\gtrsim 13.7$ days, suggesting a  $P$ decay.  $\\theta$ shows a $\\sim 100^{\\circ}$ rotation at $5.6 <  t-t_0 <13.7$ days.  We propose that the $P$ evolution can be  explained by an asymmetric SN.  We discuss several ingredients  which separately, or combined,  can explain the observed $\\theta$ rotation (in  decreasing order of relevance); i) a highly  asymmetric high-velocity  SN  ejection, followed  by a  less asymmetric SN bulk,  ii) polarization induced by the  host ISM, iii) asymmetric distribution  of $^{56}$Ni and iv) scattering  by a dusty region placed a few light-days from the progenitor.  \\begin{figure}[t] \\resizebox{\\hsize}{6.5cm}{\\includegraphics[bb=30 60 709 595]{fig3.ps}} \\caption{\\label{fig3} Our polarization  results, once the the Galactic ISM correction was considered.  {\\em Upper panel:} As shown $\\theta$ exhibited a  rotation of $\\sim100^{\\circ}$  at $5.6 < t-t_0  < 13.7$ days.   {\\em Middle  panel:} $P$  shows  a likely  decay reaching  a constant  level at  a $\\sim  1.4\\%$ level.   {\\em Lower  panel:} The intrinsic $P_{SN}$ evolution assuming that our last $P$ detection is entirely originated  by the  host ISM.  The  dashed curves  show the $\\pm 1\\sigma$ uncertainty  region of the $t^{-2}$ decay,  as seen in \\object{SN~2004dj}.  } \\end{figure}"
    },
    "0609/astro-ph0609282_arXiv.txt": {
        "abstract": "A method is presented for investigating the periodic signal content of time series in which a number of signals is present, such as arising from the observation of multiperiodic oscillating stars in observational asteroseismology. Standard Fourier analysis tends only to be effective in cases when the data are perfectly regularly sampled. During normal telescope operation it is often the case that there are large, diurnal, gaps in the data, that data are missing, or that the data are not regularly sampled at all. For this reason it is advantageous to perform the analysis as much as possible in the time domain. Furthermore, for quantitative analyses of the frequency content and power of all real signals, it is of importance to have good estimates of the errors on these parameters. This is easiest to perform if one can use linear combinations of the measurements. Here such a linear method is described. The method is based in part on well-known techniques in radio technology used in every FM radio receiver, and in part on the SOLA inverse method. ",
        "introduction": "Low-amplitude solar and stellar oscillations can be considered as a linear superposition of normal modes of resonance. In theory one indexes each by a radial order $n$ for the mode, and the degree $\\ell$, and order $m$ of the spherical harmonic functions which describe the azimuthal spatial behaviour of the normal modes. For solar-like oscillations, excited stochastically by turbulent convection, there are usually many oscillation modes present at any epoch, with comparable amplitudes. If a time series, in which there are a number of periodic signals present, is sampled perfectly regularly the problem of determining the frequencies can be solved in a straightforward manner using a Fourier transform and there is a single peak for each real frequency present in the data.  If there are gaps in the data or if the data is irregularly sampled there will be more than one peak in the Fourier spectrum for each signal frequency in the data. The shape of this response for each signal frequency that is present in the data is known as the `window function'. Clearly, if there are a large number of real signal frequencies in gapped or irregularly sampled data, the Fourier transform very quickly becomes very complicated and hence it becomes very difficult to determine which are true frequencies, and which are aliases generated by the sampling.  There exist a number of techniques for resolving the problem of determining amplitudes, frequencies and phases of periodic signals in irregularly sampled, gapped data. Some attempt to `fill gaps' using autoregressive methods (cf. \\cite{FaUl82}, \\cite{BJCD90}, \\cite{Foea99}). Other methods attempt to do a least-squares fitting of sinusoids in the time domain (cf. \\cite{Sc82}, \\cite{Sc89}, \\cite{Ko99}) or to define a best-fit spectrum in a least-squares sense, but mollified to take into account that this problem is ill-posed since it is a form of an inverse problem (\\cite{SLL00}, \\cite{VSW00}).  Apart from the problems of sampling, there is also the problem of noise in the data. In principle a star can produce `intrinsic noise', through the stellar equivalent of granulation. The telescope and instrumentation can also introduce noise into time-series which in an ideal case would amount to the photon noise due to finite integration times. Such noise might be white, i.e. independent of frequency, but in some cases it could be confined to bands of frequencies. It therefore makes some sense to `chop up' the full frequency  range available, set by the sampling and length of the time series, into  different regimes. \\cite{ORPPN} refer to this as a `multiresolution approach'.  \\section {The method} In the approach of \\cite{ORPPN} the series is convolved with every member of a set of Gaussians, with successively increasing widths $\\Delta_{j+1}=2\\Delta_j$ and the differences between these successive convolutions are computed. If the data were regularly sampled each of these differences would correspond in the frequency domain to having applied a band-pass filter to the data. If all the different multiresolution levels were to be co-added the original data would be reproduced. The convolution is carried out in the time domain by defining weights $w_i$ for each datum $Y_i$ measured at time $t_i$ and then summing $w_i Y_i$ for all $i$.  This method of separating frequency space into broad bands has the convenience of simplicity since convolving with Gaussians in the time domain corresponds to multiplying by Gaussians in the Fourier domain and the filter parameters are easy to determine. However, the shape of the filter corresponding to the differences between levels is awkward since it is skewed and it reduces the amplitude of the signal everywhere except at one off-center frequency in the band. It is preferable to have a filter that is flat over the band, at least in the limit that errors $\\sigma_i$ are equal for all $i$, with a symmetric smooth drop-off at the lower and upper limit of the band. One example of that is a filter defined by the following convolution~: \\begin{eqnarray} H(\\omega) &=& \\int\\limits_{-\\infty}^{\\infty} \\textrm{d}\\omega' {1\\over\\delta_\\omega\\sqrt{\\pi}}\\exp\\left[{-(\\omega-\\omega')^2 \\over\\delta_\\omega^2}\\right] \\Pi(\\omega') \\nonumber \\\\ \\Pi(\\omega) &=& 1\\ \\ \\textrm{for}\\ \\ \\vert\\omega\\vert < \\Delta_\\omega \\nonumber\\\\ &=&0\\ \\ \\textrm{for}\\ \\ \\vert\\omega\\vert >\\Delta_\\omega \\label{bandfiltfreq} \\end{eqnarray} Band-pass filtering, i.e. multiplying in the Fourier domain by this filter function $H(\\omega-\\omega_c)$, centered on a frequency $\\omega_c$, corresponds in the time domain by convolving with the function~: \\begin{equation} h(t) = {\\sin\\Delta_\\omega t\\over \\pi t} \\exp\\left[-{1\\over 4}\\delta_\\omega^2 t^2\\right]\\cos\\omega_c t \\label{bandfilttime} \\end{equation} One can create successive bands that are orthogonal by choosing for instance~: \\begin{eqnarray} {\\delta_\\omega\\over\\Delta_\\omega} &\\equiv& r\\ll 1 \\ \\ \\textrm{fixed}\\nonumber\\\\ \\Delta_{\\omega,j+1}&=& \\alpha_j\\Delta_{\\omega,j} \\nonumber\\\\ {\\omega_{c,j+1}\\over\\Delta_{\\omega,j+1}} &\\equiv& \\gamma_{j+1}= 1 + {\\gamma_j+1\\over\\alpha_j}\\ \\ \\textrm{with}\\ \\ \\gamma_1 = 0 \\label{filtpardefs} \\end{eqnarray} The weights $w_{m,ij}$ are therefore in this case set as~: \\begin{equation} w_{m,ij} = {\\Delta_{\\omega,j}\\over W_j\\pi} h_j(x_{m,ij}) \\end{equation} where $W_j$ is again a normalisation factor and the $x_{m,ij}$ and $h_j$ are~: \\begin{eqnarray} x_{m,ij} &=& \\Delta_{\\omega,j} (t_m - t_i) \\nonumber\\\\ h_j(x) &=& {\\sin x\\over x}\\exp\\left[-{1\\over 4}r^2 x^2\\right]\\cos\\left( \\gamma_j x\\right) \\label{xwdefs} \\end{eqnarray} The band-pass filtered data $R$ are now calculated as~: \\begin{equation} R_{i,j} \\equiv R_{t_i,j} = \\sum_m w_{m,ij} Y_m \\label{broadfilter} \\end{equation}  \\subsection{local oscillator step} In order to obtain the frequency content of the time series, without performing Fourier transforms on unevenly sampled data, one can employ what is known as a Hilbert transform. In radio receiver technology this would be referred to as mixing the signal with a local oscillator, and then applying a low-pass filter. The frequency $f$ of the local oscillator is `tunable' : the width of the band-pass set by the previous step is covered, sampling regularly in frequency with some sampling step $\\Delta f$. This sampling distance $\\Delta f$ is chosen in combination with the width of the low-pass filter subsequently applied, so that some oversampling is performed.  The low pass filter to apply to the `mixed signal' is chosen to be a Gaussian. This is a convenient choice because multiplying by a (narrow) Gaussian filter in the frequency domain, corresponds to convolving with a (wide) Gaussian in the time domain. The width $\\Delta_{LPF}$ in the time domain of what corresponds to a low-pass filter, is linearly related to the parameter $\\Delta$ of the multiresolution level within which the signal is to be frequency analysed.  One proceeds as follows. For each frequency $f_k$ define sets of `local oscillator' weights $q_{i,k}$, $p_{i,k}$~: \\begin{eqnarray} q_{i,k} &=& 2\\cos(2\\pi f_k t_i)\\nonumber\\\\ p_{i,k} &=& -2\\sin(2\\pi f_k t_i) \\end{eqnarray} and also define a set of low-pass filter weights $z_{i,l}$~: \\begin{equation} z_{i,l} = {1\\over\\Delta_{LPF}\\sqrt{\\pi}}\\exp\\left[-\\left({t_i -T_l\\over \\Delta_{LPF}}\\right)^2\\right] \\label{lowpassfilter} \\end{equation} The central times $T_l$ are spaced by some factor of order unity times $\\Delta_{LPF}$. $\\Delta_{LPF}$ is large in order to only let through signal in a narrow band, and therefore the number of times $T_l$ is small~: typically of the order of the number of nights of observations. One can define a cosine-weighted average $\\langle R_{\\cos}\\rangle$ and a sine-weighted average $\\langle R_{\\sin}\\rangle$ as follows~: \\begin{eqnarray} \\langle R_{\\cos}\\rangle_{lkj} = \\sum_i z_{i,l}q_{i,k}R_{i,j}\\nonumber\\\\ \\langle R_{\\sin}\\rangle_{lkj} = \\sum_i z_{i,l}p_{i,k}R_{i,j} \\label{Hilbertterms} \\end{eqnarray} The three indices for these two averaged quantities refer to~: \\begin{enumerate} \\item{} $j$ the `resolution level' of the time series, which corresponds to a broad range in frequencies between a lower and upper bound set by a smoothing width $\\Delta_\\omega$. \\item{} $k$ the central frequency $f_k$ of the tunable narrow-band filter runs from the lower to the upper limit of the band $j$ to explore the amplitude of signal in the time series at/around each $f_k$. \\item{} $l$ the central time on which a broad Gaussian is centered, providing a narrow band filter. By taking a finite value for the width of this Gaussian (rather than infinite) and stepping through the time series, a time-frequency analysis is done, since it provides a mechanism to consider the frequency content per night, or per few successive nights. \\end{enumerate} For data covering only a few nights with a large gap during the day-time, the window function has quite strong daily sidelobes which complicates analysis of a spectrum if there are a number of peaks present, as is common in multi-periodic variable stars. It is therefore useful to consider improvements on the method with the aim of reducing sidelobe structure. This may be achieved by choosing the low-pass filtering weights $z_{i,l}$ in a different manner than described by (\\ref{lowpassfilter}).  \\subsection{Strategy} Consider again the two summations of Eq. (\\ref{Hilbertterms}) which are defined to be functions of frequency $f$. The weights are now $\\zeta_i$ and in principle one wishes to be free to choose weights differently depending on the frequency $f_k$ of the local oscillator. The time $T_l$ on which the low-pass filter is centered in Eq. (\\ref{lowpassfilter}) now enters the discussion by explicitly setting the reference time for the  phase $\\phi$, which means replacing $t_i$ with $t_i-T_l$. With the weights $\\zeta$ yet to be determined two functions of $f$ are defined as follows~: \\begin{eqnarray} \\sum_i \\zeta_{i,kl}\\cos\\left(2\\pi (f-f_k) (t_i -T_l) \\right) &\\equiv&\\Xi (f) \\nonumber\\\\ \\sum_i \\zeta_{i,kl}\\sin\\left(2\\pi (f-f_k) (t_i -T_l) \\right) &\\equiv&X (f) \\label{KsiXdefs} \\end{eqnarray} The method has three separate aims to achieve with its choice for the weights. These aims may well not be perfectly compatible and therefore require a form of compromise~: \\begin{itemize} \\item{} The first function $\\Xi(f)$ needs to be a function peaked at $f=0$ and smoothly dropping away, with as little side-lobe structure as possible. For instance~: \\begin{equation} \\Xi(f) \\approx {1\\over\\sqrt{\\pi}\\Delta_f}\\exp\\left[-\\left({f-f_k\\over\\Delta_f} \\right)^2\\right] \\label{ksitarget} \\end{equation} with as small a width $\\Delta_f$ as achievable. \\item{} The second function $X(f)$ needs to be as close to $0$ as possible everywhere \\item{} The errors in the resulting amplitudes $A^2$ need to be kept as small as possible. \\end{itemize}  \\begin{figure} \\centerline{\\psfig{file=FPijpersFig1.eps,width=8.0cm}} \\caption[]{A histogram of the distribution of intervals between successive sampling times for an artificial time series, divided by the smallest of these intervals. The vertical dashed line indicates the mean time interval for the entire set. The inset shows the part of the distribution at small intervals which is off-scale for the larger figure.} \\label{Thistfig} \\end{figure}  In the framework of the SOLA method (cf. \\cite{PT92}, \\cite{PT94}) the best compromise between these three aims can be achieved by minimising for $\\zeta_{i,k}$ the quantity ${\\cal W}(\\zeta)$ defined by~: \\begin{eqnarray} {\\cal W}(\\zeta) \\equiv (1-\\mu)\\int\\limits_{f_{\\rm low}}^{f_{\\rm high}} {\\rm d}f \\left[\\Xi(f) - {\\cal T}(f)\\right]^2  \\nonumber\\\\ +\\mu \\int\\limits_{f_{\\rm low}}^{f_{\\rm high}}{\\rm d}f \\left[X(f)\\right]^2 + \\lambda\\sum_{i,j} \\zeta_{i,k}\\zeta_{j,k}\\Sigma_{ij} \\label{SOLAmin} \\end{eqnarray} while taking into account the normalisation constraint~: \\begin{equation} \\sum_i \\zeta_{i,k}\\left[\\int\\limits_{f_{\\rm low}}^{f_{\\rm high}} {\\rm d}f \\cos\\left(2\\pi (f-f_k) (t_i -T_l) \\right)\\right] \\equiv 1 \\end{equation} Here the target function ${\\cal T}(f)$ is taken to be a Gaussian as in (\\ref{ksitarget}). The `target function' for $X(f)$ is $0$. The matrix $\\Sigma_{ij}$ is the variance-covariance matrix of the data errors. It is usual to assume that the errors on the data are Gaussian distributed and independent, so that $\\Sigma$ is a diagonal matrix, but this is not essential for the algorithm. The methods described in \\cite{Ko99} to estimate or model error properties from residuals using ARIMA models can be applied here as well. The parameters $\\mu\\in[0,1]$ and $\\lambda\\in[0,\\infty)$ weight the relative importance of the three terms. Their value needs to be set by trial and error. A very low choice of $\\lambda$ normally leads to unacceptably high errors on the inferred $A^2$, whereas a very high value generally produces the undesirable effect of increasing width or structure in $\\Xi$. Similar arguments hold for $\\mu$.  \\begin{figure*} \\centerline{\\psfig{file=FPijpersFig2.eps,width=11.5cm}} \\caption[]{The response to a monochromatic signal of $2.5\\ {\\rm mHz}$ sampled with the sampling shown in Fig. 1, which is the equivalent of the window function for this method. The left and middle columns show the functions $\\Xi$ and $X$, with in addition the target functions (dashed linmes). The right hand column shows the sum of the square of these which is the equivalent of the window function for power. The rows are in order of increasing spectral resolution with the $\\Delta_f = 8,\\ 4,\\ 2,\\ 0.69\\ \\mu{\\rm Hz}$, from top to bottom. The lower and upper frequency limit of the band are set to $f_{\\rm low} = 2.5\\ {\\rm mHz}$ and $f_{\\rm high} = 6\\ {\\rm mHz}$.} \\label{tableauFig} \\end{figure*}  A point that {\\bf is} crucial for this version of the SOLA algorithm, is that $f_{\\rm low}$ and $f_{\\rm high}$ must be finite. This is necessary because the integrations in Eq. (\\ref{SOLAmin}) are over products of cosines, which causes problems if they are to be integrated over an infinite range. The algorithm is therefore limited in application to bandwidth-limited data. In practice this is not a severe problem, since by construction of the multiresolution steps, bandwidth limitation is achieved. Also, given a finite minimum interval between measurements it is known that signals with frequencies above the Nyquist frequency $\\propto 1/\\min(t_{i+1}-t_i)$ cannot be resolved by such data, which also sets an upper limit to the frequency. ",
        "conclusions": "The method described in this paper, for the analysis of multiperiodic, irregularly sampled time series, has several potential uses. The first point to note is that the method is a time-frequency method formulated in the time domain. This means that the frequencies of the signals in an irregularly sampled time series can be monitored as a function of epoch. This allows for the possibility to detect changes in oscillation frequencies with epoch. Such changes could arise in asteroseismology through similar mechanisms that produce changes in oscillation frequencies of the Sun over the solar cycle and are therefore of interest to determine.  The second point to note is that the method is linear. The algorithm produces a set of linear coefficients on the basis of the sampling times alone. These coefficients can then trivially be combined with data, and with measurement errors, in order to provide the power or amplitude of the time series and also an error estimate for this power, and even a phase if required. Evidently this linearity is convenient for standard asteroseismic data acquisition which provides photometry or velocity. However, it is arguably even more useful if the data obtained is more complex such as a time series of images. A particular application lies in addressing another fundamental problem in asteroseismology, which is the identification of the surface node line pattern associated with any particular frequency~: the mode identification problem. For stars with a single dominant mode, long baseline optical interferometry has been suggested as a means to image the flux variation over the surface, thus allowing the identification of the mode (cf. \\cite{Jan+01} for a description of the relevant technique). In particular the closure phase of interferometric signals between three telescopes is sensitive to symmetries in the flux distribution over the surface of stars (cf. \\cite{DdS+05}). For multiperiodic stars such a technique would suffer since the superposition of patterns would reduce the closure phase response. However, the acquisition of interferometric fringes is already done repeatedly as a function of time for operational reasons. By combining closure phase signals at the appropriate phase in time one could select a single frequency to image, thus restoring the closure phase signal strength. This technique might be referred to as `stroboscopic interferometry'. The closure phase is a somewhat more complex form of data than time series of velocities or fluxes, and carrying out stroboscopic interferometry in practice is facilitated by being able to carry out the time-filtering using the coefficients generated with this SOLA based algorithm. Appropriate tests of the algorithm for this application are in progress (Pijpers, in preparation).  A third point to note is that this method can also be adapted to streamline procedures for model fitting in asteroseismology. The Gaussian form of the target function chosen here is intended to obtain measures of the oscillation power or amplitude as a function of frequency. However, in fitting models to a given star it is not just a single frequency but the entire pattern of frequencies for which a best match needs to be found. This can be achieved more directly by calculating the expected oscillation spectrum for each model and taking for the target function ${\\cal T}(f)$ not a single Gaussian, but a sum of Gaussians centered on model frequencies $f_{nl}$, in which the $nl$ are the radial mode order $n$ and azimuthal degree $l$, instead of the single $f_k$. The coefficients obtained in this way, when combined with the data, will produce a maximal response for that model which has a set of frequencies that is the closest match to the data. Further tests of this are necessary to establish the error propagation properties of such an approach, including the influence of systematic effects (Pijpers, in preparation)."
    },
    "0609/astro-ph0609557_arXiv.txt": {
        "abstract": "This paper presents results of a survey search for bright compact radio sources at 22~GHz with the VERA radio-interferometer. Each source from a list of 2494 objects was observed in one scan for 2 minutes. The purpose of this survey was to find compact extragalactic sources bright enough at 22~GHz to be useful as phase calibrators. Observed sources were either a) within 6 degrees of the Galactic plane, or b) within 11 degrees from the Galactic center; or c) within 2 degrees from known water masers. Among the observed sources, 549 were detected, including 180 extragalactic objects which were not previously observed with the very long baseline interferometry technique. Estimates of the correlated flux densities of the detected sources are presented. ",
        "introduction": "\\label{s:introduction}  VLBI Exploration of Radio Astrometry (hereafter VERA) is a Japanese VLBI array dedicated to phase-referencing astrometry to explore the Milky Way Galaxy \\citep{honma2000a,kobayashi2004}. VERA observes Galactic H$_2$O and SiO maser sources and determines their parallax and proper motions at the 10~microarcsec accuracy level. Based on the high-precision astrometry of $\\sim$1000 maser sources, VERA will study the three-dimensional structure and dynamics of the Galaxy's disk and bulge, revealing the true shape of the bulge and spiral arms, its precise rotation curve and the distribution of dark matter.  The VERA array consists of four 20-m diameter radio telescopes spread over Japan, covering baselines ranging from 1000 km to 2300 km. Unlike other VLBI stations, VERA telescopes are equipped with a dual-beam system \\citep{kawaguchi2000}, with which one can simultaneously observe two sources separated from each other by up to a hundred times the telescope's primary beam size. With such a system, VERA can observe a Galactic maser source and an extra-galactic calibrator source at the same time and perform relative astrometry of the Galactic maser source with respect to the extra-galactic calibrator, which serves as a position reference. Simultaneous observations of two sources, rather than switching observations with single-beam telescopes, are expected to be more effective in canceling out the tropospheric fluctuations, which are the main source of position error in differential VLBI astrometry at 22 GHz or higher frequency. In fact, preliminary results from VERA's dual-beam system observations have already shown its high capability of phase-referencing to reach down to 10~microarcsecond level accuracy \\citep{honma2003}, and now monitoring of several bright maser sources is on-going to produce initial results from the VERA.  \\begin{figure}[ht] {\\includegraphics[width=0.48\\textwidth,clip]{petrov.fig1.eps} } \\caption{VERA network.} \\label{f:vera_network} \\end{figure}  In order to observe a pair of two adjacent sources with VERA, the separation angle of the pair must be less than $2.\\!{}^\\circ2$. This is not only a mechanical limit of VERA's dual-beam system, but also a practical limit for effective phase referencing to achieve 10~microarcsec accuracy, i.e., the reference source needs to be close enough to the target maser to remove the tropospheric fluctuations sufficiently. Therefore, the existence of calibrators close to target sources is crucial to successful VERA observations. If we assume a uniform density of calibrators in the sky plane, $\\sim$ 2700 calibrators are required in the whole sky to find a calibrator within $2.\\!{}^\\circ2 $ of any sky position (a circle with radius of $2.\\!{}^\\circ2 $ covers 0.0046 steradians in the sky; thus $4\\pi / 0.0046 \\approx 2700$ are necessary). Practically though, more sources are needed because the distributions over the sky is not uniform.  To date, there have been many efforts to find good calibrators observable with VLBI. The pioneering work of \\citet{first-cat} provided the first catalogue of source positions determined with VLBI, which contained 35 strong sources. In 1998 the International Celestial Reference Frame (ICRF) catalogue of 667 sources produced by analyzing VLBI observations made in the framework of space geodesy programs was published by \\citet{icrf98}. At present, the most massive survey of compact calibrators is the VLBA Calibrator Survey, hereafter VCS, performed in 1994--2005 \\citep{vcs1,vcs2,vcs3,vcs4}, and Kovalev~et~al. (submitted to \\aj,\\ 2006\\footnote{Available at {\\tt http://arxiv.org/astro-ph/0607524}}). After the fifth VCS campaign, the NASA catalogue 2005f\\_astro\\footnote{Available at {\\tt http://vlbi.gsfc.nasa.gov/astro}} of all sources detected in the absolute astrometry mode at 2.2 GHz and/or 8.6 GHz frequency bands (S and X) contains 3481 sources. Among them, 85\\% are acceptable as calibrators for the VLBA. Nevertheless, in the sky area at $\\delta \\ge -40^\\circ$, which is observable with the VERA, the probability of finding an X/S calibrator with correlated flux density over 60~mJy within $2.\\!{}^\\circ2 $ of any position is $\\sim$70\\% \\citep{vcs4}, indicating that more calibrators are needed in order to observe all maser sources with VERA. Moreover, maser sources tend to be confined to the Galactic plane, because they are basically located in the Galaxy's disk, but the density of known calibrators near the Galactic plane is even lower than that in other sky regions. This is due to the strategies used to find compact calibrators: the list of candidates was based on low resolution sky surveys that severely suffered from strong radio foreground in the Galactic plane. Besides, some past surveys, JVAS \\citep{jvas}, VCS1, VCS3, VCS4, specifically avoided the Galactic plane. To overcome this problem, a calibrator survey in the Galactic plane has been done by \\citet{honma2000b} finding more than 50 new calibrators at $|b| \\le 5^\\circ$, but still many more calibrators in this region are needed.  For these reasons, we have been performing the Galactic Plane Survey of compact radio sources, hereafter GaPS, based on the VERA and VLBA observations. The objective of this campaign is to provide a list of calibrators for dual-beam VERA observations of H$_2$O masers at 22~GHz. Our strategy for GaPS is two-fold: first, we observe a wide list of candidate sources with the VERA, and then observe a narrow list of detected sources with the VLBA in order to determine their positions at the 1--10~nrad level of accuracy and to obtain high quality maps. In this paper, we present results from the first part of the GaPS project, the fringe search observations with the VERA, reporting the detection of 180 new compact sources. In section~\\S\\ref{s:observations} we describe our approach to the selection of candidates, scheduling and observations. In section~\\S\\ref{s:anal} we describe analysis of the observations. Results are discussed in section~\\S\\ref{s:results} and concluding remarks are given in section~\\S\\ref{s:summary}. ",
        "conclusions": "\\label{s:summary}  In the 2005/2006 winter campaign $\\sim 2500$ sources were observed with the VERA during 148 hours. The target sources were either within $6^\\circ$ of the Galactic plane, or within $11^\\circ$ of the Galactic center, or within $2.\\!{}^\\circ 2$ of SiO or H$_2$O masers. Among them, 533 continuum spectrum and 16 maser sources were detected, including 180 continuum spectrum sources not previously observed with VLBI. The estimates of the correlated flux densities in the ranges [0, 70], [70, 100], and [100, 200] megawavelengths were evaluated. For 1945 non-detected sources the upper limits of their correlated flux densities were evaluated.  The distribution of detected sources over the sky is shown in figure~\\ref{f:sou_map}. All detected sources are scheduled for observations with the VLBA at 22~GHz in June-October 2006 for imaging and determination of their positions with 1--10~nanoradian accuracies. Tables of detected and non-detected sources, plots of system temperatures, fringe plots, and other auxiliary information related to this campaign can be found on the Web at {\\tt http://lacerta.gsfc.nasa.gov/vlbi/fss}.  \\begin{figure*}[ht] {\\includegraphics[width=\\textwidth,clip]{petrov.fig11.eps} } \\caption{The distribution of detected sources over the sky. Disks designate new compact source not observed before with VLBI, circles designate known sources observed before at X-band with the VLBA.} \\label{f:sou_map} \\end{figure*}  Analysis of the VERA fringe survey data shows that the probability of detecting a 100~mJy source with 120~s of integration time is about 10\\%; the probability of detecting a 200~mJy source is 60\\%; and the probability of detecting a 300~mJy source is 80\\%. Future development a next generation water vapor radiometer (Kawaguchi (2006), private communication) promises to allow increasing the integration time and, thus, to improve VERA sensitivity."
    },
    "0609/astro-ph0609627_arXiv.txt": {
        "abstract": "We address the occurrence of narrow planetary rings and some of their structural properties, in particular when the rings are shepherded. We consider the problem as Hamiltonian {\\it scattering} of a large number of non-interacting massless point particles in an effective potential. Using the existence of stable motion in scattering regions in this set up, we describe a mechanism in phase space for the occurrence of narrow rings and some consequences in their structure. We illustrate our approach with three examples. We find eccentric narrow rings displaying sharp edges, variable width and the appearance of distinct ring components (strands) which are spatially organized and entangled (braids). We discuss the relevance of our approach for narrow planetary rings. ",
        "introduction": "The accidental discovery of the rings of Uranus in 1977 by stellar-occultation measurements~\\citep{Elliot77} led to a renewed interest in planetary ring systems. The main reason was that the Uranian rings turned to be extremely different from those of Saturn: They are narrow, opaque, sharp-edged, inclined and eccentric~\\citep{Elliot84,Esposito02}. To quote some figures~\\citep{Murray99}, the widest ring of Uranus, the $\\epsilon$ ring, is 20--96 km wide with a nominal semi-major axis at 51,149 km; in comparison, the main rings of Saturn are a few thousand kilometers wide. The eccentricity of the $\\epsilon$ ring is $0.0079$; Saturn's F ring has an eccentricity $0.0026$. The structural properties have raised a number of questions, most of which remain unanswered~\\citep{Esposito02,Sicardy05}. For instance, an eccentric inclined narrow ring like the $\\epsilon$ ring is expected to spread in rather short time scales, $t_{\\rm max} = 6\\times 10^{8}$ years, which is ``considerably smaller than the age of the solar system''~\\citep{Esposito02}. This is due to interparticle collisions, drag and differential precession. This suggests an efficient confinement mechanism that maintains these structural properties of the ring over longer time scales~\\citep{Esposito02}.  To explain these structural features, new models were introduced where the confinement was induced by nearby moons. Among these models we mention in particular the shepherding model introduced by \\citet{Goldreich79}, where two moons around the ring were proposed to bound it. Soon afterwards, the Pioneer and Voyager missions detected the shepherd moons around Saturn's F ring, and later the shepherds of the outermost $\\epsilon$ ring of Uranus. The shepherding confinement involves angular momentum transfer between the shepherd moons and the ring particles, self-gravity and viscous damping due to interparticle collisions~\\citep{Borderies83}. While the full scenario for shepherding has not been fully understood~\\citep{Esposito02,Sicardy05}, the presence of dissipation seems to be essential. In addition, the formulation assumes the ring boundary located at a lower-order resonance.  While the discoveries of the shepherd moons around the F Saturn ring and the $\\epsilon$ ring of Uranus represented some confirmation for the shepherd theory, other questions still remain unanswered. For instance, most Uranian rings have no associated shepherd moons around them~\\citep{Murray90}, nor the narrow eccentric rings of Saturn, which among others would provide an explanation for their sharp-edges~\\citep{Murray99}. Thus, either the shepherds are there but are too small to be detected, or ``some physics is yet to be understood''~\\citep{Sicardy05}. Other problems are related with the effects of collisions, drag or keplerian differential velocity, which define too short life-times for the narrow rings, and cast doubts on the confinement mechanism which maintains the eccentricity and sharp edges of these rings. In addition, Saturn's F ring turned out to have a very rich dynamical structure~\\citep{Smith81,Smith82,Murray97}: besides the non-zero eccentricity, it displays multiple components entangled in a complicated way, known as strands and braids, showing further puzzling features like kinks and clumps.  Numerical simulations have investigated a variety of physical interactions, like the gravitational perturbations of shepherd moons on circular and eccentric orbits~\\citep{ShowalterBurns1982,GiuliattiWinter00}, effects due to the action of embedded moonlets~\\citep{LissauerPeale1986}, and ring interparticle collision effects~\\citep{Hanninen1993,LewisStewart2000}. The central questions investigated have been the formation of structure (strands, braids, clumps) and their short-term stability. While these studies are extremely valuable and have led to interesting predictions, e.g. the formation of channels and streamers~\\citep{GiuliattiWinter00} which have been recently observed~\\citep{Murray05}, there is no self-consistent approach for the confinement of narrow rings and their radial and azimuthal structure. The F Saturn ring remains as the most fascinating example of these possibilities.  The present paper addresses this point, namely, a self-consistent scenario for the occurrence of narrow rings and the appearance of structure. The main idea is to consider phase space regions where {\\it scattering} dominates the dynamics in systems with some {\\it intrinsic rotation}. Scattering provides a simple scenario where confinement and escape are clearly distingushed. Existence of {\\it stable} periodic orbits or tori in such scattering regions implies effective (over long times) dynamically bounded motion where otherwise there would be escaping trajectories. The intrinsic rotation, which is associated to an external potential due to the shepherds or by the full many--body problem, defines a mechanism for creating the ring: the trapped region, if it exists, rotates around the central planet, thus the confined particles form a ring. While at first sight it may appear a strange idea to focus on scattering dynamics, close approaches to shepherds illustrate this situation. In plain words, what does not escape to infinity (due to scattering) builds a ring. The bifurcation scenario yielding the stable periodic orbits or tori, and other considerations supported by theorems of Hamiltonian systems with many degrees of freedom, allow us to understand the narrowness and the appearance of strands and braids. Within this approach, which at this stage is mainly qualitative, we succeed in explaining the confinement and also some of the questions described above, such as the occurrence of narrow eccentric rings with sharp edges, apse alignment, the appearance of different ring components (strands) and their spatial entanglement (braids). The narrow rings we obtain for the examples treated are stable although they may display dynamical evolution.  To illustrate our approach we consider three examples. The first, a planar billiard system (a disk rotating on a circular or Kepler ellipse) with an arbitrary large number of non-interacting massless particles, is very simple, yet the relevant steps can be followed to a good extent analytically. So it is an interesting system that allows a pedagogical insight on many properties of the scattering approach to narrow rings. The second example in addition includes a central attractive $1/r$ interaction; it helps to establish the connection to the full gravitational problem, showing the genericity and robustness of our approach. The third example is the circular restricted three-body problem with an extremely small mass parameter. With these examples we show the occurrence of rings, the appearance of radial and azimuthal structures, and the genericity and robustness of the whole approach.  The paper is organized as follows: in the next section we describe our approach in rather general terms, avoiding formalities. In Sec.~\\ref{ex1} we illustrate our theory with the first example, a scattering billiard system. The main steps in the construction of a (narrow) ring are described in detail, discussing separately circular and elliptic motion of the billiard disk. For the elliptic case we present new results related to the appearance of structure within the rings, in particular, strands and braids. In Sec.~\\ref{ex2} we present results for the rotating disk with the addition of a central attractive $1/r$ interaction. This example permits us to understand the connection to the gravitational problem, which we exemplify with the circular restricted three-body problem in Sec.~\\ref{rtbp}. Finally, in Sec.~\\ref{concl} we present conclusions and some outlook. ",
        "conclusions": "\\label{concl}  In this paper we have presented a general and self-consistent approach for the occurrence of shepherded narrow planetary rings, and some consequences for their structural properties. Our approach is based on the local structure of phase space around stable (periodic or quasi-periodic) solutions in regions where scattering dominates the dynamics. The corresponding structure in phase space allows for a set of positive measure to exhibit dynamical trapped motion, i.e. a non-zero probability to find such trapped motion; particles with initial conditions outside this small region escape to infinity along scattering trajectories. The intrinsic rotation provides a mechanism for dynamical confinement around the central planet. Here we have considered that such rotation is induced by the perturbations of a shepherd moon, though our approach is more general since the intrinsic rotation can also be associated with interactions with other, perhaps more massive, moons, or with the combined evolution of the interacting many-body problem. A possible consequence is the explanation of the non-shepherded narrow rings of Uranus: the many-body problem is such that it defines {\\it several} regions of trapped motion, one for each specific narrow ring. An ensemble of (non-interacting) particles will form a well-defined ring. Focusing in regions where scattering dominates the dynamics allows to sharply distinguish between trapped motion related to the occurrence of the ring, and open motion manifested as regions with essentially no material due to the escape of the particles.  Within our approach, the rings will be sharp-edged, as a consequence of the scattering dynamics. The size of the stability region in the extended phase space determines how narrow the rings are. For comparatively small regions of trapped motion, the narrow ring will typically be eccentric, since the particles will be close to the organizing center, which in general displays such eccentricity. Semi-analytical estimates of the edges of the ring can be obtained by analyzing the bifurcation conditions that define the region of bounded motion in the extended phase space (within the sea of scattering trajectories) around the organizing center, at least for models with two-degrees of freedom. Typically, the bifurcation scenario that yields the regions of dynamically trapped motion corresponds to the saddle--center bifurcation, and by consequence of its fast development, the rings are narrow. We have illustrated our considerations with three distinct examples, aiming to show explicitly the robustness of the approach and its extension to and consequences for a case with many degrees of freedom, and have proved that our considerations hold also in the gravitational case as illustrated by the RTBP without physical collisions. More realistic interactions, e.g. including the oblateness of the planetary, can be easily incorporated within our theory. While this does not yet explicitly prove the connection to real narrow planetary rings, which are complicated collisional many-body systems with gravitational and other interactions, such an extension seems to us feasible, since our ideas are based on general phase space considerations, and on results and theorems which apply to Hamiltonian systems with more than two degrees of freedom.  We emphasize that our construction of the rings is self-consistent, in the sense that we do not introduce any ``cell'' or periodic boundary conditions on a segment of the ring to make the problem amenable in any sense. Therefore, radial and azimuthal structures do appear. Moreover, the theoretical framework developed is not restricted to short time scales, at least as long as the model considered is believed to be applicable for longer time scales. It is in this scope that our results may be relevant and contribute to the understanding of real planetary narrow rings: the mechanism of confinement is very efficient. We have explicitly showed which phase space structures allow for the occurrence of such narrow rings. Moreover, we have showed how rings develop further radial and azimuthal structure: rings may display variable width and many components which are entangled in a way that, at least qualitatively, resembles the observed braids of Saturn's F rings. Within our approach, the appearance of different strands and braids is not related to the ring interparticle collisions, as proposed~\\citep{Hanninen1993}, nor to the existence of other moonlets~\\citep{LissauerPeale1986}; they are related to the quasi-periodic character of the perturbing force of the rotating interaction, as induced by the shepherd's explicitly eccentric motion. This non-circular motion makes the Hamiltonian system effectively have many degrees of freedom, which complicates its analysis, but allows for other bifurcations as well as resonances. As far as we are aware, this is a new explanation for the occurrence of strands and braids, although it is somewhat in agreement with the idea of associating periodic clumps with the eccentricity of shepherd moons~\\citep{ShowalterBurns1982}. The strands and braids that we obtain display a periodic dynamical evolution, where the relevant time scale is related to the time for repeated interactions, or close approaches, with the shepherd. In our opinion, this points out the importance of the quasi-periodic perturbation, i.e. the non-circular motion of the intrinsic potential. This observation is in agreement with the results of \\citet{GiuliattiWinter00}, although these authors focus on certain perturbations induced by the repeated interactions of a group of particles with Prometheus within a cell of the ring. Whenever the organizing periodic orbit closes after more than one full turn of the secondary, the ring displays non-independent multiple components, reflecting the shape of the organizing periodic or quasi-periodic orbit in the rotating frame. Despite the fact that we have considered only planar motion, and that this is again only a qualitative comparison, the generalization of these ``higher periodic'' orbits to a non-planar case serves as an interpretation of recent observational findings~\\citep{Charnoz05}. Moreover, preliminary results~\\citep{BMpreparation} show that our theory can also yield other observed properties of narrow rings, such as arcs qualitatively similar to those of Neptune~\\citep{Namouni02}.  Our ultimate goal is to provide quantitative comparisons with observed planetary rings. The results presented here are only qualitative due to the fact that, so far, the examples are analogous to restricted three-body like problems. Actually, the most realistic example we considered is the circular (gravitational) RTBP. As we have discussed, non-collision phase-space structures of the RTBP with extremely small values of the mass parameter, which originate from the overwhelming dominance of the central planetary attraction, make such simple modeling definitely not realistic, since they yield two-body keplerian wide circular rings. We have proposed to consider the interaction of other planetary satellites and, perhaps too, other shepherds as a more realistic model. The shepherds can be included in a first step as zero-mass moonlets which break certain solutions of the many-body problem (planet including oblateness, other moons and a particle of the ring). Following the ideas of \\citet{HitzlHenon77} for the RTBP with $\\mu=0$, periodic or quasi-periodic collision solutions with respect to the shepherds can be generating orbits for the case of small masses of the shepherds. Considering the stability regions in the extended phase space when the shepherds have finite mass allows a straightforward construction of the associated rings. Work illustrating a many-body problem is in progress.  Our considerations may also be used to extend the applicability of current models to longer time scales, once it has been showed that the ring particles considered do belong to some regions of dynamically bounded motion. Conversely, they may allow to qualify certain rings or some of their structural properties as temporary phenomena, if a realistic model does not have dynamically trapped particles. An interesting question is whether chaotic motion of the shepherds may still yield long-lived rings within our framework, that is, whether they allow for bounded motion, and what the consequences are for their structure. This is related to recent investigations where the motion of Prometheus and Pandora seems to be chaotic~\\citep{Goldreich03a,Goldreich03b,Renner05}.  The strongest assumption within our approach is the non-interacting character of the ring particles. This essentially allows us to separate the many-body problem into a collection of independent one-particle problems. This assumption must be taken as a first approximation in the series expansion of the full many-body Hamiltonian in terms of the mass of the ring particles. Here, the particle--particle interactions among the ring particles will be orders of magnitude smaller than other interactions, except maybe when they collide. It is precisely the inelastic collisions among the ring particles what we neglect. A possibility to include these collisions in a simple way is with a random-walk model, taking place in all the coordinates of the full phase space of the ring particles. Such a model could also be used to describe dissipative effects like drag, depending on the particular implementation. This random-walk approach would certainly introduce an erosion mechanism of the ring, i.e., the possibility of escape from the ring and a typical life-time; yet it will preserve the properties that we discuss, such as the sharp-edges, the eccentricity and the many components of the ring. This is because, even when collisions are included, once a particle is outside the region of trapped motion it escapes along a scattering trajectory. Work in this direction is also in progress."
    },
    "0609/astro-ph0609411_arXiv.txt": {
        "abstract": "The power spectrum of the two-degree Field Galaxy Redshift Survey (2dFGRS) sample is estimated with the discrete wavelet transform (DWT) method. The DWT power spectra within $0.04 <k< 2.3 h$Mpc$^{-1}$ are measured for three volume-limited samples defined in connective absolute magnitude bins $-19 \\sim -18$, $-20 \\sim -19$ and $-21 \\sim -20$. We show that the DWT power spectrum can effectively distinguish $\\Lambda$CDM models of $\\sigma_8=0.84$ and $\\sigma_8=0.74$. We adopt maximum likelihood method to perform three-parameter fitting with bias parameter $b$, pairwise velocity dispersion $\\sigma_{pv}$ and redshift distortion parameter $\\beta=\\Omega_m^{0.6}/b$ to the measured DWT power spectrum. Fitting results denotes that in a $\\sigma_8=0.84$ universe the best fitted $\\Omega_m$ given by the three samples are consistent in the range $0.28 \\sim 0.36$, and the best fitted $\\sigma_{pv}$ are $398^{+35}_{-27}$, $475^{+37}_{-29}$ and $550 \\pm 20$km/s for the three samples, respectively. However in the model of $\\sigma_8=0.74$, our three samples give very different values of $\\Omega_m$. We repeat the fitting by using empirical formula of redshift distortion. The result of the model of low $\\sigma_8$ is still poor, especially, one of the best value $\\sigma_{pv}$ is as large as $10^3$km/s. The power spectrum of 2dFGRS seems in disfavor of models with low amplitude of density fluctuations. ",
        "introduction": "The present clumpy structures indicated by galaxies on large scales are evolved from the very small density fluctuations in the early era of the universe. The amplitude of the fluctuation is fundamental to understand the structure formation. A remarkable success of modern cosmology is that the amplitude of mass fluctuations detected by the anisotropy of cosmic microwave background radiation is in excellent agreement with the analysis of the galaxy clustering at low redshifts. Recently released WMAP third year data (WMAP3) refines most results of cosmological parameters given by the WMAP 1st year data. However, the fluctuation amplitude smoothed in a spherical top hat window of radius of $8\\rm{h^{-1} Mpc}$ is found as small as $\\sigma_8=0.74 ^{+0.05}_{-0.06}$ \\citep{Spergel2006}, which is significantly lower than $\\sigma_8=0.84\\pm 0.04$ of the WMAP 1st year data. The new result of $\\sigma_8$ is a challenge to the cosmological parameter determinations from samples of galaxies and galaxy clusters, most of which yield $\\sigma_8 \\simeq 0.9-1$ if the matter content of the universe $\\Omega_m \\leq 0.3$ \\citep[e.g.][]{2002ApJ...567..716R, 2002ApJ...577..595H, 2002ApJ...572L.131R,2002A&A...393..369V,2003MNRAS.344..673B, 2003ApJ...588L...1B, 2005PhRvD..71j3515S,2006MNRAS.365..231V}.  The problem motivates us to revisit the constrains on $\\sigma_8$ from the power spectrum of the sample of the two-degree Field Galaxy Redshift Survey (2dFGRS). The final released spectroscopic catalog of the 2dFGRS contains 221414 galaxies with good redshift quality $Q\\geq 3$ and covers approximately 1800 square degrees of the sky. It is a good sample for studying the fluctuations of cosmic mass field on large scales in linear regime as well as on scales in nonlinear range. Moreover, the 2dFGRS team has made detailed analysis of the Fourier power spectrum, the two point correlation functions and relevant cosmological parameter fitting \\citep{2001MNRAS,2001MNRAS.328...64N, 2001Natur.410..169P, 2001MNRAS.328...64N, 2003MNRAS.346...78H, 2004MNRAS.353.1201P, 2005MNRAS.362..505C}. They found that $\\Omega_m \\sim 0.3$ or less \\citep{2001Natur.410..169P,2005MNRAS.362..505C}, and the best fitting value of $\\sigma_8$ is $\\sim 0.95$ if one takes $\\Omega_m \\sim 0.3$ \\citep{2004MNRAS.353.1201P}, which is substantially different from WMAP3.  In the linear regime $\\sigma_8$ only controls the overall amplitude of the power spectrum, and it is degenerated with the linear bias parameter $b$. Power spectrum on small scales is more effective to constrain $\\sigma_8$ than on large scales since the nonlinearity of the power spectrum is directly reflected by the value of $\\sigma_8$. Most measurements of the 2dFGRS power spectrum are on scales of $k < 1h$Mpc$^{-1}$ while cosmological parameter estimation is performed on scale of $k<\\sim 0.2h$Mpc$^{-1}$ \\citep{2001Natur.410..169P,2001MNRAS, 2002MNRAS.335..887T,2005MNRAS.362..505C}. With the estimator based on the discrete wavelet transformation (DWT) \\citep{2001PABei..19S..37Y, YangEtal2001a, 2001ApJ...560..549Y, 2002ApJ...566..630Y}, we are to analyze the power spectrum of the 2dFGRS sample on scales down to $k \\simeq 2h$Mpc$^{-1}$. In this scale range, the redshift distortion of DWT diagonal mode power spectrum can be easily approximated and the aliasing effect is exactly eliminated by the DWT algorithm \\citep{2000ApJ...539....5F}.  The paper is organized as follows. In Section 2, the DWT power spectrum estimator is briefly introduced. Section 3 describes the construction of samples. The Section 4 demonstrates robustness and accuracy tests on the DWT power spectrum estimator. Section 5 lists our fitting results. Conclusions are stated in Section 6. ",
        "conclusions": "DWT power spectra of 2dFGRS samples are measured on scales equivalent to $0.04 <k< 2.3h$Mpc$^{-1}$. We show that these power spectra are efficient to test $\\Lambda$CDM models with high and low amplitudes of mass density fluctuations. The model with $\\sigma_8=0.84$ finds good support from the 2dFGRS sample, all the best fitting parameters, $b$, $\\beta$, $\\sigma_{pv}$ are consistent with other works on 2dFGRS. Especially, three volume-limited samples gives different $b$ and $\\beta$, but approximately the same values of $\\Omega_m=(b\\beta)^{1/0.6}= 0.28-0.36$. On the other hand, the model with $\\sigma_8=0.74$ cannot give such consistent fitting result, the best fitted $\\Omega_m$ from the three volume-limited samples are significantly different, and in deviation from $\\sim0.3$ at least at $1\\sigma$ level. Moreover, the fitting results of $\\sigma_{pv}$ are generically large, even reach $10^3$ km s$^{-1}$. Our studies suggest that the power spectrum of 2dFGRS disfavors models with low amplitude of mass fluctuations, $\\sigma_8=0.74$, if other cosmological parameters are given by the WMAP3.  It is found that $\\sigma_{pv}$ increases with luminosity, which is basically consistent with the observation of \\citet{2004ApJ...617..782J} though our estimations are lower than theirs.  Another parameter that will affect the shape of the DWT power spectrum is the slope of the primordial fluctuation spectrum $n_s$. We have been using the scale-invariant spectrum, or the Zeldovich spectrum with $n_s=1$. However, we notice that WMAP3 gives $n_s=0.95^{+0.015}_{-0.019}$, which is 5\\% smaller. In order to check the influence of a lower value of $n_s$ on our fitting results, we repeat our fitting for the sub-sample $-20 \\sim -21$ with $n_s=0.95$. We find $b=1.07^{+0.01}_{-0.01}$, $\\beta=0.44^{+0.02}_{-0.02}$ ($\\Omega_m=0.29^{+0.02}_{-0.02}$), and $\\sigma_{pv}=550^{+20}_{-20}$km s$^{-1}$. Clearly, the change of $n_s$ from 1 to 0.95 results in little modification to $b$ and $\\sigma_{pv}$, only slightly alters (less than 5\\%) the value of $\\beta$, or $\\Omega_m$.  In this paper, we used only the diagonal modes in term of vector ${\\bf j}$, $P_{\\bf j}=\\langle \\tilde{\\epsilon}_{\\bf j,l} \\tilde{\\epsilon}_{\\bf j,l}\\rangle$, and ${\\bf j}=(j,j,j)$. Even in the second order statistics of the DWT, $\\langle \\tilde{\\epsilon}_{\\bf j,l}\\tilde{\\epsilon}_{\\bf j,l'}\\rangle$, we can have the power of off-diagonal modes $P_{\\bf j}=\\langle \\tilde{\\epsilon}_{\\bf j,l} \\tilde{\\epsilon}_{\\bf j,l}\\rangle$, and ${\\bf j}=(j_1,j_2,j_3)$. In addition, we have the correlation between modes of $({\\bf j,l})$ and (${\\bf j,l'})$ (${\\bf l \\neq l'}$). It has been shown that different parts of the second order statistics of DWT contains different information of the random field \\citep{2001PABei..19S..37Y,YangEtal2001a}. Possible constrains on parameters given by various parts of the second order DWT statistics deserve further study."
    },
    "0609/astro-ph0609503.txt": {
        "abstract": "Identifying the two physical mechanisms behind the production and sustenance of the quiescent solar corona and solar wind poses two of the outstanding problems in solar physics today. We present analysis of spectroscopic observations from the {\\em Solar and Heliospheric Observatory} that are consistent with a single physical mechanism being responsible for a significant portion of the heat supplied to the lower solar corona and the initial acceleration of the solar wind; the ubiquitous action of magnetoconvection-driven reprocessing and exchange reconnection of the Sun's magnetic field on the supergranular scale. We deduce that while the net magnetic flux on the scale of a supergranule controls the injection rate of mass and energy into the transition region plasma it is the global magnetic topology of the plasma that dictates whether the released ejecta provides thermal input to the quiet solar corona or becomes a tributary that feeds the solar wind. ",
        "introduction": "The relentless, convection-driven reprocessing of magnetic fields on the Sun has been dubbed the ``Magnetic Carpet'' \\citep[e.g.,][]{Schrijver1997, Hagenaar1997, Title1998, Hagenaar1999}. Detailed observations of the Sun's photospheric magnetic field have been used to deduce that the turbulent motion of the ionized gas in the Solar interior causes small bipolar magnetic flux elements to be carried towards the photosphere. There, through their annihilation, they contribute the global magnetic structure of the outer solar atmosphere and are a source of energy to the corona and solar wind \\citep[][]{Schrijver1998, PriestSchrijver1999, Priest2002}.  Numerical simulations of magnetoconvection \\citep[][]{Cattaneo2003} and tectonic models of the magnetic carpet \\citep[][]{Priest2002} show that the emerging bipolar elements are advected towards the boundaries of large convective cells with typical diameters of 15-35Mm \\citep[e.g.,][]{Hagenaar1997}. These large supergranular convection cells are relatively easy to observe in the solar photosphere and chromosphere, and comprise the ``chromospheric network''. Over periods of many hours to days, the constant convective dredging of the magnetic field imposes a net magnetic polarity on a supergranule. This determines the polarity of the emerging magnetic dipoles which are annihilated through magnetic reconnection \\citep[e.g.,][]{Parker1988, Parker1994, PriestForbes2000}. It has been postulated that energy stored in, and eventually released by, these small annihilated magnetic dipoles is the simplest and most efficient means of providing the 100 W/m$^{2}$ required to produce the Sun's ambient 2x$10^{6}$K corona \\citep[][]{PriestSchrijver1999, Priest2002, Parker1994} and to load and accelerate mass into the solar wind \\citep[][]{PriestForbes2000,Wang1998,Wang2004}.  Recently, observations from the Solar Ultraviolet Measurement of Emitted Radiation \\citep[SUMER;][]{Wilhelm+others1995}, the Extreme-ultraviolet Imaging Telescope \\citep[EIT;][]{Boudine1995} and the Michelson Doppler Imager \\citep[MDI;][]{Scherrer1995} instruments on the Solar and Heliospheric Observatory \\citep[\\soho;][]{Fleck+others1995} have provided insight into the origins of the fast solar wind in coronal holes \\citep[][]{Hassler+others1999, Dammasch+others1999, Xia2003, Xia2004, Tu2005a, Tu2005b}. SUMER ``spectroheliogram'', or raster, observations were employed to quantitatively examine the origins of the solar wind by correlating Doppler velocity measurements of \\ion{Ne}{8} \\citep[formed at about 600,000K in the upper solar transition region;][]{Mazzotta1998} to a proxy for the gross super-granulation pattern of chromospheric network structure: bright \\ion{Si}{2} emission (formed at about 10,000K).  Results of these prior studies suggested a clear correlation between the observed \\ion{Ne}{8} blue shift\\footnote{In this paper we will use the convention that a blue shift is associated with a negative velocity and as such a blue Doppler shift indicates the presence of plasma moving towards the observer, possibly indicating an out-flow from the Sun.} and the chromospheric network pattern seen in \\ion{Si}{2}. We show that while there is indeed a strong correlation between bright \\ion{Si}{2} emission and \\ion{Ne}{8} blue Doppler-shifts, it is limited to regions that occur at the intersections of three or more supergranules (see, e.g., Fig.~\\pref{fig3}). We call these locations ``bright network vertices''.  We demonstrate that a substantial fraction of the blue shifted \\ion{Ne}{8} plasma in coronal holes is not correlated with the bright \\ion{Si}{2} emission that forms the supergranular boundaries or bright network vertices, consistent with the recent analysis of \\citet[][]{Aiouaz2005}. We investigate the remaining portion of the blue shifted \\ion{Ne}{8} plasma in the coronal hole to examine the mechanism behind the energetic processes taking place in the outer solar atmosphere.  We find that a critical - and largely overlooked - step towards understanding the excess \\ion{Ne}{8} Doppler shift lies in the interpretation of the \\ion{C}{4} emission line, formed in the transition region at roughly 100,000K between \\ion{Si}{2} and \\ion{Ne}{8}. Data available at this thermodynamic mid-point allows us to physically couple the spatial variation of the lower temperature \\ion{Si}{2} and higher temperature \\ion{Ne}{8} plasmas to form an entirely new interpretation of these well studied data.  We present the self-consistent analysis, comparison, and interpretation of the results derived from spectroscopic observations of coronal hole and quiet Sun regions which point to the magnetic carpet's constant recycling of magnetic flux as a significant energy supply mechanism to the outer solar atmosphere. Indeed, we deduce that the difference in the form of the energy provided to the plasma stems, rather simply, from whether or not its magnetic topology is ``open'' (coronal hole; largely kinetic) or ``closed'' (quiet Sun; largely thermal) to interplanetary space. ",
        "conclusions": "% Together, the \\soho{} (SUMER, EIT and MDI) observations, statistical relationships and cartoon representations lead us to the conclusion that the observed Doppler velocity and emission patterning of the upper transition region and low solar corona is consistent with the action of convection-driven magnetic field emergence and reconnection: the magnetic carpet \\citep[][]{Schrijver1997, Schrijver1998, PriestSchrijver1999, Priest2002, PriestForbes2000, Wang1998}.  We have demonstrated that the driven reconnection events largely neighbor the supergranular boundaries and propose that they contain plasma that is loaded and driven by the magnetic carpet's relentless stirring and destruction of the injected and advected magnetic flux \\citep[][]{Schrijver2003}. We have deduced that while the net magnetic flux on the scale of a supergranule controls the injection rate of mass and energy into the transition region plasma it is the global magnetic topology of the plasma that dictates whether the released ejecta provides thermal input to the quiet solar corona or becomes a tributary to the solar wind.  The magnetoconvection-driven reconnection and resulting ejecta that we observe and have discussed has significant impact on the energetics of the outer solar atmosphere. The same must be true for stellar atmospheres for which dynamo-driven, cyclic magnetic field behavior has been observed in their optical/UV/EUV radiative output. The hypothesis presented in this Paper can be directly tested by the instruments on the upcoming Solar-B satellite and should provide motivating science for the Solar Probe and Solar Orbiter Missions currently in the pre-proposal phase."
    },
    "0609/astro-ph0609405_arXiv.txt": {
        "abstract": "We present a single epoch of high signal-to-noise ratio spectropolarimetry of the Type Ia supernova (SN Ia) 2004S taken nine days after maximum light.  The flux spectrum is normal, but with the additional presence of high-velocity (HV) line features in both \\ion{Ca}{2} and \\ion{Fe}{2}.  The object shows continuum polarization at the 0.4\\% level in the red, a value which appears to be typical of SNe Ia.  The continuum data are consistent with a $\\sim$10\\% global asphericity in an axisymmetric geometry.  Unlike previous observations of other SNe Ia with HV features, the HV features in SN~2004S show no strong polarimetric signature, though this may be due to the timing of our observations.  Instead, the object shows line polarization features ($P\\leq0.5$\\%) that are rotated with respect to the axis of symmetry of the continuum.  The line features are visible in \\ion{Si}{2}, \\ion{Fe}{2}, and \\ion{Ca}{2}, and appear to be narrowly confined in velocity space just above the photosphere. These polarization features are a result of compositional inhomogeneities in the ejecta.  They may represent newly synthesized elements whose clumpy spatial distribution within the ejecta is distinct from that of the globally aspherical ejecta as a whole. ",
        "introduction": "Carbon-oxygen white dwarfs that accrete enough mass to approach the Chandrasekhar limit can explode as Type Ia supernovae (SNe~Ia).  The first successful models of the explosion process invoked parameterized one-dimensional (1D) flame fronts propagating outward from the center of the star as a subsonic deflagration \\citep{nom84}. However, the explosion process is fundamentally three-dimensional (3D) in nature as the propagation of a deflagration front is marked by large-scale fluid motions and mixing due to Rayleigh-Taylor instabilities.  Recent 3D pure deflagration models are not as successful at reproducing the observations because the simulated explosions are underluminous, underenergetic, inefficient in burning the available fuel, and result in ejecta that are too well-mixed \\citep{rhn02,gam03,koz05}.  Some modelers interpret these results as requiring the burning front to make a transition from a subsonic deflagration to a supersonic detonation as it becomes turbulent \\citep{kh91}, though the details of how this might happen are unclear \\citep{hn00}. Simulations, such as those of \\citet{gam05}, demonstrate that models with different prescriptions for the evolution of the burning front produce quite different ejecta structures.  Observational probes of the geometry and structure of the ejecta are therefore necessary in light of the lack of an \\emph{a priori} theoretical understanding of the explosion.  The deviations from spherical symmetry seen in the simulations discussed above are a result of the ignition process and the random physics of turbulent flame propagation.  More systematic departures from sphericity could be a signature of the progenitor systems, which are largely unknown \\citep{br95,liv01,nom03}, or their environments. Rotational flattening of the progenitor white dwarf, interaction with circumstellar material, or the process of merging two white dwarfs could lead to SNe~Ia with large-scale asphericities \\citep{ho01,wa03}.  Spectropolarimetry provides a unique opportunity to probe the geometry of SNe~Ia.  Outward streaming radiation in a supernova atmosphere acquires a polarization when it is scattered by electrons.  Distant supernovae are unresolved by the observer, so the net polarization of the received light represents an integration of emitted flux over the visible surface of the object.  In the case of a spherical supernova, symmetry demands that the total polarization be zero.  Non-zero net polarization is therefore a sign of some deviation from spherical symmetry in the system \\citep{ss82}.  Global asphericities of the photosphere, such as ellipsoidal deformations, are manifested as net continuum polarization \\citep{je91,ho91}. Past investigations of SNe~Ia have found typical intrinsic polarizations of order a few tenths of a percent in the continuum. The object with the highest measured intrinsic continuum polarization percentage is still the subluminous SN~1999by \\citep{ho01}, which showed at least 0.7\\% polarization.  \\citet{ho01} found that the spectropolarimetric data could be modeled by an oblate spheroid with an axial length ratio of 1.17, seen nearly equator-on.  The lower continuum polarizations measured in the half-dozen other SNe~Ia with good data are consistent with global asphericities at the $\\sim$10\\% level \\citep{wa03,wa06,leo05,me06}.  Line polarization allows other approaches to the study of asphericity. Polarization modulations and angle rotations across a spectral feature could be a sign of an aspherical distribution of that ionic species.  \\citet{leo05} found that the angle rotations seen in spectral features of some SNe~Ia could be due to clumps or ionization asymmetries in the ejecta.  Even if the ejecta lack such compositional inhomogeneities, line polarization can give insight into the global asphericity of the ejecta.  Photons scattered in a line are generally depolarized \\citep{ho96}, but in a purely ellipsoidal geometry we should expect strong, isolated lines to be associated with ``inverted'' P-Cygni polarization line profiles \\citep{jef89}.  The effect of many overlapping lines is to produce a decrease in the polarization at the same angle as the continuum, as was seen in SN~1999by \\citep{ho01}.  These signatures of asphericity are largely independent of the uncertainties in the removal of interstellar polarization (ISP) that plague estimates of the continuum polarization.  One exciting recent development in SN~Ia research is the association of ``high-velocity'' (HV) features in the flux spectra with strong spectropolarimetric features.  These HV features are so-named because they are signatures of a separate line-formation region at higher velocities than the normal photospheric component.  Separate HV material was first noticed in the \\ion{Fe}{2} lines of SN~1994D by \\citet{ha99}, but it may be ubiquitous in SNe~Ia.  Every object in the sample of \\citet{maz05b} showed evidence of HV absorption in the \\ion{Ca}{2} near-infrared (NIR) triplet if the spectra were taken at sufficiently  early epochs.  The HV \\ion{Ca}{2} NIR triplet feature in the otherwise-normal SN~2001el was responsible for a 0.7\\% spectropolarimetric feature that was rotated with respect to the continuum \\citep{wa03,ka03}.  This feature was modeled as forming in a spatially separate region that did not share an axis of symmetry with the continuum \\citep{ka03}.  Subsequent observations of the SNe~Ia 2004dt and 2002bf \\citep{leo05,wa06} also showed strong polarization modulations across the HV lines.  The composition and origin of this HV material remain mysterious.  Suggestions in the literature have included clumps of material newly synthesized in the explosion \\citep{wa03}, signatures of a gravitationally confined detonation \\citep{kp05}, and interaction with circumstellar material \\citep{wa03,ger04,maz05a}.  In all cases, viable models for the origin of the HV material must recognize its intrinsically 3D nature, as revealed by the spectropolarimetric line features and angle rotations.  Inverting the observed polarization data for one object to construct a model is highly underconstrained.  A more fruitful approach is to collect data for a sample of objects to see which, if any, model can best match the range of observed polarimetric behaviors.  Though observations are still sparse, \\citet{leo05} have shown that SNe~Ia exhibit diverse polarimetric characteristics that may correlate with other photometric and spectroscopic properties.  For example, the two SNe~Ia with the highest observed intrinsic continuum polarizations are both subluminous, and the objects with HV features have shown the largest line polarizations.  In addition, \\citet{wa07} have claimed to find a correlation between the strength of the spectropolarimetric modulation across the \\ion{Si}{2} $\\lambda$6355 feature and the light-curve shape (a luminosity indicator) for a sample of SNe~Ia.  In this paper we add to the small sample of SNe~Ia observed spectropolarimetrically and present a single epoch of high signal-to-noise-ratio (S/N) observations of \\sn, a SN Ia with HV features that is a virtual twin of SN~2001el.  The data were taken 9~d after maximum light and the observed continuum polarization provides evidence for asphericity in the ejecta. Deviations from axisymmetry, seen as spectropolarimetric modulations across prominent spectral lines, are the main result and focus of the discussion below. ",
        "conclusions": "We have presented a single epoch of high S/N spectropolarimetry of the Type Ia \\sn, which had an intrinsic continuum polarization of 0.4\\%.  This observation adds to a growing body of evidence that normal SNe~Ia have global asphericities of order 10--20\\%.  The data points from the continuum are consistent with an axisymmetric geometry such as an ellipsoid.  Polarization angle modulations across spectral features due to \\ion{Si}{2}, \\ion{Ca}{2}, and \\ion{Fe}{2} are evidence that the distributions of some elements within the supernova ejecta differ from that of the continuum.  One natural interpretation of these data is that we are seeing evidence for clumps of newly synthesized elements.  These clumps could be a diagnostic of turbulent burning during the explosion process (e.g., Gamezo et al. 2005).  Previous work has shown non-axisymmetric structures in HV material \\citep{ka03}, but here the deviations from axisymmetry are present close to the photosphere.  We argued that the clumps in the ejecta are sufficiently narrowly confined in velocity space to prevent the \\ion{Ca}{2} NIR triplet polarization features from being blended together.  This conclusion could be tested with a time series of spectropolarimetric observations of a future SN~Ia starting well before maximum light.  At early times these clumps might be hidden because the lines are optically thick at higher velocities.  As the absorption minima in the flux spectra recede through the ejecta, the clumps could become visible polarimetrically at later times.  In particular, the spectroscopic sequence of \\sn\\ presented by \\citet{kk06} shows that at early times the flux minimum of the \\ion{Ca}{2} NIR triplet line was at higher velocity than the polarimetric features we see at this later epoch.  Spectropolarimetric data of the similar SN~2001el before maximum light \\citep{wa03} show no features at photospheric velocities in the \\ion{Ca}{2} NIR triplet (which has a strong HV component), even though photospheric polarimetric features are visible in the \\ion{Si}{2} $\\lambda$6355 line (which does not).  A time series of spectropolarimetric data would also be useful in understanding the HV material.  \\sn\\ is the first SN~Ia observed polarimetrically with prominent HV features in the total-flux spectrum that lacked counterparts in the polarization.  We argued above that this was likely a consequence of the timing of our observations, as the HV polarization features in the similar SN~2001el also disappeared a couple of weeks after maximum light \\citep{wa03}.  The HV material must decrease in optical depth as it expands, so determining the epoch at which its effects disappear in both the total-flux spectrum and the polarization data has the potential to constrain the optical depth or covering fraction of the HV material, key parameters in any model.  Alternatively, we could be viewing the HV material from a special orientation, such as along an axis of symmetry.  The deviations from axisymmetry in \\sn\\ are responsible for subtle variations in the polarization across the profiles of some spectral lines and would have been unobservable at lower S/N. Spectropolarimetric observations have the potential to be a powerful probe of SN~Ia explosion models; however, quantitative conclusions will have to wait for detailed radiative transfer calculations of current 3D explosion models (e.g., Kasen et al. 2006).  In addition, interpretation of the continuum observations is hampered by uncertainties in the removal of ISP, even for objects like \\sn\\ suffering small extinction.  Progress can be made with observations of a larger set of objects, preferably behind minimal dust columns."
    },
    "0609/astro-ph0609319_arXiv.txt": {
        "abstract": "We present a systematic analysis of all the {\\rm BeppoSAX} data of SGR1900+14. The observations spanning five years show that the source was brighter than usual on two occasions: $\\sim$20 days after the August 1998 giant flare and during the 10$^5$ s long X--ray afterglow following the April 2001 intermediate flare. In the latter case, we explore the possibility of describing the observed short term spectral evolution only with a change of the temperature of the blackbody component. In the only {\\rm BeppoSAX} observation performed before the giant flare, the spectrum of the SGR1900+14 persistent emission was significantly harder and detected also above \\mbox{10 keV} with the PDS instrument. In the last {\\rm BeppoSAX} observation (April 2002) the flux was at least a factor 1.2 below the historical level, suggesting that the source was entering a quiescent period. ",
        "introduction": "\\label{intro} The soft gamma-ray repeater (SGR) \\mbox{SGR\\,1900+14} was discovered in 1979 through a series of short and soft gamma-ray bursts  \\citep{mazets79}. Many years later, its persistent pulsating X--ray counterpart was discovered in the \\mbox{2--10 keV} energy band \\citep{hur99}. More recently, it was also detected in the hard X--ray range (20--100 keV) with the INTEGRAL satellite, becoming the second SGR with a measurable hard X--ray tail in the spectrum of its persistent emission \\citep{gotz06}.  The bursting activity of \\mbox{SGR\\,1900+14} is rather discontinuous (see Figure \\ref{fig:1}, bottom panel) and culminated on 1998 August 27 with the emission of a Giant Flare, when more than 10$^{44}$ ergs of $\\gamma$--rays were emitted in less than one second \\citep{hurley99gf,maz99,feroci01gf}. This was one of the three Giant Flares detected up to now from three different SGRs and was interpreted as a new evidence in favor of the {\\it magnetar} model. In this model \\citep{thompson95,thompson96}, the SGRs and the \\linebreak Anomalous X--ray Pulsars (AXPs, another class of X--ray sources with similar properties, \\citealt{mereghetti02}) are believed to be neutron stars powered by the decay of their extremely intense magnetic field (\\mbox{$B\\sim10^{14}-10^{15}$ G}).  Here we present the analysis of the persistent emission of \\mbox{SGR\\,1900+14} both in the soft and hard X--ray range and its evolution across the Giant Flare and in relation to its bursting activity. ",
        "conclusions": "\\label{sec:3} We have studied the variability of \\mbox{SGR\\,1900+14}, both in the hard and in the soft X--ray range, finding the following results:  \\begin{itemize}  \\item Except for the observations immediately following exceptional flares, the flux level in soft X--rays was stable while the source was moderately active and progressively decreased when it entered a 3 years long quiescent period.  \\item The Intermediate Flare of 2001 April 18 was followed by an X--ray afterglow that can be successfully interpreted as due to the heating of a significant fraction of the neutron star surface, that then cools down in $\\sim$1 day. This is consistent with the interpretation of similar events in other magnetar candidates \\citep{woods2004}.  \\item The soft X--ray spectrum during the only available pre-flare observation was harder than in the following quiescent observations. This is similar to what observed in SGR~1806--20, the only other SGR that could be monitored before and after a Giant Flare \\citep{mte05,rea05,tiengo05}.  \\item Comparing the hard X-ray spectrum of \\mbox{SGR\\,1900+14} recently observed with INTEGRAL to that observed with the PDS instrument in 1997, we find evidence for variations in flux and spectral slope.  \\item The reduction of the X--ray tail in coincidence with the Giant Flare is supported by the count rates detected in the PDS instrument above 50 keV during the different {\\rm BeppoSAX} observations, that indicate how they significantly decreased already in the first post-flare observation. Since the hard X--ray tail in the spectrum of \\linebreak \\mbox{SGR\\,1900+14} might contain most of its total emitted energy, its variability in relation to the bursting activity is a key point to try to understand the SGR emission processes.   \\end{itemize}"
    },
    "0609/astro-ph0609633_arXiv.txt": {
        "abstract": "We show that  isotropic component of circular polarization of stochastic  gravitational wave background can be explored by breaking two dimensional configuration of multiple laser interferometers for correlation analysis. By appropriately selecting orbital parameters for the proposed BBO mission, the circular polarization degree $\\Pi$ can be measured down to $\\Pi\\sim 0.08 (\\Omega_{GW}/10^{-15})^{-1}(SNR/5)$ with slightly ($\\sim 10\\%$) sacrificing  the detection limit for  the total intensity $\\Omega_{GW}$  compared to the standard plane symmetric configuration.   This might allow us to detect signature of parity violation in the very early universe. ",
        "introduction": "Due to extreme penetrating power of gravitational waves, observation of the waves may enable us to study the very early universe in a way that cannot be achieved with other  methods \\cite{Allen:1996vm}. Recently, follow-on missions to the Laser Interferometer Space Antenna (LISA) \\cite{lisa} have been actively discussed to directly detect stochastic gravitational wave background from the early universe around $\\sim 1$Hz. The primary target for the proposed missions, such as, the Big Bang Observer (BBO) \\cite{bbo} and DECIGO \\cite{Seto:2001qf}, is the gravitational wave background produced at inflation. With  various observational supports for inflation, existence of the background is plausible, while it is currently difficult to predict its amplitude.  Meanwhile, considering the facts that observation of gravitational waves will be a truly new frontier of  cosmology and our understanding of physics is limited at very high energy scale, it would be quite meaningful to prepare flexibly to various observational possibilities. Actually,   detection of  something unexpected with odd properties will be generally more exciting than confirmation of something widely anticipated.  For this purpose it is preferable that we can study gravitational wave background  beyond its simple spectral information, and model independent approach  would be effective with regard to the sources of the background.   One of  fundamental as well as interesting properties of the background is its circular polarization. Circular polarization characterizes asymmetry of  amplitudes of  right- and left-handed waves, and its generation is expected to be closely related to parity violation (see {\\it e.g.} \\cite{Lue:1998mq}). In a recent paper \\cite{p1} it was shown that LISA can measure the dipole ($l=1$) and octupole ($l=3$) anisotropic patterns of circular polarization of the background in a relatively clean manner, but, at the same time, LISA cannot capture  its monopole ($l=0$) mode due to a symmetric reason. Since our observed universe is highly homogeneous and isotropic at cosmological scales, it would be essential to have sensitivity to the monopole mode of circular polarization of cosmological background. The  proposed missions like BBO or DECIGO are planed to use multiple sets of interferometers to perform correlation analysis to measure the total intensity $\\Omega_{GW}$ of cosmological  background by beating out  detector noises with a long-term signal integration. The standard configuration of these missions is to put the multiple interferometers on a same plane. This is advantageous to get a good sensitivity to the total intensity $\\Omega_{GW}$, as a larger spatial separation between interferometers results in reducing their overlapped responses to gravitational waves \\cite{Flanagan:1993ix}. However, with this plane symmetric configuration, we are totally blind to the monopole mode of  circular polarization, as in the case of LISA. This means that even if the isotropic background is circularly polarized by  100\\%, we will not be able to discriminate its extreme nature. In this paper we study how well we can detect the monopole mode of circular polarization by breaking symmetry of the detector configuration. ",
        "conclusions": ""
    },
    "0609/astro-ph0609069_arXiv.txt": {
        "abstract": "We present results from cryogenic tests of Volume-Phase Holographic (VPH) gratings at $\\sim$ 100 $K$. The aims of these tests are to see whether the diffraction efficiency as a function of wavelength is significantly different at a low temperature from that at room temperature and to see how the performance of a VPH grating is affected by a number of thermal cycles. We have completed 10 cycles between room temperature and 100 $K$, and find no clear evidence that the diffraction efficiency changes with temperature or with successive thermal cycle. ",
        "introduction": "Volume-Phase Holographic (VPH) gratings potentially have many advantages over classical surface-relief gratings (Barden, Arns \\& Colburn 1998; Barden et al. 2000). They are already in operation in some existing astronomical spectrographs (Kashiwagi et al. 2004) and their use is also planned for a number of forthcoming instruments (Smith et al. 2004). While the main applications of VPH gratings are currently for optical spectrographs, they will also be useful for near-infrared (NIR) spectrographs if the performance at low temperatures is satisfactory. In particular, diffraction efficiency and its dependency on wavelength should be confirmed. Contraction of dichromated gelatin with decreasing temperature could cause variations in line density and the diffraction efficiency profile (the thickness of the gelatin layer is one of the parameters that define diffraction efficiency). Since thermal cycling may cause some deterioration of a VPH grating and reduce its operational life time, we also need to investigate whether the characteristics vary with successive thermal cycles. Previously, we tested a VPH grating at 200 $K$ and confirmed that the performance is nearly independent of temperature during 5 thermal cycles (Tamura et al. 2003, 2004). While cooling to 200 $K$ can be sufficient for a spectrograph operating at wavelengths out to $\\sim$ 1.8 $\\mu$m ($H$ band) such as Fiber Multi Object Spectrograph (FMOS; e.g. Kimura et al. 2003), a much lower temperature (e.g., 100 $K$) is required to extend the spectral coverage of a spectrograph out to $\\sim$ 2.4 $\\mu$m ($K$ band).  In this paper, measurements of diffraction efficiency of VPH gratings at 100 $K$ and at room temperature are reported. Pictures of the gratings investigated are shown in Fig. \\ref{vphpics}, both of which were manufactured by Ralcon Development Lab. The one in the left panel of Fig. \\ref{vphpics} has a line density of 385 lines/mm and the peak of diffraction efficiency exists around 1.3 $\\mu$m at the Bragg condition when the incident angle of an input beam to the normal of the grating surface is 15$^{\\circ}$. The measurements of this grating are performed at wavelengths from 0.9 $\\mu$m to 1.6 $\\mu$m. Since it is important to see whether the performance at low temperatures is different from grating to grating (Bianco et al. 2003; Blais-Ouellette et al. 2004; Branche et al. 2004) we also investigate a different VPH grating, which is shown in the right panel of Fig.  \\ref{vphpics}. (This grating was provided as a free sample for demonstration purposes only.) The line density is 300 lines/mm and the peak efficiency is obtained around 0.7 $\\mu$m for an incident angle of 6$^{\\circ}$. The measurements of this grating are performed at wavelengths from 0.5 $\\mu$m to 0.9 $\\mu$m. ",
        "conclusions": "In this paper, we present results from cryogenic tests of VPH gratings at $\\sim$ 100 $K$. The aims of these tests are to see whether the diffraction efficiency at a low temperature as a function of wavelength is significantly different from that at room temperature and to see whether the grating can withstand a number of thermal cycles. Having exposed VPH gratings to 10 cycles between room temperature and 100 $K$, we find that diffraction efficiency measured at 100 $K$ agrees with that at room temperature within the errors. We also find no clear evidence that the performance changes with the successive thermal cycles. These results were found for both of the two different VPH gratings investigated here, which may imply that VPH gratings can withstand such cryogenic temperatures in general. Ideally, an investigation of more gratings, in particular from different manufacturers and with different substrate materials, should be carried out to confirm this point.  It needs to be emphasized that we have only confirmed the performance of a low dispersion VPH grating at cryogenic temperature. It would be useful to repeat the same experiments for high dispersion VPH gratings. Since the band-width of a throughput curve is narrower for high dispersion gratings, some changes of physical properties (e.g., gelatin thickness) due to temperature variations are expected to be revealed much more clearly in the form of a shift of the throughput peak and/or a global decrease of the throughput. In this case, one would have to predict the change of characteristics due to temperature variations and take them into account in the design and fabrication of a VPH grating so that it could work with the optimal performance at the operating temperature."
    },
    "0609/astro-ph0609296_arXiv.txt": {
        "abstract": "We present the results of a spectroscopic survey with the Low Resolution Imaging Spectrograph on the Keck~I telescope of more than $280$ star-forming galaxies and AGN at redshifts $1.4\\la z\\la 3.0$ in the GOODS-North field.  Candidates are selected by their $\\ugr$ colors using the ``BM'' and ``BX'' criteria to target redshift $1.4\\la z\\la 2.5$ galaxies and the Lyman break criteria to target redshift $z\\sim 3$ galaxies; combined these samples account for $\\sim 25-30\\%$ of the {\\cal R} and $\\ks$-band counts to ${\\cal R}=25.5$ and $\\ks{\\rm (AB)}=24.4$, respectively.  The sample of $212$ BM/BX galaxies and $74$ LBGs is presently the largest spectroscopic sample of galaxies at $z>1.4$ in GOODS-N, and includes $19$ spectroscopically-confirmed distant red galaxies (DRGs) with $\\jmk>2.3$ (Vega).  Extensive multi-wavelength data, including our very deep ground-based near-IR imaging to $\\ks{\\rm (AB)}=24.4$, allow us to investigate the stellar populations, stellar masses, bolometric luminosities ($L_{\\rm bol}$), and extinction of $z\\sim 2$ galaxies.  Deep {\\it Chandra} X-ray and {\\it Spitzer} IRAC and MIPS data indicate that the sample includes galaxies with a wide range in $L_{\\rm bol}$, from $\\simeq 10^{10}$~L$_{\\odot}$ to $>10^{12}$~L$_{\\odot}$, and covering 4 orders of magnitude in dust obscuration ($L_{\\rm bol}/L_{\\rm UV}$).  The sample includes galaxies with a large dynamic range in evolutionary state, from very young galaxies (ages $\\simeq 50$~Myr) with small stellar masses ($M^{\\ast}\\simeq 10^{9}$~M$_{\\odot}$) to evolved galaxies (ages $>2$~Gyr) with stellar masses comparable to the most massive galaxies at these redshifts ($M^{\\ast}>10^{11}$~M$_{\\odot}$).  {\\it Spitzer} data indicate that the optical sample includes some fraction of the obscured AGN population at high redshifts: at least 3 of 11 AGN in the $z>1.4$ sample are undetected in the deep X-ray data but exhibit power-law SEDs longward of $\\sim 2$~$\\mu$m (rest-frame) indicative of obscured AGN.  The results of our survey indicate that rest-frame UV selection and spectroscopy presently constitute the most time-wise {\\it efficient} method of culling large samples of high redshift galaxies with a wide range in intrinsic properties, and the data presented here will add significantly to the multi-wavelength legacy of the GOODS survey. ",
        "introduction": "\\label{sec:intro}  Rapid advances in our understanding of galaxy evolution have been prompted by the recognition that observations covering the full spectrum are necessary to adequately interpret the physical nature of galaxies. Multi-color {\\it Hubble Space Telescope} ({\\it HST}) imaging of two otherwise inconspicuous fields in the high Galactic latitude sky \\citep{williams96, williams00} marked the inception of a decade dominated by large-scale multi-wavelength surveys.  The two {\\it Hubble} Deep Fields are now encompassed or supplanted by other areas of the sky which are the focus of a number of space and ground-based observations both within and peripheral to the {\\it Great Observatories Origins Deep Survey} (GOODS; \\citealt{dickinsongoods03}).  Included among these data are the deepest {\\it Chandra} X-ray observations \\citep{alexander03}, {\\it HST} ACS optical imaging \\citep{giavalisco04}, {\\it Spitzer} IR to far-IR imaging (Dickinson et~al., in prep; Chary et~al., in prep), {\\it GALEX} far and near-UV imaging \\citep{schiminovich03}, ground-based optical and near-IR imaging and spectroscopy \\citep{capak04,cowie04,vanzella05}, and radio/submillimeter observations (\\citealt{richards00,pope05}).  Despite the easy access to broadband photometry and subsequent insights into galaxy evolution made possible by multi-wavelength surveys such as GOODS, important issues regarding survey completeness and the physical conditions in galaxies and their surrounding intergalactic medium can only be investigated spectroscopically. Spectroscopy of galaxies in blind flux-limited surveys can be quite inefficient and expensive, particularly if one only wants to study galaxies at certain cosmological epochs.  However, we have shown that the technique of photometric pre-selection can allow one to cull large samples of galaxies in particular redshift ranges over a large range in redshift $1.0\\la z\\la 4$ (e.g., \\citealt{adelberger04, steidel04,steidel03,steidel95,steidel93}), which can then be efficiently followed up using multi-object optical spectrographs such as the Low Resolution Imaging Spectrograph (LRIS) on the 10~m Keck telescope.  Near-UV sensitive spectrographs such as the blue arm of LRIS (LRIS-B) on Keck and the Focal Reducer/low dispersion Spectrograph (FORS) on the VLT have significantly extended our capabilities by allowing for spectroscopy of key features that fall shortward of the OH emission forest for redshifts $1.4\\la z\\la 3$, a particularly active epoch in the context of galaxy evolution and the buildup of stellar and black hole mass.  To take advantage of extensive multi-wavelength data, we included the GOODS-North (GOODS-N, hereafter) field in our ongoing survey to select and spectroscopically followup large samples of galaxies at redshifts $1.4\\la z\\la 3.0$ \\citep{steidel04}.  In the interest of public dissemination of data, we present in this paper the results of our spectroscopic survey of $1.4\\la z\\la 3.0$ star-forming galaxies in the GOODS-N field including associated photometry and spectroscopic redshifts.  Information on the galaxies, including their photometric measurements and errors and stellar population fits, are available at the following public website: http://www.astro.caltech.edu/$\\sim$drlaw/GOODS/.  The outline of this paper is as follows. In \\S~\\ref{sec:data} we briefly describe the optical imaging, photometry, and spectroscopy. To supplement these, we have also obtained the deepest wide-area near-IR $J$ and $K$-band imaging in the GOODS-North field, and these data are also presented in \\S~\\ref{sec:data}.  The spectroscopic results and associated catalog are presented in \\S~\\ref{sec:specres}. We describe the {\\it Spitzer} IRAC and MIPS data (taken from the GOODS-N public release; Dickinson et~al. in prep. and Chary et~al. in prep.) for our spectroscopic sample of galaxies in \\S~\\ref{sec:multiwave}.  Ground-based photometry and {\\it Spitzer} IRAC data, together with spectroscopic redshifts, enable the modeling of the stellar populations of galaxies given certain simplifying assumptions.  Our modeling procedure and results are discussed in \\S~\\ref{sec:seds}.  In \\S~\\ref{sec:diversity} we describe a few characteristics of the sample of star-forming galaxies and AGN to demonstrate the wide range in intrinsic properties of UV-selected galaxies at high redshift.  A flat $\\Lambda$CDM cosmology is assumed with $H_{0}=70$~km~s$^{-1}$~Mpc$^{-1}$ and $\\Omega_{\\Lambda}=0.7$. All magnitudes are on the AB \\citep{oke83} system. ",
        "conclusions": "\\label{sec:summary}  We have presented the results of a spectroscopic survey of redshift $1.4\\la z\\la 3.0$ star-forming galaxies in the GOODS-North field, made possible by efficient UV ($\\ugr$) color selection and the unique multi-object capabilities of the LRIS instrument on the Keck I telescope.  Our sample consists of 212 redshifts for galaxies at redshifts $1.4\\la z\\la 2.5$ selected using the BM and BX criteria of \\citet{adelberger04} and \\citet{steidel04}, and 30 new redshifts (of a total of 74) for Lyman break galaxies at redshifts $2.5\\la z\\la 3.5$. Our deep optical and near-IR imaging, supplemented by publicly available {\\it Spitzer} IRAC and MIPS data (Dickinson et~al., in prep; Chary et~al., in prep), allow us to measure the stellar populations, stellar masses, star formation rates, and dust extinction for galaxies in our sample (e.g., \\citealt{erb06b, shapley03, shapley05, reddy04, reddy05a, reddy06a, steidel04}).  These analyses indicate that the $\\ugr$-selected sample consists of galaxies which span two orders of magnitude in age and stellar mass, and 4 orders of magnitude in dust obscuration ($L_{\\rm bol}/L_{\\rm UV}$).  Included are galaxies with bolometric star formation rates ranging from $\\sim 5$~$\\sfr$ to $>1000$~$\\sfr$.  We further identify at least 3 of 11 AGN in our sample which appear to be heavily dust-obscured based on their power-law SEDs longward of $2$~$\\mu$m (rest-frame) and lack of detection in the deep {\\it Chandra} 2~Ms data \\citep{alexander03}.  A compilation of the multi-wavelength data for these 11 AGN indicates that optical and {\\it Spitzer} data are able to more efficiently (in terms of integration time) select AGN at $z>1.4$ than X-ray data, but optical spectra and {\\it Spitzer} and {\\it Chandra} data are all required to fully account for the census of AGN at high redshifts. The photometry and SED fitting results for galaxies in our sample are available at http://www.astro.caltech.edu/$\\sim$drlaw/GOODS/.  Large spectroscopic samples at high redshifts allow for a number of other detailed investigations such as the galaxy and AGN/QSO luminosity functions (\\citealt{steidel99, adel00, shapley01, hunt04}, Reddy et~al. in prep.); metallicities \\citep{pettini98, pettini01, shapley04, erb06a}; signatures of galaxy feedback and IGM metal enrichment \\citep{adel03}; and accurate clustering analyses \\citep{adelberger05a, adelberger05b}.  This large range in galaxy evolution studies highlights the versatility and efficiency of optically-selected samples in addressing many fundamental issues in cosmology."
    },
    "0609/astro-ph0609775_arXiv.txt": {
        "abstract": "Recent experiments suggest that polarized photons may couple significantly to pseudoscalar particles such as axions.  We study the possible observational signatures of axion-photon coupling for radiation from magnetic stars, with particular focus on neutron stars. We present general methods for calculating the axion-photon conversion probability during propagation through a varying magnetized vacuum as well as across an inhomogeneous atmosphere. Partial axion-photon conversion may take place in the vacuum region outside the neutron star. Strong axion-photon mixing occurs due to a resonance in the atmosphere, and depending on the axion coupling strength and other parameters, significant axion-photon conversion can take place at the resonance. Such conversions may produce observable effects on the radiation spectra and polarization signals from the star.  We also apply our results to axion-photon propagation in the Sun and in magnetic white dwarfs. We find that there is no appreciable conversion of solar axions to photons during the propagation. ",
        "introduction": "The axion is a hypothesized pseudoscalar particle arising from the breaking of a $U(1)$ Peccei-Quinn symmetry, introduced to explain the absence of strong CP violation\\cite{peccei,weinberg,wilczek}. The axion is also an ideal candidate for cold dark matter that makes up five-sixths of all matter in the universe. The allowed axion mass $m_a$ is in the range of $10^{-6}\\lo m_a\\lo 10^{-3}$~eV\\cite{raffelt04}.  A general property of the axion is that it can couple to two photons (real or virtual) via the interaction \\be {\\cal L}_{a\\gamma\\gamma}=-{1\\over 4}\\,g\\,a\\,F_{\\mu\\nu}{\\tilde F}^{\\mu\\nu}= g\\,a\\,\\bE\\cdot\\bB, \\ee where $a$ is the axion field, $F_{\\mu\\nu}$ (${\\tilde F}^{\\mu\\nu}$) is the (dual) electromagnetic field strength tensor, and $g$ is the photon-axion coupling constant. Accordingly, in the presence of a magnetic field ${\\bf B}$, a photon may oscillate into an axion and vice versa.  Exploiting such photon-axion oscillation, experiments based on ``photon regeneration'' (``invisible light shining through walls'') set the constraint $g<6.7\\times 10^{-7}~\\GeV1$ ($95\\%$ CL) for axions with $m_a<10^{-3}$~eV\\cite{cameron93}. The first results from the CERN Axion Solar Telescope (CAST) experiment imply an upper limit $g\\lo 10^{-10}~\\GeV1$ for $m_a\\lo 0.02$~eV \\cite{zioutas05}. Other experimental and astrophysical constraints on the axion parameters (mass $m_a$ and $g$) are reviewed in Refs.~\\cite{raffelt04,raffelt05,vanBibber06}.  Recently, the PVLAS collaboration has reported measurements of the rotation of the polarization of photons in passing through a vacuum cavity in a magnetic field~\\cite{zavattini06}. If interpreted as due to the conversion of photons to axions, this may imply $g\\simeq (1.6-5) \\times 10^{-6}~\\GeV1$ for $m_a\\sim 10^{-3}$~eV. Unfortunately, this result contradicts the constraint from the CAST experiment and the previous energy-loss limit from globular cluster stars. Thus, either there remain systematic effects in the PVLAS experiment or nonminimal models of psuedoscalar-photon coupling are required \\cite{masso05}.  In this paper, we study the effects of axion-photon coupling on radiation from magnetic stars (see \\cite{raffelt88} for previous works).  We will consider generic psuedoscalar parameters with $m_a\\lo 10^{-3}$~eV and $g\\lo 10^{-6}~{\\rm GeV}^{-1}$, not necessarily the QCD axion (for which there exists a unique relation between $m_a$ and $g$). Our goal is to understand and quantify what might be the observational signatures of axion-photon coupling for various axion parameters.  Our method and basic results (such as the axion-photon conversion probability) can be applied for all axion parameters and in different astrophysical environments. Most of the paper will deal with magnetic neutron stars, for which most significant effects are possible, but we also discuss applications of our results in the contexts of magnetic white dwarfs and the Sun.  Our paper is organized as follows. In Sect.~II we study axion-photon propagation in the vacuum region outside a magnetic neutron star, taking into account of the variation of the magnetic field.  Section III examines the propagation in a magnetized plasma characteristic of a neutron star atmosphere. We study in detail the properties of the axion-photon resonance as the photon-axion propagates in an inhomogeneous medium. Depending on the axion coupling strength and other parameters, significant axion-photon conversion may take place and such conversion may leave an imprint on the radiation spectrum and polarization from the star.  In Sect.~IV we consider axion-photon propagation in the Sun and in magnetic white dwarfs, and Sect.~V contains a brief discussion of our results. ",
        "conclusions": "In this paper we have studied the propagation of the axion-photon system in the atmospheres and the near vicinity of magnetic stars. Because of their strong magnetic fields, we have focused on neutron stars, but our main results/methods can be similarly applied to other astrophysical bodies. Our study goes beyond previous work~\\cite{raffelt88} in that we quantify the various parameter regimes for which axion-photon conversion is important, we calculate the conversion probabilities for propagation in varying magnetic fields and varying plasma densities, and we present examples to illustrate how axion-photon coupling may affect observed photon spectra and polarizations.  The axion-photon resonance (where maximum mixing occurs) is always present in magnetized neutron star atmospheres/magnetospheres.  This resonance [see Eq.~(\\ref{eq:veq})] is at a higher density than the previously-studied vacuum resonance. Complete axion-photon conversion is possible only when the adiabatic condition [$\\gamma_{\\rm res}\\gg 1$; see Eq.~(\\ref{eq:gammmares2})] is satisfied. A high axion-photon coupling strength, gradual density gradient and lower photon energy tend to make such conversion a possibility.  Note that in this paper we have treated the plasma dielectric tensor in the cold-plasma approximation. For low-energy photons (e.g. optical), the resonance occurs in the neutron star's magnetosphere and the cold plasma treatment is no longer valid. It would be straightforward to to generalize our results to more general plasma dielectric tensors.  Even without the axion-photon resonance, partial axion-photon conversion may take place during the propagation in the vacuum region with spatially varying magnetic fields (see Sect.~II).  Applying our result to the axions produced at the center of the Sun, we find that there is no possibility for appreciable axion-photon conversion during propagation (see Sect.~IV). Thus the nondetection of axions in the CAST experiment cannot be explained by such oscillation effect. Our analysis also shows that with the axion parameters allowed by the PVLAS experiment, significant photon-axion resonant conversion is possible in highly magnetized white dwarfs. This may produce interesting spectral and polarization signatures in the observed radiation from the white dwarf.  \\bigskip"
    },
    "0609/astro-ph0609822_arXiv.txt": {
        "abstract": "{} {We investigate the cosmological test based on the evolution of the gas fraction in X-ray galaxy clusters and the stability of the cosmological parameters derived from it.} { Using a sample of distant clusters observed by \\emph{XMM-Newton} and \\emph{Chandra}, and comparing their gas fraction at different radii to the gas fraction observed for nearby clusters, we have determined the likelihood functions for $\\Omega_m$ in a flat universe and the confidence contours in the $\\Omega_m-\\Omega_{\\Lambda}$ plane.} { Results obtained at the virial radius point to a high matter density Universe, while for inner radii the $\\Omega_m$ parameter obtained tends to decrease, reaching values compatible with the concordance model. The analysis allows us to conclude that this test provides ambiguous results due to the complex structure of the ICM that induces a dependence of the gas fraction on temperature, radius, and redshift, which cannot be accounted for by the self-similar picture expected from pure gravitational heating of the ICM.} {The use of gas fraction in X-ray clusters to constrain cosmological parameters seems therefore to be compromised until a better understanding of the ICM physics and the ability to obtain observations of better quality up to the virial radius are achieved.} ",
        "introduction": "Physical proprieties of galaxy clusters have been considered a promising tool for cosmological probes. A test based on the assumption of a constant baryon fraction in these objects was proposed by \\cite{sasaki} and \\cite{pen}. These authors showed that under this assumption, the gas fraction of clusters can, in principle, be used as a standard ``ruler'' to determine cosmological parameters. The main idea is that if the only heating source of the gas is gravitational collapse, the baryon fraction should be constant and independent of redshift. Since the measured baryon fractions depend on the cosmology through the angular distance to the clusters, the ``true'' cosmological model would be one where all the measured baryon fractions at different redshifts would be equal. This method has been applied in the past years (\\cite{Rines}; \\cite{ettori1}; \\cite{Allen2002}; \\cite{Ettori}; \\cite{Allen2004}) using clusters observed by \\emph{Chandra} and \\emph{ROSAT PSPC}, with the results pointing to a low value of $\\Omega_{m}$, consistent with the concordance model ($\\Omega_m\\approx0.3,\\Omega_{\\Lambda}\\approx0.7$). However, these authors have often restricted the method to a single radius, defined by a given high overdensity relative to the critical density of the Universe. To check the validity of this cosmological test, it is desirable to perform a good study on the behavior of the gas fraction distribution (the main contribution of baryons) inside clusters and its evolution with redshift. Such a study was done recently in \\cite{sadat05}, hereafter S05, who claim that clusters are more complex objects than it was expected in self-similar models, as they found that gas fraction presents a dependence on radius, temperature, and redshift.  The goal of the present work is to give a quantitative analysis on the results obtained in S05, by determining the likelihood functions of the cosmological test based on the gas fraction for the sample of S05 and study the effect of using radii of different overdensities to perform it. The paper is organized as follows. In Sect. 2, we present the cluster samples that was used in the test. In Sect. 3, we present the method used to determine the gas fractions for all clusters in the sample. The cosmological results are presented in Sect. 4 and finally, in Sect. 5, we draw our conclusions. We used a Hubble constant of 50 km/s/Mpc, unless the dependence is explicitly given (with $H_0$=100$h$km/s/Mpc). ",
        "conclusions": "We have used combined data from \\emph{Chandra} and \\emph{XMM-Newton} to apply the cosmological test based on the gas fraction evolution considering that this quantity varies with the temperature of the cluster and that this dependence should evolve with redshift through a simple scaling argument based on the gravitational growth of structure in an expanding Universe. The obtained results allow us to conclude that this kind of test is highly dependent on the radius considered to implement it, since we found significantly different values for $\\Omega_m$ when the test is performed at different normalized radii. The complexities in cluster's internal structures which were evidenced in S05 are at the origin of this observed ambiguity. It seems that the gas fraction observed in low and high-redshift clusters cannot be accounted for by a purely gravitational heating of the gas, which in other words means that non-gravitational effects present in clusters induce a significant deviation of the gas fraction profile from standard scaling. We found no significant differences in the results when using two different definitions for normalizing the radius, and it is interesting to point out that the results obtained at R$_{2500}$ are compatible with those obtained by Allen et al.(2002) and Allen et al.(2004) at the same radius. Note although that these authors did not apply the same method, as they assumed a constant gas fraction at their working radius.  The main question raised by this work is whether there is a radius at which scaling relations can be used with more confidence. At first one is tempted to say that the virial radius should provide a positive answer, since at this radius non-gravitational processes are thought to be less important. However, there are two aspects that do not allow us to make definite conclusions on this point: a) the observations of high-redshift clusters do not extend up to the virial radius and so the gas masses at this region were obtained by extrapolating the $\\beta$-model fitted on the observed parts of the clusters, and b) the maximum likelihood values for $\\Omega_m$ at this radius are too high ($\\Omega_m \\sim 0.7$), which translates to a significant deviation from the concordance model supported by other observational data such as CMB, Type Ia supernovae, and baryonic oscillations (\\cite{Spergel}, \\cite{astier}, \\cite{eisenstein}, \\cite{Blanchard}). Observations of X-ray emissivity in distant clusters up to the virial radius will be critical to providing definitive answers on the first aspect. Despite this, the use of gas mass fraction evolution (or lack of it) in clusters as a precision method for constraining cosmological parameters seems to be seriously compromised."
    },
    "0609/astro-ph0609013_arXiv.txt": {
        "abstract": "Using a statistically rigorous analysis method, we place limits on the existence of an isotropic stochastic gravitational wave background using pulsar timing observations. We consider backgrounds whose characteristic strain spectra may be described as a power-law dependence with frequency. Such backgrounds include an astrophysical background produced by coalescing supermassive black-hole binary systems and cosmological backgrounds due to relic gravitational waves and cosmic strings.  Using the best available data, we obtain an upper limit on the energy density per unit logarithmic frequency interval of $\\Omega^{\\rm SMBH}_g(1/8\\mbox{ yr}) h^2 \\leq 1.9\\times 10^{-8}$ for an astrophysical background which is five times more stringent than the earlier Kaspi et al. (1994) limit of $1.1\\times 10^{-7}$. We also provide limits on a background due to relic gravitational waves and cosmic strings of $\\Omega^{\\rm relic}_g(1/8\\mbox{ yr}) h^2 \\leq 2.0\\times10^{-8}$ and $\\Omega^{\\rm cs}_g(1/8\\mbox{ yr}) h^2 \\leq 1.9\\times10^{-8}$ respectively. All of the quoted upper limits s correspond to a 0.1\\% false alarm rate together with a 95\\% detection rate. We discuss the physical implications of these results and highlight the future possibilities of the Parkes Pulsar Timing Array project.  We find that our current results can 1) constrain the merger rate of supermassive binary black hole systems at high red shift, 2) rule out some relationships between the black hole mass and the galactic halo mass, 3) constrain the rate of expansion in the inflationary era and 4) provide an upper bound on the dimensionless tension of a cosmic string background. ",
        "introduction": "Pulsar timing observations (see Lorimer \\& Kramer 2005\\nocite{lk05}, Edwards, Hobbs \\& Manchester 2006\\nocite{ehm06} for a review of the techniques) provide a unique opportunity to study low-frequency ($10^{-9}-10^{-7}$ Hz) gravitational waves (GWs) (e.g., Sazhin 1978\\nocite{saz78}, Detweiler 1979\\nocite{det79}, Bertotti, Carr, \\& Rees 1983\\nocite{bcr83}, Foster \\& Backer 1990\\nocite{fb90}, Kaspi, Taylor \\& Ryba 1994\\nocite{ktr94}, Jenet et al. 2005\\nocite{jhlm05}). Sources in this low-frequency band include binary supermassive black holes, cosmic super-strings, and relic gravitational waves from the Big Bang (Jaffe \\& Backer 2003, Maggiore 2000\\nocite{mag00}).  An isotropic stochastic background can be described by its characteristic strain spectrum $h_c(f)$ which, for most models of interest, may be written as a power-law dependence on frequency, $f$: \\begin{equation} h_c(f) = A \\left(\\frac{f}{{\\rm yr}^{-1}}\\right)^\\alpha. \\label{hc} \\end{equation}  Table~\\ref{table1} shows the expected values of $A$ and $\\alpha$ for different types of stochastic backgrounds that have been addressed in the literature. The characteristic strain is related to the one-sided power spectrum of the induced timing residuals, $P(f)$, as \\begin{equation} P(f) = \\frac{1}{12 \\pi^2}\\frac{1}{f^3}h_c(f)^2, \\label{resspec} \\end{equation} and to $\\Omega_{gw}(f)$, the energy density of the background per unit logarithmic frequency interval, as \\begin{equation} \\Omega_{gw}(f) = \\frac{2}{3}\\frac{\\pi^2}{H_0^2} f^2 h_c(f)^2 \\label{omega} \\end{equation} where $H_0$ is the Hubble constant. Note that the one-sided power spectrum, $P(f)$, is defined so that \\begin{equation} \\int_0^\\infty P(f) df = \\sigma^2 \\end{equation} where $\\sigma^2$ is the variance of the arrival time fluctuations, or timing residuals, generated by the presence of the GW background. Since $\\sigma^2$ has the physical units of s$^2$, P(f) has the units of s$^2$/Hz or s$^3$.      Jenet et al. (2005)\\nocite{jhlm05} developed a technique to make a definitive detection of a stochastic background of GWs by looking for correlations between pulsar observations. It was shown that approximately 20 pulsars would need to be observed with a timing precision of $\\sim$100\\,ns over a period of $5$\\,years in order to make such a detection if the GW background is at the currently predicted level (Jaffe \\& Backer 2003, Wyithe \\& Loeb 2003, Enoki et al. 2004, Sesana et al. 2004)\\nocite{shmv04}. The Parkes Pulsar Timing Array (PPTA) project (Hobbs 2005\\nocite{hob05}) is trying to achieve these ambitious goals, but the currently available data-sets do not provide the required sensitivity for a detection. In this paper, we introduce a method to place an upper bound on the power of a specified stochastic GW background using observations of multiple pulsars. Full technical details of our implementation will be published in Hobbs et al. (2006).  Here, this method is applied to data (see Section 2) from seven pulsars observed for the PPTA project combined with an earlier publicly available data-set.    \\begin{table*} \\caption{The expected parameters for predicted stochastic backgrounds}\\label{table1} \\begin{tabular}{llll}\\hline \\hline Model & A & $\\alpha$ & References \\\\ \\hline  Supermassive black holes & $10^{-15} - 10^{-14}$ & $-$2/3 & Jaffe \\& Backer (2003)\\\\ & & & Wyithe \\& Loeb (2003)\\\\ & & & Enoki et al. (2004) \\\\ Relic GWs & $10^{-17} - 10^{-15}$ & $-1$ -- $-0.8$ & Grishchuk (2005) \\\\ Cosmic String  & $10^{-16} - 10^{-14}$ & $-$7/6  & Maggiore (2000) \\\\ \\hline\\end{tabular} \\end{table*}\\nocite{jb03}\\nocite{wl03a}\\nocite{eins04}  Upper limits have already been placed on the amplitude of any such background of GWs. Using eight years of observations for PSR~B1855$+$09, Kaspi et al. (1994) obtained a limit of $\\Omega_g h^2 \\le 1.1 \\times 10^{-7}$, where $H_0 = 100h$ km~s$^{-1}$~Mpc$^{-1}$, at the $95\\%$ confidence level\\footnote{Their more stringent constraint of $\\Omega_g h^2 \\le 6 \\times 10^{-8}$ was obtained when data from PSRs~B1855$+$09 and B1937$+$21 were combined. Since the data from PSR~B1937+21 is far from white, we believe this limit is artificially low and therefore restrict our discussion to the PSR~B1855+09 data only.} for the case when $\\alpha = -1$ (i.e. $\\Omega_{gw}$ is independent of frequency). This work was continued by Lommen (2002)\\nocite{lom02} who used 17\\,yrs of observations to obtain $\\Omega_g h^2 < 2\\times10^{-9}$.  However, the statistical method used for both of these analyses has been criticized in the literature (see, for instance, Thorsett \\& Dewey 1996\\nocite{td96}, McHugh et al. 1996\\nocite{mzvl96}, Damour \\& Vilenkin 2005). In this paper, we develop a frequentist technique, similar to that used by the LIGO science collaboration (Abbott et al. 2006\\nocite{aaa+06}), to place an upper bound on $A$, given $\\alpha$. The technique makes use of a statistic, $\\Upsilon$, defined below, which is sensitive to red noise in the pulsar timing residual data. Upper bounds on $A$ are determined using $\\Upsilon$ together with a specified false alarm rate, ${\\cal P}_f$, and detection rate, ${\\cal P}_d$.  Monte-Carlo simulations are used to determine these probabilities by generating pulsar pulse times-of-arrival consistent with a GW background. All of the upper limits quoted in this paper correspond to ${\\cal P}_f = 0.1\\%$ and ${\\cal P}_d = 95\\%$. ",
        "conclusions": "The upper bound on the stochastic background can be used to probe several aspects of the Universe. Precisely what is being constrained depends on the physics of the particular background in question. Here, both the measured upper bounds using the currently available data and the expected upper bounds using the full five-year PPTA data-set are discussed in the context of several GW backgrounds.    \\subsection{Supermassive black holes} A GW background generated by an ensemble of supermassive black holes distributed throughout the universe has been investigated by several authors (Jaffe \\& Backer 2003, Wyithe \\& Loeb 2003, Enoki et al. 2004). In general, the characteristic strain spectrum for this background can be written as: \\begin{equation} h_c(f) = 2.5 10^{-16} h \\left(\\frac{f}{\\mbox{yr}^{-1}}\\right)^{-2/3}\\left\\langle\\left(\\frac{M_c}{10^{7}\\Msolar}\\right)^{5/3}\\right\\rangle^{1/2} \\left(\\frac{N_0}{\\mbox{Mpc}^{-3}}\\right)^{1/2} I^{1/2} \\end{equation} where \\begin{equation} I = \\int \\frac{N(z)}{N_0} H_0 \\frac{a(z)}{\\dot{a}(z)} \\frac{dz}{(1+z)^{4/3}}, \\end{equation} a(z) is the cosmological scale factor written in terms of red shift, $z$, $\\dot{a}(z)$ is the derivative of $a(z)$ with respect to cosmic time, $H_0$ is the Hubble constant, the chirp mass $M_c = [M_1 M_2 (M_1 + M_2)^{-1/3}]^{3/5}$ of a given binary system, $\\langle \\rangle$ represents ensemble averaging over all the systems generating the background, $N(z)$ is the galaxy merger rate as a function of red shift, and $N_0$ is the present day number density of merged galaxies that created a black-hole binary system. The values of each of these physical quantities are currently poorly constrained and each investigator has chosen a different parameterization. Under the framework described by Jaffe and Backer (2003), $\\langle M_c^{5/3}\\rangle$ and $N_0$ are constrained by observations at the current epoch to be $\\langle M_c\\rangle \\approx 2.3 \\times 10^7 \\Msolar$ and $N_0 \\approx 1/$\\,Mpc$^3$. They parameterized the galaxy merger rate such that $R(z)$ goes as $(1 + z)^\\gamma$. Hence, $I$ depends on $\\gamma$. Combining the estimates of $\\langle M_c\\rangle$ and $N_0$ together with our measured upper bound of $A_{\\rm upper} = 1.1\\times 10^{-14}$, one finds that $I \\leq 3$. Using the full PPTA after 5 years, one expects $I \\leq 0.8$. These constraints together with the calculations of Jaffe \\& Backer (2003) (see Fig.~4 in their work) constrain $\\gamma$. Currently the limit on $\\gamma$ is $<2.6$ and with the full PPTA $\\gamma < 0.4$. This value is expected to lie somewhere between $-0.4$ and 2.3 (Carlberg et al. 2000\\nocite{ccp+00}, Patton et al. 2002\\nocite{ppc+02}). Current PPTA sensitivity (i.e. using the data presented in this paper) is just above the expected range, while the full PPTA should be able to place meaningful constraints on this exponent.  In the Wyithe \\& Loeb (2003) work, both $\\langle M_c^{5/3}\\rangle$ and $I$ depend strongly on the black-hole versus galactic-halo mass ($M_{\\rm BH}-M_{\\rm HM}$) relationship. They discuss several different scenarios which yield different $M_{\\rm BH}-M_{\\rm HM}$ relationships and hence different levels of the background. For the case of an $M_{\\rm BH}-M_{\\rm HM}$ relationship determined by Ferrarese (2002)\\nocite{fer02}, the expected value of $A$ is $2 \\times 10^{-15}$. For the $M_{\\rm BH}-M_{\\rm HM}$ relationship derived from Navarro, Frenk, \\& White (1997)\\nocite{nfw97}, $A = 5 \\times 10^{-15}$. Using an $M_{\\rm BH}-M_{\\rm HM}$ relationship derived from simple considerations of BH growth by feedback from quasar activity (Wyithe \\& Loeb 2003, Haehnelt, Natarajan \\& Rees 1998\\nocite{hnr98}, Silk \\& Rees 1998\\nocite{sr98}), $A \\approx 10^{-15}$.  Our measured upper limit for $\\alpha = -2/3$ cannot rule out any of these models. However, if only a limit is obtained from the full PPTA observations, it will rule out all of the models described above.  \\subsection{Relic gravitational waves}  A relic gravitational wave background is generated by the interaction between the large-scale dynamic cosmological metric and quantum fluctuations of the metric perturbations occurring in the early universe \\citep{gri05}. In the nano-Hz frequency regime, the background takes the following form: \\begin{equation} h_c(f) =  h_c(f_h)\\left(\\frac{f}{H_0}\\right)^{\\alpha} \\left(\\frac{a_2}{a_H}\\right)^{\\frac{1}{2}} \\label{relicgw} \\end{equation} where $h_c(f_h)$ is the magnitude of the characteristic strain spectrum at $f=H_0$, $a_H$ is the current value of the cosmological expansion factor, and $a_2$ is the value of the expansion factor at the start of the matter-dominated epoch.  Note that this expression is not valid in the ultra-low frequency regime where $f\\approx H_0$. The notation used here is consistent with \\citet{gri05} except for $\\alpha$ which is related to Grishchuk's parameter $\\beta$ by $\\alpha = 1 + \\beta$. The exponent determines the evolution of the inflationary epoch which starts the GW amplification process. When $\\alpha = -1$, the scale factor grows exponentially with global cosmic time. The ratio $a_2/a_H$ is believed to be about $10^{-4}$. $h_c(f_h)$ is constrained by cosmic microwave background measurements to be about $10^{-5}$. Using these values and assuming the validity of the amplification scenario described in Grishchuk (2005), the upper bound on $A$ may be used to constrain $\\alpha$. The upper bound on $\\alpha$ is given by the solution to the following equation: \\begin{equation} h_c(f_h)\\left(\\frac{1/\\mbox{1 yr}}{f_h}\\right)^{\\alpha} \\left(\\frac{a_2}{a_H}\\right)^{\\frac{1}{2}} = A(\\alpha). \\end{equation} The above equation yields $\\alpha \\leq -0.7$ for the current PPTA and $\\alpha \\leq -0.84$ for the full PPTA. Within the theoretical framework described by \\citet{gri05}, if $\\alpha$ is larger than $-0.80$, small scale gravitational waves will effect primordial nucleosynthesis, while an $\\alpha$ less than $-1.0$ will result in an infinitely large energy density in small scale gravitational waves. Hence, the full PPTA will be able to place useful constraints on the relic gravitational wave background. Since $\\alpha$ determines the rate of expansion in the inflationary epoch, it turns out that it is related to the equation of state of the ``matter'' in that epoch by \\begin{equation} \\frac{P}{\\epsilon} = w = \\frac{2 - \\alpha}{3 \\alpha}. \\end{equation} where $P$ is the pressure and $\\epsilon$ the energy density. The full PPTA will constrain $w$ in the early universe to be greater than $-1.17$. This would limit inflationary models based on ``quintessence'' and ``phantom energy'' \\citep{nog06,pad05}.  \\subsection{Cosmic strings}  It has been proposed that oscillating cosmic string loops will produce gravitational wave radiation (Vilenkin 1981\\nocite{vil81}). Recently, \\citet{dv05} discussed the possibility of generating a stochastic GW background using a network of cosmic superstrings. Using a semi-analytical approach, they derived the following characteristic strain spectrum valid in the pulsar timing frequency range (see their equation 4.8): \\begin{equation} h_c(f) = 1.6 \\times 10^{-14} c ^{1/2} p^{-1/2} \\epsilon_{eff}^{-1/6} (h/.65)^{7/6}\\left(\\frac{G \\mu}{10^{-6}}\\right)^{1/3} \\left(\\frac{f}{\\mbox{yr}^{-1}}\\right)^{-7/6} \\end{equation} where $\\mu$ is the string tension, $G$ is Newton's constant, $c$ is the average number of cusps per loop oscillation, $p$ is the reconnection probability, $\\epsilon_{eff}$ is a loop length scale factor, and $h$ is the Hubble constant in units of 100 km/s/Mpc. Note that for the above estimate, $h$ was evaluated at $0.65$ in order to be consistent with \\citet{dv05}. The combination $G \\mu$ is the dimensionless string tension which characterizes the gravitational interaction of the strings. The predicted string tensions are $10^{-11} \\leq G \\mu \\leq 10^{-6}$ \\citep{dv05}. Using the above spectrum together with the measured upper bound on $h_c$ for $\\alpha = -7/6$, an upper bound may be placed on the dimensionless string tension: \\begin{equation} G \\mu \\leq 1.5 \\times 10^{-8} c^{-3/2} p^{3/2} \\epsilon_{eff} ^{1/2} (h/.65)^{-7/2}. \\end{equation} As emphasized by \\citet{dv05}, the above expression for the upperbound may be simplified using the fact that both $p$ and $\\epsilon_{eff}$ are less then one and $h$ is expected to be greater than .65: \\begin{equation} G \\mu \\leq 1.5 \\times 10^{-8} c^{-3/2}. \\end{equation} Using standard model assumption where $c=1$, the upperbound becomes $G\\mu \\leq 1.5 \\times 10^{-8}$. This is already limiting the parameter space of the cosmic string model of Sarangi and Tye (2002).  With the full PPTA, the limit will become $G \\mu \\leq 5.36 \\times 10^{-12}$. Hence, the full PPTA will either detect GWs from cosmic strings or rule out most current models."
    },
    "0609/astro-ph0609539_arXiv.txt": {
        "abstract": "% We present the first results of our study of stellar populations in the Large and Small Magellanic Clouds based on multiband WFPC2 observations of ``random\" fields taken as part of the ``pure parallel\" program carried out by the Space Telescope Science Institute as a service to the community. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609225_arXiv.txt": {
        "abstract": "{We present the results of modelling the 0.45--1\\mic\\ spectral energy distribution of \\vmon\\ for 2002 November. Synthetic spectra were calculated using the NextGen model atmospheres of Hauschildt et al. (1999), which incorporate line lists for H$_2$O, TiO, CrH, FeH, CO, and MgH, as well as the VALD atomic line list. Fits to the observed spectra show that, in 2002 November, the effective temperature of \\vmon\\ was approximately $2000\\pm100$~K. Our theoretical spectra show a comparatively weak dependence on $\\log g$. Preliminary analysis of the hot star observed together with \\vmon\\ shows it to be a normal B3V dwarf. ",
        "introduction": "\\vmon\\ (also known as  EQJ0704.0-0350, GSC~04822-00039, IRAS~07015-0346, USNO-A2.0 0825-03833116, AAVSO 0659-03) has been the subject of intense interest since its discovery as an eruptive variable by \\cite{brown} on January 6, 2002. Two further  maxima on the light curve subsequently occurred, in 2002 February (\\citealt{Munari2002, Kimeswenger2002, crause2003}) and the optical flux in the $V$ band increased by $\\sim 9$~mag. The luminosity peaked at $V \\approx 7$ in 2002 February. Later, a gradual decline in the $V$ flux began which, by 2003 January, had faded by 8 mag.  \\cite{Henden2002} reported the presence of a light echo, subsequently studied in detail by \\cite{Bond2003} and \\cite{crause2005}. The reflecting dust may be interstellar \\citep{Tylenda2005a} or may have originated in the envelope of \\vmon\\ and ejected in the past \\citep{loon2004}. Van Loon et al. (2004) reported the discovery of the multiple dust shells around \\vmon.  Alternatively, these shells arose through heavy mass loss during the evolution of a massive progenitor (Munari et al. 2005).  The distance to \\vmon\\ can be estimated from the evolution of the light echo \\citep{Henden2002}. According to works based on Hubble Space Telescope data, its distance is $> 6$~kpc \\citep{Bond2003}, $8\\pm 2$~kpc \\citep{Tylenda2004}, or as high as 12 kpc (Crause et al. 2005). However, according to recent work based on HST data, its distance is 5.9 kpc (Sparks et al. 2006). If these estimates are correct, then at the time of maximum brightness \\vmon\\ was the most luminous star in the Galaxy.  The cause of the eruption of \\vmon\\ and the nature of its progenitor are unclear. The integrated colors of the progenitor looked that of an unreddened F-type star \\citep{Munari2002}. \\cite{desidera2002} discovered spectroscopically a hot companion to the outbursting star, later confirmed by Wagner et al. (2002), and classified as a B3V star by \\cite{munari05}.   The spectral, photometric, and polarimetric evolution of \\vmon\\ has been described in several studies (\\citealt{Munari2002, Kimeswenger2002, Kolev2002, Wisniewski2003, Osiwala2003, crause2003, Kipper2004, rushton05}). By mid-April 2002, bands of TiO had appeared in the spectrum and by May, the spectrum had evolved to a ``very cold'' M giant \\citep{banerjee02}. Its infrared spectrum in 2002 October was characterized as an ``L-supergiant'' by \\cite{evans03}.  The complexity and the rarity of V838 Mon optical spectrum resulted in only a few sporadic attempts to derive the star's atmospheric properties. \\cite{Kipper2004} found that, for the iron group of elements, $\\mbox{[M/H]} = -0.4$, while abundances of lithium and of some s-process elements are clearly enhanced. Later, \\cite{kaminsky05} provided fits to echelle data and obtained a similar result, $\\mbox{[Fe/H]} = -0.3\\pm0.2$. These results were obtained using a static LTE model and are therefore very dependent on the model atmosphere and spectrum synthesis assumptions.  In this paper we discuss fits to the spectral energy distribution of \\vmon\\ using data obtained in 2002 November. The observational data used in this paper are described in Section 2. Section 3 explains some of the background to our work and some details of the procedure and input parameters used. The results are discussed in Section 5. ",
        "conclusions": "It seems that the pseudo\\-photosphere of \\vmon\\ was evolving sufficiently slowly that its optical spectrum could be approximated by a sequence of classical,  i.e. static model atmospheres \\citep{kaminsky05}; we find that the standard computational procedures for SEDs, applied to \\vmon, give results that are reliable enough to determine the effective temperature. It is worth noting that the presence of P Cyg profiles indicates expansion velocities up to 200\\vunit. However the emission components of the P Cyg profiles form at the outer boundary of the envelope, well above the region where the observed continuum and absorption lines form.  There are distinct differences between the SED of \\vmon\\ and of late-type dwarfs with similar spectral type, which obviously arise from their differing physical nature and the input physical data. In the case of \\vmon\\ we have modelled its SED at an advanced stage of its evolution and find that its optical spectrum is formed by absorption of the saturated bands of VO and TiO.  We note that our models do not satisfactorily reproduce the absorption band located around 0.75\\mic; this absorption cannot be identified as TiO or VO bands (see Fig.~\\ref{opacities}). Hypothetically, the strong band of CH$_4$ could be formed in a low-temperature regime at these wavelengths. This band, as well as the bands at 0.619 and 0.890\\mic, are known to be present in the spectra of Jupiter and other gas giants. Possibly the band may be formed in the outermost layers of the extended envelope of \\vmon. We plan to provide a more detailed study of the molecular features in a forthcoming paper.   Nonetheless, we have obtained a qualitatively good fit of the spectrum of \\vmon\\ over a wide spectral range (0.4--1\\mic) with a model atmosphere having $\\Tef=2000 \\pm 100$~K and C/O $<$ 1, and a hot component consistent with Teff=18000 - 23000 K and log g =4.0 $\\pm$ 0.5; this agrees well with previous studies \\citep{tylenda05}. The cool component is slightly better reproduced with a moderately sub-solar metallicity model."
    },
    "0609/astro-ph0609155_arXiv.txt": {
        "abstract": "}[2]{{\\footnotesize\\begin{center}ABSTRACT\\end{center} \\vspace{1mm}\\par#1\\par \\noindent {~}{\\it #2}}}  \\newcommand{\\TabCap}[2]{\\begin{center}\\parbox[t]{#1}{\\begin{center} \\small {\\spaceskip 2pt plus 1pt minus 1pt T a b l e} \\refstepcounter{table}\\thetable \\\\[2mm] \\footnotesize #2 \\end{center}}\\end{center}}  \\newcommand{\\TableSep}[2]{\\begin{table}[p]\\vspace{#1} \\TabCap{#2}\\end{table}}  \\newcommand{\\FigCap}[1]{\\footnotesize\\par\\noindent Fig.\\  % \\refstepcounter{figure}\\thefigure. #1\\par}  \\newcommand{\\TableFont}{\\footnotesize} \\newcommand{\\TableFontIt}{\\ttit} \\newcommand{\\SetTableFont}[1]{\\renewcommand{\\TableFont}{#1}}  \\newcommand{\\MakeTable}[4]{\\begin{table}[htb]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newcommand{\\MakeTableSep}[4]{\\begin{table}[p]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newenvironment{references}% { \\footnotesize \\frenchspacing \\renewcommand{\\thesection}{} \\renewcommand{\\in}{{\\rm in }} \\renewcommand{\\AA}{Astron.\\ Astrophys.} \\newcommand{\\AAS}{Astron.~Astrophys.~Suppl.~Ser.} \\newcommand{\\ApJ}{Astrophys.\\ J.} \\newcommand{\\ApJS}{Astrophys.\\ J.~Suppl.~Ser.} \\newcommand{\\ApJL}{Astrophys.\\ J.~Letters} \\newcommand{\\AJ}{Astron.\\ J.} \\newcommand{\\IBVS}{IBVS} \\newcommand{\\PASP}{P.A.S.P.} \\newcommand{\\Acta}{Acta Astron.} \\newcommand{\\MNRAS}{MNRAS} \\renewcommand{\\and}{{\\rm and }} {  We present $VI$ photometry for the open cluster NGC 637 which is located in the Cassiopeia region. Morphology of cluster color-magnitude diagram indicates that it is a young object with age of a few million years. The apparent distance modulus of the cluster is $13.9<(m-M)_V<14.3$, while reddening is $0.69<E(V-I)<0.73$. We estimated the heliocentric distance as $2.6<d<3.3$ kpc. We also report the identification of two variable stars in NGC 637. One of the variables is a non-radial pulsating $\\beta$ Cep-type star. Other one is a likely ellipsoidal variable, however its pulsating nature can not be excluded. }  {Hertzsprung-Russell (HR) and C-M diagrams -- Stars: variables : beta Cep -- open clusters and associations: individual: NGC 637} ",
        "introduction": "CURious Variables Experiment (CURVE) is a long-term project focused on observations of open clusters, globular clusters and cataclysmic variable stars in the northern hemisphere (Olech \\etal 2003a, Olech \\etal 2003b). In stellar clusters we principally search for variable objects. However, our data also allow to estimate basic parameters of observed clusters, such as distances and ages.  NGC 637, also known as C 0139+637 or OCl 329, is a not very abundant open cluster situated in the Cassiopeia region of the Perseus spiral arm of our Galaxy ($l$=128\\zdot\\arcd55, $b$=+1\\zdot\\arcd73). The first photometric studies of this cluster were done by Grubissich (1975) in $RGU$ system. He estimated reddening and distance to the cluster. Later, Huestamendia \\etal (1991) published photoelectric $UBV$ observations and assessed age of the cluster as 15 Myr. NGC 637 was also observed by Phelps and Janes (1994) as one of 23 open clusters to explore the star formation history of the Cassiopeia region. They found that the previous authors had overestimated the age of the cluster and suggested NGC 637 to be 0-4 Myr old.  This paper presents CCD $VI$ photometry and results of a year-long monitoring of NGC 637 in search for variable objects. We also give characteristic parameters of the cluster. ",
        "conclusions": "We have presented the CCD $VI$ photometry and results of searches for variable stars in the open cluster NGC 637. The analysis of the derived color-magnitude diagram allowed us to assess some basic parameters. The apparent distance modulus for the cluster is $13.9<(m-M)_V<14.3$, the reddening is $0.69<E(V-I)<0.73$, and the estimated heliocentric distance of $2.6<d<3.3$kpc. The morphology of the CMD indicates that it is a young object with age of a few million years. Based on the comparison of the CMD with the theoretical isochrones, we were also not able to firmly establish the metallicity of the cluster. Spectroscopic observations would help to clarify it.  The search for variability led to discovery of two variable stars in NGC 637. One of them, V1, is a bona fide $\\beta$ Cep-type variable. It has two close periodicities what clearly indicates non-radial pulsations. The light curve of object V2 has nearly sinusoidal shape and small amplitude and therefore we classified it as an ellipsoidal variable. However, one can not excluded a possibility that V2 is a pulsating star.  \\Acknow{  The authors would like to thank Drs. A. Pamyatnykh and W. Dziembowski for helpful comments on detected variables and Dr. W. Pych for providing some useful software which was used in the analysis. This work was supported by Polish KBN grant number 1~P03D~006~27 to A. Olech. }"
    },
    "0609/astro-ph0609680_arXiv.txt": {
        "abstract": "{Determination of ages of powerful radio sources is crucial in understanding galaxy evolution, the activity cycle of galactic nuclei and their influence on the surrounding intergalactic medium. So far, several different methods for age estimates of classical double radio galaxies have been proposed and widely used in the literature, although each of them faces some difficulties due to observational limitations and/or freedom in choosing the underlying model assumptions.}{We propose a new approach in determining ages of FR~II type radio sources, on one hand exploiting a dynamical model developed for these objects by Kaiser et al. (1997), and on the other hand using multifrequency radio observations not necessarily restricted to the high-resolution ones.}{In particular, we apply the assumed dynamical model to a number of FR~II type radio galaxies observed at different radio frequencies, and fit --- for each frequency separately --- the model free parameters to the observed sources' quantities. Such a procedure, using enlarged in fact a number of observables, enables us to determine relatively precise ages and other crucial characteristics (like the jets' kinetic power) of the analyzed sources.}{The resulting age estimates agree very well with those obtained by means of `classical' spectral ageing method for objects not older than 10 Myr, for which good-quality spectral data are available. The presented method is however also applicable in the case of the sources older than this, and/or the ones for which the only available low-resolution radio data do not allow for detailed spectral ageing studies. Interestingly, the estimated ages always correspond to the realistic values of the jets' advance velocity $\\sim 0.01\\,c - 0.1\\,c$}{Our analysis indicates that the main factor precluding precise age determination for FR~II type radio galaxies regards the poorly known shape of the initial electron energy distribution injected by the jet terminal shocks to the expanding lobes/cocoons. We briefly consider this issue, and conclude that the broad-band single power-law form assumed here may be accurate enough for the presented estimates, although most likely it does not strictly correspond to some well-defined realistic particle acceleration process. Instead, it should be considered as a simplest model approximation of the initial electron continuum, averaged over a very broad energy range and over the age of the source, with the \\emph{effective} spectral index which may be different for different sources, however within the relatively narrow range $p = 2.0-2.4$ suggested by our modeling.} ",
        "introduction": "The dynamical age of extragalactic radio sources is of a great importance for studying and understanding details of the physical processes taking place in both Active Galactic Nuclei (AGNs) emitting energy in forms of jets boring through the ambient intergalactic medium (IGM), and also lobes or cocoon consisting of shocked (backflowing) jet material and shocked IGM. The `standard model' of classical double radio sources (Blandford \\& Rees 1974; Scheuer 1974) implies that the axial lengthening of the cocoon is governed by the balance between the jet's thrust and the ram pressure of the IGM, whilst the width of the cocoon is determined by its internal pressure. This scenario has formed the basis for all further analytical models of the dynamics and radio-emission properties of powerful double radio sources (e.g., Begelman \\& Cioffi 1989; Falle 1991; Nath 1995; Kaiser et al. 1997; Chy\\.{z}y 1997; Blundell et al. 1999; Manolakou \\& Kirk 2002, Kino \\& Kawakatu 2005).  All such dynamical models allow one to predict a number of observational parameters of the sources (e.g., linear size, width of the cocoon, radio power at different wavelengths) for their given ages and redshifts, assuming values of the models' basic free parameters (e.g., jet power, external medium density and its gradient, etc.). A comparison of the predictions of the three main scenarios, i.e. the ones by Kaiser et al. (1997), Blundell et al. (1999) and Manolakou \\& Kirk (2002), was recently undertaken by Barai \\& Wiita (2005). They argued that {\\sl no} existing model could give acceptable fits to the common properties of radio sources from the survey considered (3CRR, 6CE, 7CRS), though the Manolakou \\& Kirk model gave `better overall results than do either of the other two'.  The age of a classical double radio source can, in principle, be determined from its apparent radio spectrum. There is no doubt that a radio continuum spectrum in different parts of a radio source should contain crucial information about its history, depending on various energy losses (and gains) of the radiating particles. This simple idea, relying on the determination of a single characteristic break frequency in the radio spectrum (at a given part of the source) caused by the above time-dependent losses, was used to estimate the radiative ages and expansion speeds in large samples of powerful 3CR sources (Myers \\& Spangler 1985; Alexander \\& Leahy 1987; Leahy et al. 1989; Carilli et al. 1991; Liu et al. 1992), as well as in samples of low-luminosity and intermediate-luminosity radio galaxies (Klein et al. 1995; Parma et al. 1999). However, the spectral break and steepening of the spectrum beyond this break need not be entirely due to radiative ageing. A possible role of the local magnetic field structure and evolution, details of the backflow process of the cocoon material, or the difficulties in disentangling various energy losses of the radiating particles have been raised in a number of papers (e.g. Rudnick et al. 1994; Eilek \\& Arendt 1996; Jones et al. 1999).  The task of resolving some of the above difficulties was undertaken by Kaiser (2000; hereafter referred to as K2000) who has constructed a 3-dimensional model of the synchrotron emissivity for a cocoon of a FR~II-type radio source, and applied it to the self-similar model of the cocoon's dynamics developed by Kaiser \\& Alexander (1997; hereafter referred to as KA) and its extension by Kaiser et at. (1997; hereafter KDA). In the K2000 model, the cocoon is split into small volume elements whose surface-brightness evolution is traced individually. The author has argued that the proposed careful analysis of the bulk backflow process and of the electron energy losses (both radiative and adiabatic) can provide an accurate estimate of the sources' spectral ages, being in addition in agreement with the estimates of their dynamical ages. Blundell \\& Rawlings (2000) have contended however that this could only be the case if these ages were much less than 10 Myr.  Analyzing the K2000 model, we realized that in the majority of the extended FR~II-type radio sources --- even those without distorted lobe structures --- the surface-brightness profiles are far from the expected smooth shapes, causing the fitted free parameters to be highly uncertain. In addition, the K2000 method requires rather high resolution observations of the radio lobes. But usually high-resolution observations of low-brightness sources cause a serious loss of the flux density, so that application of the K2000 method is in fact limited to a few strongest sources such as Cyg\\,A.  The KDA model has been successively used by Machalski et al. (2004a, b) to derive the basic physical parameters of the FR~II-type radio sources with known age, namely the jet power, the central density of the ambient medium, the cocoon's energy density and pressure, and the total energy deposited in the radio source. All of these values were derived by means of fitting the model free parameters to the observed parameters of the source: its redshift, monochromatic radio luminosity at 1.4 GHz, projected linear size, and the cocoon's axial ratio. In this paper we propose a simplified method of the age determination for powerful radio sources based on a further exploitation of the KDA model, which avoids the limitations present in the K2000 method. In particular, we show that performing the above fitting procedure for a {\\sl few} observing frequencies, i.e. enlarging the number of the observables used by adding the radio spectrum data of the entire cocoon, a {\\sl unique} set of the age and the values for the two fundamental model parameters --- the jet power and the central density of the ambient medium --- can be found. This, in turn, enables us to determine relatively precisely the age of the analyzed source, even if the available (low-resolution) radio data do not allow for detailed spectral ageing studies.  The paper is organized as follows. First, a brief presentation of the analytical KDA model together with its application to the dynamical ageing studies are given in Section 2. In Section 3 we present the dataset used in the subsequent analysis. Exemplary results of the dynamical age determination for the radio galaxy Cygnus~A and its dependence on the assumed model parameters are given in Section 4. Application of the proposed method to a variety of other FR~II type radio sources is summarized in Section 5. Remarks on the implications of the obtained results for modeling the electron injection spectrum in lobes of powerful radio sources are given in Section 6. Final conclusions are presented in Section 7. ",
        "conclusions": "Here we propose a new approach to determining the ages of powerful radio sources, on the one hand exploiting a dynamical model developed for these objects by Kaiser et al. (1997; KDA), and on the other hand using multifrequency radio observations not necessarily restricted to the high-resolution ones. The KDA model is chosen for its minimal number of free parameters when compared to the more sophisticated scenarios by Blundell et al. (1999) or Manolakou \\& Kirk (2002).  We apply the assumed dynamical model to a number of FR~II type radio galaxies observed at different radio frequencies, and fit --- for each frequency separately --- the model free parameters to the sources' observed quantities. Such a procedure, enlarging in fact a number of observables, enables us to determine, relatively precisely, ages and other crucial characteristics (such as the jet kinetic power) of the analyzed sources. We show, that the resulting age estimates extend from about 1 Myr for the small-size but powerful high-redshift radio galaxies to over 100 Myr for the giant-sized radio galaxies. The estimated ages agree very well with the ones obtained by means of `classical' spectral ageing method for objects not older than 10 Myr, for which good-quality data are available. Meanwhile, in the case of sources older than 10 Myr, our age estimates are systematically larger (by a factor of $2-4$) than the ages obtained by means of spectral ageing analysis. Such a behavior is in fact consistent with the conclusions of Blundell \\& Rawlings (2000), who argued that the spectral ageing analysis is most likely erroneous for the lobes older than 10 Myr.  Interestingly, the ages of FRII-type sources estimated by using the presented method correspond always to `realistic' values of the jet advance velocity $\\leq 0.3\\,c$, i.e. consistent with those measured directly with the VLBI technique in many of the CSO-type sources (see section 5.2.). Thus, we believe that the method presented here is applicable also in the case of the oldest radio galaxies, and/or ones for which the only available low-resolution radio data do not allow for detailed spectral ageing studies (i.e., for majority of powerful radio sources). This implies that the applied procedure of favoring the value of the injection spectral index which corresponds to the minimum of the cocoons' total energy (see section 4.3), although to some level arbitrary, nevertheless gives very reasonable results.  Our analysis indicates that the main factor precluding precise age determination for FR~II type radio galaxies regards the poorly known shape of the initial electron energy distribution injected by the jet terminal shocks to the expanding lobes/cocoons. We consider briefly this issue, and conclude that the broad-band single power-law form assumed here (and in all the other analogous models) may be accurate enough for the presented estimates, although most likely it does not strictly correspond to some well-defined realistic particle acceleration process. Instead, it should be considered as a simplest model approximation of the initial electron continuum, averaged over a very broad energy range and over the age of the source. The \\emph{effective} spectral index of such a distribution may be different for different sources, however within the relatively narrow range $p = 2.0-2.4$ suggested by our modeling.  In a subsequent paper we indend to apply the presented method for the age determination in a large sample of FR~II-type (but also younger, MSO-type) radio sources."
    },
    "0609/astro-ph0609363_arXiv.txt": {
        "abstract": "Our aim is to study the evolution of tidal dwarf galaxies.  The first step is to understand whether a model galaxy without Dark Matter can sustain the feedback of the ongoing star formation.  We present tests of the evolution of models in which star formation efficiency, temperature threshold, initial distribution of gas and infall are varied.  We conclude that it is feasible to keep a fraction of gas bound for several hundreds of Myr and that the development of galactic winds does not necessarily stop continuous star formation. ",
        "introduction": "Tidal tails, resulting from the interactions of galaxies, can create self-gravitating structures which are the seeds of dwarf galaxies.  In these galaxies, commonly known as {\\it Tidal Dwarf Galaxies} (TDGs), the Dark Matter (DM) content is very small, which make them more vulnerable to stellar feedback.  1-D DM-free models of dwarf galaxies have already demonstrated that they can remain bound even with a large fraction of supernova-expelled gas (Hensler et al. \\cite{htg04}). ",
        "conclusions": "From our models we found that TDGs are not necessarily destroyed rapidly and can sustain SF for times longer than 200 Myr and that galactic winds do not always suppress SF. Further work is necessary to sharpen our understanding of these issues."
    },
    "0609/astro-ph0609649_arXiv.txt": {
        "abstract": "The variational method with constraints recently developed by Verkley and Gerkema to describe maximum-entropy atmospheric profiles is generalized to ideal gases but with temperature-dependent specific heats. In so doing, an extended and non standard  potential temperature is introduced that is well suited for tackling the problem under consideration. This new formalism is successfully applied to the atmosphere of Venus. Three well defined regions emerge in this atmosphere up to a height of $100\\, km$ from the surface: the lowest one up to about $35\\, km$  is adiabatic, a transition layer located at the height of the cloud deck  and finally a third region which  is practically isothermal. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/hep-th0609155_arXiv.txt": {
        "abstract": "We investigate a string-inspired dark energy scenario featuring a scalar field with a coupling to the Gauss-Bonnet invariant. Such coupling can trigger the onset of late dark energy domination after a scaling matter era. The universe may then cross the phantom divide and perhaps also exit from the acceleration. We discuss extensively the cosmological and astrophysical implications of the coupled scalar field. Data from the Solar system, supernovae Ia, cosmic microwave background radiation, large scale structure and big bang nucleosynthesis is used to constrain the parameters of the model. A good Newtonian limit may require to fix the coupling. With all the data combined, there appears to be some tension with the nucleosynthesis bound, and the baryon oscillation scale seems to strongly disfavor the model. These possible problems might be overcome in more elaborate models. In addition, the validity of these constraints in the present context is not strictly established. Evolution of fluctuations in the scalar field and their impact to clustering of matter is studied in detail and more model-independently. Small scale limit is derived for the perturbations and their stability is addressed. A divergence is found and discussed. The general equations for scalar perturbations are also presented and solved numerically, confirming that the Gauss-Bonnet coupling can be compatible with the observed spectrum of cosmic microwave background radiation as well as the matter power spectrum inferred from large scale surveys. ",
        "introduction": "General relativity predicts singularities. Therefore, in spite of its being a highly successful as a classical theory of gravitation, its modification seems inevitable at high energy scales. The low energy action could then also feature corrections to the Einstein-Hilbert term from additional fields and curvature invariants. Among the possible corrections, a particular combination of the quadratic Riemann invariants, the Gauss-Bonnet term, is of special interest. It appears in extensions of gravitational actions, whether motivated by the form of most general of a scalar-tensor theory, uniqueness of the Lagrangian in higher dimensions or the leading order corrections from string theory. An interesting possibility then arises that the modifications could change the way the universe itself gravitates at large scales.  Indeed, the current cosmological observations cannot be explained by the standard model of particle physics and general relativity. Precise cosmological experiments have confirmed the standard Big Bang scenario of a flat universe undergoing an inflation in its earliest stages, where the perturbations are generated that eventually form into structure in matter. Most of this matter must be non-baryonic, dark matter. Even more curiously, the universe has presently entered into another period of acceleration. Such a result is inferred from observations of extra-galactic supernovae of type Ia (SNeIa) \\cite{Riess:2004nr,Astier:2005qq} and is independently supported by the cosmic microwave background radiation (CMBR) \\cite{Spergel:2006hy} and large scale structure \\cite{sloan} data. It seems that some dark energy, with its negative pressure that speeds up the universal expansion, dominates the density of the universe \\cite{Copeland:2006wr}. This concordance model agrees very well with the astrophysical data, but features inflation, dark matter and dark energy as phenomenological ingredients of undisclosed nature. Perhaps dark energy is the most enigmatic of the latter two, as at least there are reports of tentative discoveries of dark matter \\cite{bertone}.  The problem of dark energy reduces to the questions 1) what is the culprit (perhaps in fundamental physics) for the acceleration, and 2) why did the speed-up begin just at the present stage of the cosmological evolution. In recent years dark energy and its cosmological and astrophysical signatures has been addressed in many papers considering both modifications of the energy-momentum tensor, with the inclusion of scalar fields \\cite{Caldwella,brookfield2,bertolami,carsten} or imperfect fluids \\cite{tomicard,koivisto} and of gravitational physics at large scales \\cite{Capozziello02,Nojiri:2006ri,Straumann:2006tv,Borowiec:2006qr}. The latter approach has proven to be perhaps surprisingly difficult. Extending the action with curvature invariants rather generically results in fourth-order gravity, which is problematical because of the implied instabilities\\cite{Woodard:2006nt}. An $f(R)$ extension, although not generically, can be stable\\cite{Faraoni:2006sy}. However, though some $f(R)$ -type modifications might be compatible with the acceleration following a standard matter dominated phase\\cite{Nojiri:2006gh,Nojiri:2006su}, the simplest modifications as proposed this far seem not to be able to generate a viable cosmology\\cite{Amendola:2006kh}. On the other hand, within the Palatini formulation, the corresponding extensions result in second order field equations, but it turns out that the observationally allowed modifications are then practically indistinguishable from a cosmological constant \\cite{tomimod2,morad,tomimod3,Li:2006ag}.  Scalar-tensor theories of gravity are interesting alternatives to the concordance model and seem to have potential to provide a linkage between the acceleration and fundamental physics \\cite{st1,st2,Albrecht:2001xt,Kainulainen:2004vk}. They have a desirable feature which is their quasi-linearity: the property that the highest derivatives of the metric appear in the field equations only linearly, so as to make the theory ghost free. Interestingly, there is a particular combination of the curvature squared terms with such behavior, known as the Gauss-Bonnet (GB) integrand. It is constructed from the metric as \\be R^2_{GB} \\equiv R^{\\mu\\nu\\rho\\sigma}R_{\\mu\\nu\\rho\\sigma} - 4R^{\\mu\\nu}R_{\\mu\\nu} + R^2. \\ee As mentioned in the beginning, this term appears frequently in attempts at a quantum gravity, especially in stringy set-ups. In fact, all versions of string theory (except Type II) in 10 dimensions ($D=10$) include this term as the leading order $\\alpha^\\prime$ correction\\cite{Callan,gross}. In $D=4$ the Gauss-Bonnet term is a topological invariant. It's appearance alone in the action can then be neglected as a total divergence (though even then this topological term can have interesting role in view of the boundary terms and regularization of the $D=4$ action in asymptotically AdS spacetimes\\cite{Aros:1999id,Olea:2005gb}). However, if coupled, its presence may lead to contributions to the field equations also in $D=4$. The low energy string action typically features scalar fields with such couplings. Namely the moduli, associated to the internal geometry of the hidden dimensions, become non-minimally coupled to the curvature terms. Hence, in a general $D=4$ scalar-tensor theory one might couple the scalar field non-minimally not only to the Ricci curvature $R$ but to the Gauss-Bonnet invariant $R_{GB}^2$ as well, and in $D>4$ it is necessary to include this term in order to preserve the uniqueness of the gravitational action.  Such a coupling could thus be useful in modeling both the early inflation and the late acceleration\\cite{Carter:2005fu,Neupane:2006dp,Tsujikawa:2006is,Carter:2005sd,Neupane:2006ip}. Here we will concentrate on the post-inflationary epochs in such a universe, and more specifically, within the scenario we recently studied\\cite{Koivisto:2006xf}. There the scalar field lives in an exponential potential, and the coupling can likewise be exponential. This kind of model was originally proposed by Nojiri, Odintsov and Sasaki\\cite{Nojiri:2005vv}. More recently it was shown that if the coupling grows steeper than the potential decays with the field, acceleration can occur with a canonic scalar, after a transition from a scaling era\\cite{Koivisto:2006xf}. Then the interaction with the Gauss-Bonnet curvature term causes the universe to enter from a scaling matter era to an accelerating era. Existence of scaling solutions with more general parameterizations has then been taken under investigation\\cite{Tsujikawa:2006ph}, noting that requiring exact scaling is otherwise possible only with a tachyon field which is non-minimally coupled with matter. Other dark energy solutions in a broad variety of second order string cosmologies, taking into account both coupled and uncoupled matter, have been explored also previously \\cite{Sami:2005zc,Calcagni:2005im}, higher order terms have been also incorporated and a method to reconstruct the coupling and the potential has been found\\cite{Nojiri:2006je}. If the scalar field is non-dynamical, the model is equivalent with a modified Gauss-Bonnet gravity, featuring a function of $R^2_{GB}$ the action \\cite{Cognola:2006eg,Nojiri:2005am}.  The coupling mechanism can also momentarily push the universe to a phantom era, until the scalar field potential begins to dominate. Therefore it is possible to get a very good accordance with the SNeIa data. Dark energy models other than of the Gauss-Bonnet form but featuring transition to phantom expansion and having possible origin in string theory have been considered also \\cite{Aref'eva:2005fu,Aref'eva:2006et}. In addition, one should clearly distinguish the models investigated here from those featuring a coupling to the Ricci scalar \\cite{Baccigalupi:2000je}, since for them the so-called R-boost \\cite{Pettorino:2005pv} plays a role in the early universe, while the potential has to be shallow in order to have accelerating behaviour at the present. On the other hand, one can couple the self-tuning scaling field to matter \\cite{Amendola:1999er,manera,brookfield1}, and thereby indeed set the acceleration ongoing and thereby perhaps alleviate the coincidence problem \\cite{Bean:2000zm,Huey:2004qv}. However, implications for the large-scale structure formation seem to make this possibility problematical \\cite{Koivisto:2005nr,nunes}. Here we will not find the same, since modifications to the growth of matter perturbations are much less drastic as the scalar field remains minimally coupled to the matter sector.  In this paper, besides investigating the cosmological evolution of such models, we also derive quantitative constraints on the scenario of Gauss-Bonnet dark energy mentioned above. Our aim is to present in more detail this scenario and to generally uncover in more depth what are the possible effects of the Gauss-Bonnet combination on dark energy cosmologies, in order to assess whether they could be compatible with present days observations, and to estimate the amount of fine-tuning that this might require. To this end the the model in Ref.\\cite{Koivisto:2006xf} is subjected to further constraints. In addition to tightening the limits for this particular parameterization, we will discuss in more detail and generality the phenomenology of Gauss-Bonnet dark energy at cosmological, astrophysical and Solar system scales during the cosmological evolution. Especially, we consider the impact of the coupling to the CMBR and large scale structure. In view of the remarks above, this is a crucial aspect of a dark energy model with any non-minimal couplings or modifications of gravity. Since the class of models studied here features both and these can lead to completely different predictions for the expansion rate of the Universe than without the $R^2_{GB}$ term, it is rather non-trivial to find that the perturbation evolution in these models is in some sense only modestly altered. Namely, the shape of the matter power spectrum is retained during the presence of the Gauss-Bonnet gravity and the effects one expects to see in the large multipoles of the CMBR spectrum are typically small. Still, there are distinguishable features peculiar to these models that can lead to essential constraints for the model parameters.  The article is organized as follows. The cosmological dynamics for these models is described in Section \\ref{cs} for both the background expansion and for linear perturbations. In Section \\ref{constraints} we discuss the implications and derive constraints of the Gauss-Bonnet coupling to nucleosynthesis, Solar system, large scale structure, CMBR and SNeIa. Section \\ref{conclusions} contains our conclusions. Some technical details are confined to the Appendixes. ",
        "conclusions": "\\label{conclusions}  We studied cosmological phenomenology of dark energy based upon a low-energy effective string theory action featuring a compactification modulus and taking into account the leading order curvature corrections. One is then lead to consider a generalized scalar-tensor theory which includes a coupling to the Gauss-Bonnet invariant. Specifically, we investigated the evolution of perturbations. There the simple closed form equation (\\ref{d_evol}) for the linear matter inhomogeneities, together with its effective gravitional constant (\\ref{g_eff}) could be highlighted as one of the main results of the paper, since it enabled to find several new {\\it model-independent results of perturbation evolution} in Gauss-Bonnet dark energy, these including the following. \\begin{itemize} \\item The evolution of matter perturbations is scale-invariant at small scales in the presence of the Gauss-Bonnet term, and thus the shape of the matter power spectrum is retained. \\item The growth rate of matter perturbations, which is easily extracted from Eq.(\\ref{d_evol}), can be compatible with observations even in the presence of significant contribution from the Gauss-Bonnet interaction. \\item The stability of perturbations can be read off from the expression for the effective gravitational constant. A divergence is possible. \\end{itemize} This gravitational constant cannot be deduced by matching Schwarzchild and FRW -type metrics for subhorizon evolution of spherical matter overdensities, as the Jebsen-Birkhoff theorem is not in general respected in these models. The Newtonian limit, as discussed in Section \\ref{solar}, does feature this effective strength of gravity, which might be used to efficiently constrain the coupling. However, the relevance of this limit to the experiments within Solar system is an open question to be addressed more carefully in forthcoming studies.  In addition to such general considerations mentioned above, we chose a particular model and scrutinized its predictions for cosmology. Specifically, we characterized the potential and the coupling by exponential functions, as with the Nojiri-Odintsov-Sasaki modulus \\cite{Nojiri:2005vv}. This parameterization is appealing in its minimality, as within it one can describe Gauss-Bonnet dark energy with just two extra degrees of freedom compared to the $\\Lambda$CDM model. Moreover, the exponential forms are well motivated on fundamental grounds, and allow dark energy solutions without introducing unnatural scales into the Lagrangian. Many of the previous cosmological studies of this type of low-energy string theory have been concerned with asymptotic solutions and confined to cases where the Gauss-Bonnet term is subdominant, and thus only introducing small deviations to standard quintessence scenarios. Here allowed a crucial role for the Gauss-Bonnet term and examined its consequences in detail. Such approach also proved to be very useful, as it revealed several {\\it possibilities of cosmological evolution present already in this simple model}, these including the following. \\begin{itemize} \\item The Gauss-Bonnet coupling can act to switch the decelerating expansion into acceleration after a scaling matter era, thus perhaps alleviating the coincidence problem. \\item The curvature interaction may momentarily push the universe into a phantom era, which will then not lead into a Big Rip singularity in the future. Hence the model can explain superacceleration. \\item However, the linear perturbations might diverge in the future. This might terminate the de Sitter phase, possibly helping to make contact with string theory. \\end{itemize} Let us briefly comment each of these possibilities.   The field is drawn into the scaling attractor virtually regardless of initial conditions, and at late times acceleration can be onset provided simply $\\alpha>\\lambda$. This can be seen as a possible way towards solution of the coincidence problem, though it is occasionally considered that in a preferable solution the future evolution of the Universe would be characterized by a constant ratio of $\\Omega_\\phi$ and $\\Omega_m$ different from 1. However, a transition from a scaling matter era to an accelerating scaling era has been shown to be impossible for a very general class of scalar field models \\cite{Amendola:2006qi}. Another remark is that the question of why dark energy is beginning its domination today is now translated into the question that why the coupling comes into play just recently. The same holds for any other coupling mechanism proposed so far, but it is still fair to say that such mechanism can provide an approach to tackle the coincidence problem. The relevant dynamics are not always obvious when only the asymptotic behaviour is considered, and therefore it was only very recently \\cite{Koivisto:2006xf, Tsujikawa:2006ph} found that the simple model considered here can (without flipping the field to a phantom, i.e. putting $\\gamma=-1$ in the action (\\ref{action})) feature a transition from a matter dominated to an accelerating phase.  We also showed that during this transition, the dynamics of the coupling can cause a phantom expansion. Again, this phenomenon is unaccessible when only asymptotic behaviour or vacuum is taken into account. Several authors have shown that phantom dark energy does not necessarily lead a Big Rip singularity in the presence of the Gauss-Bonnet term. Here we found that this term can in fact ally with a canonical scalar field into an effective phantom energy. This is clearly a more appealing way to produce $w_{eff}<-1$, since thereby the phantom era can take place without the introduction strong energy condition violating components, while the Big Rip singularity is still avoided.  However, one of the peculiarities of this class of models is that their linear perturbations can diverge, even while the background solution is not singular. We argued that such a possibility does not necessarily indicate an inconsistency of the model. For the specific model considered here, the divergence typically awaits in the future when the evolution would begin to asymptote to the de Sitter expansion. The dynamics is such that the (linearized) inhomogeneities grow explosively as the $s_{SC}<0$ limit is approached. Hence the linear approximation perhaps together with the whole FRW description breaks down. It is tempting to believe that the de Sitter phase will not then be reached but the inhomogeneities take over. This would enable to reconcile the accelerating universe with string theory in these models in yet another way. Since if the acceleration is transient one can consistently define the set of observables in string theories\\cite{Hellerman:2001yi}. Their present S-matrix formulations seem to be in odds with an eternally accelerating universe, like in the concordance model with the cosmological constant. Would the universe however somehow pass through $c^2_{S}=0$, one would have to invoke higher-order curvature corrections or modifications of the exponential parameterization to escape the possible singularity.  We calculated the implications of Gauss-Bonnet dark energy to various different cosmological and astrophysical observations and matched them with several data sets. The general conclusion we arrive at is that the quadratic curvature coupling has interesting and non-trivial signatures while being in a good agreement with a large amount of data. We checked how the Gauss-Bonnet interaction affects the evolution of scalar field fluctuations and how these in turn impact the large scale structure and CMBR observations. Using combined datasets related to the CMBR shift parameter, baryon oscillation scale, supernovae Ia luminosity distance and variations of the effective gravitational constant in the Solar system, we derived confidence limits on the parameters of the model. The study presented not only shows the requirements to be imposed on models of Gauss-Bonnet dark energy for them to exhibit completely realistic cosmology, but also quantifies the degree to which these requirements can be within the most minimalistic set-ups. Due to its simplicity, matching with observations, capability to shed light on the coincidence problem, possibility incorporate phantom (and perhaps also transient) acceleration, and not least its theoretical motivation, the Gauss-Bonnet quintessence seems to be a promising alternative to a cosmological constant.  We end by mentioning some issues yet to be pursued. Among these are the nature of the divergence and the Solar System limit in these models. Also, we found that there is some contrast with the amount of early quintessence and the nucleosynthesis bound and with the present matter density and the baryon oscillation scale. Whether these enforce one to resort to more elaborate models depends on several assumptions, like the validity of the baryon oscillation constraint and the initial conditions of the scalar field after inflation. One caveat is also that these contrasts can disappear if the Gauss-Bonnet modification affects the intrinsic evolution of supernovae and thus changes the effective luminosity-distance relationships of these objects. This does render some of our conclusions indecisive, but in an interesting way. Suppose that the coupling Gauss-Bonnet modification does influence small scale physics considerably. Then one has to consider $f(\\phi)$ such that it evolves to a constant as dark energy domination begins, so as to comply with the Solar system experiments (which is very much the approach in standard scalar-tensor theories). Then one could match all the observations, while the expansion of the universe at small $z$ could be something considerably different than commonly established from the SNeIa luminosity-redshift curves."
    },
    "0609/astro-ph0609612_arXiv.txt": {
        "abstract": "We analyze $N$-body simulations of halo mergers to investigate the mechanisms responsible for driving mixing in phase-space and the evolution to dynamical equilibrium. We focus on mixing in energy and angular momentum and show that mixing occurs in step-like fashion following pericenter passages of the halos.  This makes mixing during a merger unlike other well known mixing processes such as phase mixing and chaotic mixing whose rates scale with local dynamical time.  We conclude that the mixing process that drives the system to equilibrium is primarily a response to energy and angular momentum redistribution that occurs due to impulsive tidal shocking and dynamical friction rather than a result of chaotic mixing in a changing potential. We also analyze the merger remnants to determine the degree of mixing at various radii by monitoring changes in radius, energy and angular momentum of particles. We confirm previous findings that show that the majority of particles retain strong memory of their original kinetic energies and angular momenta but do experience changes in their potential energies owing to the tidal shocks they experience during pericenter passages. Finally, we show that a significant fraction of mass ($\\approx 40\\%$) in the merger remnant lies outside its formal virial radius and that this matter is ejected roughly uniformly from all radii outside the inner regions. This highlights the fact that mass, in its standard virial definition, is not additive in mergers. We discuss the implications of these results for our understanding of relaxation in collisionless dynamical systems.  \\end {abstract} ",
        "introduction": "Almost four decades have passed since ``violent relaxation'' was first proposed by \\citet{lyndenbell67} as a mechanism to explain the remarkable regularity and smoothness of the light distributions in elliptical galaxies. Several authors had noted that the time scale for two-body relaxation \\citep{chandrasekhar43a} was several orders of magnitude too long to be relevant for the evolution of galaxies \\citep{zwicky39}. Lynden-Bell derived the distribution function that would result from rapid fluctuations in the potential of the galaxy, such as those that might occur during gravitational collapse or following a merger. The process of violent relaxation outlined by Lynden-Bell produces a smooth distribution function as a result of the gravitational scattering of particles by the time dependent global gravitational potential. He argued that this process could result in a mass-independent and smooth final distribution of particles in phase-space.  Another process that was considered by \\citet{lyndenbell67} is ``phase mixing'' - a process by which an initially compact ensemble of non-interacting phase-space points in a time-independent potential can be stretched into a long ribbon in phase-space as a result of small differences in the initial energies (or angular momenta) of particles in the ensemble.  Phase mixing conserves the fine grained distribution function but in a coarse grained sense it produces smooth/relaxed looking distributions on timescales that depend only on the range of values of the orbital integrals (or actions) in the ensemble \\citep{BT}.  In the last decade, a new type of mixing process, which occurs due to the presence of large fractions of chaotic orbits in a galactic potential (such as when box orbits are scattered by a central supermassive black hole in a triaxial galaxy), has been investigated \\citep{kandrup_mahon94,MV96,habib97,KPS,kandrupSiopis03}. These studies have shown that microscopic ensembles of iso-energetic test particles which occupy an infinitesimal volume of phase-space surrounding a chaotic orbit can evolve rapidly to a near-invariant distribution that uniformly covers the energy surface available to the ensemble.  This process of diffusion to a near-invariant distribution is termed ``chaotic mixing'' and has been shown to occur in a variety of galaxy-type potentials. These authors argued that in a triaxial galaxy with a supermassive central black hole, a significant fraction of orbits are chaotic, and consequently, this mixing process could drive secular evolution to more axially symmetric distributions.  The physically interesting aspect of this mixing process is that the evolution to a near-invariant distribution has been shown to occur on timescales of order 5 to 100 orbital crossing times in strongly chaotic systems making it a potentially important process for the evolution of galaxies.  This was demonstrated by \\citet{merritt_quinlan98} who employed $N$-body simulations of triaxial elliptical galaxies containing growing supermassive central black holes and found the triaxial figures to evolve to more axisymmetric shapes.  Like phase mixing, chaotic mixing conserves the fine-grained phase-space density.  Unlike phase mixing, its rate does not depend on the initial spread in the energy distribution of particles in the ensemble, but on how strongly chaotic the orbits in the ensemble are \\citep{MV96}.  However two concerns remain regarding these earlier experiments. First the rate of evolution to the invariant measure depends on the properties of the ensemble of particles. For ensembles near a ``sticky'' or weakly chaotic orbit, a near-invariant distribution may not be attained in the lifetime of the galaxy \\citep{MV96}. Second, the majority of the numerical studies which found strong chaotic mixing and short time scales for evolution were not carried out in self-consistent potentials leading to concerns that in self-consistent systems such mixing could be self-limiting.  Since all orbits in a time-dependent potential with long-lived oscillations can, in principle, exchange energy with the potential, a large fraction of the orbits will not conserve the energy integral. In such systems previous authors have found that a significant fraction of orbits can be chaotic \\citep{KVS}. This raises the possibility that the underlying physical process driving relaxation in violently fluctuating potentials could be a consequence of chaotic mixing. The first step in making this connection was studied by \\citet{KVS} and \\citet{TK}, who demonstrated that large fractions of chaotic orbits arise in potentials subjected to long-duration (damped) periodic oscillations. These authors found that the fraction of chaotic orbits depended on the frequency as well as the amplitude of the oscillations in the potential and attributed the chaos to a ``broad parametric resonance''.  In this paper we analyze  the mechanisms driving the approach to equilibrium following the collisionless merger of two self-gravitating systems.  Our principal interest is in understanding whether changes in the fine-grained distribution function (e.g., changes in the angular momenta and energies of nearby particles) occur as a result of exponential instabilities of orbits or some more straightforward energy exchange process which transfers kinetic energy of the merging halos to the internal energies of particles.  We  consider two dynamical heating effects that are not usually discussed in this context.  The first is the effect of compressive tides that arise when one collisionless gravitating system passes through another on a time scale that is short compared to the internal dynamical time of the infalling system. Such tides can impulsively heat particles as a result of the transient deepening of the net potential. Impulsive compressive tidal shocks have been studied previously and are well known to produce large changes in the internal structure and evolution of globular clusters \\citep{spitzerChevalier73,gnedin99}, galaxies orbiting within a cluster potential \\citep{valluri93}, and DM subhalos orbiting within larger hosts \\citep{kravtsov_gnedin_etal04,kazantzidis_etal04b}.  The second process is dynamical friction \\citep{chandrasekhar43a}, which is the primary mechanism that causes the formation of a single remnant when two collisionless halos merge. When collisionless halos merge, dynamical friction transfers their relative linear and angular orbital momenta to the individual particles. The original version of the \\citet{chandrasekhar43a} dynamical friction theory applied to the deceleration of a massive star  moving in an infinite homogenous isotropic stellar distribution. However several studies have shown that with minor modifications the theory is applicable in systems that resemble real galaxies and clusters \\citep{White76,binney77,velazques_white99}. Dynamical friction is found to be enhanced in systems with inhomogenous density profiles (such as cuspy profiles \\citeauthor{Delpopolo03} \\citeyear{Delpopolo03}), as compared with systems with uniform homogenous distributions.  Recent $N$-body studies of live satellites in extended $N$-body galaxies \\citep{ jiang_binney00, fujii_etal06} have shown that the timescale for dynamical friction is shorter than that predicted by Chandrasekhar due to two additional components of drag force: the tidal torque from the leading tidal tail and an enhanced wake due to particles that are no longer bound to the satellite but still trail behind it in the same orbit.   It is important to emphasize at the outset, that the terms ``mixing'' and ``relaxation'' are often used rather loosely (and often interchangeably) in the astrophysics literature to mean the approach of a gravitating system to global as well as small-scale equilibrium. The most restrictive definition of ``relaxation'' is a process which leads to the complete ``loss of memory of initial conditions'' or loss of correlations between initially nearby particles in phase space. The large changes in integrals of motion required for such relaxation only occur as a result of two-body interactions \\citep{chandrasekhar43a} and as a result of ``violent relaxation''. Both processes cause changes in the fine-grained distribution function. In contrast phase mixing and chaotic mixing in time-independent potentials, under the right conditions, can lead to large changes in correlations but conserve integrals of motions (if they exist) and conserve the fine-grained distribution function. In this paper we focus our attention on loss of correlations in phase space accompanying a collisionless merger and use the term ``mixing'' to connote the approach to macroscopic equilibrium and the accompanying change in the fine-grained distribution function (even though this is quite small).   According to the prevailing cosmological model for structure formation, galaxies and dark matter (DM) halos in the Universe form hierarchically via multiple mergers. Numerous minor mergers and several major mergers of subhalos eventually form the dark halos at $z=0$. Thus, the merger history of cosmological structures is complex, and DM halos are predicted to abound in substructure and have mass profiles that may still be evolving today.  A detailed analysis of the physics of gravitational relaxation, with the specific goal of understanding what processes drive the DM and stellar component to equilibrium, can be more readily done using dissipationless controlled merger simulations.  Several aspects of the universality in the structure and properties of dark matter halos have lead to new interest in the detailed physical processes driving gravitational relaxation. The discovery of ``universal'' dark matter halo profiles resulting from cosmological $N$-body simulations \\citep[e.g.,][]{dubinski_carlberg91,NFW,navarro_frenk_white97}, the more recent realization that these profiles have only a ``nearly universal'' form with a correlation between the density profile's shape parameter and the halo mass \\citep{merritt_etal05, merritt_etal06} as well as the finding that phase-space density profiles of these halos have power-law distributions over about 2.5 orders of magnitude in radial extent \\citep{taylorN01,Arad04,ascasibar_binney05} all point to a universality in the processes driving structure formation. In particular, violent relaxation and the statistical distribution functions that result from this process are being actively investigated \\citep{aradLB05, aradJ05, dehnen05}. Several studies of the remnant density profiles have shown that the remnants retain a surprisingly strong memory of their initial density profiles despite the strength of the relaxation process \\citep{barnes96,boylan-kolchin_ma04, kazantzidis_etal06}.  Several authors have criticized Lynden-Bell's \\citeyear{lyndenbell67} original suggestion that it is potential fluctuations that occur during gravitational collapse that drive the system to equilibrium, on the grounds that time dependence of the potential alone merely causes particles to be relabeled in energy and causes no mixing of nearby particles in energy and angular momentum.  It has been therefore argued that a process like chaotic mixing is actually responsible (at the microscopic level) for driving the evolution of the system to equilibrium \\citep{MV96,KVS}.  An analysis of the merging process can help us better understand the physical mechanisms that drive the evolution to equilibrium and cause the redistribution of nearby particles in phase-space. In this investigation we restrict ourselves to the study of mixing in positions, velocities,  energies and angular momenta of particles. The goals of this study are: (a) to determine whether the relaxation of a remnant to an equilibrium distribution following a collisionless  $N$-body merger is driven by chaotic mixing that is believed to take place in time-independent potentials (b) to quantify the degree of relaxation/mixing in various phase-space parameters that occurs during the merger by examining the remnants of major mergers of DM halos, and (c) to obtain a better understanding of why steep cusps are as robust as they are.  Both the time-dependent nature of the problem, as well as the fact that $N$-body systems exhibit exponential sensitivity to initial conditions (as discussed in greater detail in \\S~\\ref{sec:micro}), make the first of the above goals particularly challenging.  The paper is organized as follows. In \\S~2 we describe the $N$-body experiments that were conducted and define several new measures that we use to quantify the rate of separation of nearby orbits and the degree of ``mixing'' or ``relaxation'' in phase-space.  In \\S~3 we discuss the results of the application of our measures of mixing to the merger experiments by considering the separation in various phase-space variables of pairs of nearest neighbors in phase-space (\\S~\\ref{sec:separation}) as well as by monitoring the rate of mixing of ensembles of nearby particles (\\S~\\ref{sec:mixing}).  In \\S~4 we examine how particles are redistributed in phase-space following the merger. In \\S~5 we summarize our results and discuss the implications of our findings for understanding relaxation during mergers and the mixing that it produces. ",
        "conclusions": "The principal goal of this study was to investigate the processes that drive the evolution of self-gravitating systems to equilibrium and cause mixing in phase-space. The goal was also to determine whether the possible presence of large fractions of chaotic orbits arising from the time-dependent potential was responsible for driving the system to equilibrium.  We summarize the main results below.\\\\    \\begin{figure} \\epsscale{0.9} \\plotone{f14.eps} \\caption{Fraction of particles that originated in each of the radial shells (1-11) in one of the original NFW halos involved in the merger that end up outside the fiducial virial radius of the remnant. (All particles beyond the virial radius are assigned to shell 11.)} \\label{fig:vir-out} \\end{figure}  \\begin{enumerate}  \\item{All orbits in the $N$-body simulation exhibit the well-documented ``Miller instability'' of the $N$-body problem.  This instability results in an initial nearly exponential separation of nearby orbits in configuration space and energy, and a decrease in velocity separation. The $e$-folding time ($t_e$) for the separation increases with increasing softening length and decreasing particle number. The initial exponential separation saturates on a short timescale ($\\sim 0.4 t_c$). The $e$-folding time and the number of $e$-folds after which the instability saturates are identical in the merging and isolated halos indicating that this is purely a consequence of the $N$-body nature of the problem and not a result of the merger. To this end, we use control simulations of isolated halos as a baseline for comparison with merging halos to identify the mixing processes occurring in mergers.}  \\item{After the initial growth in $\\ln(\\Delta r)$ and decrease in $\\ln(\\Delta |v|)$, there are periods of time for which these quantities are almost constant, separated by several sharp decreases in $\\ln(\\Delta r)$ and correspondingly sharp peaks in $\\ln(\\Delta |v|)$ that are not seen in the isolated halo.  The transient changes in $\\ln(\\Delta r)$ and $\\ln(\\Delta |v|)$ are coincident with each other and correspond to the pericenter passages of the centers of the merging halos.  Pericenter passages also cause sharp transient step-like changes in the energy and angular momentum separation of pairs of particles. Between pericentric passages there is very little evidence for for mixing in the innermost shell,  but   at larger radii$\\ln(\\Delta r)$ and  $\\ln(\\Delta |v|)$ continue to grow slowly with time due to the overlap of the merging halos.}  \\item{The changes in $\\Delta E$ and $\\Delta J$ are more strongly correlated  with pericenter passages at all radii pointing to the importance of compressive tidal shocks. At larger radii there are is also some growth in separation between pericentric passages, possibly pointing to the continued action of chaotic mixing. If we assume that separations of nearby particles in phase-space are linked to ``mixing'', this implies that the principal drivers of strong mixing are the compressive tidal shock and increased dynamical friction that occurs during the pericentric passages.}  \\item{Mixing experiments with ensembles of 1000 nearest particles in phase-space showed that mixing in the phase-space variables $\\mathcal E, \\mathcal J$ also occurs primarily during pericenter passages. Mixing leads  ensembles of nearby self-gravitating particles to spread and fill a triangular region in $\\mathcal E, \\mathcal J$ space. Mixing in these parameters occurs at the same rate at all radii, largely independent of the characteristic crossing time of the ensemble. This is further evidence that the mixing process is driven by pericentric passages and is unrelated to mixing processes that scale with the local crossing time. }   \\item{We confirm the findings of several other authors that cusps of DM halos are remarkably robust. The robustness of cusps is a consequence of the fact that 80\\% of the particles in the first shell have apocenters that lie within the first shell. During the merger only a small fraction (20\\%) of particles with apocenters within the original cusp are ejected to larger radii. The majority of the ejected particles do not get beyond shell 4. A much larger fraction ($\\sim$ 45\\%) of particles in the central regions of core-like profiles is ejected to radial shells directly outside the cusp during a merger.}  \\item{Particles in a given radial shell outside the cusp are redistributed almost uniformly with radius primarily due to phase mixing. This result holds for both the isolated spherical halo and the merging halos. The merger helps to redistribute an additional $\\sim $20\\% of orbits that have predominantly tangential velocity distributions in the original halos. }  \\item{Particles within the scale radius of the merging NFW halos experience a net decrease in total energy following the merger. This is a result of the overall decrease in the potential energy of the particles, resulting from the ($\\sim 60$\\%) increase in the mass of the cusp in the final remnant.}  \\item{The majority of particles in the merging halos experience only a small change in their kinetic energies and moderate changes in their potential energies. The change in total energy of particles is driven predominantly by the changes in their potential energies during pericenter passages. }  \\item{At smaller radii the majority of particles experience only small change in angular momentum evidenced by highly peaked distributions of $\\Delta J$. The angular momentum of the merger is absorbed primarily by particles in the outermost radial shells. }  \\item{The largest fraction of particles lying beyond the virial radius of the final remnant (25\\%) were located beyond the virial radius of the original halos. The inner 2 shells are most robust to the ejection of particles beyond the virial radius. All shells between shell 4 and the virial radius (shell 10) contribute roughly equally to the particles ejected beyond $R_{\\rm vir}$ of the remnant.} \\end{enumerate}  While this paper focuses on simulations of the merger of two dark matter halos with NFW potentials, most of the results are generic to mergers of collisionless gravitating systems and therefore applicable to dissipationless (``dry'') mergers of elliptical galaxies.  In particular, the absence of large amounts of mixing in radius over and above that expected from phase mixing in an isolated halo supports previous work that indicates that radial gradients in stellar properties (such as metallicity gradients) are can to survive intact even in major mergers \\citep{white80,barnes96,boylan-kolchin_ma04}. These results are also expected to be generically applicable to cosmological mergers. In fact soon after the submission of this paper we became aware of the work of \\citet{faltenbacher_etal06} - a study of the relaxation processes that operate in hierarchical mergers of dark matter halos in cosmological $N$-body simulations. These authors independently arrived at the same conclusions that we do, namely that the primary driver for mixing and relaxation in cosmological mergers is tidal shocking.  Finally, we find strong evidence that the processes that drive the mixing in phase space and evolution to a dynamical equilibrium do not occur continuously due to the time-dependent potential of the system. In fact strong mixing occurs primarily following episodic injection of energy and angular momentum into the internal energies of particles during pericenter passages. The injection of energy and angular momentum causes nearest particles in phase-space to separate in $E, {\\mathbf J}$.  Previous studies of orbital evolution in time-dependent and time-independent potentials have argued that the presence of large fractions of chaotic orbits could be the principal driver of the mixing and evolution to an equilibrium distribution.  While chaotic mixing could well be occurring, we do not find evidence that such chaotic mixing is driving the relaxation. This is probably a consequence of the fact that the timescales required for chaotic mixing to be relevant are much longer than the duration of the merger.  Two  interesting questions remain regarding ``violent relaxation'' 1) does the mixing that occurs during collapse of an isolated cold  gravitating system give rise to similar mixing in phase-space? 2) what gives rise to the universal power-law phase-space density profiles seen in cosmological mergers? We defer addressing both questions to future work."
    },
    "0609/astro-ph0609338_arXiv.txt": {
        "abstract": "Lensing tomography with multi-color imaging surveys can probe dark energy and the cosmological power spectrum. However accurate photometric redshifts for tomography out to high redshift require imaging in five or more bands, which is expensive to carry out over thousands of square degrees.  Since lensing makes coarse, statistical use of redshift information, we explore the prospects for tomography using limited color information from two or three band imaging. With an appropriate calibration sample, we find that it is feasible to create up to four redshift bins using imaging data in just the $g$, $r$ and $i$ bands. We construct such redshift sub-samples from mock catalogs by clustering galaxies in color space and discarding regions with poorly-defined redshift distributions. The loss of galaxy number density decreases the accuracy of lensing measurements, but even losing half or more of the galaxies is not a severe loss for large area surveys. We estimate the errors on lensing power spectra and dark energy parameters with color tomography and discuss trade-offs in survey area and filter choice. We discuss the systematic errors that may change our conclusions, especially the information needed to tackle intrinsic alignments. ",
        "introduction": "Over the last decade a concordance model has emerged in cosmology in which about two-thirds of the energy density of the universe today may be in the form of ``dark energy''. This explains the observation that we reside in an accelerating universe \\citep{Riess,Perl}.  Despite its importance to the formation and evolution of the universe there are no compelling theories that explain the energy density nor the properties of the dark energy.  To address questions about the nature of dark energy a number of ambitious wide-field optical and infrared imaging surveys have been proposed. These range from space-based missions in the optical and infrared, such as the Supernova Acceleration Probe (SNAP\\footnote{http://www.snap.lbl.gov}, proposed as the space-based Joint Dark Energy Mission (JDEM)), to ground-based surveys such as the Panoramic Survey Telescope \\& Rapid Response System (Pan-STARRS\\footnote{http://pan-starrs.ifa.hawaii.edu}), the Dark Energy Survey (DES\\footnote{ http://www.darkenergysurvey.org}), the Large Synoptic Sky Survey (LSST\\footnote{http://www.lsst.org}) and others. Each of these missions approaches the study of dark energy using multiple, complementary observational probes: gravitational weak lensing (WL) to study the growth of structure and geometry, baryon oscillations to measure the angular diameter distance vs. redshift relation and Type Ia supernovae to measure the  luminosity distance vs. redshift relation.  In this paper we focus on one of these probes, weak lensing, and how the design of a survey might impact its scientific value for constraining dark energy. Observations to date have succeeded in measuring the amplitude $\\sigma_8$ of the $z\\approx0$ dark matter power spectrum to $\\approx10\\%$ accuracy ({\\it cf.} \\citet{Jarvis05,CFHLS05} and references therein) by surveying $\\simeq 100$~deg$^2$ of sky. To constrain the equation of state of dark energy requires that we improve the accuracy of these measures by close to an order of magnitude. To achieve this deep imaging surveys have been proposed that cover surveys of 1,000-20,000 square degrees in 5-6 filters. There is, however, a natural trade off in the design of these surveys; depth vs number of filters vs area surveyed. To date the appropriate weighting of these components is not fully understood. For example, given the time and cost of a survey, does a large area, shallow survey in a small number of filters provide more scientific return than a deep survey in a large number of filters but sampling only a few hundred square degrees?  One of the primary requirements for any lensing application is the need for redshift information on the source galaxies. In the first studies of lensing only the statistical redshift distribution was available, but now surveys aim to get photometric redshifts (Connolly \\etal \\ 1995; hereafter photo-z's) for individual galaxies. Accurate photo-z information enables new qualitative and quantitative advances in lensing. Conversely photo-z errors can be a limiting systematic in the use of lensing for precision cosmology.   Lensing tomography refers to the use of depth information in the source galaxies to get three-dimensional information on the lensing mass \\citep{Hu99}. By binning source galaxies in photo-z bins, the evolution of the lensing power spectrum can be measured. This greatly improves the sensitivity of lensing to dark energy in cosmological applications. The lensing power spectra measured from two redshift bins is shown in Figure \\ref{fig:cl}. The two auto-spectra and the cross-spectrum, along with expected statistical errors, are shown.  The relative shift in the amplitudes of the lensing spectra is sensitive to the properties of dark energy.  It depends on both distances and the growth of structure, thus enabling tests of dark energy or modified gravity explanations for the cosmic acceleration. In Figure \\ref{fig:cl}, the amplitude shifts of the three spectra (for a given cosmological model) depend on the mean redshifts and widths of the two photo-z bins; clearly any errors in the bin redshifts will degrade the ability to discriminate cosmological models.  Thus the use of broadband multicolor photometry to estimate the distances of galaxies (i.e.\\ photometric redshifts; henceforth abbreviated as photo-z's) is critical to the use of lensing for dark energy studies. The capability of these surveys to meet their scientific goals will depend on our ability to characterize the uncertainties present within photometric redshift estimates, i.e. the scatter, bias and fraction of outliers. Photometric redshifts must be calibrated with an appropriate sample of spectroscopic redshifts (Huterer \\etal \\ 2005; Ma \\etal \\ 2005). This may be done more cheaply by using auto- and cross-correlations of photometric and spectroscopic redshifts samples \\citep{Newman2006,Schneider2006}, which can also be used to estimate the redshift distribution for a galaxy sample where the calibration data is incomplete (see also Zhan \\& Knox 2006).  In this paper we study the possibility of doing lensing tomography from limited color information. Imaging in five or more bands is needed for well-characterized photo-z's of galaxies extending beyond $z\\sim 1$. We consider here whether lensing tomography can be carried out from a wide-area imaging survey in just two or three filters, along with a relatively small calibration sample that provides the statistical redshift distributions in all parts of the color space. These can be used to create a few sub-samples of the full galaxy sample that occupy distinct redshift bins, while discarding galaxies with colors that lead to badly defined redshift distributions. Since lensing does not require redshift bins much narrower than $0.2$-$0.4$ in redshift, and does not need a full or fair sample of the galaxy population, there is reason to expect that limited color information may be sufficient for tomography. We use a mock catalog of galaxies with known redshifts to see how well one can make cosmological measurements with this approach, taking into account statistical errors in lensing measurements. In a future paper we will expand upon this analysis to include various sources of systematic errors in colors/photo-z's and in the lensing measurements.  In Section 2 we describe the formalism for computing lensing power spectra given the redshift distribution of a galaxy sample. In Section 3 the mock catalog used for our study is described. Section 4 contains the results on how well one can do with color cuts in making well separated redshift distributions for galaxy sub-samples. The errors on the power spectra and dark energy parameters are compared with the forecasts for idealized photo-z's. In Section 5 we discuss the prospects for tomography with color cuts from imaging in a limited number of filters -- which we call color tomography. ",
        "conclusions": "We have studied the prospects for lensing tomography with a wide area imaging survey that has limited color information. While it is never desirable to limit color information, it may be that a full complement of filters is not available at a survey telescope, that it is un-realistic to image the full survey area in all filters, or that additional science requirements for a survey dictate an observing strategy that is not optimal for lensing tomography. The questions we address in this study are whether lensing tomography is at all feasible with two or three band imaging data, and if so, what is the trade-off between filter choice and sky coverage.  We use the fact that lensing tomography makes only statistical use of estimated redshifts and, given the width of the lensing kernel, the redshift bins can be quite broad (though they must be known accurately). Further, for planned wide-area surveys lensing measurements are not shot noise limited -- hence a significant fraction of galaxies can be discarded if their redshifts are ill-defined without severely degrading parameter accuracy. With these considerations in mind, we generate a mock catalog of galaxies extending to high redshift with known types, redshifts and colors.  This provides an estimate of the redshift distributions for each part of $g-r$ and $r-i$ color space. From this we select regions in color space that produce sub-samples with redshift distributions well suited for tomography.  Errors on the power spectra in different redshift bins from this color tomography are then compared with what would be expected with idealized photo-z's.  The resulting degradation in the accuracy of cosmological parameters is shown in Tables \\ref{tab:lens} and \\ref{tab:lens+cmb}, and Figures \\ref{fig:fisher1} and \\ref{fig:fisher2} for dark energy parameters from color tomography. With three-band imaging, the errors on dark energy parameters are only modestly degraded compared to idealized six-band photo-z's. Even with two-band imaging, lensing tomography may be feasible for the case of high signal-to-noise data and an adequate calibration sample. This result may have implications for how to trade-off filter choice and sky coverage for large area surveys and how to optimize survey strategy to maximize the scientific returns in its initial stages. For example, imaging in three filters instead of six is more than twice as efficient in terms of survey time since the $u$ and $y$ filters (on the blue and red end) are substantially less efficient. Thus, for a fixed survey duration, the survey area for a six band program would be less than half that of the three band survey, with all parameter errors scaling as $f_{\\rm sky}^{-1/2}$. Such trade-off studies must be carried out for the specific parameters and available choices in observing strategy for a given survey; the fiducial parameters we have considered do not correspond to any real or planned survey known to us.  For color tomography an adequate calibration sample is essential. The calibration will likely require a two-step approach: a sample of spectroscopic redshifts of over $\\sim 10^4$ galaxies (for a $\\sim 1000$ square degree survey), and a larger sample of multi-band imaging with a full set of optical and possibly infra-red filters.  Imaging to a depth equivalent to $r \\sim 25$ in say six bands over $10$-$20$ square degrees would provide a sample of well measured photo-z's for over $1$ million galaxies. This would enable us to map the redshift distributions over the desired color space: e.g., it could provide photo-z's of over $10^4$ galaxies in each bin of $\\sim 0.1\\times 0.1$ magnitude in two-colors (assumed to be available for the full survey area).  It has been recognized that even with five or six band imaging data, photo-z's would need to be carefully calibrated with spectroscopic redshifts so that the means and widths of redshift bins are known to high accuracy (Bernstein \\& Jain 2004; Huterer \\etal \\ 2006; Ma, Hu and Huterer 2006), and proposals for calibration are being developed (Newman 2006). For color tomography the calibration sample is even more important, as it forms the basis for grouping and discarding regions in color space. We leave for future work the detailed requirements and strategy for obtaining the calibration sample.  Our current study is based on estimates of statistical errors in the lensing power spectra -- in future work we will include systematic photometric uncertainties and other systematic errors that affect lensing. These will include a realistic modeling of photometric errors, filter design, and will address whether the spectral templates are adequate for describing high-redshift galaxies. The inclusion of these realistic errors is essential before one can make detailed trade-off studies of survey duration and filter choice. Several caveats are in order until such a study is done: systematic errors may lower the signal-to-noise and this may depend on scale, testing for systematics is more difficult with coarser redshift bins, and marginalizing over intrinsic ellipticity contributions may be harder (see below). Finally, the inclusion of other statistics such as higher order correlations and cross-correlations with the galaxy distribution can alter the results on parameter accuracies. Some of these issues can be tested on ongoing multicolor surveys, such as the CFHT Legacy Survey.  Our results on the power spectrum errors in Figure \\ref{fig:power_ratio} show that the degradation in errors is worse at high $\\ell$, because the shot noise term dominates on small scales. This raises the question: should one choose different galaxy samples at different $\\ell$: be more conservative at low $\\ell$, since number density matters less, to minimize redshift overlap and bias? This merits a detailed study as it has the potential to impact different survey strategies.  Such a study must include non-Gaussian contributions that increase the sample variance at high $\\ell$; this in fact lowers the degradation for color tomography as the shot noise regime shifts to higher $\\ell$ (e.g. Kilbinger \\& Schneider 2005).  Finally, intrinsic ellipticity alignments (e.g. Heymans \\etal \\ 2006) and ellipticity-shear alignments (Hirata \\& Seljak 2004) must be considered. While we did choose as our fiducial sub-samples ones with well separated redshift distributions, it may be necessary to only use cross-spectra to eliminate intrinsic ellipticity contributions (e.g. Takada \\& White 2004).  This would increase the parameter errors from having a smaller number of redshift bins. Similarly ellipticity-shear correlations may need to be measured from the data and marginalized over, which is easier to do with finer redshift binning. Four redshift bins are the minimum needed to separately fit for both kinds of intrinsic ellipticity correlations, this may impose a minimal requirement on the needed color information. The impact of intrinsic alignments is likely to be the most important issue for future studies related to color tomography."
    },
    "0609/astro-ph0609424_arXiv.txt": {
        "abstract": "With recent and archival {\\em Rossi X-Ray Timing Explorer} ({\\em RXTE}) X-ray measurements of the heavily obscured X-ray pulsar EXO 1722-363 (IGR J17252-3616), we carried out a pulse timing analysis to determine the orbital solution for the first time. Using a single pulse period derivative, we connected datasets separated by over 2.4 years without pulse number ambiguity ($> 180,000$ pulses), implying continuous spin-up of the pulsar from 2003 August (and possibly earlier) until 2006 February (and possibly later). The orbital solution also shows that a torque reversal occurred sometime between 1998 November and 2003 August. The binary system is characterized by $a_{x} \\sin{i} = 101 \\pm 3$ lt-s and $P_{\\rm orb} = 9.7403 \\pm 0.0004$ days (90\\% confidence), with the precision of the orbital period being obtained by connecting datasets separated by more than 7 years (272 orbital cycles). The orbit is consistent with circular, and $e < 0.19$ at the 90\\% confidence level. The mass function is 11.7 $\\pm$ 1.2 $M_{\\sun}$ and confirms that this source is a High Mass X-ray Binary (HMXB) system. The orbital period, along with the previously known $\\sim$414 s pulse period, places this system in the part of the Corbet diagram populated by supergiant wind accretors. Using previous eclipse time measurements by Corbet et al. and our orbital solution, combined with the assumption that the primary underfills its Roche lobe, we find $i > 61\\degr$ at the 99\\% confidence level, the radius of the primary is between 21 $R_{\\sun}$ and 37 $R_{\\sun}$, and its mass is less than about 22 $M_{\\sun}$. The acceptable range of radius and mass shows that the primary is probably a supergiant of spectral type B0I--B5I. Photometric measurements of its likely counterpart are consistent with the spectral type and luminosity if the distance to the system is between 5.3 kpc and 8.7 kpc. Spectral analysis of the pulsar as a function of orbital phase reveals an evolution of the hydrogen column density suggestive of dense filaments of gas in the downstream wake of the pulsar, with higher levels of absorption seen at orbital phases 0.5--1.0, as well as a variable Fe K$\\alpha$ line. ",
        "introduction": "EXO 1722-363 (IGR J17252-3616) was first discovered by {\\em EXOSAT} scan observations in 1984 along with 7 other sources in the Galactic ridge and 14 sources outside the Galactic ridge (Warwick et al. 1988). A pointed {\\em Ginga} observation in 1987 October for a total of 80 ks discovered pulsations with a 413.9 $\\pm$ 0.2 s period (Tawara et al. 1989). The source exhibited significant variability, with the 6--21 keV flux decreasing from about 2 mCrab to 0.2--0.3 mCrab over 8 hrs. Using an absorbed power-law model, they found $N_{\\rm H} = 1.26_{-0.47}^{+0.74} \\times 10^{24}$ cm$^{-2}$, suggesting a dense envelope of circumstellar matter surrounding the pulsar. Additional {\\em Ginga} observations were performed in 1988 March--April by Takeuchi et al. (1990). Using pulse timing analysis, lower limits on the orbital period and mass of the companion were set to 9 days and 15 $M_{\\sun}$, assuming a canonical neutron star mass of 1.4 $M_{\\sun}$. Takeuchi et al. (1990) also measured the spectrum but included a high energy cut-off, and found a lower column density $N_{\\rm H} = 1.26_{-0.26}^{+0.33} \\times 10^{23}$ cm$^{-2}$.  Using {\\em RXTE} Proportional Counter Array (PCA) scans of the Galactic center region from 1999 February to 2003 October, Corbet el al. (2005) measured a 9.741 $\\pm$ 0.004 day orbital period through periodic changes in the source flux. By folding the scan observations on this orbital period, eclipses were detected lasting 1.7 $\\pm$ 0.1 days (eclipse half-angle $\\theta_{e} = 31\\fdg8 \\pm 1\\fdg8$). Corbet et al. (2005) also measured the spectrum using pointed PCA observations from 1998 October--November, and found that the spectrum varies but not dramatically, although the spectrum was sampled over limited ranges of orbital phase.  EXO 1722-363 has also been observed by {\\em INTEGRAL} during surveys of the Galactic plane. Zurita Heras et al. (2006) found that the source is persistent in the 20--60 keV band but is not detected above 60 keV. On a time scale of $\\sim$3 days, the 20--60 keV source flux varied by a factor of about 4. Flaring activity was also observed, with the source flux increasing by a factor of 5 and returning to its base level in about a day. Using {\\em INTEGRAL}, the pulse period was measured to be 413.7 $\\pm$ 0.3 s in 2003 August. Thus, aside from slight ($\\sim$1 s) changes, the neutron star pulse period has been stable for over a decade. Phase-resolved spectroscopy did not show any variation of the continuum emission with pulse phase. Using a follow-up {\\em XMM-Newton} observation from 2004 March, a 414.8 $\\pm$ 0.5 s pulse period was measured. More importantly, however, {\\em XMM-Newton} measured the source position to 4\\arcsec~accuracy, and an IR counterpart 2MASS J17251139-3616575 is located 1\\arcsec~away from the best position. The possible IR counterpart is not detected in the 2MASS survey in the {\\em J}-band but appears in the {\\it H}-band and {\\it K}$_{s}$-band with 11.8 mag and 10.7 mag respectively.  In this work, we use recent and archival {\\em RXTE} PCA observations to determine the orbital parameters accurately for \\igr~using pulse arrival time analysis. Since the orbital period for this system was previously known, the new observations were timed to provide even coverage in orbital phase, allowing us to obtain firm constraints on the orbital parameters, and especially the eccentricity. The observations are described in \\S~2, as well as the pulse arrival time measurements and error analysis. The results of the pulse timing analysis are presented in \\S~3.1. We confirm the 9.74 day orbital period and increase its precision by an order of magnitude. The orbital solution is then used to observe orbital modulation of the flux, pulse fraction, and pulse profile in \\S~3.2, to constrain the mass and radius of the mass-donating primary in \\S~3.3, and to study the orbital modulation of the source spectrum in \\S~3.4. We discuss the results of the spectral analysis in \\S~4. Finally, in \\S~5 we summarize our main results.  \\begin{deluxetable*}{cclcccccc} \\tablenum{1} \\tabletypesize{\\scriptsize} \\tablecolumns{9} \\tablewidth{0pt} \\tablecaption{\\sc{RXTE Observation Log of EXO 1722-363}} \\tablehead{ \\colhead{ObsID\\tablenotemark{a}} & \\colhead{Epoch} & \\colhead{Date} & \\colhead{Exp. Time} & \\colhead{Count Rate\\tablenotemark{b}} & \\colhead{Pulse Min.\\tablenotemark{c}} & \\colhead{Orb. Phase\\tablenotemark{d}} & \\colhead{Spec. Num.\\tablenotemark{e}} \\\\ \\colhead{} & \\colhead{} & \\colhead{(U.T.)} & \\colhead{(ks)} & \\colhead{(Counts/s/PCU)} & \\colhead{} & \\colhead{} & \\colhead{} } \\startdata                                            % 30142-01-02-01 & 1 & 1998 Oct 23.08--23.25 & 14.80 & 24.77 & 10 & 0.67 & \\nodata \\\\ 30142-01-02-00 & 1 & 1998 Oct 27.14--27.31 & 15.14 & 6.30 & 4 & 0.09 & \\nodata \\\\ 30142-01-03-00 & 1 & 1998 Oct 31.93--01.10\\tablenotemark{f} & 14.22 & 13.13 & 7 & 0.58 & \\nodata \\\\ 30142-01-04-00 & 1 & 1998 Nov 05.00--05.17 & 14.16 & 4.18 & 0 & 0.00 & \\nodata \\\\ 30142-01-05-00 & 1 & 1998 Nov 09.07--09.25 & 15.84 & 7.89 & 4 & 0.42 & \\nodata \\\\ \\hline 80424-01-01-00 & 2 & 2003 Aug 22.21--22.22 & 0.96 & 9.40 & 1 & 0.78 & \\nodata \\\\ 80424-01-02-00 & 2 & 2003 Aug 29.58--29.61 & 2.10 & 14.30 & 1 & 0.54 & \\nodata \\\\ \\hline 91080-02-04-00 & 3 & 2006 Jan 16.92--17.00 & 4.92 & 4.49 & 0 & 0.00 & 1 \\\\ % 91080-02-06-00 & 3 & 2006 Jan 17.90--17.99 & 5.04 & 5.21 & 0 & 0.10 & 3 \\\\ % 91080-02-14-00 & 3 & 2006 Jan 21.84--21.92 & 4.86 & 13.73 & 4 & 0.50 & 11 \\\\ % 91080-02-08-00 & 3 & 2006 Jan 28.63--28.71 & 4.98 & 7.43 & 1 & 0.20 & 5 \\\\ % 91080-02-01-00 & 3 & 2006 Feb 04.05--04.13 & 4.74 & 7.94 & 2 & 0.86 & 18 \\\\ % 91080-02-02-00 & 3 & 2006 Feb 04.52--04.55 & 2.82 & 5.22 & 0 & 0.91 & 19a \\\\ % 91080-02-01-01 & 3 & 2006 Feb 04.57--04.60 & 2.10 & 5.28 & 0 & 0.91 & 19b \\\\ % 91080-02-10-** & 3 & 2006 Feb 08.38--08.46 & 4.86 & 25.32 & 4 & 0.30 & 7 \\\\ % 91080-02-16-00 & 3 & 2006 Feb 11.26--11.30 & 3.54 & 13.50 & 3 & 0.60 & 13 \\\\ % 91080-02-12-00 & 3 & 2006 Feb 19.18--19.22 & 3.36 & 20.90 & 3 & 0.41 & 9 \\\\ % 91080-02-15-** & 3 & 2006 Feb 20.56--20.65 & 3.18 & 13.51 & 3 & 0.55 & 12 \\\\ % 91080-02-18-00 & 3 & 2006 Feb 22.00--22.04 & 3.42 & 13.29 & 3 & 0.70 & 15 \\\\ % 91080-02-20-00 & 3 & 2006 Feb 23.05--23.09 & 3.18 & 9.82 & 3 & 0.81 & 17 \\\\ % 91080-02-17-** & 3 & 2006 Mar 03.26--03.40 & 3.18 & 8.76 & 1 & 0.65 & 14 \\\\ % 91080-02-19-** & 3 & 2006 Mar 04.24--04.39 & 3.24 & 6.90 & 1 & 0.76 & 16 \\\\ % 91080-02-03-** & 3 & 2006 Mar 06.20--06.29 & 2.88 & 4.18 & 0 & 0.95 & 20 \\\\ % 91080-02-05-** & 3 & 2006 Mar 07.12--07.20 & 3.30 & 4.12 & 0 & 0.05 & 2 \\\\ % 91080-02-07-** & 3 & 2006 Mar 08.10--08.19 & 3.42 & 6.05 & 0 & 0.15 & 4 \\\\ % 91080-02-09-** & 3 & 2006 Mar 09.08--09.17 & 3.48 & 8.27 & 1 & 0.25 & 6 \\\\ % 91080-02-11-** & 3 & 2006 Mar 10.13--10.22 & 3.42 & 8.68 & 2 & 0.36 & 8 \\\\ % 91080-02-13-** & 3 & 2006 Mar 11.04--11.13 & 4.02 & 13.05 & 3 & 0.45 & 10 \\\\ % \\enddata \\tablecomments{Epochs are separated by horizontal lines in the table.} \\tablenotetext{a}{Observation IDs ending with asterisks (**) represent two or more merged observations.} \\tablenotetext{b}{For 3--24 keV photons after background subtraction.} \\tablenotetext{c}{The number of pulse minima obtained after excluding light segments that are consistent with a constant flux (see text).} \\tablenotetext{d}{The orbital phase at the middle of the observation, inferred from the fit presented in \\S~3.1. The phases typically extend about $\\pm$0.01 in either direction. } \\tablenotetext{e}{Epoch 3 spectra number labels for Tables 4 \\& 5, ordered by orbital phase. Observations 91080-02-02-00 and 91080-02-01-01 were summed for spectral analysis.} \\tablenotetext{f}{This observation ended in November.} \\end{deluxetable*} ",
        "conclusions": "We have found the orbital solution to the eclipsing HMXB \\igr~by applying pulse timing analysis to twenty recent {\\em RXTE} observations (nearly evenly spaced in orbital phase) and two sets of observations occurring 2.4 years and 7.3 years earlier.  With the orbital solution, we constrained the nature of the mass-donating companion star, and we investigated the evolution of the spectrum through the pulsar's orbit. Our primary results are summarized as follows:  \\begin{enumerate}  \\item The circular orbital solution is characterized by $a_{x} \\sin{i} = 101 \\pm 3$ lt-s, $P_{\\rm orb} = 9.7403 \\pm 0.0004$ days, and an epoch of mid-eclipse at HJD 53761.68 $\\pm$ 0.04. The eccentricity is less than 0.19 at the 90\\% confidence level. Our orbital solution also implies that a reversal of the torque on the neutron star occurred between 1998 November and 2003 August.  \\item Using the orbital solution and the eclipse half-angle measurement of $\\theta_{e} = 31\\fdg8 \\pm 1\\fdg8$ from Corbet et al. (2005), we constrained the companion star's mass and radius as a function of inclination. Assuming a neutron star mass of 1.4 $M_{\\sun}$ and that the primary underfills its Roche lobe and, the system must have an inclination greater than 61\\degr~with 99\\% confidence. With the lower limit on the inclination, the primary's mass is between about 11 $M_{\\sun}$ and 22 $M_{\\sun}$, and its radius is between 21 $R_{\\sun}$ amd 37 $R_{\\sun}$. The range of acceptable masses and radii are consistent with a supergiant star with a spectral type between B0 and B5.  \\item There is only one 2MASS IR counterpart in the 4\\arcsec~{\\em XMM-Newton} error circle. Assuming this source is the primary, we found that the photometric measurements from 2MASS and DENIS are consistent with the inferred spectral type if $A_{V}=20.1 \\pm 0.6$ and the distance to the system is between 5.3 kpc and 8.7 kpc. In the 2--24 keV band, the unabsorbed X-ray luminosity is $L_{x} \\sim$ 0.2--4.4 $\\times 10^{36}$ ergs s$^{-1}$.  \\item The spectrum of \\igr~shows variable flux, absorbing column density, power-law photon index, and iron line equivalent width through the course of the orbit. Between orbital phases 0.25--0.45, $N_{\\rm H}$ is consistent at about $10^{23}$ cm$^{-2}$, but apparently increases for the second half of the orbit.  \\end{enumerate}"
    },
    "0609/astro-ph0609754_arXiv.txt": {
        "abstract": "% In preparation for James Webb Space Telescope (JWST), a set of cross calibration programs with HST and Spitzer for suitable primary photometric standards and astrometric fields were developed.  NICMOS/HST and IRAC/Spitzer photometry observations of new solar analog standards in NGC\\ 6791 and Melotte\\ 66 were executed. These new photometric standards will provide $\\sim 5$\\% photometric precision at V~$\\sim$~19 from the near-IR to the mid-IR wavelength range for efficient on-orbit calibration and measuring of photometric stability of the JWST complement.  For the astrometric calibration, a field in the LMC has been selected. This field falls within the JWST continuous viewing zone, within 5\\deg from the ecliptic poles, and has the stellar density necessary to achieve accuracies better than 1 mas with HST/ACS. These independent observations will play a key role in meeting the mission requirements and will allow a fast commissioning of the observatory. ",
        "introduction": "The James Webb Space Telescope (JWST), led by NASA's Goddard Space Flight Center  in collaboration with the European Space Agency (ESA) and the Canadian Space Agency (CSA), is a 6-M class observatory optimized for the near and mid IR. It has a  three-mirror anastigmatic design with a primary constructed from 18 hexagonal deployable berylium mirror segments. A large sunshield will protect the Optical Telescope Element (OTE) and instruments from the sun and earth radiation and will enable passive cooling of the telescope and instruments. The imaging and spectroscopic capabilities will span the 0.6-28 $\\mu m$ region with a diffraction-limited image quality at 2 and 4 $\\mu m$. To minimize station-keeping maneuvers, the observatory will have a second Lagrange point (L2) orbit with a period of about 6 months about the L2 point \\citep{Stal04}.  The main scientific goals of the observatory span many orders in time, space, and size, from the first light objects, reonization, and the assembly of galaxies  to the birth of stars, proto-planetary systems, planetary systems and the origins of life \\citep{Gaal06}. JWST will address these themes with its four main scientific instruments: the Near Infrared Camera (NIRcam), the Near Infrared Spectrograph (NIRSpec), the near-infrared Tunable Filter Imager (TFI), and the Mid Infrared Instrument (MIRI). Table \\ref{tab-1} summarizes the main instrument characteristics.  \\begin{table} \\caption{Science Instrument Characteristics \\label{tab-1}} \\smallskip {\\footnotesize \\begin{tabular}{lccccc} \\tableline \\noalign{\\smallskip} Instrument & Wavelength & Optical Elements & FOV & Plate Scale & Pixels Format\\\\ & ($\\mu m$) & & & (mas/pixel)  \\\\ \\noalign{\\smallskip} \\tableline \\tableline \\noalign{\\smallskip} NIRcam & 0.6-2.3 &fixed filters& 2\\farcm2 x4\\farcm4 & 34\t& Two 4096$^2$ \\\\ (Short Wave) &  &  R$\\sim$ 4, 10,100 & &\t& \\\\ &  & coronagraphic spots  & & & \\\\ (Long Wave)& 2.3-5    & & & 68 & Two 2048$^2$ \\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} NIRspec & 0.6-5 & Grating (R=1000) & 3\\farcm4x3\\farcm1 & 100& Two 2048$^2$ \\\\ &  & Prisim (R=100) & &  & \\\\ &  & slits, MSA & &  & \\\\ \\tableline \\noalign{\\smallskip} MIRI & 5-27 & filters &1\\farcm4x1\\farcm9 & 110 &1024$^2$   \\\\ (imaging) &  & coronagraphic spots &(26\\arcsec x26\\arcsec & &   \\\\ &  &  phase masks &  Chor.)& &   \\\\ (spectroscopic) &   5-10 & prism (R $\\sim$ 100) & & & \\\\ &5-27 &  IFU (R$\\sim$ 3000) & 3\\farcs6x3\\farcs6 to& 200 to &  Two 1024$^2$  \\\\ &           &                                        & 7\\farcs5x7\\farcs5 & 470 &    \\\\ \\tableline \\noalign{\\smallskip} FGS-TF &1.2-4.5 & filters + & 2\\farcm3x2\\farcm3 & 68 & 2048$^2$ \\\\ & &  etalon (R$\\sim$100) & & & \\\\ \\noalign{\\smallskip} \\tableline \\end{tabular} } \\end{table}  To accomplish these objectives, we have designed a set of cross calibration programs with the HST and the Spitzer observatories. These will provide accurate absolute photometric and spectrophotometric calibrators, as well as the required observatory astrometric accuracies. This paper presents the motivation behind the current  cross-calibration programs and sumarizes the targets and the strategy followed in order to achieve the required accuracies. Here, we discuss only the photometric and astrometry cross-calibration programs. For the details of the spectrophotometric cross-calibration, please refer to the paper of R. Bohlin elsewhere in these proceedings. ",
        "conclusions": ""
    },
    "0609/astro-ph0609081_arXiv.txt": {
        "abstract": "HII regions surrounding supermassive black holes (BHs) in an otherwise still neutral intergalactic medium (IGM) are likely to be the most easily detectable sources by future 21cm experiments like LOFAR. We have made predictions for the size distribution of such HII regions for several physically motivated models for BH growth at high redshift and compared this to the expected LOFAR sensitivity to these sources.  The number of potentially detectable HII regions does not only depend on the ionisation state of the intergalactic medium and the decoupling of the spin temperature of the neutral hydrogen from the cosmic microwave background (CMB) temperature, but is also strongly sensitive to the rate of growth of BHs at high redshift. If the supermassive BHs at redshift $6$ were built up via continuous Eddington-limited accretion from low mass seed BHs at high redshift, then LOFAR is not expected to detect isolated QSO HII regions at redshifts much larger than 6, and only if the IGM is still significantly neutral.  If the high-redshift growth of BHs starts with massive seed BHs and is driven by short-lived accretion events following the merging of BH hosting galaxies then the detection of HII regions surrounding supermassive BHs may extend to redshifts as large as $8-9$ but is still very sensitive to the redshift to which the IGM remains significantly neutral. The most optimistic predictions are for a model where the supermassive BHs at $z>6$ have grown slowly.  HII regions around supermassive BHs may then be detected to significantly larger redshifts. ",
        "introduction": "Reionisation - the change of state of the intergalactic medium (IGM) from a neutral state at recombination to the present day ionised state - is an epoch which is currently under intense study. Analyses of the spectra of high redshift QSOs indicate a smooth increase in the optical depth of the IGM with redshift which accelerates at $z>6.1$ (Fan et al. 2002, Songaila 2004, Fan et al. 2006).  Quasars at $z>6.1$ show almost complete absorption of the flux blue-ward of the Lyman$-\\alpha$ resonance [the Gunn-Peterson (1965) effect].  Inferring the ionisation state of the IGM at these high redshifts from QSO absorption spectra is difficult and ambiguous (e.g. Oh \\& Furlanetto 2005, Bolton \\& Haehnelt 2006).  If the small amounts of flux transmitted just blue-ward of the Lyman$-\\alpha$ centre are attributed to the presence of a locally ionised region embedded within a significantly neutral medium, then lower limits on the neutral fraction of $x_{\\rm HII}$ may be inferred (Wyithe \\& Loeb~2004a; Mesinger \\& Haiman~2004). Malhotra \\& Rhoads~(2005), however, argue that the number density of Lyman-$\\alpha$ emitters can be used to place a lower limit on the ionised fraction of gas at $z \\sim 6$ at $20 - 50$~per cent. Fan et al. (2006) also come to the conclusion that the IGM probed by the highest redshift QSOs is predominantly ionised.  The optical depth to Thompson scattering inferred from the recent WMAP measurements offers an independent constraint on the integrated reionisation history. Interpreted as a sudden reionisation at a redshift $z_{\\rm reion}$ the year three WMAP results give $z_{\\rm reion} = 10^{+2.7}_{-2.3}$ (Spergel et al. 2006).  The emerging picture is consistent with the reionisation of cosmic hydrogen being complete by $z \\sim 6 - 13$ and therefore significant contribution to reionisation by sources beyond current detection limits.  Experiments to detect the redshifted 21cm signature of high redshift neutral hydrogen will potentially provide copious amounts of information on the epoch  of reionisation. The European design for a low frequency interferometer, LOFAR, may be in operation as early as 2007 and other low-frequency experiments (see \\textsection \\ref{prospects}) are in various stages of planning and development.  However, mapping the high redshift neutral gas distribution in 21cm will be extremely challenging.  The redshifted 21cm signal due to ionised regions produced by bright, high redshift QSOs surrounded by an otherwise still neutral intergalactic medium may be our best bet for an detection of individual objects due to their 21cm signature (see Tozzi et al. 2000 for a detailed discussion).  Kohler et al.~(2005) used numerical simulations to show that ionised regions around high redshift QSOs are indeed likely to be the most prominent features in the high redshift neutral gas distribution.  Wyithe, Barnes \\& Loeb~(2005) use a semi-analytic model for the high redshift QSO HII region distribution and predict detectable HII regions at $z \\sim 6 - 8$.  However, the growth history of high redshift supermassive BHs is very uncertain. In this paper we will investigate a range of models for high redshift BH growth which reproduce the bright end of the luminosity function of redshift $6$ QSOs and sample the range of plausible growth histories at higher redshift.  We begin by reviewing the physics of detecting a 21cm signal due to neutral hydrogen (\\textsection \\ref{21cm}). In section \\ref{QLF_models} we describe our models for the evolution of the QSO luminosity function (QLF).  We discuss the expected size of HII regions due to QSOs in \\textsection \\ref{HII_regions}.  Finally, in \\textsection \\ref{prospects}, we compare our computed size distributions to the expected sensitivity of LOFAR\\footnote{http://www.lofar.org/} to extended sources and discuss observational prospects.  Throughout this paper we adopt a cosmological matter density $\\Omega_m = 0.27$, baryonic matter density $\\Omega_b = 0.044$, cosmological constant $\\Omega_{\\Lambda} = 0.73$, present day Hubble constant $H_o \\equiv  100h$~km~s$^{-1}$~Mpc$^{-1} = 71$km~s$^{-1}$~Mpc$^{-1}$, a mass variance on scales of $8 h^{-1}$~Mpc $\\sigma_8 = 0.84$ and a scale invariant primordial power spectrum (slope $n=1$). ",
        "conclusions": "We have investigated the prospects of using the 21cm signature of the regions around supermassive black holes at $z>6$ to probe the growth history of these supermassive black holes and the ionisation history of their surrounding IGM.   We have presented three simple, physically motivated models for the high redshift growth history of supermassive black holes in which the growth is either  simple Eddington limited accretion or due to short-lived accretion phases triggered by the merger of galaxies.  We have used a simple analytical model for the size of HII regions in order to generate predictions for the size distribution of HII regions for each of our model QLFs. Our results show that models which are all consistent with the luminosity function of bright QSOs at $z = 6$ but differ in the growth history of the supermassive black holes give very different results for the predicted number of detectable HII  regions.   Our estimates for the number of detectable HII regions at $z \\sim 6$ are promising, up to several hundred detectable sources per LOFAR field around active QSOs and possibly many more fossil regions. However, the number of detectable HII regions depends sensitively on unknown details of SMBH growth and the ionisation state of the IGM. If the QLF evolves to $z >6$ as predicted in our merger-driven model, then isolated HII regions will be prominent features of the high redshift neutral hydrogen distribution providing reionisation was completed late ($z < 7$).  However if the SMBH population is growing rapidly due to continuous accretion, or reionisation occurred much earlier than $z \\sim 6$, then the detection of QSO HII regions may prove elusive.  The most optimistic picture emerges if supermassive black holes grow more slowly over an extended period of time as in our Eddington limited accretion model with a smaller duty cycle ($f_{\\rm duty}= 0.1$). In this case detectable sources may be present to very large redshifts and the prospects for detection are good even if reionisation is completed relatively early.   Observation of future 21cm experiments like LOFAR should  not only provide a wealth of information on how the Universe was reionised but may also provide important clues for the growth history of supermassive black holes at very high redshift."
    },
    "0609/astro-ph0609562_arXiv.txt": {
        "abstract": "The diversity of Type Ia supernova (SN Ia) photometry is explored using a grid of 130 one-dimensional models. It is shown that the observable properties of SNe~Ia resulting from Chandrasekhar-mass explosions are chiefly determined by their final composition and some measure of ``mixing'' in the explosion. A grid of final compositions is explored including essentially all combinations of $^{56}$Ni, stable ``iron'', and intermediate mass elements that result in an unbound white dwarf. Light curves (and in some cases spectra) are calculated for each model using two different approaches to the radiation transport problem. Within the resulting templates are models that provide good photometric matches to essentially the entire range of observed SNe~Ia.  On the whole, the grid of models spans a wide range in $B$-band peak magnitudes and decline rates, and does not obey a Phillips relation. In particular, models with the same mass of $^{56}$Ni show large variations in their light curve decline rates. We identify and quantify the additional physical parameters responsible for this dispersion, and consider physically motivated ``cuts'' of the models that agree better with the Phillips relation, discussing why nature may have preferred these solutions. For example, models that produce a constant total mass of burned material of 1.1 $\\pm 0.1$ \\msun\\ do give a crude Phillips relation, albeit with much scatter. If one further restricts that set to models that make 0.1 to 0.3 \\msun\\ of stable iron and nickel isotopes, and then mix the ejecta strongly between the center and 0.8 \\msun, reasonable agreement with the Phillips relation results, though still with considerable spread. We conclude that the supernovae that occur most frequently in nature are highly constrained by the Phillips relation and that a large part of the currently observed scatter in the relation is likely a consequence of the intrinsic diversity of these objects. ",
        "introduction": "\\lSect{intro}  The application of Type Ia supernovae (SNe~Ia) to cosmological distance determination has yielded revolutionary insights into the structure and composition of the universe \\citep{Per99,Rie98}. The utility of such explosions is based upon two empirical observations. First, that Type Ia supernovae are, for the most part, approximate standard candles, even without any corrections being applied. This probably reflects their common origin in the explosion of a white dwarf of standard mass \\citep[see review by][] {Hil99}. Second, a large part of the residual diversity in peak brightness can be corrected for by use of either light curve template fitting or observed correlations between light curve decline rate and peak brightness - the so called ``Phillips Relation''\\citep{Phi93}, or ``width-luminosity relation'' (henceforth WLR)\\footnote[1]{It should be noted that a correlation of the peak luminosity and the decline rate - with the correct sign - was already established by \\citet{Psk77}. See the history and other references in \\citet{Phi05}.}.  While the empirical relation between brightness and width has worked well for most purposes, as we move into an era of ``precision cosmology'', one must feel increasing unease at the lack of an agreed upon standard model for how Type Ia supernovae explode and the possibility that evolutionary or environmental factors may erode accuracy at large distances. If the mass of $^{56}$Ni made in a supernova is the dominant physical parameter affecting both brightness and light curve shape, what are the magnitude and direction of other possible parameters of the explosion such as kinetic energy, intermediate mass element production, stable iron production and mixing?  And if, as we shall find, not all physically plausible models obey a WLR, why has nature chosen to realize frequently a particular subset? That is, what does the WLR tell us, not just about cosmology, but about supernova models? If one can make progress on both these questions, then it may be possible to derive more rigorous and quantitative tools for using Type Ia supernovae for distance determination. That is the long term goal of our investigation.  In this paper, we introduce a simple parameterized approach to computing SN~Ia explosions and use it to generate a large grid of one-dimensional models.  For each model, we calculate synthetic broadband light curves and examine the relationship between peak $B$ magnitude and $B$-band decline rate.  Such an approach allows us to study the physical parameters affecting the light curves of SNe~Ia and the WLR.  Several previous theoretical studies have performed similar investigations for a given set of models \\citep{Khokhlov_93,Hoe96a,Hoe98,Pin01, Hoeflich_99by,Mazzali_WLR}.  These studies, however, were typically confined to a certain class (or classes) of theoretical explosion paradigms.  In this study, rather than adopt a specific theoretical framework, we take a general and exhaustive approach by simulating the full range of final ejecta structures conceivably arising from the disruption of a Chandrasekhar-mass carbon-oxygen white dwarf. ",
        "conclusions": "The multi-band photometry (and, in some cases, the spectra) of a large set of Type Ia supernova models have been studied using two approaches to the radiation transport problem. These models were constructed in a simple fashion (\\Sect{toys}) that allowed complete control over the major parameters affecting the outcome - the explosion energy and the masses of \\Nifs, stable iron, and IME produced. The models were exploded by depositing, uniformly, an amount of energy corresponding to the change in composition in a standard white dwarf model. All sensitivity to the uncertain physics of the actual burning was thus absorbed into the three model parameters, plus some prescription for mixing. Comparison of the light curves and spectra of sample models with those of more complicated models and to observations of SN Ia gave good agreement.  Using these models, which consisted essentially of every explosion one could produce starting from a 1.38 \\msun \\ carbon-oxygen white dwarf, the necessary conditions were determined for reproducing the observed relation between decline rate and peak magnitude (both in the $B$-band). The physical parameters leading to intrinsic dispersion in that relation were also identified and quantified.  The relation between peak magnitude and color at peak magnitude was studied as well, and found to be more robust to parameter variation.  The set of all explosions with positive kinetic energy did {\\sl not} give a WLR. Instead a large range of decline rates, quite inconsistent with observations, was found for SN Ia with the same mass of \\Nifs. If SN Ia light curves are a single parameter family, that parameter cannot be as simple as just the mass of \\Nifs. Other possible cuts of the model set were thus considered based upon total kinetic energy (or mass burned), and various restrictions on the final composition and mixing.  In order to satisfy the WLR, the ejecta structures of SNe~Ia must obey certain constraints.  First, all SNe~Ia must have a common total mass burned (and hence kinetic energy). In the present calculations, good agreement with observations was found if that common burned mass was $1.1 \\pm 0.1$ \\msun\\ and the kinetic energy at infinity was $1.2 \\pm 0.2 \\times 10^{51}$ erg.  Second, the radial distribution of iron group elements (including \\Nifs) in the ejecta must be fairly uniform among all SNe~Ia.  We found the best agreement with the observed WLR among those models where the distribution of iron group elements extended from near the center of the ejecta to $\\sim$0.8 \\msun.  For example, one can consider the subset of models in which the total mass of iron-group elements is nearly constant at $\\sim 0.8 \\pm 0.1$~\\Msun, with variations in \\Mni\\ arising from difference in the ratio of \\Nifs\\ to stable iron group species produced in the explosion.  Differences in this ratio are indeed predicted to arise from metallicity and electron-capture effects, although it is not clear that the entire range of SN~Ia luminosities can be so explained given the constraints provided by the underlying explosion physics and Galactic nucleosynthetic measurements.  Various arguments lead to the conclusion that the mass of stable iron (mostly $^{54,56}$Fe and $^{58,60}$Ni) is generally restricted to the range 0.1 to 0.3 \\msun \\ with about half of that range due to metallicity effects.  Alternatively, one could vary the amount of \\Nifs\\ at the expense of IME.  In this case, good agreement with the observed WLR is found as long as the inner layers of ejecta ($\\la 0.8$~\\Msun) are fairly well mixed. In the present calculations, well-stratified ejecta structures do not reproduce the observed slope of the WLR, as the lower \\Mni\\ models have \\Nifs\\ concentrated relatively closer to the ejecta center, which tends to slow their light curve decline rate. To populate the entire observed span of the WLR, especially at faint luminosities, the mass of these IME is quite large in some instances, arguing for a delayed detonation.  The present calculations adopted an LTE approximation for both the level populations and the line source functions (i.e., the thermalization parameter, $\\epsilon_{\\rm th}$ = 1).  Although a common approach in many previous transfer studies, this (and other simplifying approximations adopted) can be expected to lead to errors in the model observables on the order of 0.1 to 0.3 mag.  The relative differences between models, however, may be known more accurately, and so the general trends obeyed by sets of models and the level of {\\sl dispersion} in such trends may be considered more robust.  For example, our choice $\\epth = 1$ overestimates the true redistribution probabilities and leads to maximum light colors that are too red by $\\sim 0.06$~mag relative to the observations in the color-peak-magnitude plot of, e.g., Fig.~\\ref{Fig:bmv_constrain}. Nevertheless, the clustering of the points in this plot is quite tight and the slope is correct.  The possible refined calibration of the radiative transfer calculations and its impact on the study of the WLR will be explored further in a subsequent paper.  Given the large dispersion found, even in selected subsets of our models (e.g., Fig.~\\ref{Fig:all_constrain}), the narrow spread in the observed relation between \\Mb\\ and \\dmfb\\ is surprising. No single physical parameter yields it, unless several other parameters are highly constrained. Part of that spread can be reduced by a judicious choice of mixing and a tight constraint on the mass burned (Fig.~\\ref{Fig:nw_all}). Even so, it is difficult to avoid the conclusion that a major fraction of the observed scatter of the WLR, $\\gtaprx 0.1$~mag, reflects intrinsic physical diversity, and not observational effects. This conclusion has important implications as one plans for future studies using SN Ia as calibrated standard candles to ever higher precision.  Given the dispersion expected from the models themselves, a larger sample of supernovae will have to be studied to get a strong signal to noise ratio for cosmological effects.  A primary benefit of the approach adopted in this paper is the ability to construct an expansive grid of SNe~Ia explosion models without restricting ourselves to any specific theoretical explosion paradigm. Such a large and general database of model light curves and spectra nicely compliments the large sample of SNe~Ia data sets currently being acquired by ongoing observational programs.  A direct comparison of the models to individual observations would be useful in interpreting the physical properties of any given SN~Ia event.  In addition, the model grid should be helpful in developing refined techniques for calibrating the luminosities of SNe~Ia so as to limit and reduce intrinsic sources of dispersion or evolution.  For example, our studies suggest that an improved strategy might be {\\sl spectroscopic} template fitting in which as much data as possible from each individual SN Ia is compared to a template of model spectra and colors. The models employed could be the ones computed here or some derivative of that set."
    },
    "0609/astro-ph0609032_arXiv.txt": {
        "abstract": "{ {Accretion rates from disks around pre-main sequence stars are of importance for our understanding of planetary formation and disk evolution. We provide in this paper estimates of the mass accretion rates in the disks around a large sample of Herbig Ae stars.} {We obtained medium resolution 2 $\\mu$m spectra  and used the results to compute values of   \\Macc\\  from the measured luminosity of the \\Brg\\ emission line, using a well established   correlation between L(\\Brg) and the accretion luminosity \\Lacc.} {We find that 80\\% of the stars, all of which have evidence of an associated circumstellar disk, are accreting matter, with rates $3\\times 10^{-9}\\simless \\dot M_{acc} \\simless 10^{-6}$ \\Msun/yr; for 7 objects, 6 of which are located on the ZAMS in the HR diagram, we do not detect any line emission. Few HAe stars (25\\%) have \\Macc$>10^{-7}$ \\Msun/yr. } {In most HAe stars the accretion rate is sufficiently low that the gas in the inner disk, inside the dust evaporation radius, is optically thin and does not prevent the formation of a  puffed-up rim, where dust is directly exposed to the stellar radiation. When compared to the \\Macc\\ values found for lower-mass stars in the star forming regions Taurus and Ophiuchus, HAe stars have on average higher accretion rates than solar-mass stars; however, there is a lack of very strong accretors among them, probably due to the fact that they are on average older.} } ",
        "introduction": "Pre-main sequence stars accrete matter from their circumstellar disks for a large fraction of their life, until the  disk is accreted entirely onto the star or is dissipated by a number of potentially competing processes, such as dynamical perturbations by close companions, photoevaporation by the host or nearby  stars, or planet formation. Although the amount of matter accreted is small compared to the mass of the central object, accretion of matter through the disk has important consequences on the physics of the disk itself and on the properties of the planetary systems that the disk may form.  The mass accretion rate through the disk is reasonably well known for pre-mains sequence stars with masses from about one solar mass (T Tauri stars or TTS), down to brown dwarfs (BDs). Results can be found, e.g., in Gullbring et al.~\\cite{Gea98}, Muzerolle et al.~(\\cite{Mea98}, \\cite{Mea03}, \\cite{Mea05}), White and Ghez (\\cite{WG01}), Natta et al.~(\\cite{Nea04}, \\cite{Nea06}), Mohanty et al.~(\\cite{Subu05}).    It has been much more difficult to  determine \\Macc\\  in stars more massive than about 1 \\Msun. Recently Calvet et al.~(\\cite{Cea04}) obtained  the accretion rate in nine young pre-main sequence stars with masses in the interval 1.5--3.7 \\Msun\\ (intermediate mass TTS or IMTTS) using HST and optical spectra. They found an average accretion rate of $\\sim 3\\times 10^{-8}$ \\Myr, with a spread of an order of magnitude at least. Muzerolle et al.~(\\cite{Mea04}) derived a similar value of \\Macc\\ for the $\\sim 2$ \\Msun\\ star UX Ori by fitting  its \\Ha\\ profile with magnetospheric accretion models.  We report in this paper measurements of the accretion rate in a large sample of Herbig AeBe objects (intermediate mass pre-main sequence stars), mostly of spectral type  A (HAe). We  derive the accretion rate  from the luminosity of the \\Brg\\ line seen in emission in pre-main sequence accreting objects. Calvet et al.~(\\cite{Cea04}) have shown that the correlation between the hydrogen IR line luminosity and the accretion luminosity found in  TTS  and BDs (Muzerolle et al.~\\cite{Mea98}; Natta et al.~\\cite{Nea04}) extends to their sample of IMTTS (see also van den Ancker ~\\cite{vda04}). Establishing the luminosity of the hydrogen near-IR recombination lines as proxy of the accretion luminosity has been extremely important, since it has made possible to derive  accretion rates with reasonably good accuracy in objects where measuring the UV veiling is challenging. This is the case, for example, of  very young star-forming regions, such as \\Roph, where the high extinction is the limiting factor (e.g., Natta et al.~\\cite{Nea06}), but also of the Herbig AeBe stars, where it is difficult to separate the UV excess emission from the hot photospheric continuum.  The main purpose of this work  is to estimate \\Macc\\ in  objects with disks for which we are currently obtaining near-IR interferometric data with the ESO VLTI (AMBER GTO\\footnote{ http://www.eso.org/observing/proposals/gto/amber/}; see also Malbet et al.~\\cite{Malb06}, Tatulli et al.~\\cite{Tea06}). Knowledge of \\Macc\\ is crucial for  understanding  the disk geometry, in particular   the nature of the disk inner rim (e.g., Natta et al.~\\cite{Nea01}; Dullemond et al.~\\cite{DDN01}; Isella \\& Natta \\cite{IA05}), which can only form if the accretion rate is lower than a critical value of the order of $10^{-7}$ \\Msun/yr  (Muzerolle et al.~\\cite{Mea04}). For low \\Macc, the inner gaseous disk is optically thin and does not prevent the direct stellar irradiation of the inner dust edge of the disk, thus giving raise to the puffed-up rim. Secondly, we  are interested in  comparing the results to the accretion properties of the lower-mass objects studied in Taurus and Ophiuchus. In both regions it has been found that the mass accretion rate \\Macc\\ is a strong function of the mass of the central object, increasing roughly as \\Mstar$^{1.8-2.0}$ (e.g., Natta et al.~\\cite{Nea06} and references therein; but see also Clarke \\& Pringle \\cite{CP06}). The causes of this behaviour are not understood, and are matter of debate (Muzerolle et al.~\\cite{Mea05}; Natta et al.~\\cite{Nea06}; Alexander \\& Armitage \\cite{AA06}; Dullemond et al.~\\cite{Dea06}). The HAe stars are scattered in different  regions of the sky and are, on average, older than objects in Ophiuchus and Taurus; they may provide interesting information on the dependence of the disk accretion on age and environment.    The outline of this paper is as follows. The observations and data reduction are described in \\S2; method and  results are presented in \\S3 and discussed in \\S4. Conclusions follow in \\S 5.  \\section {Observations and data reduction}  Medium resolution (R$\\sim$9000) spectra of all stars in our sample were obtained with the ISAAC spectrograph at the Antu 8.2m VLT unit telescope in service mode in spring 2004. The spectral region  covered by our spectra included the Br$\\gamma$ and the H$_2$ (1--0)S(1) lines, which was not detected in any of the objects of the sample. All spectra were obtained in moderately poor conditions (seeing $\\sim$1.5 arcsec and non photometric transparency), as these were adequate for the scope of the project. As a consequence, we did not attempt to obtain an accurate flux calibration of the spectra.  The raw data were wavelength calibrated using daytime arcs and rectified using data from the ISAAC calibration plan. After aligning and coadding the dithered integrations on each target, the spectra were extracted and telluric absorption and instrumental throughput were corrected using observations of telluric standards obtained by the observatory staff during the same nights as the science observations. The details of the ISAAC calibration plan, the reduction  process and the instrument characteristics can be found on the ESO web pages\\footnote{http://www.eso.org/instruments/isaac}.  The portion of the spectra around  \\Brg\\ are shown and briefly discussed in the appendix. In each spectrum we measured the equivalent width of the Br$\\gamma$ line; the measurements are given in Tables~1 and~2, together with the 1$\\sigma$ uncertainties.  Tables~1 and 2 report also the star name and properties. Spectral type and distance are compiled from the literature (e.g., van den Ancker et al.~\\cite{vda97}, \\cite{vda98}; Testi et al.~\\cite{Tea98}; Meeus et al.~\\cite{Meeus01}; Rodgers \\cite{Rod01}; Hernandez et al.~\\cite{Hea05},  and references therein). Extinction and luminosity are computed from the V, (B-V) magnitudes as described in Testi et al.~(\\cite{Tea98}); stellar masses are derived by comparing the location of each star on the HR diagram to the evolutionary tracks of Palla and Stahler (\\cite{PS93}).  \\section {Method and results}   We  derive the accretion luminosity \\Lacc\\ of each star from the luminosity of \\Brg, following the relation derived by Calvet et al.~(\\cite{Cea04}) for a sample of objects ranging in mass from about 4 \\Msun\\ to BDs: \\begin{equation} \\log L_{acc}/L_\\odot = 0.9 \\times \\large(\\log (L(Br_\\gamma)/L_\\odot) +4\\large) -0.7 \\end{equation} where \\Lbg\\ is the luminosity of the \\Brg\\ line emitted by the circumstellar gas. Note that Eq.(1) is an empirical relation, where \\Lacc\\ is measured from the observed veiling, i.e., independently of any assumption on the  origin of the \\Brg\\ emission. It has been derived from objects with L(\\Brg) between $\\sim 3\\times 10^{-6}$ and $\\sim 3\\times 10^{-3}$ \\Lsun; the scatter of \\Lacc\\ for any given L(\\Brg) is typically a factor $\\pm$3 (see also van den Ancker \\cite{vda04}).  In order to compute \\Lbg, we need to extract the equivalent width of the circumstellar component of the line from the observed one. In HAe stars, the observed flux is the sum of the stellar photospheric emission, which for A stars shows broad and rather strong absorption in the hydrogen lines, of the emission of the circumstellar gas and of the emission of the disk. Following Rodgers (\\cite{Rod01}), we  assume that the star dominates the observed V emission. The equivalent width of the photospheric component, corrected for the K-band veiling, is computed for each object with spectral type earlier than A8 using the template spectra of Rodgers (\\cite{Rod01}). For objects of later spectral type, for which we did not have spectra of template stars of suitable resolution,  we compute \\Wcirc\\ from the expression: \\begin{equation} W(Br_\\gamma)_{circ}= W(Br_\\gamma)_{obs} -W(Br_\\gamma)_{phot}\\> 10^{-0.4 \\Delta K} \\end{equation} where \\Wobs\\ is the observed equivalent width, \\Wphot\\ the equivalent width of the photospheric \\Brg\\ line of a template with the  same spectral type (see Table~\\ref{table_templates}; values from Rodgers \\cite{Rod01} and Wallace \\& Hinkle \\cite{WH97}) and $\\Delta K$ the disk continuum emission, computed subtracting from the observed K magnitude  the photospheric one. All the relevant parameters used in the derivation are given in Table~\\ref{table_hae}. K magnitudes are from 2MASS \\footnote{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. }; note that none of the quantities used in eq.(2) has been measured simultaneously.  Of the 32 targets, 4 are very luminous objects of spectral type early B. Three of them have \\Brg\\ in emission, one in absorption. Table~\\ref{table_hbe} refers in Column 3 a lower limit to the line  luminosity, derived from the observed \\Brg\\ equivalent width, i.e., without correcting for the photospheric absorption. For the three stars where it is seen in emission L(\\Brg)$\\simgreat 0.1-0.2$ \\Lsun. The validity of Eq.(1) for these very luminous objects is very doubtful, and we will not consider them further in the following of this paper.  After  subtraction of the photospheric component, 24 of the 28 HAe stars have circumstellar \\Brg\\  emission; for these, we compute \\Lbg\\ from \\Wcirc, the continuum flux derived from the extinction-corrected K magnitude and the distance of the star. \\Lacc\\ is then obtained from Eq.(1). For the remaining 4, where no circumstellar  emission is detected, we estimate that \\Wcirc $\\simless$ 2 \\AA, and  compute upper limits to \\Lacc\\ accordingly. For HD~98922 and HD~144432, \\Lacc\\ refers to the quoted distance lower limit, and is therefore also a lower limit to the true value. The results are given in Table~\\ref{table_hae} and shown in Fig.~\\ref{fig_lacc} as a function of the stellar luminosity. Note that Eq.(1) has been derived for objects with L(\\Brg)$\\simless 3\\times 10^{-3}$ \\Lsun; 2 objects (HD~98922 and VV Ser) have larger L(\\Brg),  and the use of Eq.(1) should be viewed with some caution.  The mass accretion rate  follows from the relation $\\dot M_{acc}=L_{acc} \\, R_\\star/(GM_\\star)$, and is shown in Table~\\ref{table_hae}, Column (13).  The uncertainties on \\Lacc\\ and \\Macc\\ are dominated in most cases by the uncertainties deriving from the subtraction of the veiled photospheric absorption component from the observed spectrum. The error bars shown in Fig.~\\ref{fig_lacc} correspond to an uncertainty on the resulting \\Wcirc\\ of  $\\sim\\pm$2 \\AA, which we consider realistic in the majority of cases.  For any given object, variability of lines and continuum can also be large. There are no systematic studies of time variability of the \\Brg\\ equivalent width in our sample stars, but one can compare our measurements with the Rodgers (\\cite{Rod01}) results from lower resolution spectra (R$\\sim 500-2000$) for the 7 stars we have in common. The differences are within the errors for 3 of them, but larger than 2 \\AA\\ in 2, one being R CrA. This confirms the well known fact that the accretion rate in pre-main sequence stars varies and that one can only estimate the ``typical\" value  of any individual object if the results of long term time monitoring are available. However, if one is interested in the  properties of a large sample of stars, as in this work, one can safely assume that their time variability will  increase the spread of the determinations, but not bias the average values.  \\begin{figure}[h!] \\begin{center} \\leavevmode \\end{center} \\centerline{ \\psfig{file=5575fig1.eps,width=9cm,angle=0} } \\caption{Accretion luminosity as a function of \\Lstar\\ for our sample (filled squares) and the Rodgers (\\cite{Rod01}) objects (open squares); arrows show upper limits. For comparison, we show also the IMTTS \\Lacc\\ measured by Calvet et al.~\\cite{Cea04} from veiling (crosses).} \\label{fig_lacc} \\end{figure}    \\begin{figure}[ht!] \\begin{center} \\leavevmode \\end{center} \\centerline{ \\psfig{file=5575fig2.eps,width=9cm,angle=0} } \\caption{Distribution of the mass accretion rate in HAe stars (top Panel). The shadowed boxes show actual measurements, the empty boxes (thick histogram) the distribution of the upper limits. The mid and bottom panels show the distribution of \\Macc\\ for Class II objects in Ophiuchus with mass $0.3 <M_\\ast/M_\\odot <1$ and $M_\\ast/M_\\odot \\le 0.3$, respectively (data from Natta et al.~\\cite{Nea06}). } \\label{fig_hist} \\end{figure}  \\begin{figure}[ht!] \\begin{center} \\leavevmode \\centerline{ \\psfig{file=5575fig3.eps,width=9cm,angle=0} } \\end{center} \\caption{Mass accretion rate as a function of \\Mstar\\ for  HAe stars (filled squares and upper limits), Ophiuchus Class II objects (open dots and upper limits; Natta et al.~\\cite{Nea06}), Taurus CTTS and BDs (crosses; Muzerolle et al.~\\cite{Mea05}, Mohanty et al.~\\cite{Subu05}, Calvet et al.~\\cite{Cea04} and references therein), IMTTS in Orion (plus signs, Calvet et al.~\\cite{Cea04}).} The dashed line shows the best-fit to the Ophiuchus data (\\Macc $\\propto$ \\Mstar$^{1.8}$); the dotted lines are $\\pm$1 dex. \\label{fig_macc} \\end{figure}   \\begin{table*} \\caption{Herbig Ae properties -- Column 1: star name; Column 2: spectral type; Column 3: V magnitude; Column 4: K magnitude; Column 5: V extinction; Column 6: distance; Column 6: effective temperature; Column 8: stellar luminosity; Column 9: stellar mass; Column 10: observed \\Brg\\ equivalent width; Column 11: equivalent width of the circumstellar \\Brg; Column 12: accretion luminosity; Column 13: mass accretion rate.} \\label{table_hae} \\centering \\begin{tabular}{ l c c c c c c c c c c c  c}     % \\hline\\hline (1) & (2)   & (3)& (4) & (5) & (6)& (7)      & (8)      & (9)      & (10)& (11)& (12)& (13)\\\\  name  & ST &   V & K& A$_{\\rm V}$  & D  & T$_{eff}$& \\Lstar& \\Mstar& Br$_\\gamma$(obs)& Br$_\\gamma$(circ)& Log L$_{acc}$& Log $\\dot M_{acc}$\\\\ &    & (mag)& (mag)& (mag)& (pc)& (K)& (\\Lsun)& (\\Msun)& (\\AA)& (\\AA)& (\\Lsun)& (\\Msun/yr)\\\\ \\hline \\object {HD 149914} & B9 &6.75& 5.69& 0.9  &165 & 10600   & 158      & 3.5 & 6.9$\\pm$0.2& -0.8 & -0.51 & -7.99      \\\\ \\object {HD 179218} & B9 & 7.2 & 5.99& 1.8 & 240& 10600   & 500      & 4   & 1.3$\\pm$0.2 & -10.1 & 0.70 & -6.59       \\\\ \\object  {HD 176386} & B9& 7.30 & 6.69& 0.5 & 140& 10600   & 49       & 2.7  & 6.2$\\pm$0.2& -3.1    &  -0.49& -8.11 \\\\ \\object  {HD 97300}  & B9 & 9.00& 7.15& 1.2 & 190& 10600   & 35       & 2.5      & 9.1$\\pm$0.2& -- &$<$-0.5 &$<$-8.18  \\\\ \\object  {HD 98922}  & B9 & 6.77& 4.28& 0.3 & $>$540& 10600& $>$890   & 5: &-2.4$\\pm$0.2 & -3.7    &  1.50& -5.76    \\\\ \\object  {VV Ser}    & B9 &11.92 & 6.32& 3.1 &260 & 10600 & 85   & 3.1   &-9.0$\\pm$0.2& -13.0 & 1.20& -6.34        \\\\ \\object  {51 Oph}    & B9.5&4.81& 4.29& 0.1 & 131& 10200   & 245      & 4  &1.3$\\pm$0.3& -5.6 & 0.54 &-6.87        \\\\ \\object  {HD 141569} & B9.5&7.00& 6.82& 0.  &99  & 10200   & 22    & 2.3   &6.0$\\pm$0.1& -4.5  &  -0.68 & -8.37      \\\\ \\object  {NX Pup}    & A0 &9.96 &6.08& 1.4 & 450& 9840    & 42     & 2.5   &-1.9$\\pm$0.2& -2.4     &  0.58& -6.97    \\\\ \\object  {HR 4796}   & A0 &5.78 &5.77 & 0.1 & 67 & 9840   & 26     &2.4    &9.4$\\pm$0.2& --& $<$-0.9 & $<$-8.53\\\\ \\object  {HD 97048}  & A0 &8.46 &5.94& 1.0 & 180& 9840     & 41    & 2.5     &-6.3$\\pm$0.15& -8.4   & 0.39 &-7.17     \\\\ \\object  {HD 104237} & A0 &6.59 &4.58& 0.7 & 116& 9840     & 6     & 2.0     &-4.7$\\pm$0.2& -6.7   & 0.43 &-7.45     \\\\ \\object  {HD 95881}  & A1 &8.25 &5.73& 0.3 & 120& 9450    & 10     & 2       &-1.6$\\pm$0.2& -2.8   & -0.31 & -8.04     \\\\ \\object  {HD 150193} & A1 &8.88 &5.48& 1.5 & 150& 9450     & 30     & 2.3     &-3.5$\\pm$0.15& -5.5  & 0.27 & -7.29      \\\\ \\object  {HD 163296} & A1 &6.87 &4.78& 0.3 & 122& 9450    & 36      & 2.3     &-4.7$\\pm$0.1& -6.9  &  0.40 & -7.12     \\\\ \\object  {HD 144667} & A1.5&6.63&6.67&0.1&140& 9300     & 51       & 2.5     &5.8$\\pm$0.15& -7.2   & -0.17 & -7.63     \\\\ \\object  {TY Cra}    & A2 &9.42 &6.67& 1.3 & 130& 9120   & 8        & 1.8      &4.7$\\pm$0.2& --     &$<$-0.6 &$<$-8.31  \\\\ \\object  {WW Vul}    & A3 &10.51& 7.28&1.0 & 550& 8670   & 30     & 2.4     &-3.6$\\pm$0.2&-5.4     & 0.60& -6.89   \\\\ \\object  {HD 169142} & A5 &8.15 & 6.41&0.3 & 145& 8260     & 15    & 1.8     &-5.0$\\pm$0.1&-9.7    &  0.08 & -7.40    \\\\ \\object  {KK Oph}    & A6 &10.57 &5.79& 0.44 & 160& 8100  & 6.2   & 1.6     &-1.2$\\pm$0.15& -1.5    & -0.30 & -7.91   \\\\ \\object  {HD 139614} & A7 &8.25 & 6.75&0.1 & 150& 7900     & 11    & 1.8     &2.2$\\pm$0.2& -3.0& -0.48 &-7.99      \\\\ \\object  {HD 142666} & A8 &8.81 & 6.08&0.8 & 116& 7580     & 8        & 1.6     & 1.0$\\pm$0.2&-2.4 &-0.50 & -8.00        \\\\ \\object  {HD 100453} & A9 &7.79 &5.6& 0.0 & 111& 7300     & 9        & 1.7     &1.6$\\pm$0.2& -1.5   & -0.58& -8.04     \\\\ \\object  {HD 144432} & A9 &8.16 &5.89& 0.2 & $>$200&7300& $>$23      & $>$2      &-2.1$\\pm$0.2& -5.1 &0.26 & -7.07       \\\\ \\object  {HD 135344} & F3 &8.61 &5.84& 0.3 & 84 & 6590     & 8      & 1.5        &0.3$\\pm$0.2& -1.3   &-0.93 &-8.27    \\\\ \\object  {T CrA }    & F5:&12.04&6.60& 1.9 & 130& 6400     & 1.4      & 1.2     &9.6$\\pm$0.2& --   &$<$-0.6 &$<$-8.20     \\\\ \\object  {R CrA}     & F5:&11.50&2.86& 1.4 & 130& 6400     & 1.3      & 1.2     &-1.2$\\pm$0.3& -1.2 & 0.50 & -7.12       \\\\ \\vspace{0.3cm} \\object  {HD 142527} & F6&8.33&4.98&0.7 & 200& 6270     & 69       & 3.5     &0.3$\\pm$0.2&-1.2  &0.04  & -7.16\\\\ \\object {AB Aur} &  A0&7.05 & 4.2 & 0.5& 140&9840 &48& 2.4& -4.4$\\pm$0.15& -5.5&  0.63& -6.85\\\\ \\object {HD 34282}& A0&9.85 & 7.68& 0.6& 400&9840 &30& 2.2&  2.7$\\pm$0.17& --   & $<$-0.20 &$<$-7.71\\\\ \\object {RR Tau}& A0&11.0   & 7.39& 1.1& 800&9840& 37& 2.5& -2.56$\\pm$0.05& -3.5&  0.70&-6.86\\\\ \\object {HD 35187}& A2&7.78 & 5.91& 0.65&150&9120 &27& 2.3&  7.13$\\pm$0.27& --   & $<$-0.3&$<$-7.81\\\\ \\object {UX Ori}&A3&10.0 & 7.21& 0.3& 450&8670& 39& 2.3& -2.4$\\pm$0.02& -3.9&  0.32& -7.19\\\\ \\object {HK Ori}&  A5&11.66  & 7.34& 1.2 & 450&8260& 12& 2.0& -1.44$\\pm$0.08& -2.2&  0.08& -7.50 \\\\ \\object {HR 5999}& A6&7.04  & 4.39& 0.5& 210&8100& 85& 3.2& -1.74$\\pm$0.01& -6.5&  0.95& -6.37\\\\ \\object {CQ Tau} & F2& 9.42& 6.17& 1.9& 130&    &6.6& 1.5&  3.17$\\pm$0.18&   --& $<$-0.6 & $<$ -8.08\\\\ \\hline \\end{tabular} \\flushleft Note: The last 9 lines are objects from Rodgers (\\cite{Rod01}). \\end{table*}  \\begin{table} \\label{table_hbe} \\caption{Herbig Be properties -- Column 1: star name; Column 2: spectral type; Column 3: observed \\Brg\\ equivalent width; Column 4: \\Brg\\ luminosity.} \\centering \\begin{tabular}{ l c c c}     % \\hline\\hline (1) & (2)   & (3)& (4)\\\\  name  & ST &   Br$_\\gamma$(obs)& Log L(\\Brg)\\\\ &    &  (\\AA)& (L$_\\odot$)\\\\ \\hline \\object {V921 Sco}& B0& -17.3$\\pm$0.25& $>$-0.7\\\\ \\object {MWC 300}& B0-B1& -10.9$\\pm$0.25& $>$-0.8\\\\ \\object {MWC 166}& B0 IV& 2.5$\\pm$0.2& --\\\\ \\object {MWC 297}& B1.5& -16.0$\\pm$0.2& $>$-0.6\\\\ \\hline \\end{tabular} \\end{table}  \\begin{table} \\begin{flushleft} \\caption{Br$_\\gamma$ equivalent width of template stars -- Column 1: spectral type; Column 2: adopted \\Brg\\ equivalent widht.} \\label{table_templates} \\centering \\begin{tabular}{ l  c }     % \\hline\\hline (1) & (2)   \\\\  ST &   W(Br$_\\gamma$)\\\\ &  (\\AA)\\\\ \\hline B8&  11.4\\\\ A0& 9.7\\\\ A2& 13.4\\\\ A8& 12.0\\\\ F3& 6.6\\\\ F6& 5.5\\\\ \\hline \\end{tabular} \\end{flushleft} \\end{table}  \\section {Discussion}  Fig.~\\ref{fig_lacc} plots \\Lacc\\ as a function of \\Lstar\\ for the HAe observed in this paper. We have added six HAe stars (listed at the bottom of Table~1) taken from the Rodgers (\\cite{Rod01}) thesis, for a total of 36 objects. Typically, the HAe stars have \\Lacc\\ roughly between 10\\% and 1\\% \\Lstar. There  are very few objects with \\Lacc$>$0.1 \\Lstar, contrary to what happens for   lower mass objects in Taurus and Ophiuchus (see Natta et al.~\\cite{Nea06}, Calvet et al.~\\cite{Cea04}).  Fig.~\\ref{fig_hist}, top Panel, shows the \\Macc\\ distribution of the HAe stars; the values of \\Macc\\ cover the interval $\\sim 10^{-6}$ -- $\\simless 3\\times 10^{-8}$ \\Msun/yr, but few stars (25\\%) have \\Macc $>10^{-7}$ \\Msun/yr; 36\\% have \\Macc$< 10^{-8}$ \\Msun/yr. The median value is \\Macc$=3.5\\times 10^{-8}$ \\Msun/yr. These results agree with those obtained by Muzerolle et al.~(\\cite{Mea04}), who found that  intermediate-mass objects have typical \\Ha\\ profiles and/or Balmer discontinuity roughly consistent with the predictions of magnetospheric accretion models with accretion rates $\\simless 10^{-7}$ \\Msun/yr. They studied in detail only one object, UX Ori, for which they found a value \\Macc$\\sim 10^{-8}$ \\Msun/yr, comparable with the value ($\\sim 6\\times 10^{-8}$ \\Msun/yr) we derive from the low resolution spectrum of Rodgers (\\cite{Rod01}).  Our results support the conclusion of Muzerolle et al.~(\\cite{Mea04}) that in most HAe stars the accretion rate is sufficiently low that the gas in the inner disk, inside the dust evaporation radius, is optically thin, allowing the direct irradiation of the inner dust edge (the rim) by the star. One of the conditions for the formation of a puffed-up rim (Natta et al.~\\cite{Nea01}) seems therefore generally verified.    Fig.~\\ref{fig_macc} plots \\Macc\\ as a function of \\Mstar\\  for the HAe stars, together with objects in \\Roph\\ (Natta et al.~\\cite{Nea06}), Taurus (Muzerolle et al.~\\cite{Mea05} and references therein; Mohanty et al.~\\cite{Subu05}) and  IMTTS in Orion and Taurus (Calvet et al.~\\cite{Cea04}). Fig.~\\ref{fig_hist} compares the \\Macc\\ distribution of the HAe stars to that of objects in Ophiuchus in two mass ranges, as labelled. One can see that the median  \\Macc\\ for the HAe stars ($\\sim 3\\times 10^{-8}$ \\Msun/yr) is about a factor ten higher than for Ophiuchus stars with $0.3 < M_\\star < 1$ \\Msun, and a factor ~100 larger than in the very low mass  objects with \\Mstar$<0.3$ \\Msun. This is roughly consistent with the predictions of the \\Macc$\\propto$ \\Mstar$^{1.8}$ relation derived for sub-solar mass objects. A close inspection of Fig.~\\ref{fig_macc}, however, shows   that the HAe stars  lack very strong accretors, as already noted in discussing Fig.~\\ref{fig_lacc}. If the upper envelope of the distributions observed in Taurus and Ophiuchus extended to HAe stars, we would expect a much larger fraction of objects accreting at rates higher than $10^{-7}$ \\Msun/yr.  It is not easy, at this stage, to understand why this happens. The HAe stars form a sample which is not complete nor homogeneous. The objects are located in a variety of star forming regions, and, in some cases, are isolated in the sky. We think that the most likely interpretation of their behaviour is ``aging'', since the  accretion rate in viscous disks with constant $\\alpha$  parameter is expected to deacrease with time roughly as $t^{-1.5}$ (Hartmann et al.~\\cite{Hea98}). HAe stars are, on average, older than Ophiuchus objects, which cluster at ages $\\sim 0.5-1$ Myr, respectively (Palla \\& Stahler \\cite{PS00}). Their nominal ages, obtained by comparing their location on the HR diagram with Palla \\& Stahler (\\cite{PS99}) evolutionary tracks, range from $\\sim 1$ to $> 10$ Myr; several stars are close to the ZAMS, and it is in fact interesting that 6/7 stars with  non-detected \\Brg\\ emission are ZAMS objects, the exception being HD 35187, with an estimated age $\\sim 8$ Myr. However,  one should not overinterpret these findings, as ages of intermediate-mass stars derived from their location on the HR diagram are very uncertain.  It is possible that the steep dependence of \\Macc\\ on \\Mstar\\ does not extend to objects more massive than, say, $\\sim 1$ \\Msun, reflecting a change in the physical processes that control  disk evolution and/or formation, and which have been discussed in a number of recent papers (e.g., Muzerolle et al.~\\cite{Mea03}, Natta et al.~\\cite{Nea06}, Padoan et al.~\\cite{Pea04}, Alexander \\& Armitage \\cite{AA06}, Dullemond et al.~\\cite{Dea06}). Unfortunately, given the lack of completeness and the selection biases that affect it, it is difficult to establish the relevance of the HAe sample for this ungoing  discussion. It is worth to remind, in this context, the well  known bias due to  the spectral type selection, which introduces an age-mass correlation in the sample (e.g., van Boekel et al.~\\cite{vBea05}) and automatically excludes young, relatively low mass ($\\simless$ 2.5 \\Msun) objects. ",
        "conclusions": "We have reported in this paper measurements of the disk mass accretion rates in a sample of 36 Herbig Ae stars. The values of \\Macc\\ are derived from the measured luminosity of the \\Brg\\ emission line, using the  correlation between L(\\Brg) and the accretion luminosity \\Lacc, established by Muzerolle et al.~(\\cite{Mea98}) and Calvet et al.~(\\cite{Cea04}). Note that it is an empirical  correlation,  which makes no assumptions on the  origin of \\Brg.  We find that 80\\% of the stars, all of which have evidence of an associated circumstellar disk, are accreting matter, with rates $3\\times 10^{-9}\\simless \\dot M_{acc} \\simless 10^{-6}$ \\Msun/yr; for 7 objects, 6 of which are located on the ZAMS in the HR diagram, we do not detect any line emission. Few HAe stars (25\\%) have \\Macc$>10^{-7}$ \\Msun/yr. These results support the conclusion of Muzerolle et al.~(\\cite{Mea04}) that in most HAe stars the accretion rate is sufficiently low that the gas in the inner disk, inside the dust evaporation radius, does not prevent the formation of the  puffed-up rim, proposed by Natta et al.~(\\cite{Nea01}) and Dullemond et al.~(\\cite{DDN01}).  When compared to the \\Macc\\ values found for lower-mass stars in the star forming regions Taurus and Ophiuchus, HAe stars have on average higher accretion rates; however, there is a lack of very strong accretors among them, as  shown by the  very few objects with \\Lacc $>$ 10\\% \\Lstar. It seems likely that this is because on average HAe stars are older than TTS in Ophiuchus or Taurus. However, ages of intermediate-mass stars derived from the HR diagram location are very uncertain; this, and  the  various selection biases which plague the HAe sample, make the comparison very difficult.  When we consider individual stars, the largest uncertainties on  \\Macc\\  come from line variability, as in TTS, but also from  the presence of a photospheric absorption component of \\Brg\\ in early-type stars, which needs to be subtracted from the observed one, and from the K-band excess emission. A better characterization of the photospheric profiles of template stars would be very valuable."
    },
    "0609/astro-ph0609803_arXiv.txt": {
        "abstract": "We present and discuss correlations for optical and near-infrared ($5500-10030$~\\AA) line intensity measurements at many positions in the Crab Nebula. These correlations suggest the existence of gas produced by a range of nuclear processing, from material in which synthesis ended with the CNO-cycle, to some helium-burning and nitrogen depletion, to regions containing enriched products of oxygen-burning. The latter exhibit a gradual, linear rise of $[$Ni~II$]$ emission with increasing argon enrichment, whereas gas with less nuclear processing shows markedly different $[$Ni~II$]$ emission characteristics, including the highest derived abundances.  This suggests two origins for stable, neutron-rich nickel in the nebula: a type of ``alpha-rich freezeout'' in the more highly processed material, and possibly removal of ions from the neutron star in other regions.  In addition, the data indicate that anomalously strong observed $[$C~I$]$ emission comes from broad, low-ionization H$^{+}$ to H$^{0}$ transition zones.  Although the strongest He~I emission could also be enhanced in similar low-ionization gas, correlations between relevant line ratios argue against that interpretation, strengthening the case for an exceptionally high helium mass fraction in some locations. ",
        "introduction": "Young supernova remnants are excellent laboratories for investigating how stars make elements, and the Crab Nebula in particular can provide unique information about the precursor star, the supernova event, associated heavy element production, and the environment of a highly energetic pulsar.  It is the bright remnant of a core-collapse supernova observed in 1054~A.D.  Its age and location, roughly 180~pc away from the plane of the Galaxy, suggest that the ejecta are not significantly contaminated by swept-up interstellar material. Furthermore, measured electron temperatures in the gas  (Woltjer 1958; Miller 1978; Fesen \\& Kirshner 1982; MacAlpine et al. 1989, 1996) along with the lack of other possible evidence for shocks (Frail et al. 1995) imply that the line-emitting gas shines primarily because of photoionization by the locally-generated synchrotron radiation field (see also Davidson \\& Fesen 1985).  Therefore it can be analyzed using powerful numerical photoionization modeling codes.  It is generally believed that the supernova precursor star initially contained about $9-11$~M$_{\\odot}$ (e.g., Arnett 1975; Nomoto 1985), representing the important low end of the Type~II supernova mass range, below which stars could ignite carbon degenerately and above which successive nuclear reaction stages would be expected to take place through silicon burning.  As discussed by Woosley \\& Weaver (1986a), the applicable stellar models allow for a number of explosive and nucleosynthesis possibilities.  Spectroscopic and photometric investigations of the Crab Nebula to date have indicated several apparent ``gas components.''  The majority of the observed nebular gas is helium (e.g., MacAlpine et al. 1989), consisting of less than $2$ M$_{\\odot}$ (MacAlpine \\& Uomoto 1991).  It has been postulated that this represents ``helium mantle'' material from deep within the original star, ejected by the explosive event.  This is consistent with some stellar models, or scenarios involving another star, in which outer layers of the precursor were lost prior to the core-collapse event.  Most of this helium-mantle gas appears to be nitrogen-rich, confirming its origin from CNO processing (MacAlpine et al. 1996).  There also is a major component of gas, primarily in the southern part of the nebula near the pulsar, which is significantly nitrogen-poor and much of which is sulfur-rich (MacAlpine et al. 1996).  It was suggested that this gas resulted from localized oxygen-burning episodes, consistent with stellar models (Woosley \\& Weaver 1986b, 1995) that involve off-center (in a shell) oxygen-burning.  Strolger \\& MacAlpine (1996) provided a preliminary demonstration that this explanation is probably correct.  Still another significant nebular component or ``anomaly'' involves an apparent helium-rich torus viewed as an east-west band across the pulsar region, which constitutes approximately 25\\% of the visible material (Uomoto \\& MacAlpine 1987; MacAlpine et al. 1989; MacAlpine \\& Uomoto 1991). The computed helium mass fraction is about 95\\%, and there is not yet a realistic explanation for this apparent structure.  Other known apparent anomalies in the Crab Nebula include exceptionally strong $[$Ni~II$]$ and $[$C~I$]$ line emission, resulting in large spatial variations for deduced abundances of nickel (along with iron) and carbon.  Strong $[$Ni~ II$]$~$\\lambda$7378 has been reported by numerous authors (Miller 1978; Fesen \\& Kirshner 1982; Henry et al. 1984; MacAlpine et al. 1989). The latter work suggested neutron-rich nickel isotopic abundance enhancements, compared with solar, by factors of 5-50 at various locations.  Iron was also found to vary widely, together with nickel but at a much lower level; deduced nickel/iron abundance ratios are roughly 60-75 times the solar value.  Henry et al. (1984) found surprisingly high $[$C~I$]$~$\\lambda$$\\lambda$9823,9850 emission at several locations in the Crab Nebula, where it was measured to be as much as 7 times stronger than predicted by the photoionization models of Henry \\& MacAlpine (1982).  To account for this, it was postulated that the $[$C~I$]$ lines might arise from collisional excitation involving hydrogen atoms as well as electrons.  As part of an effort to develop a more consistent and accurate overall understanding of the Crab Nebula, we have measured relative line intensities of He~I~$\\lambda$5876, $[$O~I$]$~$\\lambda$6300, H$\\alpha$, $[$N~II$]$~$\\lambda$6583, $[$Ar~III$]$~$\\lambda$7136, $[$Ni~II$]$~$\\lambda$7378, $[$S~III$]$~$\\lambda$9069, $[$S~III$]$~$\\lambda$9531, and $[$C~I$]$~$\\lambda$9850 for roughly 200 well-distributed positions throughout the emitting gas.  In this paper, we address all of the above issues (apparent gas components and anomalies) using comparisons among these line measurements.  A follow-up paper, involving more in-depth photoionization analyses for deriving improved gas physical conditions and chemical abundances, is planned.  The spectroscopic observations are described in \\S~2. Then \\S~3 presents emission-line correlations and resulting inferences about the nebular gas. The correlations illustrate a broad range of nuclear processing and confirm significant enrichment with products of oxygen-burning in some areas. They also suggest two distinct nickel/argon line relationships. In addition, the importance of optical depth for enhanced $[$C~I$]$ lines is demonstrated, along with the {\\it lack} of a dominant role for optical depth in the production of strong He~I emission.  A summary discussion is given in \\S~4. ",
        "conclusions": "We have measured emission of He~I~$\\lambda$5876, $[$O~I$]$~$\\lambda$6300, $[$N~II$]$~$\\lambda$6583, $[$S~II$]$~$\\lambda$6731, $[$Ar~III$]$~$\\lambda$7136, $[$Ni~II$]$~$\\lambda$7378, $[$S~III$]$~$\\lambda$9069, $[$S~III$]$~$\\lambda$9531, and $[$C~I$]$~$\\lambda$9850 at many locations within the Crab Nebula.  The different line intensities (or subsets thereof) were plotted against each other in efforts to investigate correlations and improve our understanding of the range of nuclear processing in the gas, as well as the causes of exceptionally strong emission from $[$Ni~II$]$, $[$C~I$]$, and He~I.  We identified gas where nucleosynthesis has not progressed significantly beyond the CNO-cycle, gas in which some helium-burning and nitrogen depletion have taken place, and regions where oxygen-burning has occurred. The anomalously strong observed $[$Ni~II$]$ emission may have two sources, in one case resulting from high temperature and subsequent cooling in gas enriched with products of oxygen-burning, while in the other case possibly representing the release of nuclei from the neutron star surface.  Line correlations indicate that very strong $[$C~I$]$ emission arises in low-ionization H$^{+}$$\\rightarrow$H$^{o}$ transition regions. On the other hand, exceptionally strong He~I~$\\lambda$5876 does not show similar evidence of a low-ionization zone origin; and it does not appear to correlate with different levels of nuclear processing as represented by $[$N~II$]$ emission.  Therefore, the apparent high-helium torus around the pulsar may be a distinct component of the nebula."
    },
    "0609/astro-ph0609668_arXiv.txt": {
        "abstract": "We present our recent work on the conditions under which star formation occurs in a metal-poor environment, the Large Magellanic Cloud ([Fe/H]\\,$\\sim$\\,$-$0.4). Water masers are used as beacons of the current star formation in H\\,{\\sc ii} regions. Comparing their location with the dust morphology imaged with the Spitzer Space Telescope, and additional H$\\alpha$ imaging and groundbased near-infrared observations, we conclude that the LMC environment seems favourable to sequential star formation triggered by massive star feedback (\\cite[Oliveira et al. 2006]{oliveira06}). Good examples of this are 30\\,Doradus and N\\,113. There are also H\\,{\\sc ii} regions, such as N\\,105A, where feedback may not be responsible for the current star formation although the nature of one young stellar object (YSO) suggests that feedback may soon start making an impact. The chemistry in one YSO hints at a stronger influence from irradiation effects in a metal-poor environment where shielding by dust is suppressed (\\cite[van Loon 2005]{vanloon05}). ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609342_arXiv.txt": {
        "abstract": "In this paper two theoretical approaches for the calculation of the rate of quasi-stationary, two-dimensional magnetic reconnection with nonuniform anomalous resistivity are considered in the framework of incompressible magnetohydrodynamics (MHD). In the first, ``global'' equations approach the MHD equations are approximately solved for a whole reconnection layer, including the upstream and downstream regions and the layer center. In the second, ``local'' equations approach the equations are solved across the reconnection layer, including only the upstream region and the layer center. Both approaches give the same approximate answer for the reconnection rate. Our theoretical model is in agreement with the results of recent simulations of reconnection with spatially nonuniform resistivity by Baty, Priest and Forbes (2006), contrary to their conclusions. ",
        "introduction": "\\label{INTRODUCTION}  Magnetic reconnection is one of the most important processes of plasma physics, and is believed to play the central role in observed phenomena in laboratory and cosmic plasmas. At the same time there has been a long-standing debate about the correct theoretical model of magnetic reconnection. Most previous theoretical and numerical work focused on reconnection processes in two-dimensions, in which all physical scalars and vectors are independent of the third coordinate ($z$). There exist two original and well known models of magnetic reconnection with constant resistivity. First, the Sweet-Parker reconnection model~\\cite{sweet_1958,parker_1963}, which predicts a slow magnetic reconnection rate in hot low-density plasmas. Second, the Petschek model~\\cite{petschek_1964}, in which a fast reconnection rate is achieved by introducing switch-off magnetohydrodynamic (MHD) shocks attached to the ends of a relatively short reconnection layer in the downstream regions. Many numerical simulations have been carried out to discriminate between these two models. More recent high-resolution simulations generally favor the Sweet-Parker model of slow reconnection in the case of constant resistivity, and do not confirm the Petschek theoretical picture for the geometry of the reconnection layer with shocks~\\cite{biskamp_1986,breslau_2003,uzdensky_2000}. However, reconnection becomes much faster and Petschek-like if resistivity is not constant and is enhanced locally in the reconnection layer (numerical studies of this case were pioneered by Ugai and Tsuda~\\cite{ugai_1977,tsuda_1977}, by Hayashi and Sato~\\cite{hayashi_1978,sato_1979}, and by Scholer~\\cite{scholer_1989}).  These results have been called into question in a recent paper, Baty, Priest and Forbes (2006)~\\cite{baty_2006}, which claims that Petschek reconnection can occur when the resistivity is not absolutely constant but varies by an arbitrarily small amount. This paper has motivated us show on the basis of an earlier paper of ours, Malyshkin {\\it et~al$.$}~\\cite{malyshkin_2005}, that the Petschek shocks are produced not by the variation of the resistivity but by the rate of variation which for their case is very large and unlikely to happen naturally.  In 2001 Kulsrud~\\cite{kulsrud_2001} provided some theoretical insight into magnetic reconnection process which qualitatively explained the results of recent simulations. He considered quasi-stationary, two-dimensional magnetic reconnection in the classical Sweet-Parker-Petschek reconnection layer with zero guide field ($B_z=0$), zero plasma viscosity and an anomalous resistivity that is a piecewise linear function of the electric current. Kulsrud was first to suggest that one has to calculate the half-length of the reconnection layer $L'$ from the MHD equations and the jump conditions on the Petschek shocks, instead of treating $L'$ as a free parameter (as Petschek erroneously did when chose $L'$ to be equal to its minimal possible value under the condition of no significant disruption to the plasma flow). As a result, in the case of constant resistivity Kulsrud correctly estimated the layer half-length $L'$ to be approximately equal to the global magnetic field scale, $L'\\approx L$. In this case the Petschek reconnection rate reduces to the slow Sweet-Parker reconnection rate~\\cite{kulsrud_2001,kulsrud_2005}, in agreement with numerical simulations. The second result obtained by Kulsrud is that in the case when resistivity is non-constant but anomalous and enhanced (e.g.~by plasma instabilities), the reconnection rate can become considerably faster than the Sweet-Parker rate.  In our recent paper (Malyshkin, Linde and Kulsrud~\\cite{malyshkin_2005}), to which we will hereafter refer as MLK2005 paper, we put Kulsrud's derivations on a rigorous analytical basis and extended his model by using a new theoretical approach to calculate the reconnection rate. This approach is based on ``local'' analytical derivations across a thin reconnection layer, and it is applicable to the case when resistivity is anomalous, nonuniform and an arbitrary function of the electric current and of the spatial coordinates. We included the case of non-zero guide field $B_z\\ne0$ and non-zero plasma viscosity in our model. We found an approximate formula for the reconnection rate which confirmed Kulsrud's theoretical results.  The present paper has two goals. First, to calculate the reconnection rate, and second, to compare it with the recent simulations of Baty {\\it et~al$.$}~\\cite{baty_2006}.  In the next section we simultaneously follow two theoretical approaches to the calculation of the reconnection rate. In the first approach, which we call the ``global'' equations approach, we derive and solve approximate MHD equations for a whole reconnection layer, including the upstream and downstream regions and the layer center. These theoretical derivations are similar to those done before, except for an important difference. Namely, we take into consideration an additional important equation, the spatial homogeneity of the z-component of electric field $E_z$ along the reconnection layer. This equation, together with the jump condition on Petschek shocks, allows us to find the reconnection layer length $L'$, which must be determined consistently~\\cite{kulsrud_2001}. The second theoretical approach, which we call the ``local'' equations approach, basically coincides with the calculations done in the MLK2005 paper. In this approach we derive and solve approximate MHD equations across a thin reconnection layer, including only the upstream region and the layer center. Both approaches give the same approximate formula for the reconnection rate, which is valid in the general case of an arbitrary nonuniform anomalous resistivity [see Eq.~(\\ref{RATE})]. Using these two approaches simultaneously allows us to better understand the physics of a reconnection process.  In Sec.~\\ref{LOCALIZED_RESISTIVITY} we pursue the second goal of this paper. We consider the case treated by Baty {\\it et~al$.$}~\\cite{baty_2006}, when resistivity is a prescribed function of the two spatial coordinates $x$ and $y$ (remember that we consider two-dimensional reconnection, so that no physical quantities depend on $z$). We argue that our theoretical model is in agreement with the results of these recent simulations of reconnection by Baty, Priest and Forbes~\\cite{baty_2006}. This agreement contradicts their theoretical conclusion that Petschek shocks exist with constant resistivity, and we explain why. Finally, in Sec.~\\ref{DISCUSSION} we discuss our results. ",
        "conclusions": "\\label{DISCUSSION}  Let us now discuss our major results.  First, equation~(\\ref{RATE}) that determines the reconnection current $j_o$ and the quasi-stationary reconnection rate $\\eta_oj_o$ is derived by using the ``local'' equations theoretical approach (see the last column of Table~\\ref{TABLE}). This approach involves only ``local'' quantities and equations that are defined on the interval $OM$ across the reconnection layer, refer to Fig.~\\ref{FIGURE_LAYER}. Thus, the quasi-stationary reconnection rate in a thin two-dimensional layer is determined locally, in the layer central point~$O$ and in the upstream region outside the layer at point~$M$. In other words, the rate is {\\it fully} determined by a particular functional form of anomalous resistivity $\\eta(j_z,x,y)$ and by the local configuration of the magnetic field in the upstream region. The later determines field $B_m$ and its scale $L\\!\\equiv\\![-2B_m/(\\partial_y^2 B_y)_m]^{1/2}$ at point~$M$. As a result, the global properties of the reconnection layer (e.g.~its length $L'$, the plasma outflow velocity $V_{\\rm out}'$ and the presence or absence of Petschek shocks) do not directly matter for the quasi-stationary reconnection rate. This is a very important result because the reconnection layer global geometry can be very complicated (e.g.~in turbulent plasmas). Of course, there exists an indirect dependence of the reconnection rate on the field global properties because the local field configuration in the upstream region is determined by the field global configuration. The above statements are also true for all other ``local'' parameters -- for the layer thickness $\\delta_o$ and for the reconnection velocity $V_R$, refer to Eqs.~(\\ref{DELTA}) and~(\\ref{V_R}).  Second, the reconnection rate can also be estimated by using the ``global'' equations theoretical approach, in which the whole reconnection layer is considered, including the downstream region at point~$O'$ (refer to Fig.~\\ref{FIGURE_LAYER}). The equations used in this approach are presented in the second column of Table~\\ref{TABLE}. We would like to point out that two of these equation play the key role in correct determination of the reconnection layer length $L'$ and geometry. The first key equation is the jump condition on the Petschek shocks, $B_x'/\\sqrt{\\rho}\\approx V_R'\\approx V_R$, whose importance was first pointed out by Kulsrud~\\cite{kulsrud_2001}. The second key equation is the constancy of the electric field z-component along the reconnection layer, $\\eta_oj_o=\\eta'j_z'+V_{\\rm out}'B_x'$, which has been overlooked in the previous theoretical models (e.g.~in the Petschek model).  Third, for the case of a strong dependence of resistivity on the current, i.e.~when $(j_o/\\eta_o)(\\partial\\eta/\\partial j_z)_o\\gg\\max[1,L^2/l_\\eta^2]$, our results, given by Eqs.~(\\ref{PK_RECONNECTION}), coincide with the results obtained by Kulsrud~\\cite{kulsrud_2001}. Thus, both the ``local'' and ``global'' theoretical approaches confirm Kulsrud's results and ideas, contrary to the doubts raised in paper by Baty {\\it et~al$.$}~\\cite{baty_2006}.  Fourth, we found that in the case of uniform or quasi-uniform resistivity the magnetic reconnection rate is the slow Sweet-Parker rate and not the fast Petschek rate, see Eqs.~(\\ref{SP_RECONNECTION}). This theoretical result follows from rigorous analytical derivations and agrees with numerical simulations. At the same time it contradicts the original Petschek theoretical model. Let us consider both the ``global'' and ``local'' analytical approaches in the case of a strictly uniform resistivity $\\eta={\\rm const}=\\eta_o$, and let us explain why the Petschek reconnection layer geometry is not realized in this case. We take the ``global'' equations approach first. In the case of constant resistivity Eqs.~(\\ref{TAYLOR_EXPANSION}) and~(\\ref{J_YY}) result in equation $(\\eta_oj_o-\\eta'j_z')/\\eta_oj_o\\approx (L'/L)^2$ for the fractional drop of the $\\eta j_z$ term along the reconnection layer. On the other hand, the constancy of the electric field z-component along the reconnection layer implies that this fractional drop is $(\\eta_oj_o-\\eta'j_z')/\\eta_oj_o\\approx V_{\\rm out}'B_x'/\\eta_oj_o\\approx 1$, where we use the ``global'' equations on lines $3$ to~$6$ in the second column of Table~\\ref{TABLE}. These two equations agree only if $L'\\approx L$, which means that the geometry of the reconnection layer is Sweet-Parker and not Petschek. Next, we take the ``local'' equations approach. In this ``local'' approach we prefer not to consider the reconnection layer geometry and any ``global'' parameters, such as the layer length $L'$, for the calculation of the reconnection rate. Instead we argue as follows. Refer to the ``local'' equations in the last column of Table~\\ref{TABLE}. Equations on lines~$1$, $2$ and~$5$ result in an estimate of the plasma outflow velocity derivative $(\\partial_y V_y)_o\\approx\\eta_oj_o^2/B_m^2$, see Eq.~(\\ref{V_y_y}). Plasma acceleration equation on line~$3$ results in the upper estimate for the $B_x$ field derivative, $(\\partial_y B_x)_o\\le\\rho{(\\partial_y V_y)_o}^2/j_o$. As a result, we can find the upper estimate for the reconnection current $j_o$ from the ``local'' equation on line~$6$, which is the condition of constancy of the electric field z-component along the reconnection layer. If resistivity is uniform, this estimate turns out to be the Sweet-Parker value, $j_o\\approx (B_m/L)(V_AL/\\eta_o)^{-1/2}$, as given by Eq.~(\\ref{SP_RECONNECTION}). In addition, if one estimates the reconnection layer length as $L'\\sim B_x'/(\\partial_y B_x)_o\\sim V_A/(\\partial_y V_y)_o$, he/she would again recover the Sweet-Parker result $L'\\approx L$~\\cite{malyshkin_2005}. We conclude that in the case of constant resistivity both the ``global'' and ``local'' equations approaches consistently lead to the Sweet-Parker reconnection rate and the Sweet-Parker geometry of the reconnection layer ($L'\\approx L$).  Finally, let us point out that whether reconnection is unforced (free) or forced does not matter for our results and conclusions. Indeed, on one hand, in the case of unforced reconnection one solves Eqs.~(\\ref{RATE}) and~(\\ref{V_R}) for the current $j_o$ and for the reconnection velocity $V_R$. The solution will depend on the magnetic field $B_m$ in the upstream region at point~$M$. On the other hand, in the case of forced magnetic reconnection the reconnection velocity $V_R$ is prescribed and fixed. In this case the field $B_m$ in the upstream region should be treated as an unknown parameter, and Eqs.~(\\ref{RATE}) and~(\\ref{V_R}) are to be solved together in order to find the correct quasi-stationary values of $j_o$ and $B_m$. In other words, in the forced reconnection case an initially weak outside magnetic field $B_m$ gets piled up to higher values until the resulting current $j_o$ in the reconnection layer becomes large enough to be able to match the prescribed velocity $V_R$ of magnetic flux and energy supply in the upstream region."
    },
    "0609/astro-ph0609497_arXiv.txt": {
        "abstract": "{Numerous surveys are currently underway or planned that aim to exploit the strengths of the Lyman-alpha emission line for cosmological purposes. Today, narrowband imaging surveys are frequently used as a probe of the distant universe.} {To investigate the reliability of the results of such high-$z$ \\lya\\ studies, and the validity of the conclusions that are based upon them. To determine whether reliable \\lya\\ fluxes (\\flya) and equivalent widths (\\wlya) can be estimated from narrowband imaging surveys and whether any observational biases may be present.} {We have developed software to simulate the observed line and continuum properties of synthetic \\lya\\ galaxies in the distant universe by adopting various typical observational survey techniques. This was used to investigate how detected \\flya\\ and \\wlya\\ vary with properties of the host galaxy or intergalactic medium: internal dust reddening; intervening \\lya\\ absorption systems; the presence of underlying stellar populations.} {None of the techniques studied are greatly susceptible to underlying stellar populations or the relative contribution of nebular gas. We find that techniques that use one off-line filter on the red side of \\lya\\ result in highly inaccurate measurements of \\wlya\\ under all tests. Adopting two off-line filters to estimate continuum at \\lya\\ is an improvement but is still unreliable when dust extinction is considered. Techniques that employ single narrow- and broad-band filters with the same central wavelength are not susceptible to internal dust, but \\lya\\ absorption in the IGM can cause \\wlya\\ to be overestimated by factors of up to 2: at $z=6$, the median \\wlya\\ is overestimated by $\\sim25$\\%. The most robust approach is a SED fitting technique that fits \\ebv\\ and burst-age from synthetic models -- broadband observations are needed that sample the UV continuum slope, 2175\\AA\\ dust feature, and the 4000\\AA\\ discontinuity. We also notice a redshift-dependent incompleteness that results from DLA systems close to the target LAEs, amounting to $\\sim 10$\\% at $z=6$. } {} ",
        "introduction": "}        Galaxies hosting young, violent starbursts are expected to accommodate numerous massive, hot stars that are bright in the far ultraviolet (UV) regime. UV photons with wavelengths shortwards of the Lyman break ionise their local interstellar media (ISM), resulting in recombination nebulae bright in hydrogen lines. Two-thirds of the Lyman-continuum photons are reprocessed as Lyman-alpha (\\lya) under the assumption of case B recombination, and it could be expected that young starburst regions should be consistently \\lya-bright.  Early observations of local star-forming galaxies with the {\\em International Ultraviolet Explorer} (IUE) demonstrated this not to be the case. In small samples of relatively un-evolved galaxies (low dust and metal content) \\lya\\ was frequently shown to be absent or unexpectedly weak (eg. Meier \\& Terlevich \\cite{ref:meier82}; Hartmann et al. \\cite{ref:hartmann88}; Calzetti \\& Kinney \\cite{ref:calkin92}). Although dust could be invoked to explain the absence of this FUV emission line, Giavalisco et al. (\\cite{ref:giavalisco96}) demonstrated that pure dust extinction could not explain the weak \\lya\\ fluxes from any of the galaxies in the {\\em IUE} samples. Instead results suggest that \\lya\\ photons have decoupled from the non-resonant continuum radiation, resulting in line-photons escaping the host over significantly extended path-lengths; thereby increasing their sensitivity to dust (Neufeld \\cite{ref:neufeld90}). On the other hand, Neufeld (\\cite{ref:neufeld91}) also found that resonance scattering can cause {\\em less} attenuation of \\lya\\ photons compared to continuum if the ISM is multiphase: if dusty, neutral clouds are embedded in a dust-free, ionised ISM, resonance scattering can reflect \\lya\\ photons from the surface of the clouds, allowing them to diffuse out of the ISM relatively unattenuated. Radiative transport in such multiphase, clumpy media has more recently been studied by Hansen \\& Oh (\\cite{ref:hansen05}) who found \\lya\\ equivalent widths (\\wlya) could be `boosted' by factors of $2-3$.  In the local universe, the observational situation was further-complicated by \\lya\\ spectroscopic observations using the {\\em Goddard High-Resolution Spectrograph (GHRS)} and {\\em Space Telescope Imaging Spectrograph (STIS)} both onboard {\\em Hubble Space Telescope (HST)}. In a sample of 8 local starbursts, Kunth et al. (\\cite{ref:kunth98}) found that when \\lya\\ is seen in emission, blueshifted \\lya\\ absorption features are also common, leading to P-Cygni-like profiles. In the same study, \\ion{O}{i} and \\ion{Si}{ii} UV absorption lines were also blueshifted with respect to the ionised gas, indicating large-scale outflows of the ISM with velocities up to 200~\\kms. Mas-Hesse et al. (\\cite{ref:mas-hesse03}) demonstrated that varying \\lya\\ profiles from starbursts are well explained by an evolutionary super-shell model; over the lifetime of a starburst, absorption, pure emission, and P-Cygni emission phases will all be observed, depending on the properties of the outflow and viewing angle. If \\lya\\ escape is associated with winds and photons can diffuse through \\ion{H}{i}, low surface brightness \\lya\\ emission may be expected to span large areas, significantly more extended than the starburst itself as probed by FUV emission. Hence imaging becomes an efficient method by which the line can be isolated. This was the motivation for our HST survey using the {\\em Advanced Camera for Surveys} (ACS; Kunth et al. \\cite{ref:kunth03}). Such \\lya\\ imaging studies are extremely sensitive to how the continuum is subtracted and much supporting data is required in order to estimate the continuum flux at \\lya. We found in Hayes et al. (\\cite{ref:hayes05}) that at least three off-line imaging observations are necessary in order to estimate the continuum flux at \\lya. We determined it necessary to sample the UV continuum slope and 4000\\AA\\ break in order to disentangle the effects of age and reddening, allowing us to estimate the flux due to continuum processes in the on-line filter from synthetic models. In our sample, luminous blue compact galaxy ESO\\,338-IG04 shows regions of diffuse \\lya\\ emission spanning several kpc, and patchy structure consistent with an inhomogeneous and outflowing ISM. In the diffuse regions \\wlya\\ is shown to be very large $(>200\\AA)$ where continuum flux is weak, but integrated over the whole galaxy \\wlya\\ is $\\sim20$\\AA, an order of magnitude lower. Damped absorption with \\wlya$\\sim -40$\\AA\\ is also observed in some central regions, demonstrating the sensitivity of \\lya\\ to small-scale variations in the ISM.  Due to its sensitivity to star-forming activity, high natural equivalent widths (up to 240\\AA; Charlot \\& Fall \\cite{ref:cf93}) sensitivity to dust, and convenient spectral positioning in the FUV, \\lya\\ has long been considered a competitive tracer of primeval high-$z$ galaxies as they form their first generation of stars. Partridge \\& Peebles (\\cite{ref:pp67}) suggested that in a collapsing \\ion{H}{i} halo, as much as 7\\% of the gravitational energy could be radiated in the \\lya\\ line. While surveys that target continuum emission or utilise dropout techniques may have been highly successful in recent years, they all share one common weakness: a luminosity bias in favour of galaxies with strong continuum. Such a bias leads to incompleteness at low luminosity and these surveys would miss the numerous low-mass, star-forming objects; the dominant species in the paradigm of hierarchal galaxy formation.  The early days of high-$z$ \\lya\\ studies were largely unsuccessful (see Pritchet \\cite{ref:prit94} for a review) and it was not until the late 1990's that searches for high-$z$ \\lya-emitters (LAE) became fruitful. Despite the complications, high-$z$ \\lya-emitting objects are now routinely being discovered by imaging techniques (Cowie \\& Hu \\cite{ref:cowiehu98}; Rhoads et al. \\cite{ref:rhoads00}; Fynbo et al. \\cite{ref:fynbo03}; Kodaira et al. \\cite{ref:kodaira03}; Yamada et al. \\cite{ref:yamada05}), fields surrounding massive objects (eg. Monaco et al. \\cite{ref:monaco05}), and spectroscopic observations (eg. Kurk et al. \\cite{ref:kurk04}). After much initial disagreement, the predicted and observed luminosity functions (LF) of LAEs at high-$z$ may be converging (Le Delliou et al. \\cite{ref:ledelliou05}). Interestingly however, a fraction of the LAEs uncovered in the high-$z$ universe have very large measured equivalent widths (Cowie \\& Hu \\cite{ref:cowiehu98}; Malhotra \\& Rhoads \\cite{ref:malrhoads02}). In the case of Malhotra \\& Rhoads (\\cite{ref:malrhoads02}), their LAE population at $z=4.5$ was shown to have a median \\wlya\\ of 400\\AA, 65\\% greater than the maximum value of 240\\AA\\ (Charlot \\& Fall, \\cite{ref:cf93}) that can be generated by pure star-formation models with normal metallicities and IMFs. These authors interpreted this to imply that either these objects were type II quasars or star-forming galaxies with very top-heavy IMFs or zero-metallicity stars. Subsequent X-ray observations with the Chandra satellite ruled out the possibility that these objects could be AGN (Wang et al. \\cite{ref:wang04}). Jimenez \\& Haiman (\\cite{ref:jimenez06}) demonstrate how this (and other effects) can be explained if 10-30\\% of the stars in these galaxies are primordial. `Normal' starburst conditions could still explain such high equivalent widths but the system would have to be perfectly configured so as to allow \\lya\\ photons to leak out while dust blocks much of the stellar UV radiation. Another possibility is that this is the result of some other astrophysical effect, external to the host galaxy that may cause certain observational techniques to overestimate \\wlya\\ when such narrowband imaging techniques are used.  While \\lya\\ may provide a `clean' probe of the high-$z$ universe, the common observational trade-off is evident: spectroscopic studies may be rich in information regarding the line, while narrowband imaging studies may be efficient for detections. Future generations of high-$z$-optimised integral field units (eg. the {\\em Multi Unit Spectroscopic Explorer (MUSE)} bound for ESO's VLT) have the potential to provide significant advances in this area. However, in the cases where narrowband imaging alone is used to derive information about the line, it is essential to examine exactly how efficient the technique is in doing so. To our knowledge, no such comprehensive study has previously been performed and this article represents the first step in doing so. We here simulate how various survey techniques may estimate \\lya\\ detection properties -- line flux (\\flya) and equivalent width (\\wlya) -- of galaxies at high-$z$.  The paper is organised as follows: in Sect. \\ref{sect:sw} we describe the software; in Sect. \\ref{sect:sim} we present the simulations we have performed; in Sect. \\ref{sect:resdis} we present and discuss some of the more important results; and in Sect. \\ref{sect:conc} we present our concluding remarks. We assume a flat cosmology with $\\Omega_\\mathrm{M}=0.3$, $\\Omega_\\Lambda=0.7$, $H_0=70$~km~s$^{-1}$~Mpc$^{-1}$ throughout. ",
        "conclusions": "} Simulations of observations of high-redshift \\lya\\ emitting galaxies have been performed. Specifically, the simulations have examined how efficiently the intrinsic values of \\lya\\ flux and equivalent width are recovered by narrowband imaging observations. \\flya\\ and \\wlya\\ are determined by the four methods described in Sect. \\ref{sect:sim} and the effects of intervening \\ion{H}{i} clouds, internal dust-reddening, and underlying stellar populations, have been investigated. In summary: \\\\ -- Observing \\lya\\ with one off-line continuum filter has been shown to be highly ineffectual. The steep UV continuum slope ($\\beta$) may cause overestimates of \\wlya\\ by a factor of almost 2, while the neglection of dust reddening may cause the reverse effect: narrowband imaging observations may interpret bright \\lya\\ emitters (\\wlya$>100$\\AA) as \\lya\\ absorbing systems. \\\\ -- Using two filters to estimate the UV continuum slope and extrapolating to \\lya\\ using a power-law improves the situation to varying degrees. In the dust-free cases this technique becomes a highly competitive method of determining \\flya\\ and \\wlya. However, even modest amounts of dust render this technique highly ineffectual since dust modifies the continuum in such a way that the power-law parameterisation becomes unreliable. $E_{B-V} = 0.1$ results in 25\\% underestimate of \\wlya. \\\\ -- Economic observing strategies that utilise a narrowband and single broadband filter with the same central wavelength are far less susceptible to dust. However, the continuum flux estimation in such observations becomes suspect due to the blue half of the broadband filter being on the blue side of \\lya, and hence being susceptible to absorption of the continuum by intervening \\ion{H}{i} clouds. Such a technique can lead to overestimates of \\wlya\\ by factors of $>2$. With the addition of a simple sky-noise model, this effect becomes more pronounced with decreasing $S/N$. We suggest that in such observations, simulations are performed to estimate the biases caused by such effects. \\\\ -- SED-fitting techniques that observe only redwards of \\lya\\ are not susceptible to intervening \\ion{H}{i} absorption or dust reddening, when a single extinction law is considered at all wavelengths. Such techniques need to sample the UV continuum slope, 2175\\AA\\ dust feature, and 4000\\AA\\ break. This way reliable estimates of the continuum flux at \\lya\\ can be made and \\wlya\\ estimates are much better constrained in the presence of noise. \\\\ -- None of the techniques we considered here were hugely sensitive to the presence of underlying old stellar populations or randomly mixed populations. Additional populations have no detrimental effect on the efficiency of the technique utilising two filters centred on \\lya. The contribution of such populations to the 4000\\AA\\ break cause only minor ($<10$\\%) errors in the SED-fitting routine. Similarly increasing the relative contribution of nebular emission has no significant impact upon the estimates of \\flya\\ and \\wlya. \\\\ -- Independent of technique, we find a redshift-dependent incompleteness that results from \\ion{H}{i} systems along the line-of-sight. Such an effect would have to be allowed for in the compilation of a \\lya\\ luminostiy function. The red damping wing of DLAs close to the LAE can remove sufficient flux to completely suppress the \\lya\\ line. At $z=6$, one could expect to miss around 10\\% of the sources using our perfect filter set.   We have shown that to recover \\flya\\ and \\wlya\\ from high-$z$ LAE systems, a SED fitting method is preferable if observations can reach the required depths. Other techniques suffer badly from either \\lya\\ absorption along the line-of-sight or reddening by internal dust. If such off-band observations are not available, the use of a single narrow- and broadband filter with the same central wavelength seems to be preferable, although such observations should be complemented with some tuned statistical study of the effects of intervening \\ion{H}{i} systems."
    },
    "0609/astro-ph0609174_arXiv.txt": {
        "abstract": "We present coefficients for the calculation of the continuous emission spectra of H~{\\sc i}, He~{\\sc i} and He~{\\sc ii} due to electron-ion recombination. Coefficients are given for photon energies from the first ionization threshold for each ion to the $n=20$ threshold of hydrogen ($36.5\\mu$), and for temperatures 100~K$~\\le~T~\\le~10^5$~K. The emission coefficients for He~{\\sc i} are derived from accurate {\\it ab initio} photoionization data. The coefficients are scaled in such a way that they may be interpolated by a simple scheme with uncertainties less than 1\\% in the whole temperature and wavelength domain. The data are suitable for incorporation into photoionisation/plasma codes and should aid with the interpretation of spectra from the very cold ionised gas phase inferred to exist in a number of gaseous clouds. ",
        "introduction": "Analyses of optical recombination line (ORL) spectra of photoionised regions have suggested the existence of cold ionised gas (100~K--2000K), mixed within a warmer component at more typical nebular electron temperatures (8000~K--10000~K).  The pockets of cold ionised gas of high metal content are invoked as a possible explanation of the long-standing problem of the discrepancy between elemental abundances derived from ORLs and those derived from collisionally excited lines (CELs).  In this scenario, ORLs and CELs are preferentially emitted by the cold and warm phases, respectively (for a recent review see Liu, 2002); this problem is closely linked to the observation that Balmer jump temperatures of H~{\\sc ii} regions and Planetary Nebulae are systematically lower than those derived from the $[$O~{\\sc iii}$]$ nebular to Auroral line ratio (Liu \\& Danziger 1993).  Understanding the complex spectra arising from such regions relies on the construction of detailed photoionisation models able to account for all gas phases that may be present. Currently, one of the major limitations in such modelling is the lack of an accurate atomic data set extending to such low temperatures. The need of low-temperature effective recombination coefficient for the calculation of recombination lines for metals is apparent. However, in addition to the discrete emission line spectrum, a continuous emission is also produced by the ionised gas, mainly due to free-bound recombination processes of hydrogen and helium ions, free-free transitions in the Coulomb fields of H$^+$, He$^+$ and He$^{2+}$ and two-photon decay of the 2~$^2$S$_{1/2}$ level of H~{\\sc i} and He~{\\sc ii}, and, less importantly of the 2~$^1$S$_{0}$ level of He~{\\sc i}. Accurate continuous emission coefficients are essential for the correct prediction of the Balmer jump by photoionisation codes.  The importance of the continuum processes listed above has long been known and emission coefficients have been tabulated (see e.g. Seaton 1955, 1960; Brown \\& Mathews 1970), for a range of temperatures and wavelengths mainly aimed at the study of optical data for classical H~{\\sc ii} regions. Ferland (1980) derived H~{\\sc i} and He~{\\sc ii} continuous emission and recombination coefficients for a wider range of temperatures (500~K--2$\\cdot$10$^6$~K) and wavelengths to aid the interpretation of ultraviolet (UV) and infrared (IR) as well as optical data from nova ejecta.  In this work we present new calculations of the H~{\\sc i}, He~{\\sc i} and He~{\\sc ii} continuous emission coefficients due to free-bound recombination, over temperatures ranging from 100~K to 10$^{5}$~K, for the physical conditions thought likely to occur in chemically inhomogeneous regions, which may include pockets of cold ionised material intermixed within typical nebular gas. ",
        "conclusions": "We have presented new calculations of the free-bound continuous emission coefficients for the hydrogen and helium ions. The results for He~{\\sc i} are derived from accurate {\\it ab initio} photoionization cross-section data. The coefficients are given for a wide range of temperatures and frequencies, extending to the previously unexplored very low-temperatures regime. This is needed for the interpretation of spectra from very cold ionised gas, which has been inferred to exist, possibly in the form of density/chemical inhomogeneities, from observations of ORLs in H~{\\sc ii} regions and planetary nebulae. The data are presented under an interpolation scheme that allows estimates to be obtained over the entire temperature and energy range presented with less than 1\\% uncertainty."
    },
    "0609/astro-ph0609160_arXiv.txt": {
        "abstract": "{The Taurus Molecular Cloud (TMC) is the nearest large star-forming region, prototypical for the distributed mode of low-mass star formation. Pre-main sequence stars are luminous X-ray sources, probably mostly owing to magnetic energy release. } {The {{\\it XMM-Newton Extended Survey of the Taurus Molecular Cloud}} (XEST) presented in this paper surveys the most populated $\\approx$5 square degrees of the TMC, using the {{\\it XMM-Newton}} X-ray observatory to study the thermal structure, variability, and  long-term evolution of hot plasma, to investigate the  magnetic dynamo, and to search for new potential members of the association. Many targets are also studied in the optical, and high-resolution X-ray grating spectroscopy has been obtained for selected bright sources.   } {The X-ray spectra have been coherently analyzed with two different thermal models (2-component thermal model, and a continuous emission measure distribution model). We present overall correlations with fundamental stellar parameters that were derived from the previous literature. A few detections from {{\\it Chandra}} observations have been added.} {The present overview paper introduces the project and provides the basic results from the X-ray analysis of all sources detected in the XEST survey. Comprehensive tables summarize the stellar properties of all targets surveyed. The survey goes deeper than previous X-ray surveys of Taurus by about an order of magnitude and for the first time systematically accesses very faint and strongly absorbed TMC objects. We find a detection rate of 85\\% and 98\\% for classical and weak-line T Tau stars (CTTS resp. WTTS), and identify about half of the surveyed protostars and brown dwarfs. Overall, 136 out of 169 surveyed stellar systems are detected.  We describe an X-ray luminosity vs. mass correlation, discuss the distribution of X-ray-to-bolometric luminosity ratios, and show evidence for lower X-ray luminosities in CTTS compared to WTTS. Detailed analysis (e.g., variability, rotation-activity relations, influence of accretion on X-rays) will be discussed in a series of accompanying papers.  } {} ",
        "introduction": "Studies of star-forming regions have drawn a picture in which cool, molecular gas, contracting to protostars with surrounding accretion disks, co-exists with high-energy radiation. The latter is emitted by plasma in the stellar environment that is continuously heated to temperatures beyond one million degrees. This radiation is most prominently seen in the soft (0.1--10~keV) X-ray range and is conventionally attributed to the presence of magnetically trapped plasma in the outer stellar atmosphere. This plasma is reminiscent of, in the broadest sense, the solar coronal plasma. There is little doubt that both lower-energy (extreme ultraviolet) and higher-energy (hard X-ray and gamma-ray) radiation is amply present, but strong photoelectric absorption or, respectively, extremely low flux precludes direct detection. Further, non-thermal radio emission from relativistic electrons trapped in the magnetic fields gives unambiguous evidence for the presence of an accelerated population of electrons similar to populations seen in the solar and in stellar coronae during flares \\citep{guedel02}. Again, by implication from the solar  analogy, we expect that high-energy, non-thermal particles (electrons and ions) play a crucial role also outside the confines of the magnetic stellar atmosphere.  But in contrast to the solar case, high-energy radiation from forming stars has a profound influence on its environment as dense accretion disks and the larger-scale molecular envelope are irradiated, heated, and ionized by high-energy photons and particles. There are numerous consequences of these processes. Ultraviolet and X-ray irradiation of molecular disks heats their upper layers \\citep{alexander04, glassgold04}, making them accessible to magnetic fields, thus inducing, for example, the magnetorotational instability \\citep{balbus91}. The latter is thought to be a prime source for increased disk viscosity, which in turn is the main driver of the accretion process and thus the stellar formation process in the first place. Complex chemical networks are also put in operation by heating and ionizing the molecular environment \\citep{glassgold04}, thus fundamentally altering composition, opacity, and cooling behavior of the material. Magnetic fields are thought to attach to the inner border of  the lightly ionized accretion disk, thus  transporting material to the stellar  surface while at the same time applying torques to the star that may regulate its rotation period  \\citep{montmerle00}. If the stellar magnetic fields are generated in a solar-like dynamo in which rotation is a key parameter, then the long-term evolution of the magnetic field production and thus of the high-energy radiation is indeed directly controlled by the accretion process in the feedback loop sketched above.  Further, magnetic fields and ionization of the disk may be relevant for the launch of bipolar jets and molecular outflows \\citep{pudritz86}, a widely observed but poorly understood mechanism that seems to regulate the mass buildup of the forming star. Alternatively, these jets may originate from the magnetic fields spanning between the star and the disk \\citep{shu94}, or from plasma clouds ejected in the process of magnetic reconnection of these same fields \\citep{hayashi96}.  The origin of the X-ray emission is still not entirely clear. While the bulk of the X-rays from T Tau stars is thought to originate in a magnetically trapped corona consisting of complex magnetic loops anchored in photospheric magnetic regions like in the Sun, features unknown to the Sun may exist. Speculation has arisen that some very soft X-rays may be formed in accretion shocks near the stellar surface  (e.g., \\citealt{kastner02}). Herbig-Haro objects far from the forming stars have been detected to be very soft X-ray emitters due to shock interactions with the interstellar medium (e.g., \\citealt{pravdo01, bally03, guedel05}), but new spectral phenomenology suggests that appreciable X-ray emission may also be formed close to the stars, at the (spatially unresolved) base of the jets otherwise seen in radio and in the optical \\citep{guedel05}. Whatever the analogy to the Sun, we also note that the level of the X-ray luminosity (between a few times $10^{28}$~erg~s$^{-1}$ and $10^{31}$~erg~s$^{-1}$ for low-mass pre-main sequence stars) and the measured plasma temperatures of up to 100~MK are extreme and may imply heating mechanisms related to continuous flaring \\citep{guedel03, wolk05}.  Studies of X-ray radiation in pre-main sequence stars are thus crucial for our further understanding of star formation, accretion processes, angular momentum transport, and mass outflow mechanisms in the earliest stages of a star's life. For recent reviews of X-rays from pre-main sequence stars, we refer the reader to \\citet{feigelson99}, \\citet{guedel04}, and \\citet{feigelson06}. {\\it {\\it XMM-Newton}} and {\\it Chandra} penetrate dense molecular envelopes providing unequaled information on high-energy processes. A recent, large survey of the Orion Nebula Cluster (ONC) with {\\it Chandra} (the COUP project, \\citealt{getman05}) has monitored approximately 1400 young stars in a $17\\arcmin\\times 17\\arcmin$ field around the Trapezium. This study has provided unprecedented information on the ONC stellar population and the evolution of X-ray production. For example, \\citet{preibisch05} concluded that a significant correlation between rotation and X-ray luminosity consistently found in main-sequence stars is absent in the ONC sample, arguing for a different magnetic dynamo process or a very different internal structure of these youngest stars. X-ray radiation was found to be largely variable, with numerous flares contributing significantly to the overall X-ray production \\citep{wolk05}. We refer the interested reader to a special volume of papers from this study (\\citealt{getman05} and associated papers), or to a summary provided by \\citet{feigelson06}.  The ONC  represents the so-called ``clustered mode'' of star formation in which hundreds to thousands of low-mass stars form around massive O-type stars such as those in the Trapezium. Such environments produce extreme stellar densities, and the presence of hot stars has a decisive impact on star formation. UV radiation and the winds of the central O stars have a deteriorating effect on the molecular environment, the pre-stellar  cores and the circumstellar disks by ionizing and dispersing the molecular material, prematurely altering the accretion environment \\citep{odell98} and eventually halting the star formation process altogether \\citep{dolan01}. The ionized H~II Orion Nebula is the most vivid expression of this environment.  Star formation occurs, however, also in relatively unspectacular low-mass star forming regions devoid of the influence of radiation and outflow-related effects from luminous O stars. These so-called ``dark clouds'' typically form a smaller number of stars and often do so in a ``dispersed'' or ``isolated mode'' in which stars or very small groups of stars form in relative isolation. Stars form in a similar manner also in high-mass star formation regions, outside the clusters and at larger distances from the central OB stars \\citep{carpenter00}.  The closest and best-studied  large, low-mass star formation region is the  Taurus Molecular Cloud (TMC) complex that extends to the adjacent Auriga Clouds. TMC has served as a testbed for low-mass star-formation theory and has provided some of the best examples of phenomenology relevant to star formation. There are only a few Herbig Ae/Be stars in this vast region (V892 Tau and AB Aur in our survey). The rest are predominantly   sub-solar in mass although there have been claims that a substantial population of non-accreting B stars may be members of the Taurus population \\citep{walter91}, in agreement with expectations from the initial mass function. Owing to the near-absence of massive stars and their mechanical/radiative effects, the molecular gas is still  present in large amounts ($M_{\\rm cloud} \\approx ~3\\times 10^4~M_{\\odot}$, \\citealt{ungerechts87}) much in excess of the total mass in stars already formed ($M_{\\rm stars}$ of order  $2\\times 10^2~M_{\\odot}$).    The Taurus region is quite large (some 10--15 degrees in diameter, corresponding to about 25--35~pc at a distance of 140~pc), making comprehensive studies of the entire population difficult. We have  initiated a large X-ray study of the TMC with the {\\it XMM-Newton} X-ray observatory, concentrating on the denser cloud areas that contain the majority of the TMC stellar population. The emphasis of this {\\it XMM-Newton Extended Survey of the Taurus Molecular Cloud} (XEST henceforth)  is on a wide-field sampling rather than on long exposures. It goes systematically deeper by about an order of magnitude than previous Taurus X-ray surveys. While it is less sensitive than the COUP survey, it surveys an area about 60 times larger ($\\approx 5$~square degrees), which is required because of the low surface density of Taurus members. The survey is accompanied by a deep, large-field optical survey with the Canada-France-Hawaii Telescope (CFHT) and a mid-infrared survey of the entire cloud complex with the {\\it Spitzer Space Telescope} \\citep{padgett06}. These latter surveys will be presented separately (see summary of all three surveys in \\citealt{guedel06a}).  All studies combined provide  an unsurpassed database for the nearest major star-forming cloud complex.  The purpose of the present paper is to give an introduction to the XEST  project, to present an overview of the surveyed stellar sample, and to  discuss the strategies followed during the data reduction and analysis.  The detailed results are described in a series of associated papers.  The outline of this paper is as follows. We discuss some pivotal properties of the TMC in Sect.~\\ref{TMC}, including  results from previous studies that support the key role of this region for further star-formation studies. Sect.~\\ref{XEST} introduces the TMC X-ray survey, and Sect.~\\ref{strategy} presents the data reduction and analysis strategies. Sect.~\\ref{results} contains an overview of the basic results, presenting detected and undetected stellar populations, accompanied by comprehensive tabulations of fundamental properties of all surveyed objects and spectral X-ray results from the detections. We summarize overall statistical properties of XEST in Sect.~\\ref{conclusions}.  The subsequent series of papers related to this survey discusses X-ray properties in the context of stellar class and accretion \\citep{telleschi06a}, rotation-activity-age relations \\citep{briggs06a}, a dedicated study of the L1495E subsample \\citep{silva06}, correlated behavior of X-rays and optical/ultraviolet emission \\citep{audard06}, X-ray properties in a sample of TMC protostars \\citep{briggs06b}, X-rays from jet-driving stars in TMC \\citep{guedel06b}, high-resolution X-ray spectroscopy of classical and weak-line T Tau stars \\citep{telleschi06b}, X-rays from brown dwarfs \\citep{grosso06a}, a $U$-band survey of brown dwarfs with the {\\it XMM-Newton} Optical Monitor \\citep{grosso06b}, light-curve variability studies \\citep{stelzer06}, interpretation of X-ray flares \\citep{franciosini06}, an investigation of statistical fluctuations in X-ray light curves \\citep{arzner06a}, an investigation of spectral parameters in extremely faint sources \\citep{arzner06b}, an analysis of the gas-to-dust ratio in the TMC region \\citep{glauser06},  a case study of the accreting, prototypical CTTS T Tau \\citep{guedel06c}, a study of the supposedly single Herbig star in TMC, AB Aurigae \\citep{telleschi06c}, and a search for new TMC members based on near-infrared and X-ray properties of field sources detected in XEST \\citep{scelsi06}. ",
        "conclusions": "\\subsection{Spectral interpretation: Quality of results}   \\begin{figure}[t!] \\centerline{\\resizebox{0.95\\hsize}{!}{\\includegraphics{guedel_f7a.ps}}} \\centerline{\\resizebox{0.95\\hsize}{!}{\\includegraphics{guedel_f7b.ps}}} \\caption{Comparison of $N_{\\rm H}$ values derived from the DEM method (x axis) and from 1-$T$ or 2-$T$ fits (y axis; Tables~\\ref{tab5} and \\ref{tab6}). Only results from {\\it XMM-Newton} data are reported. Upper plot is for sources with more than 600 counts in total (after Table~\\ref{tab4}), lower plot for sources with fewer than 600 source counts. The solid lines indicate equal $N_{\\rm H}$.}\\label{NHcomp} \\end{figure}  To assess the quality of our spectral interpretations, we now compare results derived from the two fit methods (DEM vs 1- or 2-$T$ plasma). Fig.~\\ref{NHcomp} compares $N_{\\rm H}$ values derived from the DEM fits (shown along the abscissa) and from the 1- or 2-$T$ fits (along the ordinate). We show results for bright spectra ($> 600$~cts for the source counts parameter in Table~\\ref{tab4} of which typically about half are found in the PN camera) and for faint spectra ($< 600$~cts) separately. The agreement is excellent. The brighter sources show a trend toward  $N_{\\rm H}$ being slightly lower for the 1-$T$ or 2-$T$ fit than for the DEM fit. This can be  understood as follows. The DEM model assumes the presence of cool material down to $\\log T = 6$. This plasma accounts for some of the soft emission where the 1- or 2-$T$ model may be subject to an  EM deficit. To account for the same spectral flux, the trend is to lower $N_{\\rm H}$ for the spectrum with deficient cool EM. The scatter obviously becomes larger for fainter sources, but the error bars grow accordingly so that the agreement is of similar quality (lower plot).  \\begin{figure} \\centerline{\\resizebox{0.95\\hsize}{!}{\\includegraphics{guedel_f8a.ps}}} \\centerline{\\resizebox{0.95\\hsize}{!}{\\includegraphics{guedel_f8b.ps}}} \\caption{Comparison of $L_{\\rm X}$ values derived from the DEM method (x axis) and from 1-$T$ or 2-$T$ fits (y axis; Tables~\\ref{tab5} and \\ref{tab6}). Only results from {\\it XMM-Newton} data are reported. Upper plot is for sources with more than 600  counts in total (after Table~\\ref{tab4}), lower plot for sources with fewer than 600 source counts. The solid lines indicate equal $L_{\\rm X}$.}\\label{LXcomp} \\end{figure}  \\begin{figure} \\centerline{\\resizebox{0.95\\hsize}{!}{\\includegraphics{guedel_f9a.ps}}} \\centerline{\\resizebox{0.95\\hsize}{!}{\\includegraphics{guedel_f9b.ps}}} \\caption{Comparison of $L_{\\rm X}$ values derived from the DEM method (x axis) and from 1-$T$ or 2-$T$ fits (y axis; Tables~\\ref{tab5} and \\ref{tab6}). Only results from {\\it XMM-Newton} data are reported. Upper plot is for sources with $N_{\\rm H} < 5\\times 10^{21}$~cm$^{-2}$, lower plot for $N_{\\rm H} > 5\\times 10^{21}$~cm$^{-2}$ (based on the DEM fits in Table~\\ref{tab5}). The solid lines indicate equal $L_{\\rm X}$.}\\label{LXcomp2} \\end{figure}  Fig.~\\ref{LXcomp} analogously compares $L_{\\rm X}$. For most of the brighter stars, the agreement is excellent although there are a few cases for which the 1- or 2-$T$ fit produces significantly higher $L_{\\rm X}$. The reason for this deviation is the following. Soft components can be added with relatively little constraints if they are sufficiently absorbed. The temperature then becomes uncertain. A cooler plasma subject to photoelectric absorption will, however, require a much larger EM to produce the same observed flux. If the cool EM is a free fit parameter, an instability may occur with $T$ converging to low values while $N_{\\rm H}$ and EM grow excessively. This becomes more evident for fainter sources (lower plot). In several cases, we therefore fixed $N_{\\rm H}$ at the value expected from $A_{\\rm V}$ for the DEM fit, whereas the 1-$T$ fits, requiring one parameter less, were performed without this constraint. The sources that likely suffer from this numerical instability in the 1-$T$ fit are CIDA 7 (XEST-10-034), HO Tau (XEST-09-010), V410 X6 (XEST-23-061),  CFHT-Tau 7 (XEST-03-017), MHO 9 (XEST-22-013), and V410 A25 (XEST-23-029) that all show $L_{\\rm X}$ from the 1-$T$ or 2-$T$ fits at least 3 times higher compared to the DEM fit. In the former, they show $N_{\\rm H}$ considerably higher than in the latter, or much higher than expected from $A_{\\rm V}$, while the dominant electron temperature was very low, between 1~MK and 3.3~MK. The first four of those six stars  are faint, while MHO 9 is at the limit between bright and faint as defined for the plots in Fig.~\\ref{LXcomp}.  Given that the DEM model is more rigid and avoids biases toward  strongly absorbed soft emission from large amounts of cool EM, we will base most of our statistical investigations in the accompanying papers on the DEM results. This model is also more physically meaningful because work on bright spectra of nearby stars has clearly shown continuous emission measure distributions of this kind (e.g., \\citealt{telleschi05}). As shown here and below, however, the results for $L_{\\rm X}$, $N_{\\rm H}$, and $T_{\\rm av}$ from the two methods agree well for the brighter spectra. Caution is in order essentially only for spectra defined as ``faint'' here.   \\begin{figure} \\centerline{\\resizebox{0.95\\hsize}{!}{\\includegraphics{guedel_f10a.ps}}} \\centerline{\\resizebox{0.95\\hsize}{!}{\\includegraphics{guedel_f10b.ps}}} \\caption{Comparison of $T_{\\rm av}$ values derived from the DEM method (x axis) and from 1-$T$ or 2-$T$ fits (y axis; Tables~\\ref{tab5} and \\ref{tab6}). Only results from {\\it XMM-Newton} data are reported. Upper plot is for sources with more than 600 counts in total (after Table~\\ref{tab4}), lower plot for sources with fewer than 600 source counts. The solid lines indicate equal $T_{\\rm av}$.}\\label{Tavcomp} \\end{figure}  Fig.~\\ref{LXcomp2} shows the same comparison for weakly absorbed sources (upper plot, for $N_{\\rm H} < 5\\times 10^{21}$~cm$^{-2}$) and for more strongly absorbed sources (lower plot, for $N_{\\rm H} > 5\\times 10^{21}$~cm$^{-2}$).  It is evident that higher absorption introduces more uncertainty to the precise determination of $L_{\\rm X}$. Also, these sources obviously tend to be fainter (in terms of count rates).   Fig.~\\ref{Tavcomp} compares the average electron temperature, $T_{\\rm av}$, derived from the DEM and the 2-$T$ or 1-$T$ fits. Given the rather different methodology, the agreement is satisfactory for the brighter sources, while the scatter increases for faint sources, making accurate statements on electron temperatures difficult. The few outliers among the bright sources show very cool average temperatures in 1-$T$ or 2-$T$ fits probably due to the same numerical instability described above. The worst agreement is found for the very faint spectrum of LR~1 (XEST-23-048) for which $T_0$ of the DEM fit is unconstrained despite fixed $\\beta = -1$. In the presence of  severe absorption, the DEM fit finds a low best-fit $T_0$, while the 1-$T$ fit converges to a very high temperature. Both $L_{\\rm X}$ are reasonable, while the DEM fit converges to a higher $N_{\\rm H}$ which is preferred in the light of the very high $A_{\\rm J}$ of this star.   Concluding from this comparison, we are confident that $N_{\\rm H}$ values are meaningful within factors of about 1.5 for the  brighter sample and mostly within factors of 2 for the fainter sample. X-ray luminosities appear to be reasonably constrained within a factor of two (although with exceptions), which is a typical range of variability for most sources in any case.       \\subsection{Population statistics}    Fig.~\\ref{hrd} shows an Hertzsprung-Russell Diagram (HRD)  of all sources observed in the framework of this survey with isochrones and evolutionary tracks (after \\citealt{siess00}) overplotted. The source symbols indicate the source type, and the symbol size  corresponds to the X-ray luminosity. \\begin{figure*} \\centerline{\\resizebox{0.9\\hsize}{!}{\\includegraphics{guedel_f11.ps}}} \\caption{HRD of all observed TMC objects for which $L_*$ and $T_{\\rm eff}$ are known. Symbol indicates object type (see panel in lower left corner for the key), and symbol size scales with $\\log L_{\\rm X}$,  crosses marking undetected sources. Additional sources from {\\it Chandra} have been included, but note that the L1551 IRS5 protostar is marked here as a non-detection. Solid lines show evolutionary tracks toward the zero-age main sequence (ZAMS) for masses as given, in units of the solar mass (ZAMS). Dotted curves indicate isochrones at ages given (after \\citealt{siess00}). The lowest of these curves marks the ZAMS. The spectral class ranges given at the bottom of the figure refer to main-sequence stars.}\\label{hrd} \\end{figure*}  The following characteristics are noteworthy: i) The surveyed sources cover a broad range of masses, from the substellar range up to 2.5-3$M_{\\odot}$, the three most massive stars being HD~28867 and the Herbig stars AB Aur and V928 Tau (the latter being a multiple system). Most objects are concentrated at ages between $3\\times 10^5 - 3\\times 10^6$~yr, and very few at ages around $10^7$~yr. ii) As pointed out by \\citet{kenyon95}, CTTS and WTTS (Class II and III, shown in gray and white, respectively) occupy the same region in the HRD. In fact, this is also true for the detected Class-I protostars. The low-mass and substellar HRD will be more specifically discussed by \\citet{grosso06a}. iii) Most  undetected sources are in the very low-mass/substellar regimes and/or have the largest isochronal ages, although for some of the latter objects, the fundamental parameters are poorly determined, as detailed below.  A few objects are apparently located below the zero-age main sequence (ZAMS) which is likely to  be a consequence of poor photometry. This is particularly true for embedded sources where scattered light distorts photometric measurements. The stars below the ZAMS are, in order of increasing $T_{\\rm eff}$: HH~30 (XEST-22-000, an edge-on star-disk system, \\citealt{burrows96}), IRAS~S04301+261 (XEST-18-000, a highly reddened object probably primarily detected in scattered light, \\citealt{briceno02}), Haro 6-5~B = FS Tau B (XEST-11-054 = C2-1, a near-edge-on star-disk system, \\citealt{padgett99}, with $A_{\\rm V} \\approx 10$~mag), HBC 353 (XEST-27-000), and HBC 352 (XEST-27-115).  These stars will not be considered in statistical studies whenever $L_*$ is required (e.g., in $L_{\\rm X}/L_*$).  A few stars are located between the 10~Myr isochrone and the ZAMS, making them much older than the bulk population in Taurus. Several of these objects are relatively poorly studied and may suffer from inaccurate photometry. For most of them, no $A_{\\rm V}$ has been given, or the extinction is large, and indeed several of them have not been detected (crosses in Fig.~\\ref{hrd}). These objects are, in order of  increasing $T_{\\rm eff}$ above the 0.1$M_{\\odot}$ track: ITG~33A (XEST-07-000, suggested to be an  edge-on  star-disk system, \\citealt{martin00}: $A_{\\rm V} = 3.5$~mag), FS Tau  (XEST-11-057, a strongly extincted close binary of CTTS, $A_{\\rm V} \\approx 5$~mag), CoKu Tau 1 (XEST-23-000, an edge-on star-disk system, \\citealt{padgett99}, $A_{\\rm V} = 6.8$~mag), HN Tau (C5-2), LR 1 (XEST-23-048, with a large $A_{\\rm J} = 6.4$~mag), V955 Tau (XEST-10-020, $A_{\\rm V} = 3.7$~mag), V410 A20 (XEST-23-000, with a large $A_{\\rm J} = 6.57$~mag), and V410 A24 (XEST-23-000, with a large $A_{\\rm J} = 6.73$~mag - see Table~\\ref{tab9} for $A_{\\rm V}$ and $A_{\\rm J}$ values).   The detection statistics of our X-ray survey is summarized in Table~\\ref{tab12} (considering only the {\\it XMM-Newton} observations), based on our final classification scheme in Table~\\ref{tab11}. Here, the protostar L1551 IRS5 has been treated as a non-detection as the source closely associated with this binary has been interpreted as X-ray emission from the jets \\citep{favata02, bally03}. We also list the statistics (in parentheses) if {\\it Chandra} detections and non-detections are counted (note that DG Tau B and FV Tau/c have been counted as non-detections in the {\\it XMM-Newton} sample, while they are detections with {\\it Chandra}). An important point for further statistical studies is that the X-ray sample of detected CTTS and WTTS is nearly complete for the surveyed fields (as far as the population is known). Most of the few remaining, undetected objects are considerably extincted and by implication X-ray absorbed: IRAS S04301+261 = XEST-18-000  with $A_{\\rm J} \\approx 1.8$, and CoKu Tau 1 = XEST-23-000, CFHT-Tau 19 = XEST-11-000, FV Tau/c = XEST-02-000, CFHT-Tau 20 = XEST-13-000, ITG33 A = XEST-07-000,  CFHT-Tau 8 = XEST-07-000, and HH~30 = XEST-22-000  with  $A_{\\rm V} = $ 6.8, 7.3, 3.25, 3.6, 3.5, 1.8, and $\\approx 3$~mag and CFHT-Tau 12 = XEST-17-000 with $A_{\\rm J} = 3.44$~mag, where the $A_{\\rm V}$ of  IRAS S04301+261, CoKu Tau 1, HH~30, ITG 33A  may have been severely underestimated owing to scattered light and strong extinction by near-edge-on disks (see above). FV Tau/c was, however, detected in a {\\it Chandra} observation (C3-1) reported here, with an unexpectedly high photoelectric absorption corresponding to  $N_{\\rm H} \\approx 10^{23}$~cm$^{-2}$. Some of the non-detections are very-low mass stars with late spectral types (CFHT-Tau 19 = XEST-11-000, CFHT-Tau 20 = XEST-13-000, CFHT-Tau 12 = XEST-17-000, and CFHT-Tau 8 = XEST-07-000 with spectral types of  M5.25, M5.5, M6, and M5.5, respectively). Very little is known about IRAS 04108+2803~A. Its companion, a Class~I object, was detected, its light curve showing a strong flare and its spectrum revealing very strong absorption. Perhaps,  both objects are subject to similar absorption. The only formally undetected WTTS in the {\\it XMM-Newton} sample is HBC 353, but it is located on the wings of HBC 352 = XEST-27-115, falls into a CCD gap of the PN detector, and is located far off-axis, elevating the corresponding detection limit to $ \\approx 180$~cts. From Table~\\ref{tab8}, we derive that only 3 out of 45 multiple systems remain undetected  (2/47 including {\\it Chandra} results), while  29 out of 114 single stars are undetected  (31/122 including {\\it Chandra} results). In total,  45 of the 159 systems surveyed by {\\it XMM-Newton}  are multiple (47/169 including {\\it Chandra} results).  \\setcounter{table}{11}  \\begin{table}[t] \\caption{XEST X-ray detection statistics\\label{tab12} } \\begin{tabular}{rlrrl} \\hline \\hline \\multicolumn{2}{l}{Object type}  & Members     &  Detections & Detection \\\\ \\multicolumn{2}{l}{ }            & surveyed    &  \t     & fraction  \\\\ \\hline 0, 1& Protostars  \t  & 20 \\hfill (21)         &     8 \\hfill (10)     & 40\\% \\hfill (48\\%)\t\\\\ 2& CTTS                   & 65 \\hfill (70)         &    55 \\hfill (60)     & 85\\% \\hfill (86\\%)\t\\\\ 3& WTTS                   & 50 \\hfill (52)         &    49 \\hfill (50)     & 98\\% \\hfill (96\\%)\t\\\\ 4& BDs    \t\t  & 16 \\hfill (17)         &     8 \\hfill (9)      & 50\\% \\hfill (53\\%)\t\\\\ 5& Herbig\t\t  & 2  \\hfill {\\ }         &    2  \\hfill {\\ }     & 100\\% \\hfill {\\ } \t\\\\ 9& others/unident.        & 6  \\hfill (7)          &    4  \\hfill (5)      & 67\\% \\hfill (71\\%)\t\\\\ \\hline & Total\t\t          & 159 \\hfill (169)       &    126 \\hfill (136)   & 79\\% \\hfill (80\\%)\t\\\\ \\hline \\multicolumn{5}{l}{NOTES:}\\\\ \\multicolumn{5}{l}{Numbers in parentheses include {\\it Chandra} observations. Source}\\\\ \\multicolumn{5}{l}{near L1551 IRS5 not considered (non-detection for XEST-22-040)}\\\\ \\end{tabular} \\end{table}   Previous X-ray surveys of TMC did not detect the  intrinsically fainter TTS population. Those previous surveys covered a wider field in Taurus (e.g., the ROSAT surveys by \\citealt{neuhaeuser95, stelzer01}) but were about ten times less sensitive and were confined to soft sources. Stellar classes that were thus inaccessible by these earlier surveys but are systematically (albeit not completely) detected in XEST are: i) strongly absorbed, embedded protostars or T Tau stars absorbed by their own disks seen nearly edge-on; their non-detection in previous surveys is little surprising given the strong photoelectric absorption. Most of the detected protostars show no X-ray counts below 2~keV \\citep{briggs06b}. ii) Low-$L_{\\rm X}$ brown dwarfs; in XEST, the detection rate of BDs (50\\% for the {\\it XMM-Newton} sample, 53\\% if one detection from {\\it Chandra} is added) is  high; the remaining objects of this class are likely to be intrinsically fainter than our detection limit rather than being excessively absorbed by gas ($A_{\\rm V}$ of those objects typically being no more than a few magnitudes; see \\citealt{grosso06a} for further discussion). iii) Several T Tau stars with double-peaked X-ray spectra in which a strongly absorbed hard component is accompanied by a soft excess \\citep{guedel06b}.  The incomplete statistics in previous surveys potentially introduces bias into statistical correlations and population studies. On the other hand, we emphasize that the so far cataloged population of Taurus members may not be complete either. In particular, WTTS are predominantly detected by their X-rays (e.g., \\citealt{neuhaeuser95}). The WTTS statistics in Taurus may thus be incomplete in particular with regard to faint X-ray sources or WTTS located behind the TMC. Several dozen new TMC member candidates are indeed reported from XEST \\citep{scelsi06} but will require spectroscopic follow-up observations to confirm membership.  \\subsection{Statistical properties and correlations}  Table~\\ref{tab13} summarizes statistical properties of the surveyed sources, providing median, mean, and standard deviation of the sample for various parameters. The sample sizes vary (given in the last column, N) because not all parameters are available for all sources (see details in the footnotes of the table). The salient features of this compilation are the following: i) The agreement between the sample X-ray parameters derived from the DEM model (``DEM'' in column 2)  and the 1-$T$ or 2-$T$ model (``2-$T$'' in column 2) is excellent. ii) Characteristic electron temperatures are around 6-8~MK as characterized by $T_0$ in the DEM model and by $T_1$ in the 1-$T$ or 2-$T$ model. The hotter component shows a mean and median around 23~MK. iii) The large majority of the sample shows $-4 \\la \\log L_{\\rm X}/L_* \\la -3$. iv) Protostars are, as expected, significantly more extincted in the visual band ($A_{\\rm V}$) in our sample than T Tau stars; among the latter, objects of type 2 (CTTS, Class-II objects) are more strongly extincted than objects of type 3 (WTTS, Class III). The median extinction of the detected sources is $\\approx 1.4$~mag, corresponding to a hydrogen column density of $\\log N_{\\rm H} \\approx 21.44$, close to the measured median of $\\log N_{\\rm H} \\approx 21.5-21.6$. iv) As pointed out on the HRD, the $L_*$ distributions of objects of type 1, 2, and 3 strongly overlap, although a marginal trend is seen toward a decreasing sample {\\it mean} of $L_*$ with increasing type, perhaps reflecting some overall evolutionary trend. v) A wide range of accretion rates is covered (for protostars and CTTS), with a mean and median around $10^{-8}~M_{\\odot}$~yr$^{-1}$.   \\begin{table}[t!] \\caption{XEST parameter statistics\\label{tab13} } \\begin{tabular}{lrrrrr} \\hline \\hline Parameter                        &Sample$^a$& Median      & Mean     & Std.        &  N \\\\ &      &             &          & dev.$^b$    &    \\\\ \\hline \\multicolumn{6}{c}{X-ray parameters} \\\\ \\hline $\\beta^c$                            &  DEM & -1.25       & -1.49    & 0.95        &  99 \\\\ $\\log T_0^c$ [K]                     &  DEM & 6.90        & 6.90     & 0.28        & 119 \\\\ $\\log T_1^{c,d}$ [K]                 & 2-$T$ & 6.79       & 6.76     & 0.21        &  94 \\\\ $\\log T_2^{c,d}$ [K]                 & 2-$T$ & 7.37       & 7.38     & 0.25        & 105 \\\\ $\\log T_{\\rm av}^{e}$ [K]            & DEM  & 7.04        & 7.03     & 0.24        & 123 \\\\ & 2-$T$ & 7.05       & 7.08     & 0.32        & 125 \\\\ $\\log N_{\\rm H}^d$ [cm$^{-2}$]       &  DEM & 21.57       & 21.52    & 0.67        & 123 \\\\ & 2-$T$ & 21.53      & 21.50    & 0.65        & 123 \\\\ $\\log L_{\\rm X}^e$ [erg~s$^{-1}$]    & DEM  & 29.82       & 29.76    & 0.77        & 123 \\\\ & 2-$T$ & 29.80      & 29.72    & 0.78        & 125 \\\\ $\\log L_{\\rm X}/L_*^{e,f}$           & DEM  & -3.56       & -3.55    & 0.50        & 116 \\\\ & 2-$T$ & -3.56      & -3.58    & 0.56        & 118 \\\\ \\hline \\multicolumn{6}{c}{Fundamental parameters} \\\\ \\hline $A_{\\rm V}$ [cm$^{-2}$]              & all  & 1.46        & 2.80     & 3.38        & 116    \\\\ & 1    & 10.20\t  & 9.52     & 1.93\t   & 6\t \\\\ & 2    & 1.95\t  & 3.30     & 3.61\t   & 57    \\\\ & 3    & 0.87\t  & 1.36     & 1.56\t   & 38   \\\\ & det  & 1.35        & 2.60     & 3.36        & 97    \\\\ $\\log L_*^{f,g}$ [$L_{\\odot}$]       & all  & -0.34       & -0.46    & 0.82        & 146    \\\\ & 1    & 0.08        & -0.01    & 0.49        & 16    \\\\ & 2    & -0.25\t  & -0.33    & 0.61        & 60    \\\\ & 3    & -0.43\t  & -0.44    & 0.59        & 46   \\\\ & det  & -0.30       & -0.33    & 0.71        & 119    \\\\ $\\log({\\rm age})^{f,h}$ [Myr]        & all  & 0.39        & 0.38     & 0.46        & 112    \\\\ & 1    & 0.05        & 0.23     & 0.54        & 9    \\\\ & 2    & 0.43\t  & 0.44     & 0.42        & 52    \\\\ & 3    & 0.39\t  & 0.32     & 0.35        & 43   \\\\ $\\log({\\rm mass})^{f,g}$ [$M_{\\odot}$]& all & -0.29       & -0.37    & 0.43        & 123    \\\\ & 1    & -0.17       & -0.23    & 0.22        & 10   \\\\ & 2    & -0.25\t  & -0.27    & 0.31        & 54    \\\\ & 3    & -0.39\t  & -0.40    & 0.34        & 43  \\\\ $\\log(P)^{i}$ [d]                    & all  & 0.69        &  0.64    & 0.28        & 69    \\\\ & 1    & 0.63        & 0.54     & 0.32        & 10   \\\\ & 2    & 0.77\t  & 0.73     & 0.21        & 35    \\\\ & 3    & 0.58\t  & 0.56     & 0.33        & 22  \\\\ $\\log(\\dot{M})^{j}$ [$M_{\\odot}{\\rm yr}^{-1}$]  & all  & -8.08       &  -8.20      & 1.21        & 57    \\\\ \\hline \\multicolumn{6}{l}{}\\\\ \\multicolumn{6}{l}{NOTES: Only results from {\\it XMM-Newton} have been considered.}\\\\ \\multicolumn{6}{l}{For multiple detections, parameters were averaged}\\\\ \\multicolumn{6}{l}{$^a$ DEM: from DEM fit; 2-$T$: from 1-$T$ or 2-$T$ fit; all; all objects;}\\\\ \\multicolumn{6}{l}{ 1, 2, 3; object types; det = X-ray detected objects}\\\\ \\multicolumn{6}{l}{$^b$ Standard deviation of distribution}\\\\ \\multicolumn{6}{l}{$^c$ Only non-fixed fit parameters for detections considered}\\\\ \\multicolumn{6}{l}{$^d$ Only non-zero values for detections considered}\\\\ \\multicolumn{6}{l}{$^e$ Only detections considered}\\\\ \\multicolumn{6}{l}{$^f$ Only for known $L_*$ above the ZAMS}\\\\ \\multicolumn{6}{l}{$^g$ For multiples, value for primary if available}\\\\ \\multicolumn{6}{l}{$^h$ For multiples, logarithmic average of components if available}\\\\ \\multicolumn{6}{l}{$^i$ Upper limits are adopted as measured values}\\\\ \\multicolumn{6}{l}{$^j$ Average of reported range if at least one non-upper limit given}\\\\ \\end{tabular} \\end{table}  In Fig.~\\ref{lxm} we characterize the detected stellar population by plotting $L_{\\rm X}$ (from the DEM models) as a function of mass $M$  for all observed stars and substellar objects with known masses, although we exclude the following objects: DG Tau A (XEST-02-022), GV Tau (XEST-13-004), DP Tau (XEST-10-045), and CW Tau (XEST-20-046) have composite X-ray spectra that may originate from two unrelated sources \\citep{guedel06b}. L1551 IRS5 is a deeply embedded protostar, while the X-ray source seen close to it is only lightly absorbed and may be related to the jet \\citep{favata02, bally03}. Finally, the light curves of DH Tau (XEST-15-040), FS Tau (XEST-11-057), and V830 Tau (XEST-04-016) showed a decay presumably from large flares.  For objects that were observed and detected twice, the logarithmic average of $L_{\\rm X}$ has been adopted. Although for pre-main sequence stars there is no strict correlation between $L_{*}$ and $M$, it is interesting that we find a rather well-developed correlation between $L_{\\rm X}$ and $M$. The correlation coefficient for the $\\log L_{\\rm X}$ vs $\\log M$ sample is $C = 0.79$ for 99 data points. For the subsample of  type 2 (``CTTS'')  and type 3 objects (``WTTS''), we find $C = 0.75$ and $C = 0.85$ for 45 and 43  data points, respectively. All values point to a very significant correlation. We find a linear regression,   $\\log L_{\\rm X} = ( 1.54\\pm 0.12) \\log M +  30.31 \\pm 0.06$. The slopes for type 2 and type 3 objects are, respectively, $1.52\\pm 0.21$ and $1.78\\pm 0.17$. While these regressions are not significantly different, we note the larger errors for type 2 objects, and also their lower correlation coefficient, which is due to a larger scatter of $L_{\\rm X}$ at a given mass. Similar trends have been noted  in Orion \\citep{preibisch05}, with $\\log L_{\\rm X} = (1.44\\pm 0.10) \\log M + 30.37 \\pm 0.06$ for the entire sample. This coincidence between XEST and COUP reveals that in TMC and ONC, {\\it the basic X-ray production mechanism scales precisely the same way with the most basic property of the stellar objects, namely their mass}, regardless of any differences in age distribution, star-formation mode, or the presence or {\\nolinebreak (near-)\\linebreak} absence of  high-mass stars. Part of this correlation might be explained by higher-mass stars being larger, i.e., providing more surface area for coronal active regions. Assuming similar average internal stellar density, the correlation between surface area and $L_{\\rm X}$ is, however, considerably weaker  than the trend shown in Fig.~\\ref{lxm}, so that additional effects may play a role (convection zone depth, convective mass, etc).   \\begin{figure}[t!] \\centerline{\\resizebox{1.05\\hsize}{!}{\\includegraphics{guedel_f12.ps}}} \\caption{X-ray luminosity $L_{\\rm X}$ (based on DEM fits) vs. stellar mass $M$ for all detected XEST sources ({\\it Chandra} sources and Herbig stars are not included), or upper limits (arrows). For multiple systems, the primary mass has been used if available. Symbols defining the object type are given in the lower right corner. The horizontal bars for some objects show the ranges of masses derived from literature $L_*$ and $T_{\\rm eff}$, while the circle is centered at the adopted $M$. The straight line gives a linear regression using the logarithmic values for the X-ray detections, and the 1$\\sigma$ ranges for the slopes: $\\log L_{\\rm X} = (1.54\\pm 0.12) \\log M +30.31 \\pm 0.06$.}\\label{lxm} \\end{figure}  The figure shows another important feature of our survey. Almost all sources plotted are at $L_{\\rm X}$ considerably higher than our approximate detection limit of  $\\approx 10^{28}$~erg~s$^{-1}$. There are hardly any objects in the range of $(1-5)\\times 10^{28}$~erg~s$^{-1}$ although they would have been detected. Exceptions are the brown dwarfs. This again testifies to the near-completeness of our survey within the observed fields.   Fig.~\\ref{lxlb} shows the distribution of the  $L_{\\rm X}$ (derived from the DEM models) as a function of the (stellar, photospheric) bolometric luminosity $L_{*}$ for all spectrally modeled TTS and protostars,  and also including  BDs. We again exclude the peculiar spectral sources with two absorbers, the protostar L1551 IRS5, the three stars with gradually decaying light curves (DH Tau = XEST-15-040, FS Tau = XEST-11-057,  V830 Tau = XEST-04-016), and also all sources located below the ZAMS in the HRD. For stars observed and detected twice, we again plot the logarithmic average of $L_{\\rm X}$. Most stars cluster between $L_{\\rm X}/L_{*} = 10^{-4} - 10^{-3}$ as is often found in star-forming regions (see \\citealt{preibisch05} for a similar analysis for  the Orion sample). The value $L_{\\rm X}/L_{*} = 10^{-3}$ corresponds to the saturation value for rapidly rotating main-sequence stars. One key parameter that drives the X-ray output is obviously $L_{*}$.  \\begin{figure}[t!] \\centerline{\\resizebox{1.05\\hsize}{!}{\\includegraphics{guedel_f13.ps}}} \\caption{X-ray luminosity $L_{\\rm X}$ (based on DEM fits) vs. stellar luminosity $L_*$ for all X-ray detected XEST  sources ({\\it Chandra} sources, Herbig stars, and objects ``below'' the ZAMS not included), or upper limits (arrows). For multiple systems, total (``system'')  luminosity $L_*$ has been used. Symbols defining the object type are given in the lower right corner. The horizontal bars for some objects show the ranges of literature values for $L_*$, while the circle marks the adopted $L_*$. }\\label{lxlb} \\end{figure}   We finally present the X-ray luminosity function (XLF) of our Taurus sample in Fig.~\\ref{xlf}. The XLF has been calculated using the Kaplan-Meier estimator in  the ASURV software package \\citep{lavalley92} that considers upper limits to $L_{\\rm X}$ also. Non-detections for which no reasonable upper limit can be given were not included. We also dropped the same non-standard X-ray sources excluded above (double-absorber spectra, L1551 IRS 5, and sources with gradual flare decays), but include BDs and Herbig stars, thus using 136 X-ray luminosity values, including 17 upper limits. We also plot separately the XLF for type 2 (CTTS, 54 values, including 6 upper limits) and type 3 (WTTS, 49 values, 1 upper limit). We note the shift of the CTTS distribution toward lower luminosities compared to WTTS, by  a factor of about two. This is also evident for $L_{\\rm X}/L_{\\rm bol}$  \\citep{telleschi06a}. A two-sample test based on the Wilcoxon test  and the logrank test performed in ASURV indicates a probability of only 6-9\\% that the two distributions are drawn from the same parent population. A more detailed discussion on distinctions between CTTS and WTTS is given by \\citet{telleschi06a}.   \\begin{figure}[t!] \\centerline{\\resizebox{1.05\\hsize}{!}{\\includegraphics{guedel_f14.ps}}} \\caption{X-ray luminosity function for the entire sample (black), type 2 objects (red), and type 3 objects (blue).  }\\label{xlf} \\end{figure}"
    },
    "0609/astro-ph0609483_arXiv.txt": {
        "abstract": "Recent observations of the Galactic center in high-energy $\\gamma$-rays (above 0.1~TeV) have opened up new ways to study this region, from understanding the emission source of these high-energy photons to constraining the environment in which they are formed. We present a revised theoretical density model of the inner 5~pc surrounding Sgr A* based on the fact that the underlying structure of this region is dominated by the winds from the Wolf-Rayet stars orbiting Sgr A*.  An ideal probe and application of this density structure is this high energy $\\gamma$-ray emission.  We assume a proton-scattering model for the production of these $\\gamma$-rays and then determine first whether such a model is consistent with the observations and second whether we can use these observations to further constrain the density distribution in the Galactic center. ",
        "introduction": "The central 5--10~pc region surrounding Sgr A* is filled with dense molecular clouds and streamers with a range of characteristic sizes and densities; these include the circumnuclear disk (CND), the Northern Ridge, the Western and Southern Streamers, and the 20~km~s$^{-1}$ and 50~km~s$^{-1}$ giant molecular clouds (GMC) and the Molecular Ridge between them (Herrnstein \\& Ho 2004, 2005).  Some features have densities above $10^5-10^6\\ {\\rm cm^{-3}}$ while the density of other features can be as low as $10^3\\ {\\rm cm^{-3}}$.  The CND is a ring of molecular material with a well-defined inner edge at $1.5$~pc and an outer edge around $3$--$4$~pc from Sgr A*; it has a thickness of $\\sim 0.4$~pc at the inner edge and expands to a thickness of $\\sim 2$~pc at larger radii (Wright et al. 2001, Vollmer \\& Duschl 2001, Christopher et al. 2005).  Observations of molecular tracers (HCN and HCO$^{+}$) in gas cores in the CND place the typical densities of these cores at $10^7$--$10^8$~cm$^{-3}$ and the total estimated mass of the torus at $10^6$~M$_\\odot$ (Christopher et al. 2005).  The volume filling factor of dense clumps in the torus is low---approximately 1\\% (Vollmer \\& Duschl 2001).  The torus is evidently orbiting the black hole; observations by Jackson et al. (1993) show a clear pattern of radial velocities indicating that the torus is inclined by $60$--$75$ degrees from the line of sight and is orbiting around Sgr A*.  The CND and the 50~km~s$^{-1}$ GMC both contain OH (1720 MHz) masers, which suggest the presence of ongoing interactions between these two molecular features and the supernova remnant Sgr A East (Yusef-Zadeh et al. 2001).  Sgr A East is a non-thermal radio shell source centered $\\sim 2.5$~pc to the east of Sgr A*.  It is approximately 8~pc long (aligned roughly with the Galactic plane) and 6~pc wide (perpendicular to the Galactic plane).  Several studies (including Mezger et al. 1989, Maeda et al. 2001, Herrnstein \\& Ho 2005) have established that Sgr A East is interacting with other components in the Galactic center, including ionized gas streamers within the central cavity of the CND (Sgr A West) and the molecular clouds and streamers surrounding the remnant.  Mezger et al. (1989) detected a dust ring surrounding Sgr A East and suggested that this material had been swept up by an explosive event that produced Sgr A East, and both theorists and observers since then have suggested several different scenarios for the formation of this structure.    The trend and mantra among theorists are to simplify, which can often lead to incorrect conclusions about what is truly happening in the Galactic Center (GC).  For example, Mezger et al. (1989) noted that the total mass of material in the dust ring around Sgr A East ($6 \\times 10^4$~M$_\\odot$) would be comparable to the total mass contained within the volume of Sgr A East if the density there were $10^4$~cm$^{-3}$.  Forming the observed dust shell and synchrotron source via a supernova or other explosive event in such a dense medium would require a very high energy ($\\sim 10^{52}$~erg) and considerable time ($\\sim 10,000$~y or more, depending on details of the formation scenario).  Such high energies in stellar explosions have been observed (so-called hypernovae - Nomoto et al. 2004), but these events are rare ($10^{-5}$--$10^{-6}\\ {\\rm y^{-1}}$) in any conditions; and there is strong theoretical (e.g.  Heger et al.  2003) and observational (e.g.  Priddey et al. 2006) evidence showing that such events are even more rare, possibly non-existent, in stars with solar (or above) metallicities.  It is safe to conclude that the scenario proposed by Mezger et al. (1989) is incorrect, but what factors were missing in their analysis?  The critical missing piece is the presence and influence of massive, evolved stars in the central cluster surrounding Sgr A*.  The winds from these stars almost certainly pre-date the explosion that produced Sgr A East, so the explosion occurred not inside a molecular cloud or other dense environment, but inside the wind bubble associated with the central cluster (Rockefeller et al.  2005; Fryer et al.  2006a).  By including these effects, not only do the inferred age and energy of the supernova explosion decrease to less than 2000~y and approximately $10^{51}$~erg, respectively, but an additional feature in the X-ray observations can be explained (Rockefeller et al.  2005).  This is just one example showing why obtaining an accurate description of the density profile in the GC is critical to understanding this region.  In this paper, we describe the basic physical features of a wind-blown bubble density structure in the inner 5-10~pc of the GC. If a constant density profile is the ``level 0'' approximation for this region, our representation is only one level higher.  However, because we know the combined mass loss from the stars near Sgr A*, this model has no more free parameters than the constant density model.  Indeed, we argue that the only free parameter is the density into which the winds are blowing.  This density determines both the density profile and the extent of the wind blown bubble region.  With this basic density structure determined by very few observational inputs, we can then add many of the dense molecular cores that are observed in molecular maps of the CND and the GC region.  The physical assumption we make is that these clumps remain intact as the massive star winds blow past them.  With our basic model, we study a potential alternative probe of the density structure in the GC region: high-energy emission from cosmic rays. TeV $\\gamma$-rays from a point source $5^{\\prime\\prime}\\pm 10_{stat}^{\\prime\\prime}\\pm20_{sys}^{\\prime\\prime}$ from Sgr A* (HESS J1745-290) have been observed by a host of observatories: CANGAROO (Tsuchiya et al. 2004), Whipple (Kosack et al. 2004), HESS (Aharonian et al. 2004; Hinton et al. 2006) and MAGIC (Albert et al. 2006).  Both the flux and the spectral index vary from these three observations, with spectral indexes ranging from $\\alpha=2.21 \\pm 0.09$ (Aharonian et al. 2004) to $\\alpha=4.6 \\pm 0.5$ (Tsuchiya et al.  2004) and fluxes ranging from 5\\% to 40\\% that of the Crab Nebula (note these comparisons are not direct, they correspond to different energy bands).  In this paper, we adopt as our standard the HESS measurements: an above-165-GeV flux of $(1.82 \\pm 0.22) \\times 10^{-7} {\\rm ~ m^{-2} ~ s^{-1}}$ with a spectral index of $\\alpha=2.29 \\pm 0.05_{stat}\\pm0.15_{sys}$, which are also confirmed by MAGIC observations.  A number of models have been proposed to explain the emission with sources originating from the entire diffuse 10~pc region of the HESS PSF to Sgr A* itself (see Aharonian \\& Neronov 2004 for a review). The recent discovery of TeV emission from the GC ridge (Aharonian et al. 2006) favors a cosmic ray origin, and the fact that its spectral index is identical to that of HESS J1745-290 suggests that they may share the same cosmic ray source(s).  Aharonian \\& Neronov (2004), through shocks, and Liu et al.(2006), through stochastic acceleration, have argued that protons can be accelerated in the strong magnetic fields generated near the central supermassive black hole at the heart of Sgr A*.  High energy protons may also be produced in the supernova remnant shocks of Sgr A East. These protons will diffuse through the turbulent magnetic fields in the GC region.  Ultimately, these high-energy protons can scatter against protons in the wind-blown bubble region, producing pions that then decay into the gamma-rays we observe.  With our range of density profiles, we study this production mechanism for gamma-rays.  This paper includes two main thrusts.  We first develop a simple (but more sophisticated than past work) theorist's model for the density structure in the inner 5~pc around Sgr A*, based on a wind-blown bubble produced by the Wolf-Rayet stars orbiting Sgr A*.  The wind-blown bubble must exist at some level in the GC and we describe the physics behind this bubble in \\S\\ 2.  We also show the versatility of such a basic structure by adding the dense structure of the circumnuclear disk.  With this density profile, we then move to the second thrust of the paper, an attempt to understand the high energy emission arising from the region in the center of the Galaxy and to use these observations to better understand the density profile.  We assume that the high energy emission is produced by proton-proton interactions, one of the options suggested by Aharonian \\& Neronov (2005).  In \\S\\ 3, we discuss the possible sources for these protons, placing some initial constraints on this source (and the density) based on varying assumptions about the observations.  Using our model density profile, we then calculate the proton scattering through this medium and present 2-dimensional spatial distributions of the high-energy emission in \\S\\ 4.  For the most part, the spatial distribution is not severely constrained by current observations, but future instruments could place strong constraints on the density profile.  In \\S\\ 5, we study the spectra produced by our mechanism and show how measurements of the spectral index can also constrain the spatial distribution.  We conclude with a summary of our results and a discussion of the future potential of such detailed calculations. ",
        "conclusions": "A wind-blown bubble density structure provides an ideal first-order approximation of the $\\sim 5$~pc region surrounding Sgr A*.  This structure is not only a prediction of theory, but it also is supported by several observations.  It allows a more reasonable estimate of the supernova explosion energy and explains the X-ray ridge observed near Sgr A*.  With a single parameter (the density of the medium beyond the bubble), we can predict the size of the bubble.  Bear in mind, however, that this model is still only a first-order approximation; many additional molecular structures complicate this region.  We apply this simple density structure to a specific model for the source of high-energy $\\gamma$-rays near the GC: scattering of high-energy protons.  Given the uncertainties in the source (both its energy and its spatial distribution), the opacities (especially the effects of magnetic fields), and the densities in the GC region, it is not surprising that we can construct multiple solutions that match the current observations.  But our calculation does allow us to determine which uncertainties must be better understood to constrain the models further with the existing data.  In addition, we can determine which future observations will most help us constrain the models.  We have shown that detailed spatial maps can help us study the density structure in the GC region, but if the source of protons is indeed Sgr A*, the resolution required is far beyond what we can do with current telescopes.  What the observations can constrain is the source itself.  If Sgr A East were the source, the emission would be quite extended.  Figure 17 shows the normalized flux as a function of radius.  For a source near Sgr A*, the gamma-ray emission really is point-like, within the inner 2~pc.  But a Sgr A East supernova remnant source argues for an extended source of gamma-ray emission (out to 5~pc) that increases out to 5~pc and then drops precipitously.  The innermost two pixels in the current data correspond roughly to 7\\,pc and our emission profile from Sgr A East is just barely consistent with these two pixels.  The strong drop in emission at 5\\,pc would mean that a pixel centered at 7\\,pc should be a factor of 2 lower than the peak.  As data improves, this emission source might be easily ruled out.  But remember that Sgr A East is not a 5\\,pc spherical shell. If the emission is dominated by the shorter axis of Sgr A East, this emission would center around 3~pc instead of 5~pc.  The centroid of the emission also places constraints on a Sgr A East source.  Sgr A East is offset by 2~pc from Sgr A*.  If the centroid truly is exactly on Sgr A *, and the density profile of Sgr A East truly is offset (not just the emission), we can also rule out Sgr A East.  Note that for a Sgr A* source, even if the termination shock of the wind-blown bubble is offset by 2\\,pc, we found that the centroid of the high energy emission does not move by more than 0.1\\,pc.  Finally the question arises whether the extended gamma-ray emission can be used as a constraint on the source.  If we want to claim that the extended emission is produced by the same proton source as the point-like source, we can use the observations to constrain the proton source.\\footnote{Note, however, that the only reason to assume this is because the spectra are similar---but if we use shock acceleration to produce our high-energy protons, we would expect any source to produce roughly the same spectrum.}  For the Sgr A East source, the primary constraint arises because, given the new estimate for the age of the remnant by Rockefeller et al. (2006), Sgr A East has been producing high-energy protons for at most 1000~y.  For those protons to emit gamma-rays 100~pc away, there must be a ``hole'' where the density is low and the scale-length of the magnetic field is high.  The proton travel time is this distance (100\\,pc) times the effective optical depth of the protons divided by the speed of light.  To make this timescale fall below 1000\\,y, this roughly requires an optical depth of less than 2-3, implying that the scale-lenght of the magnetic field is greater $\\sim$ 50~pc.  This timescale constraint is not important for a continuous source like Sgr A*.  However, the fact that the spectral index is nearly identical suggests that either the opacities---in particular, the magnetic field effective opacity---must be energy independent or that a similar hole exists.  We can not absolutely rule out a Sgr A East source for the high energy emission, but taking all of the above constraints together, we have to introduce a number of caveats to make a Sgr A East source match the existing data.  With accurate density maps and a precise localization of the high-energy source, it is likely that this source can be ruled out in a more absolute sense.  Finally, the spectral distribution of the gamma-ray emission depends on all of our uncertainties: source, opacities, and density distribution in the Galactic center.  However, the spectral distribution as a function of space is an ideal probe of our opacities, in particular the distribution of the magnetic fields (compare the plots in Figure 16).  {\\bf Acknowledgments} This work was carried out under the auspices of the National Nuclear Security Administration of the U.S. Department of Energy at Los Alamos National Laboratory under Contract No. DE-AC52-06NA25396."
    },
    "0609/astro-ph0609356_arXiv.txt": {
        "abstract": "We present results of a high-resolution, near-infrared survey of 41 nearby, young ($\\lesssim\\!300\\mathrm{\\,Myr}$) M0--M5.0 dwarfs using the Altair natural guide star adaptive optics system at the Gemini North telescope. Twelve of the objects appear to be binaries, 7 of which are reported here for the first time. One triple system was discovered. Statistical properties are studied and compared with earlier (F to K) and later ($\\ge$ M6 very low-mass, VLM) populations. We find that the separation distribution of the binaries in this sample peaks at $13^{+14}_{-9}\\mathrm{\\,AU}$, which is consistent with previous measurements of early-M binaries. Hence, early-M binaries seem to occur in---on average---tighter systems than G binaries. At the same time they are significantly wider than field VLM binary stars. The distribution of mass ratios $q$ of primary and secondary stars was found to show an intermediate distribution between the strongly $q\\rightarrow1$ peaked distribution of field VLM systems and the almost flat distribution of earlier-type stars. Consequently, we show evidence for relatively young, early-M binaries representing a transition between the well known earlier star distributions and the recently examined field VLM population characteristics. Despite the fact that this survey was dedicated to the search for faint brown dwarf and planetary mass companions, all planetary mass candidates were background objects. We exclude the existence of physical companions with masses greater than 10 Jupiter masses ($\\mathrm{M_{Jup}}$) at separations of $\\gtrsim 40\\mathrm{\\,AU}$ and masses greater than $24\\mathrm{\\,M_{Jup}}$ for separations $\\gtrsim 10\\mathrm{\\,AU}$ around 37 of the 41 observed objects. ",
        "introduction": "\\label{obs} The current paper presents observations of a subset of the X-ray active stars from the NLTT/CNS3 and Moving M samples. All targeted objects are estimated to lie within 20 parsecs of the Sun.  All observations were taken in queue mode with the \\mbox{8-m} Gemini North telescope using the facility Near-Infrared Imager (NIRI; Hodapp et al. 2003) in conjunction with the Altair AO system. The detector is a 1024$\\times$1024 InSb Aladdin infrared hybrid array with 27$\\micron$ pixels. The f/32 camera was used providing a $0\\farcs022$/pixel platescale and 22$\\arcsec\\times22\\arcsec$ field of view. The observations were conducted with the Cassegrain rotator purposefully disabled, leading to a fixed image pupil relative to camera and detector.  For all observed targets, both short ($<\\!10\\mathrm{\\,s}$) and deep (30\\,s) exposures in the H and K$_\\mathrm{s}$ filter bands were obtained. Each of the targets were imaged in a 5-position, $7\\arcsec$-offset dithering pattern; 3 images were taken per dither position. The short frames were designed to primarily obtain unsaturated target images so as to reveal any near-equal mass companions at separations $\\gtrsim\\!0\\farcs08$. They also served to provide astrometry and PSF sources. In a few cases even the minimum integration time of $0.2\\,\\mathrm{s}$ allowed by the NIRI IR camera led to saturated short frames due to the brightness of the target star. In these cases, other unsaturated sources were used for PSF fitting. All of the discovered physical binary systems were detected in the short frames.  The deep frames were designed to saturate the target star to gain sensitivity to potential very faint companions ($\\Delta$H $>5\\,$mag) at radial distances $\\gtrsim\\!1\\arcsec$. Integrated counts were typically $\\sim\\!\\,200$--$500\\%$ in excess of full well depth at both H and K$_\\mathrm{s}$ in the primary's core. Through the use of coadds, 30\\,s integration times were obtained per image; 3 images per dither position gave 7.5\\,min of total integration time per target.  Forty-one of the 56 objects from our survey sample were observed with the Gemini North telescope during semesters 2005A, 2005B, and 2006A. Table \\ref{tab1} lists the 23 objects observed with no likely physical companions according to our sensitivity limits (examination of the detection limits is given in \\S\\ref{sen}). Four of the observed objects showed faint companions of currently undetermined nature. Follow-up observations are necessary to reveal whether the detections are consistent with proper motion objects and therefore physical companions. One additional object is believed to be a tight equal-mass binary system with an estimated separation comparable to the FWHM of $0\\farcs08$ of these observations. Poor image quality probably due to bad weather, however, makes follow-up observations necessary to confirm its nature. None of the other binary images suffered from similar aberrations. Table \\ref{tab2} shows the observable properties of the 12 binary systems observed, 7 of which are newly discovered binaries, as well as the parameters of the newly discovered triple system.  \\begin{deluxetable}{lllcc} \\tabletypesize{\\scriptsize} \\tablecaption{Young M0.5--M5.0 Stars Observed with No Physical Companion Detections Within the Sensitivity Limits\\tablenotemark{a}\\label{tab1}} \\tablewidth{\\columnwidth} \\tablehead{ \\colhead{2MASS Name} & \\colhead{Other Name} & \\colhead{K$_s$} & \\colhead{SpT} & \\colhead{Ref.} } \\startdata J00115302+2259047 & LP\\,348-40 & 8.00 & M3.5 & 1 \\\\ J01351393--0712517 &  & 8.08 & M4.0 & 2 \\\\ J01531133--2105433 & LP\\,828-89 & 7.14 & M1.0 & 1 \\\\ J02085359+4926565 & G\\,173-39 & 7.58 & M3.5 & 3 \\\\ J03472333--0158195 & HIP\\,17695 & 6.93 & M3.0 & 2 \\\\ J03494325+2419046 & II\\,Tau & 8.94 & M4.0 & 1 \\\\ J04244260--0647313 &  & 8.63 & M4.5 & 2 \\\\ J04373746--0229282 & Steph\\,497 & 6.41 & M0.5 & 1 \\\\ J08081359+2106094 & LHS 5134 & 6.52 & M2.5 & 1 \\\\ J10193634+1952122 & Gl\\,388 & 4.59 & M3.0 & 3 \\\\ J11032125+1337571 & LP\\,491-51 & 7.91 & M3.5 & 1 \\\\ J11115176+3332111 & G\\,119-62 & 7.50 & M3.5 & 3 \\\\ J14130492--1201262 & Gl\\,540.2 & 8.16 & M4.5 & 3 \\\\ J14200478+3903014 & Steph\\,1145 & 7.71 & M2.5 & 1 \\\\ J15215291+2058394 & Wo\\,9520 & 5.76 & M1.5 & 1 \\\\ J15530484+4457446 &  & 9.25 & M3.5 & 4 \\\\ J18021660+6415445 & LP\\,71-82 & 7.65 & M4.5 & 5 \\\\ J18130657+2601519 & LP\\,390-16 & 8.07 & M4.0 & 1 \\\\ J19311257+3607300 & G\\,125-15 & 8.83 & M4.5 & 1 \\\\ J20464360--1148132 & LP\\,756-3 & 8.44 & M4.0 & 4 \\\\ J22143835--2141535 & BD-22\\,5866 & 6.72 & M0.5 & 1 \\\\ J22464980+4420030 & Gl\\,873 & 5.30 & M3.5 & 3 \\\\ J22515345+3145153 & Gl\\,875.1 & 6.87 & M3.0 & 3 \\\\[-1.5ex] \\enddata \\tablenotetext{a}{No companions with mass $>$\\,$10\\mathrm{\\,M_{Jup}}$ in more than $40\\mathrm{\\,AU}$ separation or $>$\\,$24\\mathrm{\\,M_{Jup}}$ in more than $10\\mathrm{\\,AU}$ separation. Sensitivity limits are examined in section \\S\\ref{sen} and Figure \\ref{fig8}.} \\tablerefs{(1) Reid et al. 2004; (2) Reid, Cruz, \\& Allen (in prep.); (3) Reid et al. 1995a; (4) Riaz, Gizis \\& Harvin 2006; (5) Reid et al. 2003} \\end{deluxetable}  \\begin{deluxetable*}{llcr@{$\\pm$}lr@{$\\pm$}lr@{$\\pm$}lr@{$\\pm$}lcc} \\tabletypesize{\\scriptsize} \\tablecaption{The Multiple Systems \\label{tab2}} \\tablewidth{-479.16212pt} \\tablehead{ \\colhead{2MASS Name} & \\colhead{Other Name} & \\colhead{SpT} & \\multicolumn{2}{c}{$\\Delta H$} & \\multicolumn{2}{c}{$\\Delta K_s$} & \\multicolumn{2}{c}{Sep. (arcsec)} & \\multicolumn{2}{c}{P.A. (deg)} & \\colhead{Date (UT)} & \\colhead{Ref.} } \\startdata J00285391+5022330 & G\\,172-1 & M4.0\\tablenotemark{1} & $0.50$ & $0.03$ & $0.49$ & $0.02$ & $0.426$ & $0.002$ & $32.13$ & $0.07$ & 05\\,Aug\\,05 &  TP \\\\ J01242767--3355086  & G\\,274-24 & M5.0\\tablenotemark{1} & $0.20$ & $0.03$ & $0.05$ & $0.02$ & $2.065$ & $0.002$ & $ 44.80$ & $0.07$ & 24\\,Aug\\,05 & TP \\\\ &  &  & \\multicolumn{2}{c}{ } & \\multicolumn{2}{c}{ } & $2.01$ & $0.07$ & $46.38$ & $1.31$ & Jul\\,99 & 5 \\\\ J04381255+2813001 & G\\,39-29 & M4.0\\tablenotemark{1} & $0.39$ & $0.03$ & $0.37$ & $0.03$ & $0.783$ & $0.002$ & $300.58$ & $0.09$ & 15\\,Oct\\,05 & TP \\\\ &  &  & \\multicolumn{2}{c}{ } & \\multicolumn{2}{c}{ } & \\multicolumn{2}{c}{$0.54$} & \\multicolumn{2}{c}{$299$} & 13\\,Sep\\,02 &  6 \\\\ J05254166--0909123 & LP\\,717-36 & M3.5\\tablenotemark{2} & $0.46$ & $0.02$ & $0.42$ & $0.03$ & $0.537$ & $0.002$ & $69.40$ & $0.11$ & 14\\,Oct\\,05 & TP \\\\ J07285137--3014490 & GJ\\,2060 & M0.5\\tablenotemark{2} & $0.44$ & $0.30$\\tablenotemark{a} & $0.44$ & $0.42$\\tablenotemark{a} & $0.175$ & $0.011$ & $143.71$ & $1.54$ & 30\\,Oct\\,02 &  TP \\\\ &  &  & \\multicolumn{2}{c}{ } & \\multicolumn{2}{c}{ } & \\multicolumn{2}{c}{$0.3 ?$} & \\multicolumn{2}{c}{?} & ? & 7 \\\\ J10364483+1521394 &  & M4.0\\tablenotemark{3} & $1.15$ & $0.02$\\tablenotemark{b} & $1.12$ & $0.01$\\tablenotemark{b} & $1.061$ & $0.002$\\tablenotemark{b} & $181.22$ & $0.07$\\tablenotemark{b} & 18\\,May\\,06 &  TP \\\\ &  &  & $0.05$ & $0.02$\\tablenotemark{c} & $0.03$ & $0.02$\\tablenotemark{c} & $0.189$ & $0.002$\\tablenotemark{c} & $310.62$ & $0.14$\\tablenotemark{c} &  &  \\\\ J13271966--3110394  & LHS\\,2739 & M3.5\\tablenotemark{1} & $0.14$ & $0.04$ & $0.17$ & $0.05$ & $0.544$ & $0.004$ & $333.79$ & $0.09$ & 12\\,Jun\\,05 &  TP \\\\ J13345147+3746195 &  & M3.5\\tablenotemark{4} & $0.38$ & $0.03$ & $0.25$ & $0.03$ & $0.082$ & $0.003$ & $197.90$ & $0.27$ & 26\\,Mar\\,06 &  TP \\\\ J15553178+3512028 & G\\,180-11 & M4.5\\tablenotemark{1} & $1.99$ & $0.02$ & $1.97$ & $0.01$ & $1.571$ & $0.002$ & $261.23$ & $0.07$ & 27\\,May\\,05 &  TP \\\\ &  &  & \\multicolumn{2}{c}{ } & \\multicolumn{2}{c}{ } & \\multicolumn{2}{c}{$1.5$} & \\multicolumn{2}{c}{$266$} & Mar\\,98 & 8 \\\\ J17035283+3211456  & LP\\,331-57 & M2.5\\tablenotemark{1} & $1.53$ & $0.06$ & $1.53$ & $0.04$ & $1.130$ & $0.002$ & $136.53$ & $0.11$ & 22\\,Apr\\,05 & TP \\\\ J22332264--0936537  & Steph\\,2018 & M3.0\\tablenotemark{2} & $0.09$ & $0.04$ & $0.03$ & $0.03$ & $1.571$ & $0.003$ & $279.73$ & $0.14$ & 09\\,Jun\\,05 & TP \\\\ &  &  & \\multicolumn{2}{c}{ } & \\multicolumn{2}{c}{ } & \\multicolumn{2}{c}{$1.66$} & \\multicolumn{2}{c}{$272.25$} & Jul\\,97 & 8 \\\\ J23205766--0147373 & LP\\,642-48 & M4.0\\tablenotemark{4} & $0.32$ & $0.02$ & $0.27$ & $0.03$ & $0.099$ & $0.002$ & $325.06$ & $0.09$ & 12\\,Jun\\,05 & TP \\\\ J23581366--1724338 & LP\\,764-40 & M2.0\\tablenotemark{4} & $0.02$ & $0.02$ & $0.02$  & $0.02$ & $1.904$ & $0.002$ & $265.30$ & $0.88$ & 16\\,Jul\\,05 & TP\\\\[-1.5ex] \\enddata% \\tablenotetext{a}{$\\mathrm{\\Delta H}$ was estimated from artificial ghost, $\\mathrm{\\Delta K_s} = \\mathrm{\\Delta H}$ was assumed (see \\S\\ref{red} for more detail).} \\tablenotetext{b}{Numbers refer to the system consisting of A and B component of the triple system 2MASS\\,J10364483+1521394.} \\tablenotetext{c}{Numbers refer to the system consisting of B and C component of 2MASS\\,J10364483+1521394.} \\tablerefs{(TP) This paper; (1) Reid et al. 1995a; (2) Reid et al. 2003; (3) Reid, Cruz, \\& Allen (in prep.); (4) Riaz, Gizis \\& Harvin 2006; (5) Jao et al. 2003; (6) Beuzit et al. 2004; (7) Zuckerman et al. 2004; (8) McCarthy et al. 2001} \\end{deluxetable*} ",
        "conclusions": "\\label{dis} Observations of binaries and their inherent parameters are important for setting constraints on the formation mechanism of stars in general. This sample of young, early-M dwarfs allows insight into a population of stars in between the well-known G and K star population and the distinctly different VLM population \\citep{bur06}, which became analyzable only a few years ago with the development and employment of highly sensitive AO and the {\\it HST} (Close et al. 2003; Gizis et al. 2003; Bouy et al. 2003).  \\subsection{The Multiple Star Fraction} We found 12 binaries and 1 triple system in our sample of 41 objects. In order to derive a multiple star fraction free of Malmquist bias, we consider only the systems in our sample whose distances are within 20\\,pc. Eight of the observed multiples are likely within 20\\,pc of the Sun, as are--estimatedly--all of the observed single stars. However, the distances of two of the 8 objects (G\\,172-1, 2M1036) are within $1\\sigma$ of the 20\\,pc limit, leaving a rather large probability that they are not part of the defined volume. Likewise, two binaries have estimated distances $<1\\sigma$ greater than 20\\,pc (LP\\,717-36 and LP\\,642-48), so they may be part of the limited sample. This therefore leaves $8^{+2}_{-2}$ confirmed binaries within 20\\,pc. To eliminate uncertainties caused by the distance dependence of the sensitivity, we conservatively cut the sample at a mass ratio of $q=0.1$. There are 28 objects in this sub-sample other than the confirmed binaries. Four of these have candidate companions with $q<0.1$. Hence, we include them in this binary fraction as single stars. In addition, we include one possible binary system in the upper limit estimate. This results in a total of $8^{+3}_{-2}$ multiples in a sample of $35^{+2}_{-2}$ objects (upper and lower limits correspond to each other, respectively), resulting in a nominal binary fraction of $23^{+15}_{-9}\\,\\%$ for $q>0.1$ and $\\mathrm{sep.}>1.6$\\,AU (we use binomial statistics to estimate the uncertainty as described in \\citet{bur03}). The cut in separation is determined by the minimum resolution of $\\mathrm{FWHM}=0\\farcs08$ at the distance limit of 20\\,pc. We note, however, that this above examination of the binary fraction does not account for selection effects in the sample selection.  \\subsection{The Mass Ratio Distribution}\\label{massratio} Figure \\ref{fig11} shows the mass ratio $q=M_B/M_A$ distribution derived from the 13 observed multiples using the masses listed in Table \\ref{tab3}. The triple system 2M1036 is accounted for as two binaries, the binary BC and the A(BC) system. Mass ratios for these two objects are therefore $q_\\mathrm{BC}=M_\\mathrm{C}/M_\\mathrm{B}$ and $q_\\mathrm{A(BC)}=(M_\\mathrm{B}+M_\\mathrm{C})/M_\\mathrm{A}$. In the following paragraphs we sometimes refer to 14 ``binaries'', but one should keep in mind that this means 12 binaries and 1 triple system.  Twelve out of the 14 have mass ratios of $q\\gtrsim0.74$; we regard this as evidence that equal-mass systems are preferred for young, early-M dwarf primaries. Our observations are sensitive to binaries with mass ratios $q<0.1$ for separations greater than $\\sim\\!8\\mathrm{\\,AU}$ (see Figure \\ref{fig12}). Observations have shown that there is a significant lack of very low-mass objects around main sequence stars for separations smaller than 5 AU (the ``brown dwarf desert''; Marcy \\& Butler 2000). We therefore do not expect to have missed a large quantity of objects. However, we have to assume that for $q<0.1$ the sample is likely incomplete.  \\begin{figure}[tb] \\includegraphics[angle=90,width=\\columnwidth]{f11.eps} \\caption{Mass ratio $q=M_B/M_A$ distribution. The continuous line shows the distribution derived from the 14 early-M binaries in this paper (triple 2M1036 included as two binaries, see \\S \\ref{massratio}). The dashed line shows the distribution of binaries included in an early F to K star sample by \\citet{rei02} that have separations $<\\!100\\mathrm{\\,AU}$. The dotted line shows the $q$ distribution for VLM stars from \\citet{bur06}. While the F to K binary distribution is described by a rather flat distribution, slightly more populated at higher $q$, the VLM distribution shows a significant peak for equal mass binaries. The early-M binary distribution with its broad peak at $1 \\ge q > 0.7 $ might represent an intermediate distribution between the earlier and later types. Due to the high sensitivity of our measurements, only the last, $q<0.1$ bin is significantly affected by incompleteness. \\label{fig11}} \\end{figure}  \\begin{figure}[tb] \\includegraphics[angle=90,width=\\columnwidth]{f12.eps} \\caption{Mass ratio $q=M_B/M_A$ as a function of projected binary separation in astronomical units for the 12 discovered binaries (filled circles). The triple 2M1036 is included as the two binaries A(BC) and BC (squares; for further explanation see \\S \\ref{massratio}). The dashed line represents the sensitivity limit of our survey, derived from the sensitivity relation shown in Figure \\ref{fig3} and 500\\,Myr models from \\citet{bar03}. As a demonstration of detection capabilities, the open circle at $\\mathrm{sep.}=53.3\\mathrm{\\,AU}$ and $q=0.035$ is the companion seen in Figure \\ref{fig10} if it were a physical companion to G\\,173-39. \\label{fig12}} \\end{figure}  This emphasis on $q\\sim1$ is not seen in the distributions of earlier type stars. \\citet{rei02} show that for late F to K binaries with separations of less than 100\\,AU the distribution of mass ratios is rather flat with only marginal tendency towards equal mass binaries. For field VLM binaries, however, the opposite effect is observed; recent statistical evaluations of all currently known VLM binaries \\citep{bur06} show a significant peak at $q\\!\\sim\\!1$. It is therefore possible that the mass ratio distribution of our early-M binaries sample has an intermediate distribution between the flat F-K distribution and the strongly peaked field VLM distribution.  However, our sample of binaries is rather small. Further observations of a larger number of young M stars will be needed to show if this relation holds.  \\subsection{Mass Ratio vs. Separation} Mass ratios $q=M_\\mathrm{B}/M_\\mathrm{A}$ as a function of binary separation are shown in Figure \\ref{fig12}. While at separations $\\lesssim 20$\\,AU systems are distributed over a broad range of mass ratios, all three binaries with separations $> 25$\\,AU show mass ratios of $q=1$. The sensitivity of the observations would have allowed for the detection of much lower mass ratios at these separations (see Figure \\ref{fig10} for example), indicating a real lack of objects with low mass ratios and high separations. However, a sample of 14 binaries cannot fill the accessible space in the $(q,\\mathrm{separation})$ plane uniformly, limiting the significance of this result due to low number statistics.  \\subsection{The Separation Distribution} Figure \\ref{fig13} shows the separation distribution of our set of 14 binaries. The physical projected separations range from 2.2 to 49.7\\,AU and possess a distribution that shows no significant peak. However, since the observations cover separations up to 15 arcseconds, which corresponds to 300\\,AU at a distance of 20\\,pc, we conclude that the decreasing number of detected objects for separations higher than $\\sim\\!50$\\,AU is a real effect. The underlying distribution therefore probably peaks within 0 to 50\\,AU. A least-$\\chi^2$ fit of a Poisson distribution to the data suggests a peak at $13^{+14}_{-9}\\mathrm{\\,AU}$. This is consistent with earlier results, for instance by \\citet{fis92}, who report a broad peak at $4$--$30\\mathrm{\\,AU}$ for M2--M4.5 binaries.  \\begin{figure}[tb] \\includegraphics[angle=90,width=\\columnwidth]{f13.eps} \\caption{Distribution of projected binary separations. The solid line shows the separation distribution based on our early-M sample consisting of 12 binaries and a triple system counted as two binaries A(BC) and BC. The dashed line describes the projected separation/semi major axis distribution of 70 VLM binaries \\citep{bur06}. The $\\mathrm{sep.}<5\\mathrm{\\,AU}$ bin is probably incomplete due to the finite resolution of $0\\farcs08$ corresponding to 1.6\\,AU in 20\\,pc. \\label{fig13}} \\end{figure}  The separation distribution, again, shows evidence that young, early-M binaries possess transitional characteristics between earlier and later type populations. With a distribution peaking at $\\sim\\!30\\mathrm{\\,AU}$ \\citep{duq91}, G stars seem to have---on average---greater separations than early-M stars. Field VLM binaries, however, are preferably found in significantly tighter systems. The distribution of 70 binaries with spectral types later than M5.5 shows a significant peak at $3$--$10\\mathrm{\\,AU}$ \\citep{bur06} with a steep quantitative decrease of systems with higher separations.  G--K binaries are described by a mass ratio distribution and a separation distribution that are very different from the VLM binaries' characteristics; young, early-M binaries appear to show an intermediate behavior. The higher and the lower mass distributions are apparently connected through a smooth transition embodied by early-M binaries."
    },
    "0609/astro-ph0609026_arXiv.txt": {
        "abstract": "{} {The main goal of the present study is to determine the fractional SiO abundance in high-mass star-forming cores, and to investigate its dependence on the physical conditions. In this way we wish to provide constraints on the chemistry models concerning the formation of SiO in the gas phase or via grain mantle evaporation. The work addresses also CH$_3$CCH chemistry as the kinetic temperature is determined using this molecule.} {We estimate the physical conditions of 15 high-mass star-forming cores and derive the fractional SiO and CH$_3$CCH abundances in them by using spectral line and dust continuum observations with the SEST.} {The kinetic temperatures as derived from CH$_3$CCH range from 25 to 39 K, the average being 33 K. The average gas density in the cores is $4.5 \\cdot 10^6$ cm$^{-3}$. The SiO emission regions are extended and typically half of the integrated line emission comes from the velocity range traced out by CH$_3$CCH emission. The upper limit of SiO abundance in this 'quiescent' gas component is $\\sim 10^{-10}$. The average CH$_3$CCH abundance is about $7 \\cdot 10^{-9}$. It shows a shallow, positive correlation with the temperature, whereas SiO shows the opposite tendency.} {We suggest that the high CH$_3$CCH abundance and its possible increase when the clouds get warmer is related to the intensified desorption of the chemical precursors of the molecule from grain surfaces. In contrast, the observed tendency of SiO does not support the idea that the evaporation of Si-containing species from the grain mantles would be important, and it contradicts with the models where neutral reactions with activation barriers dominate the SiO production. A possible explanation for the decrease is that warmer cores represent more evolved stages of core evolution with fewer high-velocity shocks and thus less efficient SiO replenishment.} ",
        "introduction": "High-mass star-forming regions with large molecular column densities and bright molecular lines offer excellent opportunities to study interstellar chemisty. Results from chemistry models can be combined with observational data on the physical conditions and molecular abundances. This information can help us to link the properties of massive dense cores, their chemical characteristics, and phenomena like massive outflows, ultracompact (UC) \\ion{H}{ii} regions and molecular masers into a coherent evolutionary sequence. This is useful since the process of high-mass star formation is not yet well understood (see, e.g., \\cite{garay1999}; \\cite{fontani2002}; \\cite{thompson2003}).  In this paper we derive the abundances of silicon monoxide, SiO, and methyl acetylene, CH$_3$CCH, in a sample of massive molecular cloud cores, and discuss the relation of these abundances to the physical conditions and the probable evolutionary stages of the cores. The purpose is to contribute towards a better understanding of the interplay between the processes associated with massive star formation and the physical and chemical properties of the surrounding dense core.  SiO is believed to trace exclusively shocked gas as its spectral lines have wings and are often shifted relative to the emission from the ambient gas (see, e.g., \\cite{martin1992}; \\cite{schilke1997}). Current understanding is that SiO is evaporated from the dust grains when the shock velocity is greater than about 20 km s$^{-1}$. This, however, depends on the grain-mantle composition in the preshock gas (e.g., \\cite{schilke1997}). In dense and warm giant molecular cloud (GMC) cores, SiO seems to present also in the quiescent gas component. In few interferometric maps available extended SiO emission at ambient cloud velocity is either seen in the vicinity of high-velocity outflows (\\cite{hatchell2001}) or totally separate from them (\\cite{lefloch1998}; \\cite{codella1999}; \\cite{shepherd2004}; \\cite{fuente2005}). The fractional abundance of 'quiescent SiO' is estimated to be about $10^{-10}$ in GMC cores. SiO absorption measurements towards dormant GMCs, so called 'spiral arm clouds', give similar abundances (\\cite{greaves1996}). The present observational evidence is insufficient to determine whether the presence of SiO is caused by thermal evaporation, enhanced neutral-neutral production pathways, photon-induced reactions, or by shock removal. To resolve this indistinctness it seems beneficial to determine fractional SiO abundances and kinetic temperatures in a representative sample of GMC cores, and to correlate these also with the kinematic information contained in the spectral lines.  CH$_3$CCH is an organic molecule observed rather frequently towards dense cores. Its abundance is on the order of $10^{-9}$. The rotational temperature, $T_{\\rm rot}$, derived from a series of $J_K-(J-1)_K$ rotational lines of CH$_3$CCH is considered a good estimate of the gas kinetic temperature, $T_{\\rm kin}$, in molecular clouds (e.g., \\cite{askne1984}; \\cite{bergin1994}). Since CH$_3$CCH has a relatively low dipole moment ($\\mu=0.78$ D, \\cite{bauer1979}; \\cite{burrell1980}), its rotational levels are thermalized at densities $n({\\rm H_2}) \\sim 10^4 \\, {\\rm cm}^{-3}$ or higher (e.g., \\cite{kuiper1984}). This property together with the assumption that CH$_3$CCH emission is optically thin ($\\tau \\ll 1$), allow us to assume local thermodynamic equilibrium (LTE), so that $T_{\\rm kin}$ may be derived. Besides for $T_{\\rm kin}$, we use $T_{\\rm rot}$ for estimating the dust temperature, $T_{\\rm d}$, which is needed to derive the total gas column densities ($\\sim N({\\rm H_2}$)) from dust continuum observations.  The source list consists of 15 high-mass star-forming cores associated with OH, H$_2$O, and CH$_3$OH masers, UC \\ion{H}{ii} regions and bright, thermal SiO rotational line emission. The sources were selected from the SiO survey of \\cite{harju1998} (hereafter HLBZ98). The positions selected for the SiO and CH$_3$CCH line observations are listed in Table \\ref{table:sources}. These positions served as map centres in the continuum observations.  The SiO and CH$_3$CCH column densities and the kinetic temperatures in the present paper are derived from spectral line observations with the SEST. The estimates for the molecular hydrogen column densities, $N({\\rm H_2})$, and the core masses are derived from dust continuum maps obtained with the SIMBA bolometer at SEST.  The observations and the data reduction procedures are described in Sect. 2. The direct observational results are presented in Sect. 3. In Sect. 4 we describe the methods used to derive the physical/chemical properties of the GMC cores. In Sect. 5 we discuss the results, and finally, in Sect. 6 we summarize our major conclusions. ",
        "conclusions": "Spectral line observations at 3 and 2 mm, and dust continuum observations at 1.2 mm performed with the SEST were used to derive physical characteristics and fractional SiO and CH$_3$CCH abundances in 15 high-mass star-forming cores. The sample represents typical GMC cores with the following average properties $<T_{\\rm kin}> = 33$ K, $<n(\\rm H_2)> = 4.5 \\cdot 10^6 \\; {\\rm cm}^{-3}$, and $<M> = 1800$ $ \\rm M_{\\sun}$. While most cores seem to be near the virial equilibrium, for five of them the gravitational potential energy is large compared with the kinetic energy.  The fractional SiO abundances determined from the velocity range with detectable CH$_3$CCH emission were found to be $\\sim 1-7 \\cdot 10^{-10}$. CH$_3$CCH abundances were found to be $3-18 \\cdot 10^{-9}$. In consistence with earlier results towards GMC cores (e.g., \\cite{ungerechts1997}; \\cite{bergin1997}) the CH$_3$CCH abundance is found to be clearly larger than in the dense cores of dark clouds. It also shows a slight increase as a function of $T_{\\rm kin}$. As this tendency does not agree with the predictions of gas-phase chemistry models (e.g. \\cite{lee1996}), a possible explanation is that warm conditions lead to an intensified desorption of the precursor molecules of CH$_3$CCH from the icy mantles of dust grains. This suggestion is supported by the chemistry model of \\cite{canosa1997}.  The SiO abundances are midway between the upper limit from dark clouds and PDRs ($\\sim 10^{-12}$) and the values derived for powerful shocks ($\\sim 10^{-8}$). The abundance seem to decrease with rising temperature. This finding contradicts with the models where SiO production is dominated by neutral reactions with activation energies (\\cite{ziurys1989}; \\cite{langer1990}), and with models where desorption of Si-containing species from grain mantles are significant for the SiO production (\\cite{mackay1996}). As suggested by the recent observational results of \\cite{fuente2005}, warmer cores represent more evolved objects where highly energetic protostellar outlows releasing SiO into the gas phase are less frequent; this together with the rapid post-shock processing have lead to a diminished SiO abundance."
    },
    "0609/astro-ph0609210_arXiv.txt": {
        "abstract": " ",
        "introduction": "The formation of open clusters has been well studied in the past two decades (see review by Lada \\& Lada 2003). In contrast, due to the lack of well-studied young clusters as massive as a typical globular cluster, little is known about the formation of globular clusters (Larson 1992). Although it has been suggested that globular clusters share some basic formation mechanism with open clusters due to the apparent continuity in properties (Larsen 1992), the formation of two types clusters with few orders of magnitude difference in total mass might be induced by significantly different physical phenomena. It has been suggested that the young clusters found by \\textit{Hubble Space Telescope} (\\textit{HST}) are protoglobular clusters formed by merger events (Holtzman et al. 1992; and the review by Whitmore 2003). However, how the star formation is triggered by mergers and what triggers the star cluster formation in non-interacting galaxies (B{\\\" o}ker et al. 2002) are still unclear.  In order to understand the formation of globular clusters, it is essential to trace the youngest protoglobular clusters. These young and massive star clusters, or super star clusters -- ``SSC'', contain hundreds to thousands O stars in their very first few million years of life. They are often embedded in their dusty natal cocoons  (Wynn-Williams et al. 1972) and visually obscured by their surrounding birth clouds. The high optical extinction is true as well for the \\HII\\ regions which are large enough, with much of the extinction internal to the nebula itself (Kawara et al. 1989; Ho et al. 1990; Beck et al. 1996). Thus, the forming SSC usually cannot be seen in optical. However, embedded clusters can be detected through their nebulae, which glow by reprocessed UV light from O stars. The thermal free-free emission from these nebulae can be detected and identified at radio wavelengths. Hence radio continuum observations which suffer much less extinction are often useful tracers of young star forming regions.  We report the centimeter continuum study at 2 and 6 cm on the centers of four nearby and well-studied spiral galaxies, IC 342, Maffei II, NGC 2903, and NGC 6946. The centers of these galaxies are infrared-bright and molecular gas rich. The goal of this investigation is to identify young SSC candidates in the star formation regions in the centers of these galaxies using subarcsecond radio continuum imaging. Subarcsecond imaging with the \\textit{VLA} in its extended configuration maximizes sensitivity to bright and compact radio ``supernebulae'' over low brightness disk synchrotron emission. Observing at shorter wavelengths, $\\lambda<$ 2 cm, also minimizes the contribution of synchrotron emission, which falls with frequency.  The results presented here are detailed in Tsai et al. (2006)  \\begin{table*} \\caption{\\normalsize{Sample Galaxies and Identified Radio Sources}} \\begin{center} \\begin{tabular}{lcccccc} \\hline\\hline\\\\[-6pt] Galaxy    & D     & $S_{6cm}^{compact}$/$S_{6cm}^{total;a}$ & $S_{2cm}^{compact}$/$S_{2cm}^{total;a}$ & \\HII$^{b}$ & SNR$^{b}$ & Indeterminated$^{b}$ \\\\ & (Mpc) & (mJy/mJy) & (mJy/mJy) & Thick/Thin &     &              \\\\ \\hline IC 342    & 3.3   &     7.5/82          &     7.5/38          & 5/1     & 6   & 0      \\\\ Maffei II & 5.0   &    13.5/107         &    12.7/46          & 4/3     & 3   & 0      \\\\ NGC 2903  & 8.9   &     1.7/35          &     3.0/12          & 4/2     & 1   & 0      \\\\ NGC 6946  & 5.9   &     6.8/39          &     6.1/23          & 1/3     & 2   & 3      \\\\   \\hline \\\\[-6pt] \\multicolumn{7}{l}{$^{a}$The flux ratio between the summation of compact source fluxes and total flux. Total flux is the flux} \\\\ \\multicolumn{7}{l}{\\ \\ obtained from \\textit{VLA} B- and C-configuration measurements (Turner \\& Ho 1983; Turner \\& Ho 1994; } \\\\ \\multicolumn{7}{l}{\\ \\ Wynn-Williams \\& Becklin 1985).} \\\\ \\multicolumn{7}{l}{$^{b}$Number of compact sources in the category. See the definition of each category in text.} \\end{tabular} \\end{center} \\end{table*} ",
        "conclusions": "More than half of the 38 identified compact sources are \\HII\\ regions with flat or positive indices ($\\alpha \\gsim -0.1$). More than 60\\,\\% of these \\HII\\ regions have a spectrum rising from 6 to 2 cm, indicating that they are at least partly optically thick at 6~cm. The radio fluxes of these radio \\HII\\ regions require that they harbor hundreds of massive stars. For electron temperatures of $\\sim 10,000$~K, rising cm-wave spectra imply that the rms electron densities are $\\sim 10^{4}~cm^{-3}$.  The deconvolved sizes of \\HII\\ regions correspond to diameters of $\\sim$ 6 - 14~pc at the distances of the galaxies. These \\HII\\ regions are much smaller than the 30 Doradus \\HII\\ region ($\\sim$ 200~pc in diameter for region of $EM > 10^{4}~pc~cm^{-6}$; see Mills et al. 1978 and Kennicutt 1984) in the Large Magellanic Cloud (LMC). They are larger but similar in nature to the Galactic compact \\HII\\ regions in W49 (Mezger et al. 1967; Conti \\& Blum 2002) and NGC 3603 (de Pree et al. 1999; Sung \\& Bessell 2004, M{\\\" u}cke et al. 2002). One possibility for the fact that these \\HII\\ regions are so compact and dense is that they are younger than 30 Doradus or NGC 3603.  Nearly two dozen radio sources in the four galaxies of this study were found to have flat or rising spectra. These are presumably thermal nebulae excited by massive stars, for which the 2 cm fluxes give a lower limit to the required $N_{lyc}$ and corresponding $L_{H_{\\alpha}}$. We cannot detect unresolved (diameter $<$ 15 pc) nebulae with $N_{Lyc} < 4\\times 10^{50}~ s^{-1}$ due to our sensitivity limits, corresponding to 40 O stars in IC 342 and 120 stars in NGC 6946. However, we also cannot detect more luminous regions if they are large and resolved (size $\\gg$ 15~pc); in this limit we are sensitive to \\HII\\ regions with $EM > 10^4 \\rm ~cm^{-6}\\,pc$. The youngest and therefore smallest \\HII\\ regions are probably not affected by this latter criterion.  The luminosity function (LF) of the 23 thermal sources in 4 spirals shown in Figure 2 suggests a broken power law with turnover point at $L_{H_{\\alpha}} = 2 \\times 10^{39} erg/sec$ and cutoff at luminous end around $L_{H_{\\alpha}} = 10^{40} erg/sec$. The total mass of such a cluster at cutoff would be $\\sim 2 \\times 10^{5} \\; M_{\\sun}$ (assuming Salpeter initial mass function, IMF), consistent with the mass of Galactic globular clusters. Very similar feature shows in LF for 21 thermal radio sources in M 82, another nearby starburst galaxy, based on more complete radio studies of Allen (1999), with sensitivity limits of $N_{Lyc} < 2\\times 10^{50}~ s^{-1}$ (20 O7 stars) and radio nebulae with size $\\lsim$ 50 pc. The agreement of the LF shape of M 82 and our sample suggests that the broken power law may be real. However, the results based on luminosity function of 44 sources are statistically inconclusive. Larger samples would help to clarify this issue.  The steep power law LF suggests a cutoff at the high luminosity end in the young clusters associated with radio nebulae in spiral galaxies. The upper end of the luminosity function is around $L_{H_{\\alpha}} \\sim 10^{40}~ erg~ s^{-1}$. This corresponds to $\\sim$ 750 O7 stars. We observe two such clusters, in Maffei 2 and NGC 2903, and their nebulae are $\\lsim5$ pc in extent. The total mass of such a cluster would be $\\sim$ 2.2 $\\times 10^{5}~ M_{\\sun}$ (assuming Salpeter IMF), consistent with the mass of Galactic globular clusters."
    },
    "0609/astro-ph0609576_arXiv.txt": {
        "abstract": "X-Ray Polarimetry can be now performed by using a Micro Pattern Gas Chamber in the focus of a telescope. It requires large area optics for most important scientific targets. But since the technique is additive a dedicated mission with a cluster of small telescopes can perform many important measurements and bridge the 40 year gap between OSO-8 data and future big telescopes such as XEUS. POLARIX has been conceived as such a pathfinder. It is a Small Satellite based on the optics of JET-X. Two telescopes are available in flight configuration and three more can be easily produced starting from the available superpolished mandrels. We show the capabilities of such a cluster of telescopes each equipped with a focal plane photoelectric polarimeter and discuss a few alternative solutions. ",
        "introduction": "\\label{sect:intro}  % A long term theoretical analysis has foreseen the possibility to test models of X-ray sources and to derive relevant parameters of the models themselves by measuring the linear polarization of X-rays. In the first years the polarization deriving from the emission process was outlined. On 1985 Sunyaev and  and Titarchuk demonstrated that in a hot accretion disk the scattered photons have an energy much higher than the original photons and thence the Chandrasekhar limit of 11.7 $\\%$ could be exceeded. In 1988 Meszaros et al. stated \"Polarimetry would add to energy and time two further observable quantities (the amount and the angle of polarization) constraining any model and interpretation: a theoretical/observational break-through\".  This very attractive perspective pushed some investigators to develop instrumentation aimed to perform this measurements. This was mainly based on Bragg Diffraction at 45$^o$ and on Compton/Thomson scattering in a certain range of angles around 90$^o$. Bragg crystal is an excellent analyzer of polarization (100$\\%$ modulation) but is effective in a very narrow interval around the energy that fulfils the Bragg condition for that particular crystal and that particular angle. A flat Bragg crystal at 45$^o$ set before the focal plane of an X-ray telescope acts as a newtonian secondary mirror and a detector in the secondary focal plane detects the source in the thin band defined by the Bragg (plus higher orders if in the band). The whole is rotated and the source brightness is modulated depending on the amount of polarization.  The scattering polarimetry is effective only when scattering exceeds photoabsorption. A well of detectors surrounds the scatterer and the polarization is measured by the angular distribution of the out-coming photons. This technique can be used by means of large experiments with a distributed multiplicity of detectors and scatterers and many experiments based on this principle have been proposed and some are on the way\\cite{Legere05}. Scattering polarimetry can also be performed in the focus of a telescope with some substantial limitations: even with lithium the method is only effective above 5 keV and this is mismatched with the bandpass of most telescopes; the angular distribution strongly depends on the impact point and the measurement is severely affected by systematics; all the detector set (and not only a few pixels) is involved in the measurement so that the improvement of signal to noise ratio introduced by the use of the telescope is essentially lost.  The pioneering result achieved by the Columbia University team was the measurement of polarization of Crab Nebula first with a rocket\\cite{Novick72} then with a Bragg polarimeter on-board OSO-8 satellite\\cite{Weisskopf76}. This result was of extreme relevance but the overall throughput of X-ray polarimetry based on conventional techniques was considered too meagre to support the inclusion of a polarimeter onboard further X-ray missions. Moreover theoretical analysis suggested that a polarimeter capable to perform polarimetry at few $\\%$ level also on relatively faint sources was needed for most of the scientific objectives. The only mission with a focal plane polarimeter (both Bragg and Compton) was the Spectrum X-Gamma\\cite{kaaret89} that sinked in the general collapse of the soviet system.  A third physical process that can analyze the polarization is photoelectric absorption: s-photoelectrons are ejected according to a cos$^{2}$ distribution, with respect to the electric field. This is a potentially ideal analyzer of polarization and many teams tried, also in early times, to build a device based on this effect. The difficulty is that electrons penetrate much less than photons, and a finely subdivided detector is needed. Various attempts arrived to perform this measurement in one dimension\\cite{Soffitta95} or in two dimensions but at higher energies\\cite{austin94}. Eventually the Micropattern Gas Chamber was developed with a full capability to convert photons and image photoelectrons, making possible focal plane photoelectric polarimetry. ",
        "conclusions": "POLARIX is a based on a new but already reliable technology and is capable to perform a set of measurements of high astrophysic interest and open the path to polarimetry with future large aperture telescopes. An advanced study to evolve this concept to a better defined design is about to start under ASI contract."
    },
    "0609/astro-ph0609095_arXiv.txt": {
        "abstract": "We investigate the effect of blending in future gravitational microlensing surveys by carrying out simulation of Galactic bulge microlensing events to be detected from a proposed space-based lensing survey.  From this simulation, we find that the contribution of the flux from background stars to the total blended flux will be equivalent to that from the lens itself despite the greatly improved resolution from space observations, implying that characterizing lenses from the analysis of the blended flux would not be easy.  As a method to isolate events for which most of the blended flux is attributable to the lens, we propose to use astrometric information of source star image centroid motion.  For the sample of events obtained by imposing a criterion that the centroid shift should be less than three times of the astrometric uncertainty among the events for which blending is noticed with blended light fractions $f_{\\rm B}>0.2$, we estimate that the contamination of the blended flux by background stars will be less than 20\\% for most ($\\sim 90\\%$) of the sample events.  The expected rate of these events is $\\gtrsim 700$ events/yr, which is large enough for the statistical analysis of the lens populations. ",
        "introduction": "Observational experiments to detect light variations of stars induced by gravitational microlensing have been and are being carried out for more than a decade \\citep{udalski93, alcock93, aubourg93, bond01}. Originally, the experiments were initiated for the purpose of searching for dark matter in the Galactic halo \\citep{paczynski86}.  However, microlensing developed many applications in various aspects of galactic, stellar, and planetary astrophysics such as probing the structure of the inner galaxy \\citep{kiraga94, evans94, han95b, zhao95}, resolving the photosphere of distant stars \\citep{witt95, albrow99, abe03, field03}, and detecting extrasolar planets \\citep{mao91, gould92, bond04, udalski05, beaulieu06, gould06}.  Current experiments are routinely detecting more than 500 events every season and the total number of detected events is now reaching 3000.   The probability of a star being gravitationally lensed is very low (of the order of $10^{-6}$).  As events are very rare, lensing searches are conducted toward very crowded star fields such as the Galactic bulge in order to monitor as many stars as possible within a single frame.  For observations in such a dense field, blending of stellar images is inevitable \\citep{distefano95, wozniak97}.  The lightcurve of a blending-affected event is expressed as \\begin{equation} F_{\\rm obs} = AF_0 + F_{\\rm B};\\qquad A={u^2+2\\over u(u^2+4)}, \\label{eq1} \\end{equation} where $F_0$ is the baseline flux of the lensed source star, $F_{\\rm B}$ is the amount of the blended flux, $A$ is the magnification induced by lensing, and $u$ is the lens-source separation normalized by the Einstein ring radius $\\thetae$.  The normalized lens-source separation is related to the lensing parameters by \\begin{equation} u = \\left[ \\left( {t-t_0\\over \\te}\\right)^2-u_0^2 \\right]^{1/2}, \\label{eq2} \\end{equation} where $\\te$ is the Einstein timescale, $t_0$ is the time of the closest lens-source approach, and $u_0$ is the separation at that moment (impact parameter).  Among the lensing parameters, the Einstein timescale is related to the mass and location of the lens and relative lens-source transverse speed, and thus it can provide constraint about the lens system.  However, the constraint is weak because the timescale results from the combination of the physical parameters.  Blending aggravates the photometric precision due to the increased noise from the blended flux.  In addition, blending causes accurate determinations of the lensing parameters difficult due to the poorly known baseline flux of the source star.  As a result, the constraint about the lens system becomes further weaker.   Recently, next-generation microlensing experiments based in space are being seriously discussed.  {\\it Microlensing Planet Finder} ({\\it MPF}), that succeeds the original concept of {\\it Galactic Exoplanet Survey Telescope} ({\\it GEST}) mission \\citep{bennett02}, is a space mission proposed to NASA's Discovery Program  with the main goal of searching for a large sample of extrasolar planets by using gravitation lensing method \\citep{bennett04}.  By conducting observations in space, it is possible to minimize blending effect thanks to the improved resolution. Another important advantage of the space observation may be that one can directly detect and characterize the lens based on the observed flux. This is possible when a light-emitting star is responsible for a lensing event and the flux from the lens contributes a substantial fraction of the observed flux.  It is believed that majority of events detected toward the Galactic bulge are caused by stars \\citep{han03}.  In addition, the {\\it MPF} mission plans to monitor faint main-sequence stars in order to optimize detections of low-mass planets by minimizing finite-source effect.  Therefore, the light from the lens would not be negligible for a significant fraction of the event sample, and thus the existence of the lens could be revealed by the excess flux from the lens. By conducting multi-band observations, it would be possible to determine the spectral type of the lens, which would provide much tighter constraint about the lens.   However, this new method of constraining lenses can be applicable under the condition that source stars to be observed are in most cases free from blending by the flux of blends other than lenses.  If an unrelated background star is located within the resolution disk of the lensed source star, the blended flux comes not only from the lens but also from the background star.  This causes complexity in the analysis of the observed blended light, making it difficult to uniquely characterize the lens.  \\citet{han05} (hereafter Paper I) warned of the possible contamination of the blended flux by background stars.  As a method to resolve the contamination, he suggested to use the astrometric information of the source star image centroid shift.  If an event is affected by background stellar blending, the centroid will shift from the apparent position of the combined image of the source plus blended stars toward the lensed star during the event \\citep{goldberg98}.   In this paper, we further investigate the blending problem in future space-based lensing surveys following Paper I and probe the feasibility of sorting out events whose blended flux is mostly attributable to the lenses by using the additional astrometric information of the centroid shift.  For this, we carry out simulation of Galactic bulge microlensing events expected to be detected from the proposed {\\it MPF} survey by imposing realistic observational conditions and detection criteria. Based on this simulation, we investigate the expected pattern of blending and the seriousness of the contamination by background stars.  We also estimate the fraction of events that can be identified to be free from background stellar blending by using the proposed method. ",
        "conclusions": "We investigated the effect of blending in future gravitational microlensing surveys using a 1 m class space telescope.  For this, we carried out simulation of Galactic bulge microlensing events to be detected from the proposed {\\it MPF} lensing survey by imposing realistic observational conditions and detection criteria.  From this simulation, we found that the contribution of the flux from blended background stars to the total blended flux is equivalent to that of the lens itself, implying that characterizing lenses from the analysis of the blended flux would not be easy.  However, with the additional astrometric information of the source star image centroid shift, we found that it would be possible to isolate events for which most of the blended flux is attributable to the lens.  For the sample of events obtained by imposing a criterion that the centroid shift should be less than three times of the astrometric uncertainty among the events for which blending was noticed with blended light fractions $f_{\\rm B}>0.2$, we estimated that the contamination of the blended flux by background stars would be less than 20\\% for most ($\\sim 90\\%$) of the sample events.  The expected rate of these events is $\\gtrsim 700$ events/yr, which is large enough for the statistical analysis of the lens populations."
    },
    "0609/astro-ph0609740_arXiv.txt": {
        "abstract": "We consider nonlinear biasing models of galaxies with particular attention to a correlation between linear and quadratic biasing coefficients, $b_1$ and $b_2$. We first derive perturbative expressions for $b_1$ and $b_2$ in halo and peak biasing models.  Then we compute  power spectra and bispectra of dark matter particles and halos using N-body simulation data and of volume-limited subsamples of Sloan Digital Sky Survey (SDSS) galaxies, and determine their $b_1$ and $b_2$. We find that the values of those coefficients at linear regimes ($k<0.2h{\\rm Mpc}^{-1}$) are fairly insensitive to the redshift-space distortion and the survey volume shape. The resulting normalized amplitudes of bispectra, $Q$, for equilateral triangles, are insensitive to the values of $b_1$ implying that $b_2$ indeed correlates with $b_1$.  The present results explain the previous finding of Kayo et al. (2004) for the hierarchical relation of three-point correlation functions of SDSS galaxies.  While the relations between $b_1$ and $b_2$ are quantitatively different for specific biasing models, their approximately similar correlations indicate a fairly generic outcome of the biasing due to the gravity in primordial Gaussian density fields. ",
        "introduction": "One of the major uncertainties in precise cosmology is the galaxy biasing relative to the underlying mass distribution. It hampers extracting the cosmological information from the large scale structure of the universe. In particular, the biasing is sensitive to the unknown physical conditions of galaxy formation. Thus its phenomenological parametrization and understanding is crucial in advancing our knowledge of the evolution of the universe.  A reasonable approximation often adopted is a local linear biasing model, which assumes a simple scale-independent relation between the density contrast fields of galaxy and mass: \\begin{equation} \\delta_g({\\bf x},z) = b(z)\\delta({\\bf x},z). \\label{eq:dg-d} \\end{equation} In the above, the density contrast of the mass is defined as \\begin{equation} \\delta({\\bf x},z) \\equiv \\frac{\\rho({\\bf x},z)-\\bar{\\rho}(z)}{\\bar{\\rho}(z)}, \\end{equation} where the over-bar indicates the mean over the entire universe. The number density contrast of galaxy density field, $\\delta_g$, is defined similarly. Because this model contains only a single time-dependent parameter $b(z)$, it is not surprising that the model cannot describe the observed features of galaxy clustering accurately.  Furthermore, the three-point correlation function of SDSS galaxies \\citep{Kayo2004} indicates the importance of the higher-order correction to equation (\\ref{eq:dg-d}); while the galaxy two-point correlation functions are well represented by the linear biasing model (e.g., \\cite{Zehavi2005}), the normalized amplitudes $Q$ of the corresponding three-point functions for equilateral triangles do not show the expected scaling with respect to $b(z)$: \\begin{equation} Q\\propto\\frac{1}{b}. \\label{eq:Qinvb} \\end{equation} In reality, however, $Q$ for equilateral triangles calculated from SDSS galaxies in redshift space proves to be almost scale-independent, and follows the hierarchical relation approximately, $Q=0.5\\sim1.0$. Moreover its dependence on the morphology, color, and luminosity is not statistically significant. Given the robust morphological, color and luminosity dependences of the two-point correlation function, \\citet{Kayo2004} argued that galaxy biasing is complex and requires a contrived relation between the linear biasing and its higher order correction terms. Another possibility is that the observed value is significantly contaminated by the redshift-space distortion and does not properly reflect the actual clustering information in real space. Indeed the previous N-body studies of the redshift-space distortion on three- and four-point correlation functions suggest that this is the case in nonlinear regimes \\citep{Suto93,MS94,SM94}.  The interpretation is further complicated by the fact that the higher-order statistics is sensitive to relatively rare large-scale structures in particular samples \\citep{Nichol06}.  The main purpose of the present paper is to provide a physical explanation for the approximate hierarchical relation for $Q$ of SDSS galaxy clustering on the basis of perturbation analysis and numerical simulations.  This is a step toward the understanding of the nonlinear nature of biasing. Specifically we employ a local {\\it nonlinear} biasing model \\citep{FG93}: \\begin{equation} \\delta_b({\\bf x},z) = \\sum_ {n=0}^\\infty\\frac{b_n(z)}{n!} \\left[\\delta({\\bf x},z)\\right]^n. \\label{eq:db-d} \\end{equation} We consider density peaks, dark matter halos, simulated halos, and the SDSS galaxies as specific examples for the density fields $\\delta_b$ in the left hand side.  The outline of the paper is as follows; section \\ref{sec:AM} briefly summarizes the halo and peak biasing models as analytically tractable examples. We compute the linear and quadratic biasing coefficients perturbatively and find that these models roughly agree with the observed hierarchical relation.  To be more realistic, we analyze simulated halo catalogs and SDSS galaxies in section \\ref{sec:simulation}, and find that the perturbative result is valid even if we take account of a variety of selection effects and particularly redshift-space distortion.  The quantitative agreement between the two indicates a generic correlation among the biasing coefficients in the SDSS galaxies. Finally, section \\ref{sec:SD} is devoted to the summary and conclusion of the paper. ",
        "conclusions": "\\label{sec:SD} We have found a fairly generic correlation between linear and quadratic biasing coefficients, $b_1$ and $b_2/b_1$, using a variety of different methodologies; perturbative expansions of peak and halo biasing models, N-body simulations of halos, SDSS galaxy data analysis, and the corresponding halo occupation distribution predictions (analytic and mock). The presence of such correlations was suggested earlier by the previous finding that the normalized three-point correlation functions of SDSS galaxies in redshift space follow the hierarchical relation approximately, $Q=0.5\\sim1.0$, despite the robust morphological, color and luminosity dependences of the corresponding two-point correlation functions \\citep{Kayo2004}. We have derived the $b_1$ -- $b_2$ correlation explicitly and showed for the first time that it indeed explains the observed behavior of $Q$ for equilateral triangles calculated from SDSS galaxies in linear regimes.  The major findings of the present paper are summarized as follows. \\begin{itemize} \\item Even with the presence of redshift distortion and complicated survey shape effects, $b_1$ can be accurately estimated in $k<0.2h{\\rm Mpc}^{-1}$. So $b_1$ estimated from SDSS galaxies is expected to reflect the true value in real space. \\item The values of $Q_b$ for equilateral triangles from simulated halos, mock galaxies based on HOD model, and SDSS galaxies do not show any clear dependence on $b_1$. The independence on $b_1$ is not an artifact from the redshift-space distortion contamination, but is a consequence of the intrinsic correlation between $b_2/b_1$ and $b_1$. \\item The HOD model, equation (\\ref{eq:HOD}), whose parameters are determined so as to reproduce the observed two-point statistics alone, seems to be also successful in predicting the value of $b_2$. \\end{itemize}  There are a few remaining tasks which should be done following the present result.  First, it is interesting to compare the current analysis in Fourier space with that in configuration space.  As long as nonlinear effect is negligible, $Q_m$ and $Q_b$ in k-space are expected to be the same as those in configuration space.  Nevertheless we would like to make sure of it, and to further explore the connection between bispectra and the three-point correlation functions of halos and galaxies. Second, the success of HOD to reproduce the correlation between $b_2/b_1$ and $b_1$ at the current level is very encouraging. Naturally it is important to further improve the HOD model. Third, the present paper considers only equilateral configuration of Fourier space triangles for $Q$. In reality, however, the different shape of triangles has a wealth of information. While the $b_1$ dependence of $Q$ barely vanishes for equilateral triangles due to the correlation between $b_1$ and $b_2/b_1$, this is not the case for arbitrary triangle shapes. Therefore it is interesting to explore if the $b_1$ dependence of $Q$ for the other triangle shapes and to see if the results are consistent with the present analysis and/or to constrain higher order biasing coefficients more tightly. The work along the line is now in progress.    \\bigskip  We thank Akihito Shirata and Kohji Yoshikawa for useful discussions. The N-body simulations were performed at ADAC (Astronomical Data Analysis Center) of the National Astronomical Observatory, Japan. K.~Y. and C.~H. acknowledge the support from Grants-in-Aid for Japan Society for the Promotion of Science Fellows. This research was partly supported by Grant-in-Aid for Scientific Research of Japan Society for Promotion of Science (No.16340053, 17740139).   Funding for the SDSS and SDSS-II has been provided 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.  The SDSS Web Site is http://www.sdss.org/.  The SDSS is managed by the Astrophysical Research Consortium for the Participating Institutions. The Participating Institutions are the American Museum of Natural History, Astrophysical Institute Potsdam, University of Basel, Cambridge University, Case Western Reserve University, University of Chicago, Drexel University, Fermilab, the Institute for Advanced Study, the Japan Participation Group, Johns Hopkins University, the Joint Institute for Nuclear Astrophysics, the Kavli Institute for Particle Astrophysics and Cosmology, the Korean Scientist Group, the Chinese Academy of Sciences (LAMOST), Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, Ohio State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory, and the University of Washington.  \\bigskip"
    },
    "0609/astro-ph0609430_arXiv.txt": {
        "abstract": "I have recently combined HST/STIS spectrophotometry with existing photometric data to analyze the calibration of three standard optical photometry systems: Tycho-2 $B_{\\rm T}V_{\\rm T}$, Str\\\"omgren $uvby$, and Johnson $UBV$. In this contribution I summarize those results, present new ones for 2MASS $JHK_s$, and combine them with recent literature results to generate a uniform set of zero points for six photometric systems, the above mentioned plus Cousins $RI$ and SDSS $ugriz$. With the exception of the latter system, the zero points use the new Vega spectrum presented at this meeting by Ralph Bohlin. I also discuss the implementation of these results in CHORIZOS, a Bayesian photometric code that compares multi-filter observational data with spectral energy distributions to solve the inverse problem of finding the models which are compatible with the observations. ",
        "introduction": "The inverse photometric problem in astronomy can be defined in the following way: one observes a series of magnitudes $m_{{\\rm obs},p}$ in a series of filter passbands (denoted here by the $p$ index), each one of them defined by a total-system dimensionless sensitivity function $P_p(\\lambda)$, and compares them with a family of \\sindex{spectral energy distribution}s (SED, denoted here by the $s$ index) $f_{\\lambda,s}(\\lambda)$ to find out which of the SEDs is compatible with the observed magnitudes. The first required step in that process is to compute the synthetic magnitudes $m_{r,p}$ in the $r$ system for each SED in each passband $p$. For a photon-counting detector, the formula is:  \\begin{equation} m_{r,p}[f_{\\lambda,s}(\\lambda)] = -2.5\\log_{10}\\left(\\frac{\\int P_p(\\lambda)f_{\\lambda,s}(\\lambda)\\lambda\\,d\\lambda} {\\int P_p(\\lambda)f_{\\lambda,r}(\\lambda)\\lambda\\,d\\lambda}\\right) + {\\rm ZP}_{r,p}. \\label{photon} \\end{equation}  A magnitude system $r$ is defined by a reference SED $f_{\\lambda,r}(\\lambda)$ and a series of (relative) zero points ZP$_{r,p}$ for each filter. In principle, one can adjust the $f_{\\lambda,r}(\\lambda)$ in such a way that ZP$_{r,p}$ = 0.0 for all filters, which generates the default system for that SED. Examples are given in Table 1. Some magnitude systems (e.g. Str\\\"omgren+Vega), however, use non-zero ZP$_{r,p}$ by definition.  It is commonly found when one attempts to use a magnitude system with zero ZP$_{r,p}$ that, after the data has been fully processed, there are small offsets between observed and synthetic values (typically in the hundredths or thousands of magnitudes) that have to be corrected\\footnote{Given the unintended nature of this effect, some authors refer to these quantities as zero-point offsets instead of zero points.}. There are two approaches to deal with this issue: the first one, used e.g. for \\sindex{HST photometry}, is to leave the synthetic magnitudes unchanged and modify $m_{{\\rm obs},p}$ by changing the calibration parameters that are used to convert from instrumental counts to observed magnitudes. This leaves the zero points (as defined here) as exactly zero, e.g. ZP$_{{\\rm HST+Vega},p} = 0.0$ for all filters. Two problems with this approach are that the observed magnitudes of a given constant object (e.g. a calibration star) change when the reference SED is modified and that it would be hard to apply to systems with long histories (e.g. Johnson $UBV$). An alternative approach, which will be followed here, is to calculate the non-zero values of ZP$_{r,p}$, leaving the published observed magnitudes fixed and introducing changes in the synthetic ones. This is more practical way to deal with historical systems but it means that, for example, the synthetic magnitudes in the Johnson+Vega system\\footnote{That is, the values that should be compared with the observed ones.}, $m_{{\\rm Johnson+Vega},p}$, where $p=U,B,V$, are not the same as those calculated in the default \\sindex{Vega system} (also called VEGAMAG), $m_{{\\rm Vega},p}$, due to the existence of zero points. This one is also a more practical approach for large-scale surveys (Tycho-2, SDSS, 2MASS\\ldots) where the magnitudes need to be published as soon as possible and it is not until a few years later that accurate zero points are calculated. As an example, synthetic SDSS magnitudes, $m_{{\\rm SDSS+AB},p}$, where $p=u,g,r,i,z$, are reported in an AB-based system and require a non-zero ZP$_{{\\rm SDSS+AB},p}$ in order to be accurately compared with observed magnitudes.  \\bigskip  \\centerline{{\\bf Table 1.} Default magnitude systems (ZP$_{r,p} = 0.0$ for all $p$).}  \\medskip  \\centerline{ \\begin{tabular}{lllll} \\hline \\hline System     & \\multicolumn{4}{c}{Reference spectral energy distribution}                                 \\\\ \\hline Default ST & $f_{\\lambda,{\\rm ST}}$   & = & $3.63079\\cdot 10^{-9}$  & erg s$^{-1}$ cm$^{-2}$ \\AA$^{-1}$ \\\\ Default AB & $f_{\\nu,{\\rm AB}}$       & = & $3.63079\\cdot 10^{-20}$ & erg s$^{-1}$ cm$^{-2}$ Hz$^{-1}$  \\\\ VEGAMAG    & $f_{\\lambda,{\\rm Vega}}$ & = &\\multicolumn{2}{l}{Vega spectrum}                            \\\\ \\hline \\end{tabular}}  \\bigskip  The relationships between different magnitude systems are easily derived from Eqn~\\ref{photon}. We start with the magnitudes of \\oindex{Vega} in the default ST and AB systems, which are given by:  \\begin{equation} m_{{\\rm ST},p}({\\rm Vega}) = -2.5\\log_{10}\\left(\\frac{\\int P_p(\\lambda)f_{\\lambda,{\\rm Vega}}(\\lambda)\\lambda\\,d\\lambda} {\\int P_p(\\lambda)f_{\\lambda,{\\rm ST}}           \\lambda\\,d\\lambda}\\right), \\label{VegaST} \\end{equation}  \\begin{equation} m_{{\\rm AB},p}({\\rm Vega}) = -2.5\\log_{10}\\left(\\frac{\\int P_p(\\lambda)f_{\\lambda,{\\rm Vega}}(\\lambda)\\lambda\\,d\\lambda} {\\int P_p(\\lambda)(cf_{\\nu,{\\rm AB}}            /\\lambda) d\\lambda}\\right). \\label{VegaAB} \\end{equation}  $m_{{\\rm ST},p}({\\rm Vega})$ and $m_{{\\rm AB},p}({\\rm Vega})$ can then be used to transform from an X+Vega magnitude system (where X = \\sindex{Johnson}, \\sindex{Cousins}\\ldots) into ST or AB magnitudes by applying:  \\begin{equation} m_{{\\rm ST},p} = (m_{{\\rm X+Vega},p} - {\\rm ZP}_{{\\rm X+Vega},p}) + m_{{\\rm ST},p}({\\rm Vega}), \\label{mST} \\end{equation}  \\begin{equation} m_{{\\rm AB},p} = (m_{{\\rm X+Vega},p} - {\\rm ZP}_{{\\rm X+Vega},p}) + m_{{\\rm AB},p}({\\rm Vega}). \\label{mAB} \\end{equation}  Similarly, if we assume ${\\rm ZP}_{{\\rm X+Vega},p} = {\\rm ZP}_{{\\rm X+ST},p} = {\\rm ZP}_{{\\rm X+AB},p}$, we can transform from an X+ST or X+AB magnitude system into an X+Vega system by using:  \\begin{equation} m_{{\\rm X+Vega},p} = m_{{\\rm X+ST},p} - m_{{\\rm ST},p}({\\rm Vega}), \\label{mVegaST} \\end{equation}  \\begin{equation} m_{{\\rm X+Vega},p} = m_{{\\rm X+AB},p} - m_{{\\rm AB},p}({\\rm Vega}). \\label{mVegaAB} \\end{equation}  Some authors list the zero-magnitude $f_\\lambda$ or $f_\\nu$ for a given magnitude system and filter. Those quantities are easily derived from Eqns.~\\ref{mST}~and~\\ref{mAB}:  \\begin{equation} f_{\\lambda,p}(0\\; {\\rm mag}) = f_{\\lambda, {\\rm ST}}\\cdot 10^{-0.4[m_{{\\rm ST},p}({\\rm Vega})-{\\rm ZP}_{{\\rm X+Vega},p}]}, \\label{fl0} \\end{equation}  \\begin{equation} f_{\\nu,p}(0\\; {\\rm mag}) = f_{\\nu, {\\rm AB}}\\cdot 10^{-0.4[m_{{\\rm AB},p}({\\rm Vega})-{\\rm ZP}_{{\\rm X+Vega},p}]}. \\label{fn0} \\end{equation}  It has not been until recently that the quality and stability of the observational data and calibrations in astronomy has been good enough to allow for a measurement of the photometric zero points of a magnitude system with an accuracy of $\\le 1\\%$. Such an accuracy is necessary to solve the inverse photometric problem without introducing significant systematic errors. In this contribution I summarize the recent work that has attained that goal and I generate a uniform list of NIR/optical photometric zero points. ",
        "conclusions": ""
    },
    "0609/astro-ph0609606_arXiv.txt": {
        "abstract": "{\\emph{Context}. We report the discovery of a new luminous supersoft source, \\object{XMMU J005455.0$-$374117}, in the nearby spiral galaxy NGC~300, in \\xmm observations performed on 2005 May 22 and on 2005 November 25. The source is not present in \\rosat data nor in the previous \\xmm observations of 2000 December/2001 January.  The unique luminous supersoft source, \\object{XMMU J005510.7$-$373855}, detected in the 1992 May/June \\rosat data and in the 2000/2001 \\xmm data, fell below detectability.  This source already appeared highly variable in  \\rosat observations. \\emph{Aims}. We report on the temporal and spectral analysis of this new supersoft source (SSS) and compare its properties with the previous known SSS. \\emph{Methods}. We present the  light curves of the SSS, model its spectrum and estimate the corresponding flux and luminosities. \\emph{Results}. The light curve of XMMU J005455.0$-$374117  does not show large fluctuations  in any of the observations and its spectrum  can be  modelled with an absorbed blackbody with $kT\\sim60$\\,eV. The corresponding bolometric luminosity is $8.1^{+1.4}_{-4.5}\\times10^{38}\\,\\text{erg}\\,\\text{s}^{-1}$ in the first observation and drops to $2.2^{+0.5}_{-1.4}\\times10^{38}\\,\\text{erg}\\,\\text{s}^{-1}$ six months later. No optical source  brighter than $m_\\text{V}\\sim21.7\\,\\text{mag}$ is found coincident with its position. \\emph{Conclusions}. The luminosity of these two SSSs is higher than what has been found for `classical' SSSs. Their nature could be explained by beamed emission from steady nuclear burning of hydrogen onto white dwarfs, or accretion onto stellar-mass black hole with matter outflow or observed at high inclination angle. The presence of an intermediate-mass black hole seems unlikely in our case. ",
        "introduction": "\\label{sec:int}  Supersoft sources (SSSs) were first discovered in the Large Magellanic Cloud with the \\textsl{Einstein} Observatory \\citep{Long1981}, including the two prototypes \\object{CAL83} and \\object{CAL87}.  SSSs are characterised by very soft thermal spectra, with temperatures typically below 100\\,eV, and have bolometric luminosities in excess of $10^{36}\\,\\text{erg}\\,\\text{s}^{-1}$.  These sources can be divided into two groups. The first one includes the ``classical'' SSSs, which are characterized by bolometric luminosities in the range $10^{36}$--$2\\times10^{38}\\,\\text{erg}\\,\\text{s}^{-1}$.  The most promising model to explain the flux emitted from such SSSs was proposed by \\cite{vandenHeuvel1992} to be a steady nuclear burning of hydrogen accreted onto white dwarfs with masses in the range of 0.7--1.2\\,M$_{\\odot}$. These sources are fairly common: 57 sources have been catalogued by J.~Greiner\\footnote{\\url{http://www.aip.de/~jcg/sss/ssscat.html}} up to 1999 December, but since then many more were discovered in distant galaxies by \\xmm and \\chandra (see below). The second group includes SSSs for which the luminosity exceeds the Eddington limit for a $1.4\\,\\text{M}_\\odot$ compact object, therefore excluding unbeamed emission from steady nuclear burning of hydrogen accreted onto a white dwarf. These sources are much less common.  For ultraluminous SSSs, i.e., SSS with bolometric luminosities exceeding $10^{39}\\,\\text{erg}\\,\\text{s}^{-1}$, models involving intermediate-mass black holes \\citep{Kong2003, Swartz2002, Kong2004} or stellar-mass black holes with matter outflow \\citep{Mukai2005} have been invoked.  According to \\cite{Greiner2004}, 25 SSSs have been discovered in \\object{M31} of which a large fraction (30\\%) are found to be transient sources, with turnoff and turnon times of the order of a few months.  Their luminosities are in the range $10^{36}$--$10^{38}\\,\\text{erg}\\,\\text{s}^{-1}$.  \\cite{DiStefano2003} found 16 SSSs in \\object{M101}, 2--3 in \\object{M51}, 10 in \\object{M83} and 3 in \\object{NGC~4697}. Of these sources, 11 have bolometric luminosities $>10^{38}\\,\\text{erg}\\,\\text{s}^{-1}$, from which six are brighter than $10^{39}\\,\\text{erg}\\,\\text{s}^{-1}$. In \\object{M81}, \\cite{Swartz2002} found 9 SSSs, including two with bolometric luminosities $>10^{38}\\,\\text{erg}\\,\\text{s}^{-1}$ and one $>10^{39}\\,\\text{erg}\\,\\text{s}^{-1}$.  According to \\cite{DiStefano2003}, normal SSSs in spiral galaxies appear to be associated with the spiral arms. The most luminous SSSs, however, have been found either in the arms, bulge, or disk \\citep{Swartz2002}, as well as in halos \\citep{DiStefano2003b}. SSSs have been found either in spiral (e.g. M31, M101, M83, M81, \\object{M104}, \\& NGC~300), elliptical (e.g., NGC~4697), interacting (e.g., M51 and \\object{NGC~4038}/\\object{NGC~4039}, i.e., the Antennae) or irregular galaxies (e.g., \\object{LMC} and \\object{SMC}).   In this paper we report  the discovery of a luminous ($>2\\times10^{38}\\,\\text{erg}\\,\\text{s}^{-1}$) SSS, XMMU~J005455.0$-$374117, in the spiral galaxy NGC~300.  This galaxy is a normal dwarf galaxy of type SA(s)d located at a distance of $\\sim$1.88\\,Mpc \\citep{Gieren2005}. The galaxy is seen almost face-on and has a low Galactic column density \\citep[$N_\\text{H}=3.6\\times10^{20}\\,\\text{cm}^{-2}$;][]{Dickey1990}. The major axes of its $D_{25}$ optical disk are 13.3\\,kpc and 9.4\\,kpc \\citep[$22'\\times 15'$;][]{deVaucouleurs1991}.  NGC~300 was observed by \\rosat between 1991 and 1997 for a total of 46\\,ksec with the Position Sensitive Proportional Counter and 40\\,ksec with the High Resolution Imager.  One SSS, XMMU J005510.7$-$373855, was present in these observations \\citep{Read1997}. This source was visible in 1992 May and June but not in 1991 December \\citep{Read2001}.  The spectrum was well described with a thermal bremsstrahlung model with $kT \\sim 0.1$\\,keV \\citep{Read1997}.  More recently, we observed NGC~300 with \\xmm during its orbit 192 (2000 December 26; 37\\,ksec on source time) and orbit 195 (2001 January 1; 47\\,ksec on source time).  The results of these observations have been presented by \\cite{Carpano2005}.  A deep analysis of the SSS XMMU J005510.7$-$373855 as seen in these observations was also performed by \\cite{Kong2003}.  These authors report that during the 6\\,days between the two \\xmm pointings the source went from a ``high state'' to a ``low state'', and that a 5.4\\,h periodicity was found during the low state. More information about this source will be given in Sect.~\\ref{sec:discus}. Recently, \\xmm re-observed NGC~300 on 2005~May~22 (orbit 998) and on 2005~November~25 (orbit 1092), for 35\\,ksec each.  In this paper, we focus on the analysis of a new SSS, which was present in the 2005 \\xmm observations, and compare its properties with the previously known SSS.  For simplicity, we will refer to XMMU~J005510.7$-$373855 as SSS$_{1}$ and XMMU J005455.0$-$374117 as SSS$_{2}$ in this work.  Sect.~\\ref{sec:obs} describes the observations and the reduction of the \\xmm data. In Sect.~\\ref{sec:res} we present the spectral and timing analysis of SSS$_2$. We discuss  the nature of these SSSs in Sect.~\\ref{sec:discus} and conclusions are given in Sect.~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc} We report the discovery of a new luminous supersoft source in the 2005 May \\xmm observation of NGC~300. The previously known luminous supersoft source detected by \\rosat and the previous \\xmm data was below detectability.  This latter source already appeared highly variable in the \\rosat observations.  No object in the SIMBAD catalogue is associated with the new SSS, and, from the optical data, no counterpart brighter than $\\sim$21.7\\,mag ($\\sim$24.5\\,mag for the previous known SSS) has been found.  The X-ray spectrum is well described by an absorbed blackbody at temperatures of $kT$$\\sim$60\\,eV. The bolometric luminosity, in the highest observed state, is $8.1^{+1.4}_{-4.5}\\times10^{38}\\,\\text{erg}\\,\\text{s}^{-1}$ and dropped to $2.2^{+0.5}_{-1.4}\\times10^{38}\\,\\text{erg}\\,\\text{s}^{-1}$ six months later.  The SSSs in NGC~300 are brighter than ``classical'' SSSs, for which steady nuclear burning of hydrogen accreted onto white dwarfs has been suggested to explain their nature. They are too faint, however, to be classified as ultraluminous sources. We summarized several possible explanations for their nature. These involve beaming emission from a WD, intermediate-mass black hole (IMBH) or a stellar-mass black hole with matter outflow or observed at high inclination angle. Except the one including IMBH, which seems unlikely in our case, all models are consistent with our observations."
    },
    "0609/astro-ph0609520_arXiv.txt": {
        "abstract": "\\baselineskip = 11pt \\leftskip = 0.65in \\rightskip = 0.65in \\parindent=1pc {\\small The Taurus Molecular Cloud (TMC) ranks among the nearest and best-studied low-mass star formation regions. It contains numerous prototypical examples of deeply embedded protostars with massive disks and outflows, classical and weak-lined T Tauri stars, jets and Herbig-Haro objects, and a growing number of confirmed brown dwarfs. Star formation is ongoing, and the cloud covers all stages of pre-main sequence stellar evolution. We have initiated comprehensive surveys of the TMC, in particular including: (i)   a deep X-ray survey of about 5 sq. degrees with {\\it XMM-Newton}; (ii)  a near-to-mid-infrared photometric survey of $\\approx$30 sq. degrees with the Spitzer Space Telescope, mapping the entire cloud in all available photometric bands; and (iii) a deep optical  survey using the Canada-France-Hawaii Telescope. Each wavelength regime contributes to the understanding of different aspects of young stellar systems. {\\it XMM-Newton} and Spitzer mapping of the central TMC is a real breakthrough in disk characterization,  offering the most detailed studies of correlations between disk properties and high-energy magnetic processes in any low-mass star-forming region, extending also to brown dwarfs in which disk physics is largely unexplored. The optical data critically complements the other two surveys by allowing clear source identification with 0.8$^{\\prime\\prime}$ resolution, identifying substellar candidates, and, when combined with NIR data, providing the wavelength baseline to probe NIR excess emission. We  report results and correlation studies from these surveys. In particular, we address the physical interpretation of our new X-ray data, discuss the entire young stellar population from embedded protostars to weak-lined T Tau stars and their environment, and present new results on the low-mass population of the TMC, including young brown dwarfs. \\\\~\\\\~\\\\~}% ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609716_arXiv.txt": {
        "abstract": "We summarize recent numerical results on the control of the star formation efficiency (SFE), addressing the effects of turbulence and the magnetic field strength. In closed-box numerical simulations, the effect of the turbulent Mach number $\\Ms$ depends on whether the turbulence is driven or decaying: In driven regimes, increasing $\\Ms$ with all other parameters fixed decreases the SFE, while in decaying regimes the converse is true. The efficiencies in non-magnetic cases for realistic Mach numbers $\\Ms \\sim 10$ are somewhat too high compared to observed values. Including the magnetic field can bring the SFE down to levels consistent with observations, but the intensity of the magnetic field necessary to accomplish this depends again on whether the turbulence is driven or decaying. In this kind of simulations, a lifetime of the molecular cloud (MC) needs to be assumed, being typically a few free-fall times. Further progress requires determining the true nature of the turbulence driving and the lifetimes of the clouds. Simulations of MC formation by large-scale compressions in the warm neutral medium (WNM) show that the generation of the clouds' initial turbulence is built into the accumulation process that forms them, and that the turbulence is driven for as long as accumulation process lasts, producing realistic velocity dispersions and also thermal pressures in excess of the mean WNM value. In simulations including self-gravity, but neglecting the magnetic field and stellar energy feedback, the clouds never reach an equilibrium state, but rather evolve secularly, increasing their mass and gravitational energy until they engage in generalized gravitational collapse. However, local collapse events begin midways through this process, and produce enough stellar objetcs to disperse the cloud or at least halt its collapse before the latter is completed. Simulations of this kind including the missing physical ingredients should contribute to a final resolution of the MC lifetime and the origin of the low SFE problems. ",
        "introduction": "\\label{sec:intro}  Molecular clouds (MCs) are the densest regions in the interstellar medium (ISM) and also the site of all present-day star formation in the Galaxy. They are known to have masses much larger than their thermal Jeans mass, a fact that led \\cite{GK74} to propose that the clouds should be in a state of generalized gravitational collapse. However, \\cite{ZP74} readily noted that this would imply that the MCs should be forming stars at very high rates ($\\sim 30 M_\\sun$ yr$^{-1}$) if all of their mass were to be transformed into stars in roughly one free-fall time $\\tff$, while the observed rates are much lower ($\\sim 5 M_\\sun$ yr$^{-1}$; see, e.g., \\cite{SP04}), suggesting that the star formation efficiency (SFE) is reduced by some mechanism. The observed SFE ranges from a few percent when whole giant molecular complexes are considered (e.g., \\cite{Myers_etal86}), to 10--30\\% in cluster-forming cores (e.g., \\cite{LL03}).  The SFE in MCs, defined as the fraction of the clouds' mass that finally makes it into a star during their lifetime, can be simply written as SFE = ${\\rm SFR} \\times \\Delta \\tc$, where SFR is the star formation rate, and $\\Delta \\tc$ is the cloud lifetime. Thus, a low SFE can be obtained through either a small $\\Delta \\tc$ or a low SFR. Currently, there is an ongoing debate within the community on whether the cloud lifetimes are long, but the SFR is small (or even zero) over a major fraction of the cloud lifetime (\\cite[e.g.]{PS00}; \\cite{PS02}; \\cite{TM04}; \\cite{MTK06}; \\cite{TKM06}), or else the lifetimes are short, but the SFRs are relatively large (e.g., \\cite{BHV99}; \\cite{KHM00}; \\cite[Hartmann, \\BP\\ \\& Bergin 2001, herafter HBB01]{HBB01}; \\cite{Hart03}; \\cite{BBB03}; \\cite{BB06}; \\cite{BH06}; \\cite{VS_etal06b}).  MCs are also known to have supersonic linewidths, which have been attributed to turbulent motions (e.g., \\cite{ZE74}; \\cite{Larson81}; \\cite{Blitz93}). Turbulent flows are characterized by a scaling of the typical velocity difference $\\Delta v$ across points separated by a distance $\\ell$ that scales as $\\Delta v \\propto \\ell^\\alpha$, with $\\alpha >0$ (e.g., \\cite{Lesieur90}), implying that the largest velocity differences occur at the largest spatial scales of a given cloud or clump (\\cite{Larson81}). Thus, turbulence is expected to have a {\\it dual} role in the dynamics of MCs (e.g., \\cite{VP99}; \\cite{MK04}; \\cite{BKMV06}): On the one hand, with respect to regions of size $L$, supersonic compressive turbulent modes of size $\\ell > L$ will act mainly as pistons that can form a density peak (a ``cloud'', ``clump'' or ``core'') out of those regions. The timescale for clump formation is essentially the turbulent crossing time across scale $\\ell$. Since the compressions are supersonic, this is typically shorter than the free-fall or sound-crossing times. On the other hand, turbulent modes with $\\ell < L$ will provide support against the self-gravity of a clump of size $L$.  In this paper we review recent numerical results concerning the effect of the molecular cloud turbulence and the magnetic field on the regulation of the SFE, and discuss how the resolution of certain issues, such as the determination of the most appropriate set of parameters requires studying the formation of the clouds themselves. This review extends the one presented earlier by \\cite{VS05}. ",
        "conclusions": "\\label{sec:conclusions}    We conclude that numerical simulations of isolated clouds up to the present have quantitatively constrained the effect of the turbulent Mach number and the magnetic field strength on the SFE, but in turn this effect depends on the nature of the turbulence production and maintenance, and on the lifetimes of the clouds themselves. Simulations of MC formation within their diffuse environment have begun to shed light on these issues, but much parameter space exploration and inclusion of additional physics (notably, magnetic fields and stellar energy feedback) remain to be done."
    },
    "0609/astro-ph0609185_arXiv.txt": {
        "abstract": "We report on the analysis of the photometric and spectroscopic properties of a sample of 29 low redshift ($z<0.6$) QSOs for which both \\textit{HST} WFPC2 images and ultraviolet \\textit{HST} FOS spectra are available. For each object we measure the $R$ band absolute magnitude of the host galaxy, the \\cq (1550~\\AA) line width and the 1350~\\AA{} continuum luminosity. From these quantities we can estimate the black hole (BH) mass through the $M_{\\rmn{BH}}-L_{\\rmn{bulge}}$ relation for inactive galaxies, and from the virial method based on the kinematics of the regions emitting the broad lines. The comparison of the masses derived from the two methods yields information on the geometry of the gas emitting regions bound to the massive BH. The cumulative distribution of the line widths is consistent with that produced by matter laying in planes with inclinations uniformly distributed between $\\sim$10$^{\\circ}$ and $\\sim$50$^{\\circ}$, which corresponds to a geometrical factor $f\\sim1.3$. Our results are compared with those of the literature and discussed within the unified model of AGN. ",
        "introduction": "Nowadays there is wide consensus that most if not all nearby early type galaxies host a supermassive black hole (SMBH) at their centers (see for a recent review Ferrarese 2006). However a direct determination of the BH mass (\\mbh) is possible only for a limited number of inactive galaxies in the local Universe, from dynamical methods (virial theorem) based on the observations of the kinematics of stars orbiting around the SMBH. An extension of this method to higher redshift objects is viable only for active galaxies, where \\mbh{} can be derived from the dynamics of gas gravitationally bound to the SMBH: this is the usual case of regions emitting the broad emission lines (BLR).  At low redshift correlations between \\mbh{} and the global properties of the host galaxies, namely the luminosity ($L_{\\rmn{bulge}}$) and the central velocity dispersion of the bulge, have been found. Locally therefore the measure of the properties of the host galaxy is an alternative method for evaluating \\mbh.  The aim of this paper is to compare the mass determinations through the virial method and through the global parameters of the hosts of close--by quasars. For QSOs it is very difficult to measure the velocity dispersion of the stars in the bulge, so one must rely on the \\mbhlb{} relationship. For nearby ellipticals, there is increasing evidence that this relation has the same behavior for both inactive and active galaxies (see for example Merritt \\& Ferrarese 2001; McLure \\& Dunlop 2001, 2002; Bettoni et al. 2003), notwithstanding a substantial increase of the dispersion in the latter case, possibly due to the greater difficulty in the determination of all the parameters. Moreover, AGN hosts and inactive galaxies appear substantially the same objects, as indicated by direct morphological studies (e.g. Dunlop et al.~2003). This is consistent with a picture where  the AGN phenomenon, which is linked to the gas accretion on the central BH, would be a phase characterizing the life of normal galaxies in particular during major mergers (see for example Kauffmann \\& Haehnelt 2000). For these reasons, in this work we start with the assumption that the correlation established for nearby galaxies is also valid for higher redshift QSOs, provided that host galaxy luminosity is corrected for galactic evolution to $z\\sim0$.  A second set of QSO BH mass determinations can be obtained through the virial theorem, provided that an estimate of the velocity \\vblr{} and radius \\rblr{} of the BLR is available. The BLR velocity can be derived from the FWHM of the broad emission lines if one assumes a value for the geometrical factor $f$, which is an unknown parameter of order unity linked to the geometry of the BLR. The measurement of \\rblr{} may require cumbersome procedures, like reverberation mapping (Ulrich et al.~1984; Blandford \\& McKee 1982; Peterson 2001), that uses the continuum source variability to evaluate time delays between the central emission and the BLR gas response; an alternative is to infer the BLR size from the ionizing luminosity with which it is correlated (Kaspi et al.~2000; Vestergaard 2002; Pian, Falomo \\& Treves 2005; Kaspi et al.~2005; Vestergaard \\& Peterson 2006). BH mass determinations can be then derived from a measure of the broad line widths and of the ionizing continuum intensity, once the value of the geometrical factor is fixed.  In the next section we focus on the quasar sample choice and on the data collection and analysis (\\ref{secima}--\\ref{secspe}). In section \\ref{secmas} we describe the formulae used for applying the two methods of mass determination (\\ref{secpho}--\\ref{secvir}) and report on the results (\\ref{seccomp}). This analysis constrains the value of the geometrical factor. The issue of the BLR geometry and orientation is further discussed in paragraph \\ref{calcolof}. In section \\ref{secsum} our results are compared with those of literature and discussed within the unified model of AGN.  Throughout this paper, we adopt a concordant cosmology with $H_0=72$ ~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_m=0.3$ and $\\Omega_{\\Lambda}=0.7$. We have converted the results of other authors to this cosmology when adopting their relations and using their data. ",
        "conclusions": "\\label{secsum} We have evaluated the BH masses of a sample of 29 low redshift quasars using both a photometric and a virial method. This comparison suggests the value of the geometrical factor, $f=1.3$. We also found that the cumulative FWHM distribution of \\cq(1550~\\AA) is consistent with both a disc--like BLR geometry and an isotropic velocity field, provided that in the latter case there is some distribution around the mean value of the velocity. However, the virial BH masses coincide on average with the photometric ones only in the disc picture, which strongly supports the hypothesis that the broad line region velocity field lies on the plane of the accretion disc. In this scenario, one can also put some constraints on the inclination angles of QSOs, obtaining that they are uniformly oriented between $5^{\\circ}-20^{\\circ}$ and $45^{\\circ}-50^{\\circ}$. This result is in good agreement with the unified model of AGN (see Barthel 1989; Antonucci \\& Miller 1985; Antonucci 1993; Urry \\& Padovani 1995).  One possible caveat on this result is that we used \\cq{} FWHM measurements instead of H$\\beta$ ones. If we adopted the H$\\beta$ FWHMs (see Section \\ref{secspe}), we would obtain on average BH masses $0.2\\,$dex larger. This would imply a value of the geometrical factor $f\\sim1$, which is still inconsistent with the isotropic picture.  The quasar BH masses range over two decades ($10^{8}\\,-\\,10^{9.7}\\,\\mmsun$). Both the photometric and the virial approaches indicate that RLQ masses exceed $10^{8.6}\\mmsun$, while RQQs are distributed in the whole mass range. RLQs have on average BH masses greater than RQQs; the difference $\\delta$ between the mean values of $\\log\\mmbh$ of the RLQ and RQQ subsamples is: \\begin{displaymath} \\delta\\equiv\\langle\\, \\log\\mmbh(\\rmn{RLQ})\\,\\rangle -\\langle\\, \\log\\mmbh(\\rmn{RQQ})\\,\\rangle =0.26. \\end{displaymath} This is in good agreement with the results obtained by Dunlop et al.~(2003) and Falomo et al.~(2004), who found respectively $\\delta=0.34$ for 23 QSOs at $\\langle z\\rangle =0.2$ and $\\delta=0.33$ for 16 QSOs at higher redshift ($\\langle z\\rangle =1.5$).  The general agreement between the photometric and the spectroscopic methods may be considered as an indirect indication that the \\mbhlb{} correlation obtained locally can be extrapolated at higher redshift ($z\\sim0.5$).  Our results on the BLR geometry agree well with that of the seminal paper of McLure \\& Dunlop (2001), who proposed that $f=1.5$ and obtained that a sample of 30 QSOs and 15 Seyfert I galaxies is well described by a disc--like BLR geometry, provided that it is composed of two populations with inclination angles uniformly distributed within $0^{\\circ}<\\theta<20^{\\circ}$ and $0^{\\circ}<\\theta<46^{\\circ}$ respectively. Our set of quasars is comparable to theirs, but the spectroscopy is of much better quality. In fact, the scatter in terms of black hole mass around the best--fitting relation of Fig.~\\ref{massemasseLQ} ($\\sigma=0.30$) is appreciably smaller than that obtained by McLure \\& Dunlop ($\\sigma=0.59$). Our results are in good agreement also with the findings by Onken et al. (2004), who compared \\mbh{} derived by the virial method with those from the stellar velocity dispersion for a sample of 16 low redshift AGN, obtaining that the geometrical factor, following our definition, is $f\\sim1.2$.  In conclusion, we believe that there is strong evidence that, regardless of the method used to derive BH masses, the geometrical factor results in the range $1.0<f<1.3$.  \\begin{footnotesize}"
    },
    "0609/astro-ph0609466_arXiv.txt": {
        "abstract": "As part of the {\\it All Wavelength Extended Groth Strip International Survey (AEGIS),\\,} we describe the panchromatic characterization of an X-ray luminous active galactic nucleus (AGN) in a merging galaxy at $z$\\,=\\,1.15.  This object is detected at infrared (8\\micpa, 24\\micpa, 70\\micpa, 160\\micpa), submillimeter (850\\micpa) and radio wavelengths, from which we derive a bolometric luminosity L$_{\\rm bol}$\\,$\\sim$\\,9$\\times$10$^{12}$\\,\\Lsol.  We find that the AGN clearly dominates the hot dust emission below 40\\mic but its total energetic power inferred from the hard X-rays is substantially less than the bolometric output of the system. About 50\\% of the infrared luminosity is indeed produced by a cold dust component that probably originates from enshrouded star formation in the host galaxy.  In the context of a coeval growth of stellar bulges and massive black holes, this source might represent a ``transition'' object sharing properties with both quasars and luminous starbursts.  Study of such composite galaxies will help address how the star formation and disk-accretion phenomena may have regulated each other at high redshift and how this coordination may have participated to the build-up of the relationship observed locally between the masses of black holes and stellar spheroids. ",
        "introduction": "Galaxies with a bolometric luminosity exceeding 10$^{12}$\\,\\Lsol \\, are often powered by a combination of massive star formation and accretion of material around active nuclei \\citep[e.g.,][]{Genzel98}. Early in cosmic history the connection between these two phenomena may have led to a coeval growth of super massive black holes (SMBHs) and stellar spheroids \\citep[e.g.,][]{Page01}, and could thus be the foundation of the correlation observed between the masses of these two components in the local Universe \\citep[e.g.,][]{Gebhardt00}. However, the implication of such a co-evolution in the more general context of the stellar mass built-up history \\citep[e.g.,][]{Dickinson03} and SMBH formation \\citep[e.g.,][]{Barger05} has not been fully addressed.  Understanding the importance of this ``coordinated'' activity of starbursts and active galactic nuclei (AGNs) requires properly deconvolving their respective contributions {\\it within individual objects,} a challenging task particularly for distant sources.  Because of the complexity of the AGN and starburst spectral energy distributions (SEDs), this decomposition can best be achieved by combining data from wavebands that offer distinctive spectral features to characterize these phenomena.  For instance, active nuclei are strong emitters at high energy and they are usually associated with a continuum of hot dust peaking in the mid-infrared (IR). Luminous starbursts, on the other hand, are characterized by a colder dust component and emit the bulk of their luminosity in the far-IR.  To illustrate how the coexistence of these two processes at high redshift could be studied using larger samples, we present a panchromatic analysis of a luminous AGN referenced as CXO--GWS--J141741.9+522823 (hereafter CXO-J1417) by \\citet{Nandra05}. It is embedded in an ultra-luminous infrared galaxy (ULIRG) at $z$\\,=\\,1.15 and its detection across the full electromagnetic spectrum allows us to constrain the level of star formation also present in this source.  We assume a $\\Lambda$CDM cosmology with H$_0$\\,=\\,70~km~s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_m$\\,=\\,0.3 and $\\Omega_{\\lambda}\\,=\\,0.7$. ",
        "conclusions": "\\subsection{On the nature of \\cxopa}  Although the bolometric correction for luminous and strongly-absorbed X-ray sources has not been very well constrained so far, the obscuration toward \\cxo is still reasonable enough to allow a fairly secure estimate of the total luminosity of the active nucleus in this object.  Comparison of the SED shown in Fig.\\,3 with typical AGN templates \\citep[e.g.,][]{Elvis94,Silva04} indicate that this bolometric luminosity should typically range between 1$\\times$10$^{12}$\\,\\Lsol \\, and 3$\\times$10$^{12}$\\,\\Lsol.  While the AGN can thus power most of the hot dust detected in the system, it is not energetic enough to account also for the far-IR emission.  We argue that the cold dust component is produced by a deeply enshrouded starburst in the host galaxy (see also e.g., \\citealt{Waskett03}). \\cxo could be therefore a high redshift analog of some nearby ``composite'' ULIRGs where the contribution of the AGN to the bolometric output is comparable to that of the star-forming activity \\citep{Farrah03}.  It could also be similar to other distant X-ray selected sources that were detected at long wavelengths \\citep[e.g.,][]{Page01}, though this far-IR emission has been sometimes assumed to originate from the active nucleus rather than star formation \\citep{Barger05}.  Assuming the calibration from \\citet{Kennicutt98}, the IR luminosity of the cold component translates into a star-formation rate $SFR$\\,$\\sim$\\,750\\,\\Msol\\,yr$^{-1}$. Such enhanced levels of activity usually occur within embedded and very compact regions surrounding the cores of galaxies ($\\sim$\\,100--300\\,pcs, \\citealt{Soifer00}). It is consistent with the absence of direct star formation signatures as inferred from our UV/optical photometry, as well as from the ACS image taking into account the spatial resolution of the $HST$ data (i.e., $\\sim$\\,1\\,kpc at $z$\\,=\\,1.15).  \\subsection{Implications} Assuming a typical accretion efficiency $\\epsilon$\\,=\\,0.1 \\citep{Marconi04}, the luminosity of the AGN in \\cxo translates into a mass accretion rate d$M$/d$t$\\,$\\sim$\\,3.1\\,\\Msol\\,yr$^{-1}$. This is typical of quasars at $z$\\,$\\sim$\\,1 \\citep{McLure04}, but it is larger than the rates measured in sources experiencing similar levels of starburst activity such as the more distant SCUBA sources \\citep{Alexander05}.  Furthermore, the bolometric luminosity of this object and the obscuration toward its nucleus suggest that the gas fueling and the accretion are occurring quite efficiently, probably close to the Eddington limit. Under this hypothesis we would derive a black hole mass of $\\sim$\\,1.4\\,$\\times$\\,10$^{8}$\\,\\Msol. This is typically an order of magnitude larger than the mass of the SMBHs determined in the submillimeter galaxies, and it would be even larger in the case of a sub-Eddington accretion.  These properties suggest that \\cxo is an object sharing characteristics with both starburst-dominated galaxies and quasars, where violent star formation is still happening while a massive black hole has already formed.  High redshift ULIRGs showing a mixture of star formation and AGN such as \\cxo could be interesting as tests of the evolutionary sequences that have been proposed to understand the connection between the two phenomena \\citep[e.g.,][]{Sanders88a}. In such scenarios for instance, merging galaxies first trigger powerful star formation, and as material settles into the cores of these objects it feeds a supermassive black hole that eventually emerges as a luminous AGN. The latter can then produce strong winds and outflows that feed energy back into the surrounding galaxy and may either quench or reactivate star formation (\\citealt{Springel05,Hopkins05,Silk05}).  In the case of \\cxo there is a dominant contribution of the nucleus in the near-IR and it is not clear whether an underlying bulge has already formed in the host galaxy. However, sources experiencing star formation and disk-accretion that both radiate a similar amount of energy throughout their lifetime would evolve toward massive galaxies that lie significantly out of the local ``$M_{\\rm BH} - \\sigma$'' relationship \\citep{Page01}. Although the starburst and the AGN in \\cxo are characterized by roughly equal luminosities now, we infer that if these two phenomena evolve together they may occur on quite different time scales and regulate each other efficiently for the bulges and SMBHs to grow in a coordinated manner.  Such transitional cases might be rare locally \\citep[e.g.,][]{Genzel98}. At higher redshift however, their importance relative to the infrared/submillimeter or X-ray selected objects where one type of activity (i.e., star formation or accretion) largely dominates is not yet known. Interestingly, \\cxo lies at the knee of the 2--8\\,keV luminosity function derived by \\citet{Barger05} at 0.8\\,$\\leq$\\,$z$\\,$\\leq$\\,1.2 but it is much more IR-luminous than most of star-forming galaxies at this epoch of cosmic history \\citep{LeFloch05}.  Searching for similar objects at higher redshifts when ULIRGs were a major component of the starbursting population \\citep{Blain02} should allow us to explore in more detail the role that this coexistence of AGNs and starbursts within galaxies played in shaping the present-day Universe.  Even though their identification could be challenging, \\cxo points to the type of evidence required for this goal. Large data sets from existing surveys like {\\it AEGIS\\,} should provide this information for enough sources to probe the prevalence of this phase of galaxy evolution at $z$\\,$\\gtapp$\\,1."
    },
    "0609/astro-ph0609699_arXiv.txt": {
        "abstract": "When taking the holographic principle into account, the vacuum energy will acquire dynamical property that its equation of state is evolving. The current available observational data imply that the holographic vacuum energy behaves as quintom-type dark energy. We adopt the viewpoint of that the scalar field models of dark energy are effective theories of an underlying theory of dark energy. If we regard the scalar field model as an effective description of such a holographic vacuum theory, we should be capable of using the scalar field model to mimic the evolving behavior of the dynamical vacuum energy and reconstructing this scalar field model according to the fits of the observational dataset. We find the generalized ghost condensate model is a good choice for depicting the holographic vacuum energy since it can easily realize the quintom behavior. We thus reconstruct the function $h(\\phi)$ of the generalized ghost condensate model using the best-fit results of the observational data. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/hep-ph0609259_arXiv.txt": {
        "abstract": "We analyze leptogenesis in a supersymmetric triplet seesaw scenario that explains the observed neutrino masses, adopting a phenomenological approach where the decay branching ratios of the triplets and the amount of CP--violation in its different decay channels are assumed as free parameters. We find that the solutions of the relevant Boltzmann equations lead to a rich phenomenology, in particular much more complex compared to the non--supersymmetric case, mainly due to the presence of an additional Higgs doublet.  Several unexpected and counter--intuitive behaviors emerge from our analysis: the amount of CP violation in one of the decay channels can prove to be be irrelevant to the final lepton asymmetry, leading to successful leptogenesis even in scenarios with a vanishing CP violation in the leptonic sector; gauge annihilations can be the dominant effect in the determination of the evolution of the triplet density up to very high values of its mass, leading anyway to a sizeable final lepton asymmetry, which is also a growing function of the wash--out parameter $K\\equiv \\Gamma_d/H$, defined as usual as the ratio between the triplet decay amplitude $\\Gamma_d$ and the Hubble constant $H$; on the other hand, cancellations in the Boltzmann equations may lead to a vanishing lepton asymmetry if in one of the decay channels both the branching ratio and the amount of CP violation are suppressed, but not vanishing. The present analysis suggests that in the supersymmetric triplet see-saw model successful leptogenesis can be attained in a wide range of scenarios, provided that an asymmetry in the decaying triplets can act as a lepton--number reservoir. ",
        "introduction": "Leptogenesis \\cite{fy} is a very appealing explanation of the baryon asymmetry of the Universe, the more so since it can also incorporate the origin of neutrino masses and mixings. In triplet see-saw models, a cosmological lepton asymmetry is produced by the decay of a heavy Higgs triplet, in presence of the Sakharov condition of i) lepton--number violation; ii) CP violation; iii) out--of--equilibrium decay. The same Higgs triplet is also responsible of the dimension--five operator that induces neutrino masses through the see-saw mechanism in the vacuum expectation values \\cite{Tss}. The lepton asymmetry is then transformed to a baryon asymmetry through the mechanism of sphaleron conversion.  This model has already been discussed in the literature with multi Higgs triplets \\cite{Tlepto} or with additional right-handed neutrinos \\cite{hybrids}. As a general feature, it presents some interesting differences compared to the usual see-saw mechanism where the decaying particles producing leptogenesis are assumed to be right--handed neutrinos. In particular, the triplet see-saw is more predictive \\cite{anna,chun03,chun05}, since Yukawa matrices can be fixed through neutrino masses, and no unknown Majorana mass matrices are present. An important feature of the triplet leptogenesis is that it suffers from a strong wash--out effect due to gauge--triplet annihilations  \\cite{Hambye, chun_scopel} which cannot be neglected, particularly, for low mass of the triplet. Even with such a strong annihilation effect, triplets can develop an appropriate lepton asymmetry during their thermal evolution, and this, in the non-supersymmetric version of the model, has been shown to lead to the possibility of a high efficiency for leptogenesis production even for low triplet masses \\cite{Hambye}. This is an intriguing aspect that enables us to circumvent the gravitino problem requiring the upper--bound on the reheating temperature of the Universe, $T_{RH} \\lsim 10^{8}$ GeV in supergravity theories.  In this paper, we want to extend the discussion of the properties of supersymmetric triplet see--saw leptogenesis to a general phenomenological scenario, with a minimal set of theoretical assumptions, and to compare it to the non--supersymmetric version of the model.  In our study, we will cover the triplet mass down to the TeV range, for which  the model has a prospect of being tested in future colliders \\cite{chun03,tev}.   The plan of the paper is as follows. In Section \\ref{section:themodel} we introduce the Higgs triplet model and the corresponding Boltzmann equations. In Section \\ref{section:discussion} we describe our minimal assumptions and then  discuss the produced lepton asymmetry in the non--supersymmetric version of the model in Section \\ref{section:non_susy}, and  in the supersymmetric one in Section \\ref{section:susy}. We devote Section \\ref{section:conclusions} to our conclusions. ",
        "conclusions": "\\label{section:conclusions} In this paper we have analyzed the phenomenology of  leptogenesis in the supersymmetric triplet see--saw mechanism. Taking the branching ratios $B_i$ of the decay rate of the triplets as free parameters, as well as the CP--violation parameters $\\epsilon_i$, with the additional constraints $\\sum_i B_1$=1 and $\\sum_i \\epsilon_i$=0, we have calculated the amount $\\bar{\\epsilon}$ of CP violation which is needed to provide successful leptogenesis.  In the most favourable case of early out-of-equilibrium decay of the triplets to leptons, this number is of order $10^{-8}$. However, it is well known that in this scenario inverse decays and annihilations of triplets (this last effect for lower values of $M$) contribute in general to erase the asymmetry, basically keeping the triplets in thermal equilibrium until late times, when $T<M$ and their number density is exponentially suppressed. An exception to this, within the framework of the non--supersymmetric version of the model, is known to be the case when one branching ratio is much smaller than the other, in such a way that one $K_i= B_i K <<$1 even if $K>>$1. We have referred to this kind of decay channel as to a slow one, opposed to the fast ones having $K_i>>1$.  In this case it is sufficient that just one slow channel produces a sizeable asymmetry, since a corresponding asymmetry is also developed in the triplet density, which is eventually converted into an asymmetry of the fast channel when the triplets decay. In the supersymmetric version of the Model, this mechanism is still at work. However, mainly because of the interplay of three decay channels instead of two, a richer phenomenology arises:   \\begin{itemize}  \\item the Yukawa coupling $\\lambda_1$ of the additional Higgs doublet can be made arbitrarily small at high triplet masses, allowing a good efficiency also for $M\\gsim 10^{12}$ GeV. In presence of only one Higgs doublet this is not possible due to the perturbativity bound on $\\lambda_L$ implied by Eq.~(\\ref{eq:neutrino_mass}).  \\end{itemize}  A hierarchy among the CP--violation parameters $\\epsilon_i$'s is allowed.  Defining them in such a way that $max(|\\epsilon_i|)= 1$, and using the notation $\\epsilon_{slow}$ and $\\epsilon_{fast}$ for the CP--violating parameters in a slow ($K_i<<1$) and fast ($K_i>>1$) channel, respectively, this enriches the phenomenology, because different combinations are possible:  \\begin{itemize}  \\item $\\epsilon_{slow}=1$: The efficiency of leptogenesis reaches its maximal value. Inverse decays in the slow channel freeze out early, and annihilations turn out to dominate over inverse decays in the determination of the triplet density up to quite high values of the triplet mass $M$. In these cases the final asymmetry is a growing function of the $K$ parameter. Moreover, the final asymmetry is insensitive to the actual value of $\\epsilon_{fast}$. An apparently surprising example of this situation is when $\\epsilon_L=\\epsilon_{fast}=0$, since in this case even a vanishing $\\epsilon_L$ can lead to efficient leptogenesis.  An exception to this case is given by the particular situation with $\\epsilon_{slow}=1$ in {\\it two} channels, namely when $\\epsilon_L<<1$ and the decay channels to $H_1$ and to $H_2$ are both slow with $\\epsilon_1=-\\epsilon_2=1$. In fact, if $B_1$ and $B_2$ are comparable, a cancellation takes place between the asymmetries in the two channels, leading to a vanishing lepton asymmetry, $Y_L=-Y_{H_1}-Y_{H_2}\\simeq 0$.    \\item $\\epsilon_{slow}<1$ and  one slow channel: The final lepton asymmetry is suppressed, as it would be na\\\"ively expected, since $\\epsilon_i$ is small in the only available slow channel which is supposed to drive leptogenesis through out-of-equilibrium decays. In this case there are two fast channels with $\\epsilon_{fast}=\\pm 1$, and in the channel where $\\epsilon_{fast}$ has the same sign as $\\epsilon_{slow}$ this may lead to a cancellation in the Boltzmann equations, implying a vanishing final asymmetry. Moreover, inverse decays freeze out late in this case ($z_f\\sim \\ln K >>1$), and decay is typically dominant over annihilation in the determination of the triplet density, except for very light values of $M$. As a consequence of this the efficiency scales as $1/(z_f K)$ whenever $K>>1$  \\item $\\epsilon_{slow}<1$ and two slow channels:  Since only one $\\epsilon_i$ can be small, the other slow channel with unsuppressed $\\epsilon_i$ may drive leptogenesis with a good efficiency. In this case, in practice the decay channel with $\\epsilon_{slow}<1$ drops out from the Boltzmann equation, and a system with just two decay channels is recovered. However, if $\\epsilon_{slow}=\\epsilon_{L,2}$, the phenomenology is different compared to the non--supersymmetric case, because the $K$ parameter is much bigger, reducing the efficiency at high masses and improving it at lower ones. Moreover, the unitarity limit is more constraining at high $M$ compared to the non-supersymmetric case.  \\end{itemize}  In conclusion, the present analysis suggests that in the supersymmetric Triplet Seesaw Model successful leptogenesis can be attained in a wide range of scenarios, some of which appear to be non trivial or even counter--intuitive, provided that an asymmetry in the decaying triplets can develop at early times and be eventually converted into a lepton asymmetry, acting in practice as a lepton--number reservoir."
    },
    "0609/hep-ph0609290_arXiv.txt": {
        "abstract": "{We point out that if the lightest supersymmetric particle (LSP) is a Higgsino- or Wino-like neutralino, the net effect of coannihilations with sleptons is to {\\em increase} the relic abundance, rather than producing the usual suppression, which takes place if the LSP is Bino-like. The reason for the enhancement lies in the effective thermally averaged cross section at freeze-out: sleptons annihilate (and co-annihilate) less efficiently than the neutralino(s)-chargino system, therefore slepton coannihilations effectively act as {\\em parasite degrees of freedom} at freeze-out. Henceforth, the thermal relic abundance of LSP's corresponds to the cold Dark Matter abundance for {\\em smaller} values of the LSP mass, and {\\em larger} values of the neutralino pair annihilation cross section. In turn, at a given thermal neutralino relic abundance, this implies larger indirect detection rates, as a result of an increase in the fluxes of antimatter, gamma rays and neutrinos from the Sun originating from neutralino pair annihilations.} ",
        "introduction": "\\label{sec:intro}  The fundamental nature of non-baryonic Dark Matter is one of the most pressing topics in contemporary Particle Physics. While Dark Matter is invoked in most cosmological models, and its existence, essentially inferred through gravitational effects, distribution and total abundance on cosmological scales are quantitatively established with an increasing and remarkable accuracy, little is known about what is the elementary constituent of this elusive and yet so substantial component of the cosmic budget \\cite{dmreviews}. Dark Matter candidates have been proposed in several extensions of the Standard Model of particle physics, or in purely phenomenological settings, motivated, for instance, by astrophysical observations with no obvious known source counterparts; the range of particle masses proposed in the literature spans over various decades of orders of magnitude (see, {\\em e.g.}, \\cite{Baltz:2004tj}).  A sensible rationale to distinguish among different Dark Matter candidates emerges in the nature of the process invoked to explain why a given candidate should have a relic abundance close to the abundance of Dark Matter we infer from observations today. Along this line of reasoning, weakly interacting massive particles (WIMPs) stand as excellent prototypes. Once in thermal equilibrium with the primordial particles thermal bath in the Early Universe, WIMPs undergo a {\\em freeze-out} of their number density as their pair annihilation rate becomes smaller than the Universe expansion rate. After decoupling, the WIMP number density per comoving volume remains substantially constant up to the present epoch;  their final  relic abundance crucially depends upon the WIMP pair annihilation rate. Be it a fortuitous coincidence or not, when the WIMP pair annihilation cross section is comparable to a typical weak-interactions cross section, the estimated WIMP relic abundance $\\Omega_\\chi h^2$ is remarkably close to the actual inferred Dark Matter abundance, $\\Omega_{\\rm CDM}h^2\\simeq 0.110$ \\cite{Spergel:2006hy} (here and in the remainder of the paper, $\\Omega_i$ indicates the ratio of the mean density of species $i$ over the critical density, and $h$ the normalized value of the Hubble expansion rate today, in units of 100 km/s/Mpc).  The detailed dynamics of WIMPs freeze-out depends crucially, however, on the specific nature of the particle physics setup at hand. For instance, the relic abundance of the lightest neutralino in minimal supersymmetric extensions of the Standard Model (MSSM), a paradigmatic WIMP, varies over several orders of magnitude, questioning whether achieving $\\Omega_\\chi\\simeq\\Omega_{\\rm CDM}$ is in fact ``natural'' at all. Further, the chemical decoupling of WIMPs in the Early Universe can be complicated by the concomitant, and possibly {\\em coupled}, decoupling of other particle species. The occurrence of the latter scenario, known in the jargon as {\\em coannihilation} \\cite{coan,Griest:1990kh}, can lead to very significant effects on the final WIMP relic abundance.  The net effect of coannihilations crucially depends upon the size of the thermally averaged annihilation and coannihilation cross sections of the extra degrees of freedom participating to the freeze-out of the lightest, stable species, averaged over the total number of degrees of freedom (see the next section for a more detailed and quantitative description of coannihilations). On top of this, since the abundance of non-relativistic species in thermal equilibrium approximately follows a Maxwell-Boltzmann distribution, the importance of coannihilation processes is exponentially suppressed by the relative mass splitting between the coannihilating particles's masses and the stable particle's mass. Evidently, depending upon the particle physics setup, coannihilations can be responsible for both an {\\em increase} or for a {\\em decrease} in the final relic abundance of the stable species.  In the widely studied context of minimal supergravity (mSUGRA)~\\cite{msugra}, a special realization of the MSSM, coannihilations are often thought to be synonym of {\\em suppression} of the neutralino relic abundance. In mSUGRA the neutralino is Bino-like ({\\em i.e.} the lightest mass eigenstate almost coincides with the fermionic superpartner of the hypercharge gauge boson) over most of the parameter space of the theory; Binos pair annihilate rather inefficiently in the Early Universe, as they feature a suppressed coupling to gauge bosons, and the pair annihilation into fermion-antifermion states is helicity suppressed. As noted in numerous publications, one of the few regions of the mSUGRA parameter space where $\\Omega_\\chi\\simeq\\Omega_{\\rm CDM}$ is possible lies at low values of the universal scalar soft supersymmetry breaking parameter $m_0$, where the lightest neutralino is close in mass to the lightest stau \\cite{elliscoan,nihei}. There, stau coannihilations enhance the effective stau-neutralino annihilation rate around neutralino freeze-out. Since staus annihilate {\\em more efficiently} than Binos, the final Bino relic abundance is {\\em lower} than without stau coannihilations, and can be such that $\\Omega_\\chi\\simeq\\Omega_{\\rm CDM}$. Slepton coannihilations where first explicitely studied in Ref.~\\cite{elliscoan}, where the (co-)annihilation cross section was approximated, in the low-velocity expansion limit, in powers of the mass-over-temperature ratio. A more accurate calculation, that takes into account the role of slepton mixing and the exact computation of the effective pair annihilation cross section, was then presented in Ref.~\\cite{nihei}  Ref.~\\cite{Edsjo:2003us} gave a nice example, again in the context of the mSUGRA paradigm, where slepton coannihilations {\\em increase} the Bino relic abundance: if Binos resonantly annihilate through the $s$-channel resonant exchange of a heavy Higgs, adding coannihilating slepton degrees of freedom makes the total effective cross section, averaged above all degrees of freedom, smaller than without coannihilations. When Binos annihilate efficiently, slepton act as {\\em parasite degrees of freedom} at the lightest supersymmetric particle (LSP) freeze-out.  Another illustrative example of the role of coannihilations in {\\em enhancing} the final relic abundance of the stable particle comes from Universal Extra Dimensions (UED) \\cite{Appelquist:2000nn}. In UED, the particle mass spectrum of Kaluza-Klein (KK) states is naturally highly degenerate, lying around a mass scale set by the inverse compactification scale radius $R^{-1}$. Coannihilations of the stable lightest KK particle (LKP) are therefore expected to play an important role. As pointed out in Ref.~\\cite{Servant:2002aq} and in subsequent refined analyses \\cite{Kong:2005hn}, for realistic spectra the effect of coannihilations is to significantly {\\em increase} the relic abundance of the LKP, corresponding, in those setups, to the $B^{(1)}$ (to a good approximation the $n=1$ KK excitation of the hypercharge gauge boson). This results in a reduction of the value of $R^{-1}$ such that $\\Omega_{B^{(1)}}\\simeq\\Omega_{\\rm CDM}$ by factors as large as 2, depending upon the details of the KK spectrum. Coannihilations with KK states featuring a smaller annihilation cross section than that of the $B^{(1)}$ itself, and close in mass to it, such as right-handed KK leptons, are responsible for this effect.  In the present note we point out that, unlike the generic case of a Bino (barring fortuitous resonant annihilation channels), when the LSP is dominated by its Higgsino or Wino components, {\\em i.e.} when the lightest mass eigenstate approximately corresponds to the fermionic superpartner of the neutral Higgses or of the SU(2) neutral gauge boson, the effect of slepton coannihilations is to {\\em increase} the LSP relic abundance. The increase in $\\Omega_\\chi$ depends upon various circumstances (next-to-lightest neutralino and/or chargino coannihilations, sizable couplings to gauge bosons) that contribute to make the total effective Higgsino and Wino annihilation cross sections larger than that in presence of slepton coannihilations. As a result, the mass $m_\\chi$ of neutralinos such that $\\Omega_\\chi\\simeq\\Omega_{\\rm CDM}$ is pushed to smaller values, and the pair annihilation cross sections $\\langle\\sigma v\\rangle$ to larger ones. In turn, this implies {\\em larger indirect Dark Matter detection rates}, as the latter are in general proportional to the combination $\\langle\\sigma v\\rangle/m_\\chi^2$. We show that the occurrence of slepton coannihilations in scenarios where the LSP is Higgsino or Wino-like is perfectly viable, and actually takes place in several well motivated theoretical setups, where the induced degree of ``fine-tuning'' is generically not larger than that invoked in the context of the stau coannihilation region of mSUGRA.  We start our analysis with a quantitative discussion of coannihilation processes, which leads us to a guiding analytical formula. We then focus on particular phenomenological MSSM setups, motivated by several examples of GUT scale completions which can lead to similar spectra, and determine the combinations of LSP masses and mass splitting between the LSP and the coannihilating sleptons such that the thermal neutralino relic abundance saturates the Dark Matter abundance (sec.~\\ref{sec:thermal}). Finally, we make use of the models featuring $\\Omega_\\chi\\simeq\\Omega_{\\rm CDM}$ with slepton coannihilations to estimate the resulting enhancement in various indirect Dark Matter detection rates (sec.~\\ref{sec:detect}), and summarize our conclusions (sec.~\\ref{sec:conclude}). ",
        "conclusions": "\\label{sec:conclude}  In this paper we studied the effects of slepton coannihilations on the thermal relic abundance of  Higgsino- or Wino-like lightest neutralinos. We pointed out that, unlike the well known case of a Bino-like neutralino, coannihilations with sleptons yield a larger Higgsino and Wino relic abundance. The effect on the relic abundance amounts to an increase of a factor ranging from a few percent up to 5. Requiring that the neutralino relic abundance lies in the range of values preferred for the abundance of Dark Matter entails, in presence of slepton coannihilations, a reduced mass for Winos and Higgsinos, and a larger pair annihilation cross section. Quantitatively, we find that the neutralino mass can be reduced up to a factor between 2 and 3, depending on the particular setup at hand. We showed that smaller values of the neutralino mass and larger pair annihilation cross sections produce potentially very large enhancements in the rates expected in indirect Dark Matter search experiments. In some cases, we showed that the occurrence of slepton coannihilations and the resulting reduction of the neutralino mass needed to produce the right amount of relics is crucial to produce signals that might allow to indirectly probe the occurrence of galactic Dark Matter annihilations."
    },
    "0609/astro-ph0609136_arXiv.txt": {
        "abstract": "Evidence for orbital modulation of the very high energy (VHE) $\\gamma$-ray emission from the high-mass X-ray binary and microquasar LS~5039 has recently been reported by the HESS collaboration. The observed flux modulation was found to go in tandem with a change in the GeV -- TeV spectral shape, which may partially be a result of $\\gamma\\gamma$ absorption in the intense radiation field of the massive companion star. However, it was suggested that $\\gamma\\gamma$ absorption effects alone can not be the only cause of the observed spectral variability since the flux at $\\sim 200$~GeV, which is near the minimum of the expected $\\gamma\\gamma$ absorption trough, remained essentially unchanged between superior and inferior conjunction of the binary system. In this paper, a detailed parameter study of the $\\gamma\\gamma$ absorption effects in this system is presented. For a range of plausible locations of the VHE $\\gamma$-ray emission region and the allowable range of viewing angles, the de-absorbed, intrinsic VHE $\\gamma$-ray spectra and total VHE photon fluxes and luminosities are calculated and compared to luminosity constraints based on Bondi-Hoyle limited wind accretion onto the compact object in LS~5039. Based on these arguments, it is found that (1) it is impossible to choose the viewing angle and location of the VHE emission region in a way that the intrinsic (deabsorbed) fluxes and spectra in superior and inferior conjunction are identical; consequently, the intrinsic VHE luminosities and spectral shapes must be fundamentally different in different orbital phases, (2) if the VHE luminosity is limited by wind accretion from the companion star and the system is viewed at an inclination angle of $i \\gtrsim 40^o$, the emission is most likely beamed by a larger Doppler factor than inferred from the dynamics of the large-scale radio outflows, (3) the still poorly constrained viewing angle between the line of sight and the jet axis is most likely substantially smaller than the maximum of $\\sim 64^o$ inferred from the lack of eclipses. (4) Consequently, the compact object is more likely to be a black hole rather than a neutron star. (5) There is a limited range of allowed configurations for which the expected VHE neutrino flux would actually anti-correlate with the observed VHE $\\gamma$-ray emission. If hadronic models for the $\\gamma$-ray production in LS~5039 apply, a solid detection of the expected VHE neutrino flux and its orbital modulation with km$^3$ scale water-Cherenkov neutrino detectors might require the accumulation of data over more than 3 years. ",
        "introduction": "Introduction}  The recent detections of VHE ($E \\gtrsim 250$~GeV) $\\gamma$-rays from the high-mass X-ray binary jet sources LS 5039 with the High Energy Stereoscopic System \\citep[HESS;][]{aharonian05} and LS~I~+63$^o$303 with the Major Atmospheric Gamma-Ray Imaging Cherenkov Telescope \\citep[MAGIC][]{albert06} establish these sources (termed ``microquasars'' if they are accretion powered) as a new class of $\\gamma$-ray emitting sources. These results confirm the earlier tentative identification of LS~5039 with the EGRET source 3EG~J1824-1514 \\citep{paredes00} and LSI~$61^o303$ with the COS~B source 2CG~135+01 \\citep{gregory78,taylor92} and the EGRET source 3EG~J0241+6103 \\citep{kniffen97}. Both of these objects show evidence for variability of the VHE emission, suggesting an association with the orbital period of the binary system. In the case of LS~I~+63$^o$303, the association with the orbital period is not yet firmly established since the MAGIC observations covered only a few orbital periods, and the orbital period \\citep[$P = 26.5$~d;][]{gregory02} is very close to the siderial period of the moon, which also sets a natural windowing period for VHE observations \\citep{albert06}. In contrast, the HESS observations of LS~5039 provide rather unambiguous evidence for an orbital modulation of both the VHE $\\gamma$-ray flux and spectral shape with the orbital period of $P = 3.9$~d \\citep{aharonian06b}. Specifically, \\cite{aharonian06b} found that between inferior conjunction (i.e., the compact object being located in front of the companion star), the VHE $\\gamma$-ray spectrum could be well fitted with an exponentially cut-off, hard power-law of the form $\\Phi_E \\propto E^{-1.85} \\, e^{-E/E_0}$ whith a cut-off energy of $E_0 = 8.7$~TeV. In contrast, the VHE spectrum at superior conjunction is well represented by a pure power-law ($\\Phi_E \\propto E^{-2.53}$) with a much steeper slope, but identical differential flux at $E_{\\rm norm} \\approx 200$~GeV. The original data from \\cite{aharonian06b}, together with these spectral fit functions, are shown in Fig. \\ref{observed_spectrum}.  A variety of different models for the high-energy emission from high-mass X-ray binaries have been suggested. These range from high-energy processes in neutron star magnetospheres \\citep[e.g.,][]{moskalenko93,moskalenko94,bednarek97,bednarek00,chernyakova06,dubus06b}, via models with the inner regions of microquasar jets being the primary high-energy emission sites \\citep[e.g.][]{romero03,bp04,bosch05a,gupta06,db06,gb06} and interactions of microquasar jets with the ISM \\citep{bosch05b}, to models involving particle acceleration in shocks produced by colliding stellar winds \\citep[e.g.,][]{reimer06}. In addition to involving a variety of different leptonic and hadronic emission processes, these models also imply vastly different locations of the $\\gamma$-ray production site with respect to the compact object in LS~5039 and the companion star. However, independent of the emission mechanism responsible for the VHE $\\gamma$-rays, the intense radiation field of the high-mass (stellar type O6.5V) companion will lead to $\\gamma\\gamma$ absorption of VHE $\\gamma$-rays in the $\\sim 100$~GeV -- TeV photon energy range. The characteristic features of the $\\gamma\\gamma$ absorption of VHE $\\gamma$-rays by the companion star light in LS~5039 have been investigated in detail by \\cite{bd05} and \\cite{dubus06a}. It was found that, if the $\\gamma$-ray emission originates within a distance of the order of the orbital separation of the binary system ($s \\sim 2 \\times 10^{12}$~cm), this effect should lead to a pronounced $\\gamma\\gamma$ absorption trough, in particular near superior conjunction, while $\\gamma\\gamma$ absorption tends to be almost negligible near inferior conjunction. In LS~5039, the minimum of the absorption trough is expected to be located around $E_{\\rm min} \\sim 300$~GeV and should shift from higher to lower energies as the orientation of the binary system changes from inferior to superior conjunction.  Most of the relevant parameters of LS~5039, except for the inclination angle $i$ of the line of sight with respect to the normal to the orbital plane, are rather well determined (see \\S \\ref{parameters}). This allows for a detailed parameter study of the $\\gamma\\gamma$ absorption effects, leaving the inclination angle and the distance of the VHE emission site from the compact object as free parameters. As will be shown in \\S \\ref{cascades}, the effect of electromagnetic cascades will lead to a re-deposition of the absorbed $\\gtrsim 100$~GeV luminosity almost entirely at photon energies $E \\lesssim 100$~GeV. For that reason, one can easily correct for the $\\gamma\\gamma$ absorption effect at energies $E \\gtrsim 100$~GeV by multiplying the observed fluxes by a factor $e^{\\tau_{\\gamma\\gamma}(E)}$ (where $\\tau_{\\gamma\\gamma}(E)$ is the $\\gamma\\gamma$ absorption depth along the line of sight) in order to find the intrinsic, deabsorbed VHE spectra for any given choice of $i$ and the height $z_0$ of the VHE $\\gamma$-ray emission site above the compact object. Results of this procedure will be presented in \\S \\ref{deabsorbed}. In \\S \\ref{constraints}, these results will then be used to estimate the total flux and luminosity in VHE $\\gamma$-rays, which can be compared to limits on the available power under the assumption that the VHE emission is powered by Bondi-Hoyle limited wind accretion onto the compact object. This leads to important constraints on the location of the VHE emision site and the geometry of the system, which will be discussed in \\S \\ref{summary}. ",
        "conclusions": ""
    },
    "0609/astro-ph0609300_arXiv.txt": {
        "abstract": "It is common to estimate black hole abundances by using a measured correlation between black hole mass and another more easily measured observable such as the velocity dispersion or luminosity of the surrounding bulge.  The correlation is used to transform the distribution of the observable into an estimate of the distribution of black hole masses. However, different observables provide different estimates: the $M_\\bullet-\\sigma$ relation predicts fewer massive black holes than does the $M_\\bullet-L$ relation.  This is because the $\\sigma-L$ relation in black hole samples currently available is inconsistent with that in the SDSS sample, from which the distributions of $L$ or $\\sigma$ are based:  the black hole samples have smaller $L$ for a given $\\sigma$ or have larger $\\sigma$ for a given $L$.  This is true whether $L$ is estimated in the optical or in the NIR. The $\\sigma-L$ relation in the black hole samples is similarly discrepant with that in the ENEAR sample of nearby early-type galaxies.  This suggests that current black hole samples are biased towards objects with abnormally large velocity dispersions for their luminosities.  If this is a selection rather  than physical effect, then the $M_\\bullet-\\sigma$ and $M_\\bullet-L$ relations currently in the literature are also biased from their true values.  We provide a framework for describing the effect of this bias.  We then combine it with a model of the bias to make an estimate of the true intrinsic relations.  While we do not claim to have understood the source of the bias, our simple model is able to reproduce the observed trends.  If we have correctly modeled the selection effect, then our analysis suggests that the bias in the $\\langle M_\\bullet|\\sigma\\rangle$ relation is likely to be small, whereas the $\\langle M_\\bullet|L\\rangle$ relation is biased towards predicting more massive black holes for a given luminosity. In addition, it is likely that the $M_\\bullet-L$ relation is entirely a consequence of more fundamental relations between $M_\\bullet$ and $\\sigma$, and between $\\sigma$ and $L$.  The intrinsic relation we find suggests that at fixed luminosity, older galaxies tend to host more massive black holes. ",
        "introduction": "Several groups have noted that galaxy formation and supermassive black hole growth should be related, and many have modeled the joint evolution of quasars and galaxies (e.g., Monaco et al. 2000; Kauffmann \\& Haehnelt 2001; Granato et al. 2001; Cavaliere \\& Vittorini 2002; Cattaneo \\& Bernardi 2003; Haiman et al. 2004; Hopkins et al. 2004; Lapi et al. 2006; Haiman et al. 2006 and references therein).  However, the number of black hole detections to date is less than fifty; this small number prevents a direct estimate of the black hole mass function.  On the other hand, correlations between $M_\\bullet$ and other observables can be measured quite reliably, if one assumes that the $M_\\bullet-$observable relation is a single power-law.  $M_\\bullet$ is observed to correlate strongly and tightly with the velocity dispersion of the surrounding bulge (Ferrarese \\& Merritt 2000; Gebhardt et al. 2000; Tremaine et al. 2002), as well as with bulge luminosity (McLure \\& Dunlop 2002) and bulge stellar mass (Marconi \\& Hunt 2003; H\\\"aring \\& Rix 2004): in all cases, a single power law appears to be an adequate description.  Since it is considerably easier to measure bulge properties than $M_\\bullet$, black hole abundances are currently estimated by transforming the observed abundances of such secondary indicators using the observed correlations with $M_\\bullet$ (Yu \\& Tremaine 2002; Aller \\& Richstone 2002; Marconi et al. 2004; Shankar et al. 2004; McLure \\& Dunlop 2004). Unfortunately, combining the $M_\\bullet-\\sigma$ relation with the observed distribution of velocity dispersions results in an estimate of the number density of black holes with $M_\\bullet>10^9M_\\odot$ which is considerably smaller than the estimate based on combining the $M_\\bullet-L$ relation with the observed luminosity function (Tundo et al. 2006).  The fundamental cause of the discrepancy is that the luminosity and velocity dispersion functions are obtained from the SDSS (Blanton et al. 2003; Sheth et al. 2003), but the $\\sigma-L$ relation in the SDSS (Bernardi et al. 2003b) is different from that in the black hole samples (e.g. Yu \\& Tremaine 2002). In Section~\\ref{why}, we argue that the SDSS $\\sigma-L$ relation is {\\em not} to blame, suggesting that it is the black hole samples which are biased.  \\begin{figure*} \\centering \\includegraphics[width=1.9\\columnwidth]{fig1.eps} \\caption{Comparison of the $\\sigma-L$ relation in the black hole sample of H\\\"aring \\& Rix (2004) (symbols) with that from local early-type galaxy samples (lines). We provide two estimates of the bulge luminosity for the H\\\"aring \\& Rix compilation: the open circles come from rescaling the bulge luminosity to the SDSS r-band, and the filled circles from assuming a mass-to-light ratio to scale the bulge masses $M_{\\rm bulge}$ to luminosities. Dotted line shows a fit to the filled circles. Solid and short dashed lines show $\\langle\\sigma|L\\rangle$ in the SDSS-B06 and ENEAR samples. Dot-dashed line shows the biased relation obtained from the ENEAR sample if one fails to account for the fact that the velocity dispersions played an important role in determining the distances from which the luminosities were estimated. The black hole compilation is clearly offset from the SDSS-B06, the ENEAR and even the biased ENEAR relations. Appendix~A gives the fits shown, and describes exactly how the black hole sample was compiled.  } \\label{sigmaLv} \\end{figure*}  If black hole samples are indeed biased, then this bias may be physical or it may simply be a selection effect.  If physical, then only a special set of galaxies host black holes, and the entire approach above (of transforming the luminosity or velocity dispersion function) is compromised.  On the other hand, if it is a selection effect, then one is led to ask if the $M_\\bullet-$observable relations currently in the literature are biased by this selection or not.  The main purpose of this work is to provide a framework for describing these biases. Section~\\ref{biases} describes Monte-Carlo simulations of the selection effect.  It shows an example of the bias which results, and provides an analytic model of the effect. Our findings are summarized in Section~\\ref{sowhat}, where estimates of the intrinsic black hole mass function $\\phi(M_\\bullet)$, and the intrinsic $M_\\bullet-\\sigma$ and $M_\\bullet-L$ relations are provided.  Throughout, we assume a spatially flat background cosmology with $\\Omega_0=0.3$ and a cosmological constant.  All luminosities and masses have been scaled to $H_0 = 70$~km~s$^{-1}$~Mpc$^{-1}$. ",
        "conclusions": "\\label{sowhat} Compared to the ENEAR and SDSS-B06 early-type galaxy samples, the $\\sigma-L$ correlation in black hole samples currently available is biased towards large $\\sigma$ for a given $L$ (Figures~\\ref{sigmaLv}--\\ref{sigmaLK}). If this is a selection effect rather than a physical one, then current determinations of the $M_\\bullet-L$ and $M_\\bullet-\\sigma$ relations are biased.  We provided a simple toy model of the effect which shows how these biases depend on the intrinsic correlations and on the selection effect.  To illustrate its use, we constructed simple models of the intrinsic correlation and of the selection effect.  We then identified a particular set of parameters which we showed can reproduce all the observed trends:  the correlations between $M_\\bullet$ and $L$ or $\\sigma$, the $\\sigma-L$ correlation in black hole samples, as well as correlations between residuals from these fits (Figures~\\ref{mockRels} and~\\ref{SIMresids}) are all reproduced. We do not claim to have found the unique solution to this problem. However, if we have correctly modeled the selection effect, then our analysis suggests that the bias in the $\\langle M_\\bullet|\\sigma\\rangle$ relation is likely to be small, whereas the $\\langle M_\\bullet|L\\rangle$ relation is biased towards having a steeper slope (Figure~\\ref{mockRels}). This has the important consequence that estimates of black hole abundances which combine the $\\langle M_\\bullet|L\\rangle$ relation and its scatter with an external determination of the distribution of $L$ will overpredict the true abundances of massive black holes. The analogous approach based on $\\sigma$ is much less biased (Figure~\\ref{mockCounts}).  The choice of parameters which reproduces the selection-biased fits reported in Appendix~A has intrinsic distribution $\\phi(M_\\bullet)$ that is well described \\begin{equation} \\phi(M_\\bullet)\\,{\\rm d}M_\\bullet = {\\phi_*\\over \\Gamma(\\alpha/\\beta)}\\, \\mu^\\alpha\\,\\exp(-\\mu^\\beta)\\,\\beta\\,{{\\rm d}M_\\bullet\\over M_\\bullet} \\label{phiMbh} \\end{equation} where $\\phi_* = 0.002/{\\rm Mpc}^{3}$, $(\\alpha,\\beta) = (2.0,0.3)$ and $\\mu = M_\\bullet/M_\\bullet^*$, with $M_\\bullet^*$ a characteristic mass. This mass is related to the mean mass $10^{8}M_\\odot$ by $\\langle M_\\bullet\\rangle = M_\\bullet^*\\,\\Gamma[(\\alpha+1)/\\beta]/\\Gamma(\\alpha/\\beta)$, so $M_\\bullet^* = 10^{5}\\,M_\\odot$.  Since this $\\phi(M_\\bullet)$ is based on the SDSS-B06 luminosity and velocity functions, it is missing small bulges, so it is incomplete at small $M_\\bullet$.  The intrinsic correlations in the sample are \\begin{equation} \\Bigl\\langle M_\\bullet|M_r\\Bigr\\rangle = 9.16 - 0.41\\,(M_r+22)\\\\ \\end{equation} and \\begin{equation} \\Bigl\\langle M_\\bullet|\\log\\sigma\\Bigr\\rangle = 8.85 + 3.82\\, \\log\\, (\\sigma/200~{\\rm km~s}^{-1}) \\end{equation} with scatter of 0.38 and 0.23~dex, respectively. A bisector fit to the $M_\\bullet-L$ relation would have slope $(\\sigma_\\bullet/\\sigma_L) (r_{\\bullet L} + 1/r_{\\bullet L})/2 = -0.62$; a similar fit to the $M_\\bullet-V$ relation would have slope $(\\sigma_\\bullet/\\sigma_V) (r_{\\bullet V} + 1/r_{\\bullet V})/2 = 4.49$. Notice that the slopes satisfy $a_{\\bullet|L}\\approx a_{\\bullet|V}\\,a_{V|L}$. This is because, for this choice of parameters, the correlation between $M_\\bullet$ and $L$ is almost entirely a consequence of the other two correlations.  If we have correctly modelled the selection effect, and we do not claim to have done so, then our analysis suggests that the intrinsic joint distribution of $M_\\bullet$, luminosity and $\\sigma$ is similar to that between color, luminosity and $\\sigma$: of the three pairwise correlations, the two correlations involving $\\sigma$ are fundamental, whereas the third is not (Bernardi et al. 2005).  Since residuals from the $\\langle\\sigma|L\\rangle$ relation are correlated with age (Bernardi et al. 2005), the correlations shown in the top right panels of Figures~\\ref{SIMresids} and~\\ref{BHresids} suggest that, at fixed luminosity, older galaxies tend to host more massive black holes than less massive galaxies.  Understanding why this is true may provide important insight into black hole formation and evolution."
    },
    "0609/astro-ph0609593_arXiv.txt": {
        "abstract": "A cosmology calculator that computes times and distances as a function of redshift for user-defined cosmological parameters is available at http://www.astro.ucla.edu/$\\sim$wright/CosmoCalc.html. This note gives the formulae used by the cosmology calculator and discusses some of its implementation. A version of the calculator that allows one to specify the equation of state parameter $w$ and $w^\\prime$ is at ACC.html, and a version for converting the light travel times usually given in the popular press into redshifts is at DlttCalc.html. ",
        "introduction": "There are many definitions of distances in cosmology, and it is often frustrating to calculate the kind of distance needed for a given situation.  In particular, press releases (Clavin, 2006) very often use the {\\em light travel time distance}, $D_{ltt} = c(t_\\circ-t_{em})$, even though this distance has very many undesirable characteristics.  I wrote the cosmology calculator in order to be able to calculate $D_{ltt}$ along with the observable angular and luminosity distances $D_A$ and $D_L$. It also computes the proper radial distance which is the only distance that is compatible with the common definition of distance for moving objects:  the spatial separation at a common time.  The common time used is now, or $t_\\circ$.  This it is called $D_{now}$ below. $D_{now}$, also known as the comoving radial distance, is not measurable from a single position, but it can in principle be measured by a family of comoving observers.  It is the one distance that satisfies the Hubble Law exactly: $dD/dt = H(t)D$ without any corrections.  Other distances are defined using quantities taken from more than one time and generally only satisfy the Hubble Law to first order. ",
        "conclusions": "}  These cosmology calculators are suitable for interactive use, providing fairly quick answers for single cases.  Users who wish to use the code for large scale calculations should translate it into a compiled language and change the quadrature formula.  In particular, if fitting to datasets with redshifts and distance, the data should be sorted by redshift and the distance integrals evaluated for all objects in a single pass through the sorted data.  Plotting figures is an obvious case where redshifts are computed in order: Figure \\ref{fig:Dsvsz} shows the distance measures discussed in this paper as a function of $z$ for 3 different models.  The formulae presented here were translated into Postscript for this figure, resulting in a 4 kB file."
    },
    "0609/astro-ph0609070_arXiv.txt": {
        "abstract": "We have performed a wide-field imaging survey of the globular cluster (GC) populations around M87 with Suprime-Cam on the 8.2m Subaru Telescope. A $2^{\\circ} \\times 0_{\\cdot}^{\\circ}5$ field extending from the centre of M87 out to a radius of $\\sim$ 0.5 Mpc was observed through the $BVI$ filters. By investigating the GC colour distribution as a function of distance from M87 and NGC 4552, another luminous Virgo elliptical in our survey field, it is found that clear bimodality ($(V-I)_{\\rm peak} \\sim 1.0$ and 1.2) is seen only in the inner regions ($\\lesssim 10$ kpc) of the host galaxies and that it becomes weaker with radius due to the decreasing contribution of the red GC ($V-I > 1.1$) subpopulation. It is also found (both around M87 and NGC 4552) that while the spatial distribution of the red GCs is as centrally concentrated as the host galaxy halo light distribution, the distribution of the blue GCs ($V-I \\leq 1.1$) tends to be more extended. However, the blue GC distribution around M87 is not as extended as the Virgo cluster mass density profile. Based on these facts, we argue that the contribution of GCs associated with the Virgo cluster (e.g., intergalactic GCs) is not significant around M87 and is not the primary reason for the high $S_N$ value of M87. Instead, most of the blue GCs around luminous ellipticals, as well as the red GCs, are presumed to be associated with the host galaxy.  We model the radial profile of GC surface densities out to $\\sim$ 0.5 Mpc from M87 by a superposition of the GC populations associated with M87 and with NGC 4552. It is found that there are some regions where the GC surface densities are larger than expected from this model, suggesting the existence of an additional intergalactic GC (i-GC) population independent of the luminous ellipticals. By comparing the estimated i-GC surface density with the luminosity density of the intracluster stellar population inferred from the intergalactic planetary nebulae in the Virgo cluster, we obtain a crude estimate of i-GC specific frequency $S_N = 2.9^{+4.2}_{-1.5}$. If this $S_N$ value represents the stellar population tidally stripped by a massive central galaxy from other (less luminous) galaxies, the contribution of tidally captured GCs in the GC population of M87 would need to be low to be consistent with the high $S_N$ value of M87. ",
        "introduction": "Globular cluster (GC) populations are one of the key probes of the archeology of their host galaxy because GCs contain stars with a single age and metallicity and are therefore in principle simpler to interpret than the integrated light of field stars. One of the most exciting recent developments in the study of extragalactic GCs is the discovery that many luminous elliptical galaxies have bimodal or multimodal colour distributions of GCs (e.g., Gebhardt \\& Kissler-Patig 1999; Larsen et al. 2001; Brodie et al. 2005). Since this is unlikely to be produced in a model of elliptical galaxy formation with a simple monolithic collapse and subsequent star formation, more elaborate models have been proposed. One is the multiphase collapse scenario (Forbes, Brodie \\& Grillmair 1997), where all GCs are presumed to form coevally with the host galaxy in a massive burst of star formation within a short time scale at high redshift and are therefore expected to be uniformly old. Since the starburst is presumed to be split into several discrete phases, some of the GCs will form in gas clouds more polluted by earlier star formation, which could result in a red GC subpopulation. Another proposed model is the merger scenario (Ashman \\& Zepf 1992), where luminous ellipticals are presumed to have formed via gas-rich galaxy mergers which induce the formation of additional red GCs, while the blue ones originate in the progenitor galaxies. In this case, the blue and red GC subpopulations would contain predominantly old/metal-poor GCs and younger/metal-rich GCs, respectively. Beasley et al. (2002) found that GC bimodality was a general characteristic of hierarchical merging models in which the formation of blue GCs was truncated at high redshift.  However, most of the elliptical galaxies whose GC populations have so far been intensively studied reside in high density environments. The GC properties upon which the proposed scenarios are based may therefore not be intrinsic but may have been substantially modified by environmental effects. Field stars in the outer parts of galaxies can be stripped off into intergalactic space due to tidal interactions between galaxies and the cluster potential, and as a result of high-speed encounters between galaxies. In fact, observations have now provided direct evidence for such an intracluster stellar population in the form of planetary nebulae (PNe) (e.g., Theuns \\& Warren 1997; Feldmeier et al. 2004a; Arnaboldi et al. 2004), red giant branch stars (Durrell et al. 2002), and diffuse light (e.g., Feldmeier et al. 2002, 2004b). It is believed that GCs may also be stripped from galaxies by tidal interactions, suggesting that intergalactic space may harbour a potential reservoir of GCs which needs to be taken into account to understand the evolution of GC populations. In an alternative explanation of the bimodality of GCs known as the accretion scenario (C\\^{o}t\\'{e}, Marzke, \\& West 1998), the blue (metal-poor) GCs are assumed to have been captured from other (presumably less luminous) galaxies through tidal stripping and/or accretion. A population of intergalactic GCs (``i-GC'' hereafter) may naturally explain the fact that GCs tend to be extremely populous around central cluster galaxies, because they would be most likely to be surrounded by a significant population of i-GCs (West et al. 1995). By surveying directly for i-GCs and characterizing this population, it is possible to test the validity of the accretion scenario and improve our understanding of the formation of GC populations. Furthermore, the i-GC specific frequency could be estimated by comparing the luminosity of the intracluster stellar population with the total number of i-GCs. Since GC specific frequency is considered to depend on galaxy morphology, this would give an insight into the parent galaxy type and luminosity of i-GCs and the physical processes of galaxy interactions.  To date, however, there is no firm evidence for or against the existence of i-GCs (Bassino et al. 2003; Jord\\'{a}n et al. 2003; Mar\\'{i}n-Franch \\& Aparicio 2003). In fact, several technical difficulties must be overcome to find i-GCs. Firstly, the surface number density of i-GCs is inferred to be highest around the cluster centre because tidally stripped GCs are expected to be distributed like the mass distribution of the host {\\it galaxy cluster} (White 1987; Muzzio 1987; West et al. 1995), but a cD galaxy or giant elliptical is located at the cluster centre in many cases. Since luminous ellipticals are normally surrounded by thousands of GCs, a large area needs to be surveyed to see whether there are additional i-GCs and to discriminate them from GCs associated with the central galaxy. This search needs to extend over several hundred kilo-parsecs from the cluster centre, since the GC population in a luminous elliptical galaxy may extend out to $\\sim$ 100 kpc (Rhode \\& Zepf 2001; Dirsch et al. 2003). By surveying a contiguous region around the central galaxy, one may be able to detect an apparent change in the slope of the radial profile of GC surface number densities due to the presence of i-GCs. Secondly, GCs need to be probed down to a very low surface number density to investigate the spatial structure of their distribution out to such large radii. It is suggested by recent wide-field studies of GC populations in luminous ellipticals that the GC surface number density in the outer halos is $\\lesssim$ 1 arcmin$^{-2}$ (Rhode \\& Zepf 2001, 2004; Dirsch et al. 2003; Dirsch, Schuberth \\& Richtler 2005). If this is a mixture of GCs associated with the host galaxy and i-GCs, the i-GC surface number density would be even lower.  In fact, a recent study of GC populations in the Fornax cluster has suggested an i-GC surface number density of $\\sim 0.1$ arcmin$^{-2}$ (Bassino et al. 2003). Accordingly, selection of GCs and statistical subtraction of foreground and background contamination will have to be carefully performed.  In this paper and a companion paper (Tamura et al. 2006; Paper I hereafter), we report a wide-field imaging survey in the $B$, $V$, and $I$ bands of the GC populations around M87, the luminous elliptical galaxy at the centre of the Virgo cluster, conducted with Suprime-Cam on the Subaru Telescope. The survey area is $\\sim 2^{\\circ} \\times 0_{\\cdot}^{\\circ}5$ (560 kpc $\\times$ 140 kpc) which is the widest survey yet undertaken of the GC populations around luminous ellipticals (the survey depth is similar or slightly deeper than the previous ground-based surveys). We also make use of wide-field imaging data on control fields taken with Subaru/Suprime-Cam to correct for foreground and background contamination. The observations, data reduction and data analyses including the selection of GC candidates, incompleteness corrections, and subtraction of foreground and background contamination in the GC candidates are presented in Paper I, where we also present GC luminosity functions and global GC specific frequencies of M87 and NGC 4552. The present paper is dedicated to the investigation of the colour distribution and spatial distribution of GCs. Observational results are presented in the next section and are discussed further in \\S~\\ref{discussion}. We summarize this paper in \\S~\\ref{summary}. We adopt a distance of 16.1 Mpc (distance modulus of 31.03) to M87, and 15.4 Mpc (distance modulus of 30.93) to NGC 4552, based on measurements using the surface brightness fluctuation method (Tonry et al. 2001). An angular scale of 1$^{\\prime}$ corresponds to 4.7 kpc and 4.5 kpc at the distance of M87 and NGC 4552, respectively.  \\begin{figure*} \\begin{center} \\includegraphics[height=17.5cm,keepaspectratio]{ntamurafig01.ps} \\caption{$V-I$ colour distributions of GC candidates with $V \\leq 24$ mag. The dotted line shows the raw colour distribution, and the solid lines indicate those after correcting incompleteness and contamination. The thick (thin) line shows the colour distribution where the subtraction of contaminating sources is performed using the HDF (LH) data, respectively.}  \\label{allvih} \\end{center} \\end{figure*}  \\begin{figure*} \\begin{center} \\includegraphics[height=16cm,keepaspectratio]{ntamurafig02.ps} \\caption{$V-I$ colour distributions for the GC candidates with $V \\leq 24$ mag in the central ($R \\leq 5^{\\prime}$) regions of M87 ({\\it left}) and NGC 4552 ({\\it right}). The histogram is the GC colour distribution obtained after the incompleteness correction and control field subtraction are applied. Overplotted are fitted Gaussians representing the blue and red GC subpopulations (thin lines) and the sum of these (thick lines). A linear relationship between $V-I$ colour and metallicity can be derived by using simple stellar population models by Vazdekis et al. (1996) with an age of 12.6 Gyr for and a Salpeter initial mass function as follows: $[{\\rm Fe/H}] = 4.33~(V-I) - 5.33$.}  \\label{bimo} \\end{center} \\end{figure*}  \\begin{figure*} \\begin{center} \\includegraphics[height=13cm,angle=-90,keepaspectratio]{ntamurafig03.ps} \\caption{$V-I$ colours of the GC candidates with $V \\leq 24$ mag are plotted (grey points) against distance from M87 ({\\it left panels}) and NGC 4552 ({\\it right panels}). The average color of GCs at a given radius is overplotted with filled circles and the 25- and 75-percentile of the colour distribution are indicated with the error bars. The colour profiles of the host galaxies are overplotted with solid lines; the thick line indicates the measured colour profile and the thin lines show the uncertainty due to $\\pm 1 \\sigma$ errors in the background estimation.} \\label{gcscol} \\end{center} \\end{figure*} ",
        "conclusions": "\\label{summary}  We have performed a wide-field imaging survey of the globular cluster (GC) populations around M87 with Suprime-Cam on the 8.2m Subaru Telescope. A $2^{\\circ} \\times 0_{\\cdot}^{\\circ}5$ (560 kpc $\\times$ 140 kpc) field extending from M87 to the east was observed through the $BVI$ filters. GC candidates are selected not only with an extended source cut but also with a colour selection which includes almost all the Galactic GCs but minimizes the contamination of foreground stars and background galaxies. We also analyze archival imaging data of the control fields (HDF-N and Lockman hole) to assess foreground and background contamination in the GC candidates which needs to be statistically subtracted from the GC candidates.  In this paper, we investigate the colour distribution, surface number density, and local specific frequency ($S_N$) of GC candidates as a function of distance from two luminous elliptical galaxies M87 and NGC 4552. In investigating these statistical properties, we subtract foreground and background contamination using both the HDF-N data and the Lockman Hole data but we find that the results do not depend significantly on choice of the control fields. The main results are summarized as follows:  \\begin{itemize}  \\item The colour distributions of GCs in the innermost region (within $5^{\\prime}$ or $\\sim 20$ kpc from the galaxy centre) shows bimodality both in M87 and NGC 4552. This bimodality becomes less clear in the outer regions due to a decreasing contribution of the red GC subpopulation. The colour of the host galaxy halo light is slightly redder ($\\Delta(V-I) \\sim 0.1$ mag) than that of the red GC subpopulation.  \\item The average colour of GCs becomes bluer with increasing distance from the host galaxy. The colour gradient within 15$^{\\prime}$ from the host galaxy is estimated to be $d (V-I) / d \\log r = -0.17 \\pm 0.07$ for the M87 GC population and $d (V-I) / d \\log r = -0.16 \\pm 0.08$ for the NGC 4552 GC population. The GC colour gradient may be due to a combination of the decreasing contribution of red GCs and a colour gradient in each GC subpopulation. At radii larger than $\\sim 15^{\\prime}$, however, the radial gradient of GC colour seems much less significant in both of M87 and NGC 4552. The colour gradient of the host galaxy halo light is found to be steeper than the GC colour gradient in M87, while both are consistent with each other in NGC 4552.  \\item The spatial distribution of the GC population as a whole is more extended than the host galaxy halo light both in M87 and NGC 4552. This is primarily due to the extended distribution of the blue GCs, while the distribution of the red GC subpopulation is as centrally concentrated as that of the halo light. Accordingly, the local $S_N$ for the blue GCs increases with radius, whereas that for the red GCs is nearly constant.  \\item Although the spatial distribution of blue GCs around M87 is more extended than that of the host galaxy halo light and the red GCs, it appears to be less extended than the mass distribution of dark matter in the Virgo cluster. Furthermore, the extended distribution of blue GCs does not seem to be a special characteristic of giant ellipticals at the cluster centres like M87. Both of these facts do not support the idea that the high $S_N$ value of M87 is due to the superposition of intergalactic GCs (i-GCs). Instead, we suggest that most of the blue GCs as well as the red GCs are associated with the host galaxy.  \\item We find marginal evidence for the existence of intracluster GCs (i-GCs) with a surface density of $0.2^{+0.3}_{-0.1}$ arcmin$^{-2}$. Comparison with the predicted surface luminosity density of intracluster stars gives a crude estimate of $S_N$ for this i-GC population $S_N = 2.9^{+4.2}_{-1.5}$. If this $S_N$ represents the true ratio of GCs to tidally stripped stars, the fraction of GCs captured by this process would be low in the GC population of M87.  \\end{itemize}"
    },
    "0609/astro-ph0609246_arXiv.txt": {
        "abstract": "We discuss the serendipitous discovery of an extremely metal poor galaxy, SDSS~0809+1729, classified as a star in the Sloan Digital Sky Survey (SDSS). The galaxy has a redshift $z = 0.0441$ and a $B$-band absolute magnitude M$_B = -17.1$. With a metallicity of \\OH $\\sim 7.44$ or $\\sim1/20$ solar, this galaxy is among the 10 most metal poor emission-line galaxies known. SDSS~0809+1729 is a blue compact galaxy (BCG) with a stellar age of $\\sim 4.5$~Myr, a star formation rate of 0.18 \\Msun yr$^{-1}$, and a large gas-phase electron density ($\\sim367$~cm$^{-3}$). Similar values of these parameters are common among other extremely metal poor galaxies, including I~Zw~18. SDSS~0809+1729 is, however, unusual among BCGs because it lies in the same region of the luminosity-metallicity diagram as the two lowest metallicity long duration gamma ray burst (GRB) hosts.  For a given $B$-band luminosity, both nearby GRB hosts and SDSS~0809+1729 have systematically lower metallicities than dwarf irregulars, the majority of BCGs, and normal star-forming galaxies. Because the star formation properties of SDSS~0809+1729 are similar to nearby long duration GRB hosts, SDSS~0809+1729 may be a potential GRB host.  Identification of larger samples of similar extremely metal poor objects may provide important insights into the conditions required to produce long duration GRBs. ",
        "introduction": "Extremely metal poor local galaxies are key to understanding star formation and enrichment in a nearly pristine interstellar medium (ISM). Metal poor galaxies provide important constraints on the pre-enrichment of the ISM by previous episodes of star formation. Evolutionary scenarios proposed to explain the low metallicities include metal loss from supernova-driven winds \\citep{Ferrara00,Recchi04}, dilution of the ISM metallicity by infall of unenriched gas \\citep{Koppen05}, or a star formation history dominated by short bursts of star formation separated by long quiescent periods \\citep{Searle72}. The latter scenario is currently favored \\citep[see][for a review]{Kunth00}, but all of these mechanisms probably play some role.  Intriguingly, metal poor galaxies have recently been linked to long duration GRBs \\citep{Stanek06,Wolf06,Fruchter06}. \\citet{Sollerman05} compared the host galaxy properties of three nearby GRBs with the properties of local BCGs. They found that the luminosity and star formation rates of the GRB hosts are similar to those of local BCGs.  \\citet{Fruchter06} demonstrate that long duration GRBs are concentrated in the highest surface brightness regions of their extremely blue hosts.  \\citet{Stanek06} show that five nearby GRB hosts have lower metallicities than normal star-forming galaxies of similar luminosity. Because the GRB energy released decreases steeply with increasing host metallicity, Stanek et al.\\ propose an upper metallicity limit for cosmological GRBs of $\\sim 0.15$~\\Zsun.  Strong Lyman-$\\alpha$ detections in GRB hosts at $z\\gtrsim 2$ also support low host metallicities \\citep{Fynbo03}.   Extremely metal poor galaxies (XMPGs) are rare.  Fewer than 1\\% of dwarf galaxies have extremely low metallicities, defined as \\OH$\\leq 7.65$ \\citep{Kunth00,Kniazev03}.  XMPGs, however, are much more common than known GRB hosts.  Thus a larger sample of XMPGs is one route to understanding the environment required for the formation of the long duration GRBs.  Most of the known XMPGs are gas-rich BCGs with spectra dominated by emission lines \\citep[e.g.,][]{Kunth83,Thuan95}. Emission-lines provide an estimate of the gas-phase metallicity from the \\OIII~$\\lambda 4363$ doublet.  This line is particularly sensitive to the electron temperature of the gas and is strong in metal poor galaxies where \\HII~region cooling is minimal.  Many investigators have searched for XMPGs with varying degrees of success.  For more than three decades, the famous galaxy I Zw 18 was the most metal poor galaxy known, with a metallicity of \\OH$\\sim7.17$ \\citep{Searle72}. Recently, \\citet{Izotov05} showed that the galaxy SBS 0335-052W has an even lower oxygen abundance of \\OH$\\sim7.12$. Searches of SDSS spectra based on emission-lines \\citep{Izotov04,Kniazev03,Kniazev04,Papaderos06a,Izotov06b} have uncovered few XMPGs, in agreement with previous searches of emission-line galaxies \\citep[e.g.,][]{Terlevich91,Masegosa94,VanZee00}.  Surveys that select on the equivalent width of strong lines such as \\OII\\ appear to be more efficient at finding XMPGs \\citep{Ugryumov03}, but the number of known XMPGs remains small.  Here we report the serendipitous discovery of a new, very blue XMPG, SDSS~0809+1729.  This galaxy is at a redshift $z = 0.0441$ and has a metallicity of \\OH$\\sim7.44$.  This metallicity places SDSS~0809+1729 among the 10 lowest metallicity emission-line galaxies known to date.  We describe the observations and derived quantities in \\S~\\ref{obs}. We compare the properties of our galaxy with those of other XMPGs in \\S~\\ref{comp}.  We show that XMPGs and GRB hosts share similar spectral properties in \\S~\\ref{GRB}.  We discuss our results and conclusions in \\S~\\ref{Conclusions}.  Throughout this paper, we adopt the flat $\\Lambda$-dominated cosmology as measured by the WMAP experiment \\citep[$h=0.72$, $\\Omega_{m}=0.29$;][]{Spergel03}). ",
        "conclusions": "We report the serendipitous discovery of an extremely metal poor galaxy, SDSS~0809+1729, classified as a stellar point source in the SDSS catalog.  This galaxy has a metallicity of \\OH$\\sim7.44$, or $\\sim 1/20$~solar, placing it in the group of ten lowest metallicity emission-line galaxies.  We compare the spectral properties of SDSS~0809+1729 with the spectral properties of seven known XMPGs, including I Zw 18.  We apply stellar population synthesis models to derive stellar ages and we calculate SFRs, taking into account the effect of metallicity.  The young stellar population in SDSS~0809+1729 has an age of $\\sim 4.5$~Myr, corresponding to the end of the Wolf-Rayet phase.  The star formation rate in SDSS~0809+1729 ($\\sim0.18$~\\Msun yr$^{-1}$) is ten times larger than the star formation rate in I Zw18.  We compare the properties of SDSS~0809+1729 and the other XMPGs with the properties of five nearby GRB hosts for which spectral data is available.  We find that:  \\begin{itemize} \\item XMPGs and GRB hosts lie in a similar region of the luminosity-metallicity diagram: both types of objects are systematically offset from the luminosity-metallicity relation defined by dIrrs, normal BCGs, and normal star-forming galaxies.  For a given luminosity, XMPGs and GRB hosts are less enriched than dIrrs, normal BCGs, and normal star-forming galaxies.  \\item XMPGs and GRB hosts share similar SFRs and extinction levels, strengthening the link between GRB hosts and XMPGs.  \\item Our new XMPG, SDSS~0809+1729, has nearly identical properties as the lowest metallicity GRB host, GRB~030329.  The \\Te\\ metallicities of the two galaxies agree to within $\\sim 0.1$~dex (the estimated error in the metallicity diagnostics is $\\sim 0.1$~dex), and the SFRs agree to within 0.06~\\Msun yr$^{-1}$, or 50\\%.  While electron density measurements are unavailable for GRB hosts with luminosities similar to SDSS~0809+1729, it is interesting that the luminous host galaxy of GRB~031203 has an electron density (and stellar age) very similar to the one we measure for SDSS~0809+1729.  \\end{itemize}  The similarity between XMPGs and long duration GRB hosts is remarkable. Our results suggest that SDSS~0809+1729 and other XMPGs like it are potential hosts for long duration GRBs.  Identification and analysis of larger samples of XMPGs may provide important insight into the environment where long duration GRBs occur.  The serendipitous discovery of an XMPG using SDSS colors suggests that a similar color-selection method may be useful for revealing new XMPGs. Some of these objects have been missed because, like SDSS~0809+1729, they are indistinguishable from very blue stars.  We will report soon on a more extensive survey to identify these objects (Brown et al.\\ 2006, in prep.)."
    },
    "0609/astro-ph0609778_arXiv.txt": {
        "abstract": "The degree of linear polarization in solar flares has not yet been precisely determined despite multiple attempts to measure it with different missions. The high energy range in particular has very rarely been explored, due to its greater instrumental difficulties. We approached the subject using the Reuven Ramaty High Energy Spectroscopic Imager (RHESSI) satellite to study 6 X-class and 1 M-class flares in the energy range between 100 keV and 350 keV. Using RHESSI as a polarimeter requires the application of strict cuts to the event list in order to extract those photons that are Compton scattered between two detectors. Our measurements show polarization values between 2\\% and 54\\%, with errors ranging from 10\\% to 26\\% in $1\\sigma$ level. In view of the large uncertainties in both the magnitude and direction of the polarization vector, the results can only reject source models with extreme properties. ",
        "introduction": "\\label{Introduction}  Measurements of the linear X-ray polarization in solar flares can provide essential information needed to identify the processes responsible for the acceleration of particles and the emission of radiation. For photon energies in the hard X-ray region, polarization is produced either through electron Brems\\-strah\\-lung or Compton scattering. The degree of polarization is usually a complex function of the strength and topology of the magnetic field. Thermal distributions of electron velocities result in small polarization values produced mainly by the conduction-driven anisotropy of the electrons in the primary source \\cite{emsliebrown80} with some contribution from photons backscattered in the photosphere. Higher polarization levels are expected from non-thermal anisotropic distributions of electrons that are accelerated in well ordered magnetic fields. In addition, for such cases the observed polarization degree is related to the photon directivity (i.e. the anisotropy of the emitted radiation), which depends both on the electron beaming details and on the viewing angle.  Starting in the late seventies, several non-thermal models of X-ray emission from solar flares were developed (see \\opencite{mcconnell02} and references therein). Generally the authors assume a uniform magnetic field perpendicular to the solar surface and electrons being accelerated towards the chromosphere (\\opencite{elwerthaug70}; \\opencite{haug72}; \\opencite{brown72}; \\opencite{bairamaty78}; \\opencite{zharkova95}). These electrons, spiraling downwards along the magnetic field lines, produce Brems\\-strah\\-lung radiation in collisions with hot plasma. The polarization in the emitted X-rays is a function of the energy spectra of the electrons, their pitch angles (angle between their velocity vector and the magnetic field), and the column density distribution of the ambient plasma. Studying the spectral characteristics of the detected X-rays provides  information about the electron energies, while the polarization is a very sensitive tool to sample the electrons pitch angle distribution. Both for very small pitch angles, corresponding to high electron beaming, and for very large ones, the polarization degree can be equally high. The two cases can be distinguished by their polarization direction: parallel and perpendicular to the magnetic field line, respectively. The predicted polarization values can reach up to 60\\% at energies above 50 keV, being even higher at low energies \\cite{haug72}. After introducing more realistic pitch angle distributions and taking into account photon backscattering processes in the photosphere, the expected polarization is reduced down to 20\\% or 30\\% (\\opencite{brown72}; \\opencite{bairamaty78}).  Several more complex non-thermal models have been developed as well. In one of them the magnetic field structure was defined as a semicircular loop anchored in the chromosphere \\cite{leachpetrosian83}. This approach allowed studying the X-ray emissions separately from different parts of the loop. The highest polarization could be produced at the top (up to 85\\%), while the photons observed from the foot-points (in the region of the dense chromosphere) would be polarized to the level of around 20\\%.  In general, lower energies are predicted to yield stronger polarization signals (\\opencite{haug72}; \\opencite{bairamaty78}; \\opencite{leachpetrosian83}), although the strength of this relation can vary depending on the model. In some recent theories an opposite trend has also been reported \\cite{zharkova95}. Experimental verification is usually difficult because the low-energy part of the spectrum is strongly contaminated by a non-polarized thermal emission.  In all the cases the observed value of flare polarization is strongly dependent on the viewing angle. The highest polarization values are expected for large angles of view, when the line of sight is perpendicular to the magnetic field line. Thus, most theories predict higher polarization for flares located near the solar limb. Similar behavior is also expected for the directivity of the flare emission: the intensity of the emitted radiation should depend on the angle of view (\\opencite{li94}; \\opencite{li95}). Constraints related to the model assumptions favor two possible directions of the polarization vector: either parallel to the plane defined by the magnetic field lines and the line of sight or perpendicular to it (\\opencite{bairamaty78}; \\opencite{leachpetrosian83}; \\opencite{zharkova95}).   Contrary to the intense theoretical work, only a few polarization measurements have been conducted in hard X-rays. In this energy range, the commonly used technique is based on Compton scattering \\cite{lei97}.  First attempts were done at energies around 15 keV by \\inlinecite{tindo70}, \\inlinecite{tindo72}, and \\inlinecite{tindo76} using polarimeter instruments on board of several Intercosmos satellites. Although initial results \\cite{tindo70} showed a linear polarization of around 40\\% $\\pm$ 20\\%, the later studies of different flares (\\opencite{tindo72}; \\opencite{tindo76}) found their polarization degrees always compatible with zero. Moreover, the data suffered from limited photon statistics and systematic errors related with the detector calibration.  In measurements of solar flare polarization with the Reuven Ramaty High Energy Spectroscopic Imager (RHESSI), Compton scattering can occur in a specially installed beryllium scatterer. This method can only be used at low energies (20--100 keV). Its details, together with RHESSI's polarimetric features, are described in \\inlinecite{mcconnell02} and \\inlinecite{mcconnell04}.  The latest measurements at energies below 100 keV have been performed  with the SPR-N instrument on board of the Coronas-F satellite by \\inlinecite{zhitnik06}. From a sample of 25 solar flares, these authors determined upper limits on the polarization degree in the range from 8 to 40\\% ($3\\sigma$). Only for the single case of the 29 October 2003 flare they found a significant polarization degree, which increases from about 50\\% at energies $20-40~keV$, up to more than 70\\% for the energy channel $60-100~keV$.  Recently published, the only two measurements of solar flare polarization at high energies (0.2--1 MeV), show 1$\\sigma$ values of 21\\% $\\pm$ 10\\% and -11\\% $\\pm$ 5\\% for one flare close to the limb and another near the Sun center, respectively \\cite{boggs06}. In their studies, the authors applied the polarimetric capabilities of the RHESSI satellite in the coincidence mode, i.e. without using the Be-scatterer.  In this paper, we present results in the energy range from 100 keV to 350 keV, obtained for seven solar flares (X and M classes) also selected from the RHESSI instrument database. We used a method based on the scattering of photons from detector to detector of RHESSI that has previously been applied for polarization studies of gamma ray bursts (\\opencite{coburnboggs03}; \\opencite{wigger04}; \\opencite{rutledgefox04}).   In \\S \\ref{sec:method} we explain how RHESSI can be used as a Compton polarimeter to study linear X-ray polarization at energies $\\geq$ 100 keV. The flares selected for analysis, and the criteria used for their selections is detailed in \\S \\ref{sec:flselec}. Monte Carlo simulations were performed to calculate the response of the instrument, and their results are discussed in \\S \\ref{sec:simulation}. The final polarization results are described and compared with previous measurements in \\S \\ref{sec:results}. The interpretation of the results is done in \\S \\ref{sec:interpretation} by comparing with theoretical predictions. Finally, a brief summary of the conclusions of our work is given in \\S \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}   Measurements of hard X-ray polarization have been performed for six X-class and one M-class flares from the RHESSI database. Our flare sample was identified after applying selection criteria to the signal strength, background levels and level of contamination by charged particles in RHESSI detectors. The selected energy range, from 100 to 350 keV, connects the old (\\opencite{tindo70}; \\opencite{tindo72}; \\opencite{tindo76}) and new \\cite{zhitnik06} results at low energies with the only measurement reported at a high energy band \\cite{boggs06}.  We found values for the polarization degree in the range between 2\\% and 54\\%, with statistical errors from 10\\% to 26\\% at the $1\\sigma$ level. The polarization angles are distributed between $66^\\circ$ and $170^\\circ$ in heliocentric coordinates. They do not show any preferential orientation of the polarization, neither parallel nor perpendicular, with respect to the radial direction defined by the position of the flare in the Sun. In addition, no significant dependency between the orientation of the polarization and the distance of the flares to the Sun center was found.  The polarization orientation with respect to the line that joins the two major foot-points of the flare was also studied. For this purpose, images of the flares were constructed with RHESSI at different energies. The emission above $\\sim$30 keV was found to be produced around the foot-points. However, no correlation between polarization direction and foot-points orientation was found. Similarly, no relationship between solar flare intensity and polarization degree could be observed.  The polarization degree from the 23 July 2002 flare measured by \\inlinecite{boggs06} at high energies is in agreement with our results at the $1.5\\sigma$ level. Their conclusion about the orientation of polarization respect to the radial direction passing trough the flare position (perpendicular to it for flares in the limb and parallel for flares close to center) cannot be confirmed by our measurements. Our angles are distributed between $35^\\circ$ and  $85^\\circ$ independently on the flare location. Comparison with theoretical predictions is more complex, as the direction of the polarization is expected by the theory to change around 300 keV \\cite{bairamaty78}. Again, more observations with better accuracy are needed.  Regarding the low energy measurements, our data are in good agreement with the recent results from the SPR-N instruments on board of the Coronas-F satellite \\cite{zhitnik06}, typically within $2\\sigma$. Their large sample of 25 solar flares reveals low polarization degrees, providing $3\\sigma$ upper limits from 8 to 40\\%. Unfortunately, no information is given about the polarization direction. In the 29 October 2003 flare, for which the Coronas-F team claims a polarization of 70\\%, RHESSI measurements were contaminated by charged particles.  Theoretical predictions of the non-thermal Bremsstrahlung emission provide polarization levels of the order of 20\\%. However, depending on the assumptions used by different authors, the expected polarization can differ not only by its value, but also by its orientation (see, for example, the calculations at 100 keV from \\inlinecite{bairamaty78} and \\inlinecite{leachpetrosian83}). As the statistical uncertainties provided by our instrument are in the same order as the model predictions, equally good agreement is found for any of them as well as for the case of 0\\% polarization. In order to distinguish between different models, polarimetry measurements of at least 2\\% accuracy are needed. Only model number 5 from \\inlinecite{leachpetrosian83}, predicting polarization up to 85\\%,  could be rejected by our RHESSI data. In this model, the magnetic field has the same strength at the top and the bottom of the loop, and the electrons spiral at pitch angles close to $90^\\circ$.  RHESSI has made the first steps towards the understanding of the polarization phenomena in solar flares above 100 keV, where non-thermal emission dominates. Measurements with accuracy better than 10 to 20\\% were, however, hardly possible. This is due to its small effective area and high levels of flare-induced background. Continuation of such studies will require a dedicated polarimeter that must solve the problems inherent to the RHESSI design. Emphasis should be put on increasing the effective area and improving the background rejection capabilities. In particular, a better time resolution will reduce the number of accidental coincidences, and the optimization of the detector dimensions will improve the detection efficiency for the Compton scattering."
    },
    "0609/astro-ph0609252_arXiv.txt": {
        "abstract": "We present a sample of 1744 of Type 1 AGNs from the Sloan Digital Sky Survey Data Release 4 (SDSS DR4) spectroscopic catalog with  X-ray counterparts in the White-Giommi-Angelini catalog (WGACAT) of ROSAT PSPC pointed observations. Of 1744 X-ray sources, 1410 (80.9\\%) are new AGN identifications. Of 4574 SDSS DR4 AGNs  for which we found radio matches in the catalog of radio sources from the Faint Images of the Radio Sky at Twenty cm Survey (FIRST), 224 turned up in our sample of SDSS X-ray AGNs. The sample objects are given in a catalog that contains optical and X-ray parameters and supporting data, including redshifts; it also contains radio emission parameters where available. We illustrate the content of our catalog and its potential for AGN science by providing statistical relationships for the catalog data. The potential of the morphological information is emphasized by confronting the statistics of optically resolved, mostly low-redshift AGNs with unresolved AGNs that occupy a much wider redshift range. The immediate properties of the catalog objects include  significant correlation of X-ray and optical fluxes, which is consistent with expectations. Also expected is the decrease of  X-ray flux toward  higher redshifts. The X-ray--to--optical flux ratio for the unresolved AGNs exhibits a decline toward higher redshifts, in agreement with previous results.  The resolved AGNs, however, display the opposite trend. The X-ray hardness ratio shows a downward trend with increasing low-energy X-ray flux and no obvious dependence on redshift. At a given optical brightness, X-ray fluxes of radio-loud AGNs are on average higher than those of radio-quiet AGNs by a factor of 2. We caution, however, that because of the variety of selection effects present in both the WGACAT and the SDSS, the interpretation of any relationships based on our sample of X-ray AGNs requires a careful analysis of these effects. ",
        "introduction": "Active galactic nuclei (AGNs)\\footnote{Both pointlike QSOs and resolved galaxies with active nuclei are referred to in this paper as AGNs.} produce X-ray emission as the gaseous material of a galaxy is accreted onto the massive black hole in its center. This emission provides powerful diagnostics to probe the inner working of the AGN central engines, and X-ray studies of AGNs at different redshifts have always been of fundamental importance for understanding the physics, origin, and evolution of these objects. X-ray data become especially informative when combined with optical data and redshift information, and there is much effort to build and analyze such samples. Thus,  Fiore et al. (2003), presented optical identifications from the HELLAS2XMM survey for 122 sources detected with XMM-Newton in the surveyed area of 0.9 square degree. Green et al. (2004) found 335 unique optical counterparts for X-ray sources in the fields included in Chandra Multiwavelength Project (ChaMP), which is a survey of serendipitous Chandra X-ray sources in a $\\sim\\! 14$ square degree field (Kim et al. 2004a). Half of the 125 counterparts that Green et al. (2004) classified spectroscopically proved to be broad-line AGNs and 40\\% turned out to be galaxies with narrow emission lines or absorption-line galaxies.  Along with designated optical follow-ups, much work was done to identify X-ray counterparts for AGNs detected in various optical and infrared surveys. Some of these surveys, such as that presented by Wolf et al. (2004) for the Chandra Deep Field South, specifically targeted the fields observed in X rays. Others serendipitously overlap with the fields from various X-ray missions. Particularly important for AGN studies is the Sloan Digital Sky Survey, which not only increased the number of known AGNs by an order of magnitude but also furnished redshifts and spectra. Many thousands of AGNs were spatially  resolved, so a large amount of morphological information on underlying galaxies became available, too. A catalog of $\\sim 1200$  SDSS AGNs from the Early Data Release with X-ray counterparts from the ROSAT All Sky Survey (RASS; Voges et al., 1999; 2000) was presented by Anderson et al. (2003); this number nearly tripled for SDSS DR3 AGNs (Shen et al. 2005). The online system ClassX, developed by McGlynn et al. (2004; see also http://heasarc.gsfc.nasa.gov/classx), uses Virtual Observatory protocols\\footnote{In this work we have made use of data obtained from the US National Virtual Observatory, which is sponsored by the National Science Foundation} to retrieve optical information from other major surveys, such as 2MASS and USNO B1, and provides object type classifications using artificial intelligence algorithms. Most of the previously unidentified ROSAT X-ray sources are classified by ClassX as AGNs. Yet most of the X-ray sources detected by ROSAT, Chandra, XMM-Newton, and other large X-ray missions lack optical identifications.  There is a wide range of issues that can be addressed with samples of AGNs that have both X-ray and optical data, including redshifts (for recent review see, e. g., Brandt \\& Hasinger 2005). Hasinger, Miyaji, \\& Schmidt (2005) compiled a sample of $\\sim$ 1000 X-ray sources  with energies of 0.5--2~keV from ROSAT, Chandra, and XMM-Newton observations that were optically identified with AGNs. With that sample they derived for the first time reliable space densities of  X-ray selected  AGNs at cosmological distances. They quantified in much detail previous findings on the strong luminosity-dependent cosmological evolution of these objects, confirming the density peaks at $z \\sim 2$ and $z \\lesssim 1$ near the high and low limits, respectively, of the AGN X-ray luminosity range ( $\\log L_x \\sim 45$~and~42).  Barger et al. (2005) looked into cosmological evolution of AGNs using a few hundred optically identified hard Chandra X-ray sources with redshifts from spectroscopy or, in the case of too faint objects, estimated from optical photometry. The study found a decrease in the AGN X-ray luminosity by an order of magnitude since redshift $z = 1.2$. That decrease was shown to be due to AGN downsizing rather than to declining accretion rates onto the central black holes. Another result from this study suggests that the widely discussed unification model, according to which the broad-line (Seyfert~1) and narrow-line (Seyfert~2) AGNs differ only due to the line-of-sight effects (different inclinations of the black hole accreting disk) needs to be modified by including luminosity-dependent effects.  Brand et al. (2005) used a sample of $\\sim$3000 Chandra ACIS sources with  redshifts from optical photometry  to study the accretion history of active nuclei. The sources were identified with red galaxies in the Bootes area. Over the redshift range $z = 0.3 - 0.9$ the accuracy of photometric redshifts sufficed to show that the mean X-ray luminosity increases toward higher redshifts, interpreted as evidence for a substantial decline in the nuclear accretion since the epoch $z \\sim 1$.  A rapidly growing body of evidence suggests the existence of a connection   between the starburst and AGN phenomena (e.g., Kauffmann et al. 2003; Imanishi \\& Wada 2004). This fact is of fundamental significance, because it implies that the same physics underlies processes leading to the central burst of star formation and the AGN phenomenon, e.g., formation of a central molecular torus feeding both the starburst and the central black hole. Both phenomena are associated with a prodigious output of X rays and thus can be studied through the properties of the galaxy X-ray emission. Deep insights into intimate details of the AGN--starburst connections, hence the aforementioned underlying physics, can be gained from a study of combined X-ray and optical properties of a large sample of galaxies. If redshifts are available, the evolution of these connections can be tracked on a cosmological timescale. Recent efforts along these lines, which are based on the Chandra X-ray Observatory  data in the first place, have already brought significant results.   Hornschemeier et al. (2005) cross-correlated the Chandra sources from the archive of ACIS observations with SDSS galaxies. The resulting sample of 42 galaxies with redshifts within $z = 0.03 - 0.25$  was used to compare optical spectroscopic diagnostics of galaxy activity, both star formation and nuclear accretion. Although the small size of the sample limited the scope of issues that could be addressed with it, a number of substantial results were obtained. In particular,  all X-ray-luminous, X-ray-hard galaxies were found to have AGN spectroscopic signatures. This supports earlier claims that simple observational effects may explain the puzzle of  the class of X-ray--bright, optically normal galaxies (XBONGs) seen in the Chandra and XMM-Newton observations of galaxies at moderate redshifts. But the sample used is rather small, and statistically larger samples of both normal and AGN X-ray galaxies are still needed in order to either firmly rule out or confirm high X-ray luminosities that are due to mechanisms other than an AGN.  Optical spectra  and redshifts will continue to be the main challenge in achieving the full science potential of modern X-ray missions such as Chandra and XMM-Newton. This is also true for the ROSAT data, although the situation here is in some sense much better. The reason is that the ROSAT data can be supported to a much larger degree by the rich SDSS optical data. Anderson et al. (2003) emphasized that the ROSAT All Sky Survey and the SDSS survey are well matched in many respects, most importantly in terms of sensitivity. The sensitivity match is even better for the PSPC pointed observations, where a typical flux limit of $\\log f_x \\sim\\ -13.2$ samples very well the peak in the SDSS AGN brightness distribution at the AGN sample $i$-band limiting magnitude of $\\sim\\! 19$ (see the $\\log f_x - i$ relationship below). Therefore,  although less deep and having lower positional accuracy than the Chandra and XMM-Newton, and sampling only the soft ($ < 2.5$~keV) X-ray band, the ROSAT catalogs remain indispensable for many scientific issues of current interest.  There are a few  thousand ROSAT sources, mainly from the shallow All-Sky Survey, identified with resolved SDSS AGNs and/or starburst galaxies in the low-redshift Universe, $z < 0.6$. The number counts in the similarly soft samples of the Chandra  objects at these redshifts are typically limited to less than a hundred (see, e.g., Green et al. 2004; Barger et al. 2005; Hornschemeier et al. 2005). Prior to the SDSS barely more than a thousand of ROSAT sources were identified with optically known AGNs (Brinkmann, Yuan, \\& Siebert 1997; Yuan, et al. 1998). The initial catalog by Anderson et al. (2003)  of SDSS AGNs with X-ray emission from the RASS contains $\\sim 1200$ objects; this number was increased to 3366 by Shen et al. (2005) using SDSS DR3 AGNs.  In this paper we present a catalog of 1744 AGNs from SDSS DR4 with X-ray emission from ROSAT PSPC pointed observations, which are substantially deeper than those in the RASS. Half of these AGNs are in the low-redshift Universe, $z < 0.6$, similar to the AGNs in the Andersen et al. (2003) catalog. A significant fraction, $\\sim\\!20$\\%, are resolved SDSS objects, so their morphology can be studied from SDSS images. Statistically large samples of such AGNs will be especially useful to study AGN-starburst connections, the role of galaxy mergers in the AGN phenomenon, the relationships between the galaxy bulges and the central black holes (e.g., Heckman et al. 2004), and other problems for which the knowledge of the details of galaxy morphology is essential. The sample's usefulness for some other problems is limited, however, by selection effects present in both the WGACAT and SDSS. For instance, both the WGACAT sources and the SDSS QSO (AGN) sample are biased against obscured AGNs. Soft X-ray emission that can be detected with ROSAT is absorbed in such objects, and instruments operating in harder X-ray bands, such as Chandra and XMM-Newton, are needed to observe them. Obscuration also depresses AGN spectral features in the optical. Because of that the SDSS QSO (AGN) sample misses the population of heavily obscured AGNs. Due to these effects the  catalog we present is essentially limited to Type 1 AGNs. The above selection effects are not the only ones, and in this paper we  discuss more of them. ",
        "conclusions": "We presented a sample of 1744 SDSS DR4 AGNs with X-ray emission from ROSAT PSPC pointed observations. The sample was obtained by positional cross-correlation of all SDSS DR4 AGNs/QSOs  (SDSS spectroscopic type 3) with all unique X-ray sources in the catalog of ROSAT pointed observations (WGACAT). Of 1744 X-ray sources, 1,410 (80.9\\%) are new AGN identifications while the rest are the sources identified in the WGACAT with previously known AGNs. Of 4574 SDSS DR4 AGNs  for which we found radio matches in the catalog of the FIRST radio sources, 224 turned up in our sample of X-ray emitting AGNs. We provided the major parameters for the sample objects from the respective catalogs, as well as a number of derived parameters, such as X-ray broadband luminosity, 1.4 GHz radio luminosity, and X-ray hardness ratios. We presented a series of statistical relationships for the sample objects that illustrate their overall properties and demonstrate a high reliability of the X-ray counterpart identification. We plan to make use of the sample data in our subsequent studies of AGNs and starburst galaxies for which the combined optical and X-ray information is crucial."
    },
    "0609/astro-ph0609064_arXiv.txt": {
        "abstract": "The recent discoveries of nearby star clusters and associations within a few hundred pc of the Sun, as well as the order of magnitude difference in the formation rates of the embedded and open cluster populations, suggests that additional poor stellar groups are likely to be found at surprisingly close distances to the Sun. Here I describe a new nearby stellar aggregate found by virtue of the parallel proper motions, similar trigonometric parallaxes, and consistent color-magnitude distribution of its early-type members. The 120 Myr-old group lies in Ophiuchus at $d$ $\\simeq$ 170\\,pc, with its most massive member being the 4th-magnitude post-MS B8II-III star $\\mu$ Oph.  The group may have escaped previous notice due to its non-negligible extinction ($A_V$ $\\simeq$ 0.9\\,mag). If the group was born with a normal initial mass function, and the nine B- and A-type systems represent a complete system of intermediate-mass stars, then the original population was probably of order $\\sim$200 systems. The age and space motion of the new cluster are very similar to those of the Pleiades, $\\alpha$ Per cluster, and AB Dor Moving Group, suggesting that these aggregates may have formed in the same star-forming complex some $\\sim$10$^8$ yr ago. ",
        "introduction": "}  This is the first in a series of papers regarding the identification and characterization of groups of young stars within a few hundred pc of the Sun (mostly, newly discovered groups of young stars).  The second paper (Mamajek, in prep.) will detail the identification of a new $\\sim$20-Myr-old group within 100\\,pc. In this contribution, I discuss a new candidate group situated 170\\,pc away in Ophiuchus that appears to be similar in age to the Pleiades ($\\approx$120 Myr).  Very young stars with ages of $\\lesssim$3\\,Myr are usually found in embedded clusters (ECs) with tens to hundreds of members, and associated with dark molecular clouds \\citep[e.g.][]{Lada03,Porras03}. The vast majority of embedded clusters must not remain as bound structures identifiable as open clusters (OCs) after their molecular gas is removed.  This can be inferred by estimating the formation rates of the local samples of ECs and OCs within 1 kpc of the Sun.  In comparing the number of ECs to OCs within the same representative volume, and assuming a constant cluster-formation rate, \\citet{Lada03} calculate that $\\lesssim$4-7\\%\\, of embedded star clusters probably survive to the age of the Pleiades ($\\approx$130 Myr). Counting only those ECs with $\\gtrsim$35 members from the catalogs of \\citet{Lada03} and \\citet{Porras03} within 1 kpc of the Sun, the local surface density of ECs in the Galactic disk appears to be $\\sim$9 kpc$^{-2}$. The number of ECs with 10-35 members is roughly equal to the number of ECs with $>$35 members \\citep{Porras03}, so the surface density is doubled if one counts groups with $>$10 members.  Assuming a flat cluster-formation rate, and a mean EC age of $\\sim$2\\,Myr, \\citet{Lada03} estimate for their sample of local ECs with $>$35 members a formation rate of $\\sim$4 Myr$^{-1}$ kpc$^{-2}$.  If one assumes that the census of OCs with ages of 10-100\\,Myr and $d$ $\\leq$ 1\\,kpc is complete \\citep[using the 2006 update\\footnote{http://www.astro.iag.usp.br/$\\sim$wilton/} of the catalog of ][]{Dias02}, then the local density of such OCs is $\\sim$35\\,kpc$^{-2}$, and the OC formation rate is $\\sim$0.34 Myr$^{-1}$ kpc$^{-2}$ (similar to previous estimates by \\citet{Elmegreen85} and \\citet{Battinelli91}). The actual cluster formation rate may be somewhat higher as nearby clusters are still being discovered \\citep{Platais98,Mamajek99,Alessi03,Kharchenko05}. These calculations suggest that only $\\sim$9\\% of the ECs seen today would be hypothetically identified as OCs $\\sim$10-100 Myr in the future, and that $\\sim$90\\% of ECs likely evolve into unbound structures that might be identifiable as ``associations'' or ``moving groups''.  Another conclusion is that for every cataloged OC in the Sun's vicinity with age $<$100\\,Myr, there may be an {\\it order of magnitude} more stellar aggregates with $>$35 members that are {\\it not} cataloged as OCs, and there are likely to be approximately {\\it double} that number again if one tracks the evolution of ECs with $>$10 members. So where are these fossil remnants of ECs that are predicted to be in copious supply?  The nearest unbound remnants of ECs that contain any O-type stars, and/or significant numbers of B-type stars, have probably been identified as OB associations (OBAs), at least up until ages of few$\\times$10 Myr \\citep{Briceno06}.  The progenitors of OBAs were probably giant molecular cloud complexes containing ensembles of ECs (with the evolutionary descendants of the individual ECs appearing later as ``subgroups'' in the OBAs). The OB associations within 1 kpc cataloged by \\citet{deZeeuw99} contain between 6 and 85 B-type stars ($\\overline{n_B} \\, \\simeq\\, 45$).  Young stellar groups with ages of $\\sim$3-100 Myr that contain few, if any, B-type stars have only recently been identified. Given a standard initial mass function (IMF), one naively expects that for each B-type star formed, an EC will also typically produce $\\sim$1 A-star, $\\sim$1 F-star, $\\sim$1.5 G-stars, $\\sim$4 K-stars, and $\\sim$28 M-stars \\citep[where I have adopted the IMF of ][ and the spectral types are reflective of the resultant main sequence population]{Kroupa01}. Hence, if born with a \\citet{Kroupa01} IMF, an unrecognized cluster or association with $\\sim$5 B-type stars may have represented an embedded cluster of $\\sim$200\\,stars.  That such small groups can escape detection is well illustrated by that fact that at least 5 well-studied stellar aggregates with ages of $\\lesssim$30\\,Myr and distances of $\\lesssim$110 pc have been discovered over the past decade \\citep[e.g. $\\eta$ Cha, $\\epsilon$ Cha, TW Hya, $\\beta$ Pic, Tuc-Hor groups;][]{Mamajek99,Mamajek00,Feigelson03,Kastner97, Barrado99,Zuckerman00,Torres00}.  Each contains between zero (TW Hya Assn.) and four (Tuc-Hor) B-type stars, tens of low-mass stars, and no molecular gas (except for $\\epsilon$ Cha, the youngest of these). Lacking bright, hot stars from their ranks, the discovery of these groups had to await the era of large uniform astrometric databases with accurate proper motions and trigonometric parallaxes \\citep[i.e. {\\it Hipparcos};][]{Perryman97} and an X-ray all-sky survey \\citep[i.e. {\\it ROSAT};][]{Voges99}.  Any heretofore missing groups may be identifiable by virtue of their common proper motions and distances (where the higher mass stars will be preferentially conspicuous due to the shallow magnitude limits of contemporary astrometric catalogs).  The low-mass members of such groups may be identifiable by virtue of various youth signatures, with strong X-ray emission being the most conspicuous. The most unbound aggregates may even lack a distinguishable nucleus (e.g. TW Hya association), and may simply be identifiable as a slight over-density of young stars of similar age and space motion. ",
        "conclusions": "}  \\subsection{Solar Reflex Motion? \\label{reflex_motion}}  Could this aggregate be composed of unrelated B- and A-type field stars whose projected positions are clustered in the sky, but simply demonstrating solar reflex motion? While solar reflex motion is imparted on the observed proper motion of every star, the effect can not be responsible for the nearly identical proper motions of the eight early-type stars in the proposed aggregate.  The spread in proper motions is an order of magnitude smaller than that expected for an ensemble of young {\\it field} stars (even if situated at identical distances).  Suppose that the group is actually composed of field stars at similar distances (170 pc) but with a distribution of 3D space motions representative of B- and early A-type field stars. Here I simulate the proper motions of early-type field stars using the velocity dispersion tensor for stars with (\\bv) $\\in$ [-0.238, 0.139] from Table 1 of \\citet{Dehnen98}, using a simple Monte Carlo technique previously employed by the author in \\citet{Siegler03}. For the simulated data sets, I adopt the solar peculiar motion with respect to the LSR from \\citet{Dehnen98}. I create a Monte Carlo sample of 10$^4$ stars with $UVW$ vectors drawn from the \\citet{Dehnen98} velocity dispersion tensor, then assume that these stars lie at the distance and celestial position of the center of the $\\mu$ Oph aggregate, and calculate what distribution of proper motions one would see. The simulated proper motions are consistent with \\mura\\,\\, =\\, +3\\,$\\pm$\\,8 (1$\\sigma$) \\masyr\\, and \\mudec\\, =\\, --5\\,$\\pm$\\,13 (1$\\sigma$) \\masyr.  The predicted scatter in proper motion components ($\\sigma_{\\mu}$\\, $\\simeq$\\, 8--13\\, \\masyr) is roughly an order of magnitude larger than the observed dispersion in the proper motion components for the proposed cluster ($\\sim$1\\,\\masyr), even before accounting for the individual proper motion errors (mean errors $\\sigma_{\\mu}$ $\\simeq$ 1.0 \\masyr; Table \\ref{tab:muoph_astrom}).  The amount of {\\it intrinsic} velocity dispersion (unaccounted for by the observational errors) is unresolvable (\\S\\ref{dynamics}) -- the proper motion errors appear to account for all of the observed spread. The mean proper motion of the proposed cluster is also not near the predicted proper motion locus for the field population (which is within $\\sim$10\\,\\masyr\\, of zero motion). From these calculations, one can rule out solar reflex motion alone as the agent for the clumping of the proper motions observed in Figs. \\ref{fig:muoph_vector}.  \\subsection{Dynamics \\label{dynamics}}  How has the rather poor stellar nucleus of the $\\mu$ Oph cluster survived to an age of $\\sim$10$^8$ years? The answer likely lies with the high density of the nucleus. For the fiducial isochrone ($Z$ = 0.015; 120 Myr), the absolute magnitudes of the nucleus stars suggest masses for the nucleus stars of 5.2\\,\\msun\\, ($\\mu$ Oph), 3.0\\,\\msun\\, (HD 159874), 2.8\\,\\msun\\, (HD 160038), and 2.1\\,\\msun\\, (HD 160037). The deprojected half-mass radius \\citep{Spitzer87} for the observed membership is approximately 0.4\\,pc (corresponding closely to the observed nucleus radius) and within this radius the stellar density is at least $\\simeq$49\\,\\msun\\,pc$^{-3}$ \\citep[600$\\times$ the local Galactic disk density of $\\sim$0.08 \\msun\\,pc$^{-3}$;][]{Creze98}.  This density is similar to that within a similar size volume of the $\\eta$ Cha cluster \\citep[$\\simeq$56\\,\\msun\\,pc$^{-3}$; calculated from Table 2 and Fig. 7 of][]{Lyo04}. Using the formula of \\citet{King62} and the Oort constants from \\citet{Feast97}, I estimate the cluster tidal radius ($r_t$) for the observed membership to be $\\approx$4 pc.  Ideally one would like an estimate of the velocity dispersion to give some indication of the dynamical state of the group (i.e. is it bound?).  If the group were bound and conformed to a Plummer model, then given the total mass (24\\,\\msun) and half-mass radius (0.4\\,pc), the central 1D velocity dispersion should be $\\sim$0.25 \\kms\\, \\citep{Gunn88}. This is negligible compared to the proper motion uncertainties, but consistent with the lack of a detectable spread in the tangential motions. To test to see whether an intrinsic velocity dispersion was present, I use the best long-baseline proper motions available for each star (Tycho-2 for $\\mu$ Oph, and 2UCAC for the other seven stars; Table \\ref{tab:muoph_astrom}). The velocity dispersion can be estimated by the distribution of proper motions in the $\\tau$ (tangential) direction, i.e. in the direction perpendicular to the great circle joining the position of each star and the convergent point for the group \\citep[e.g.][and references therein]{Mamajek05}. The group convergent point can be estimated from the space motion vector ($\\alpha_{cvp}, \\delta_{cvp}$ = 108$^{\\circ}$.9, --30$^{\\circ}$.7; uncertainty of $\\pm$2$^{\\circ}$.7)\\footnote{The classic convergent point technique does not work well for stellar groups that are close together, like the $\\mu$ Oph group. The convergent point algorithm described in \\citet{Mamajek05} yields a group convergent point of $\\alpha_{cvp}, \\delta_{cvp}$ $\\simeq$ 117$^{\\circ}$($^{+81^{\\circ}}_{-19^{\\circ}}$), --39$^{\\circ}$($^{+24^{\\circ}}_{-15^{\\circ}}$); (68\\% confidence intervals), within a long narrow ($\\approx$1$^{\\circ}$.5 wide) error ellipse oriented NW-SE. This convergent point is within 10$^{\\circ}$ of the convergent point estimated from the space motion, which is pleasing (and perhaps surprising) considering the size of the error ellipse.}.  The mean proper motion of the eight stars towards the convergent point is $\\overline{\\mu_{\\upsilon}}$ = 23.9 \\masyr\\, (with 1$\\sigma$ dispersion of 1.2 \\masyr), and the mean proper motion in the perpendicular direction is $\\overline{\\mu_{\\tau}}$ = 0.0 \\masyr\\, (with 1$\\sigma$ dispersion of 0.9 \\masyr). The mean uncertainties in the $\\alpha$ and $\\delta$ proper motion components are both 1.0 \\masyr, so the uncertainties in the rotated components ($\\sigma_{\\mu_{\\upsilon}}$ and $\\sigma_{\\mu_{\\tau}}$) are similarly 1.0 \\masyr.  Given the observational proper motion uncertainties (1.0 \\masyr), the dispersion in $\\mu_{\\tau}$ (0.9 \\masyr) for these nine stars is statistically consistent with parallel motion with negligible (indeed {\\it zero}) velocity dispersion. For reference, a tangential velocity of 1\\,\\kms\\, \\citep[typical for the velocity dispersion observed in OB association subgroups; ][]{Briceno06} at $d$ = 173\\,pc is equivalent to a proper motion of 1.2 \\masyr. A 1D velocity dispersion of 1\\,\\kms\\, added in quadrature to the proper motion uncertainties would have resulted in an observed proper motion dispersion of 1.6\\,\\masyr, however this is easily ruled out. Hence the data are consistent with the idea that the velocity dispersion is negligible, but a value can not be derived from modern astrometric data.  The evaporation time of the nucleus can be calculated following \\citet{Binney87} for two extreme cases. In the first case, where we are dealing with a ``cluster remnant'' with only the four high mass stars in the nucleus, the evaporation time is short ($\\sim$10$^8$ yr). This is perhaps not so surprisingly similar to the {\\it isochronal} age, given the ``poor'' appearance of the nucleus. In the other extreme case, one posits that the four high-mass stars in the nucleus represent a complete census of 2-5.5\\,\\msun\\, members, and hypothesize that there is an as-yet undetected entourage of low-mass members with masses drawn from a \\citet{Kroupa01} initial mass function. In this case, the existence of 4 high-mass members implies that $\\sim$130\\,$\\pm$\\,60 low-stars might exist\\footnote{If the nine B- and A-type stars represent a complete census of the 1.5-5.5\\,\\msun\\, population that initially formed, then assuming a \\citet{Kroupa01} IMF, the initial population may have been $\\sim$200 systems. This is similar in population to nearby ECs like NGC 2024, Mon R2, and Cha I \\citep{Porras03}.}, for a total nucleus mass of $\\sim$60\\,\\msun.  In this case the evaporation time is of order $\\sim$400\\,Myr. Reality probably lies in between these two extreme cases. There are likely to be some undetected low-mass members in the nucleus, but dynamical evolution will have preferentially evaporated the low mass members \\citep[e.g.][]{delaFuenteMarcos95}.  \\subsection{ISM \\label{ISM}}  In addition to the group's ``poorness'', the moderate extinction (\\av\\, $\\simeq$\\, 0.9$^{m}$) towards the $\\mu$ Oph group may be one reason why it was previously unnoticed. There are 2 Lynds dark clouds within 3$^{\\circ}$ of $\\mu$ Oph \\citep[\\#393, \\#382;][]{Lynds62}. In the new Catalog of Dark Clouds published by \\citet{Dobashi05}, there are 14 clouds and 22 ``clumps'' within 3$^{\\circ}$ of $\\mu$ Oph. The low-opacity cloud LDN 393 which covers $\\sim$20 square degrees \\citep{Lynds62} was broken down by \\citet{Dobashi05} into $\\sim$25 clumps, with clouds DUKSKUS H189, H159, and H177 accounting for most of the cloud coverage (clouds are from their Table 7). The clumps typically have peak extinctions of \\av\\, $\\simeq$ 1-4\\,mag. They are likely in the foreground of the new cluster, and represent a part of the Aquila Rift cloud complex connecting the Sco-Oph molecular clouds \\citep[$\\ell$ $\\simeq$ 355$^{\\circ}$; $\\sim$140\\,pc;][]{deZeeuw99} with the Serpens clouds \\citep[$\\ell$ $\\simeq$ 30$^{\\circ}$; $\\sim$225\\,pc;][]{Straizys03}.  To estimate how much gas is associated with the dark clouds near $\\mu$ Oph, I take the integrated extinction-area values from \\citet{Dobashi05} for clouds within 3$^{\\circ}$ (9\\,pc, projected) of the cluster nucleus, assume a distance of $d$ $\\simeq$ 150\\,pc, and adopt the dust-gas ratio from \\citet{Savage79}.  The total gas mass of the dark clouds is then approximately 5000-13000 \\msun.  \\subsection{Spread in Distances \\label{spread}}  As stated in \\S\\ref{distance}, the {\\it Hipparcos} and Tycho parallaxes for the nine members are statistically consistent with a mean distance of $\\simeq$173\\,pc. However it is not clear whether they are consistent with lying within a range of distances similar to their projected size on the sky (i.e. $\\sim$5-10\\,pc), or the estimated tidal radius (4\\,pc). Here I estimate a plausible upper limit to the spread in the distances to the individual cluster members, {\\it assuming that the members have identical space motions with negligible velocity dispersion} (physically motivated by the results from \\S\\ref{dynamics}).  A cluster member's distance from the Sun ($d_i$ in pc), angular separation from its group convergent point ($\\lambda_i$), proper motion component pointing towards the group's convergent point ($\\mu_{\\upsilon i}$ in \\masyr), and the magnitude of the cluster's space velocity (i.e. $V$ in \\kms), are related through the classic formula \\citep[e.g.][]{Atanasijevic71}:  \\begin{equation} d_i\\,=\\,\\frac{1000\\,V\\,sin\\,\\lambda_i}{A\\,\\mu_{\\upsilon i}} \\end{equation}  \\noindent where A is the astronomical unit in useful units (4.74047 km\\,yr\\,s$^{-1}$). For cluster members with negligible velocity dispersion and spread in $\\lambda$, the predicted dispersion in the distances amongst the cluster members ($\\sigma^{int}_{d}$) is approximately:  \\begin{equation} \\sigma^{int}_d\\,\\simeq\\,\\frac{1000\\,V\\,sin\\,\\overline{\\lambda}}{A\\,\\overline{\\mu_{\\upsilon}^2}}\\,\\sigma^{int}_{\\mu_{\\upsilon}} \\end{equation}  \\noindent where I have defined $\\sigma^{int}_{\\mu_{\\upsilon}}$ as the intrinsic dispersion in $\\mu_{\\upsilon}$ values unaccounted for by observational errors (here presumed to be due to the spread in distances only), and $\\overline{\\lambda}$ = 133$^{\\circ}$.9, $\\overline{\\mu_{\\upsilon}}$ = 23.9\\,\\masyr, and $V$ = 27.2\\,\\kms\\, for the cluster sample.  The observed 1$\\sigma$ dispersion in $\\mu_{\\upsilon}$ values is $\\pm$1.2\\,\\masyr\\ (\\S\\ref{dynamics}), and the mean observational error in $\\mu_{\\upsilon}$ is $\\pm$1.0\\,\\masyr. This leaves an intrinsic scatter of $\\sigma_{\\mu_{\\upsilon}}^{int}$ $\\simeq$ 0.7\\,\\masyr\\ of unaccounted dispersion. If we attribute this spread in the $\\mu_{\\upsilon}$ proper motions solely to the dispersion in distances to the cluster members using equations (4) and (5), then the upper limit on the dispersion in distances is $\\sigma^{int}_{d}$\\, $\\simeq$\\, 5\\,pc. If all of the stars were at the same distance, and the intrinsic spread in $\\mu_{\\upsilon}$ values were attributed to a 1D velocity dispersion, then the data would imply a velocity dispersion of 0.6\\,\\kms. However, a 1D velocity dispersion this large can be ruled out since the spread in $\\mu_{\\tau}$ values (the component perpendicular to $\\mu_{\\upsilon}$) is negligible compared to the observational errors (\\S\\ref{dynamics}). Based on this analysis, it appears likely that the 1$\\sigma$ dispersion in the distances to the individual cluster members is of order $\\pm$5\\,pc, consistent with the idea that the depth of the group is similar in size to the projected diameter of the group, and similar in magnitude to the estimated tidal radius."
    },
    "0609/astro-ph0609587_arXiv.txt": {
        "abstract": "We present new spatially-resolved Keck spectroscopy of early-type galaxies in three galaxy clusters at $z\\approx 0.5$. In two companion papers (van der Marel \\& van Dokkum 2006a,b) we construct dynamical models of the galaxies and compare their modeled $M/L$ ratios and dynamical structure to local samples. Here we focus on the fundamental plane (FP) relation, and combine the kinematics with structural parameters determined from {\\em Hubble Space Telescope} ({\\it HST}) images. The galaxies obey clear FP relations, which are offset from the FP of the nearby Coma cluster due to passive evolution of the stellar populations.  The $z\\approx 0.5$ data are combined with published data for eleven additional clusters at $0.18\\leq z \\leq 1.28$, to determine the evolution of the mean $M/L_B$ ratio of cluster galaxies with masses $M\\gtrsim 10^{11}\\,$M$_{\\odot}$, as implied by the FP.  We find $d \\log (M/L_B)/dz = -0.555 \\pm 0.042$, stronger evolution than was previously inferred from smaller samples. The observed evolution depends on the luminosity-weighted mean age of the stars in the galaxies, the initial mass function (IMF), selection effects due to progenitor bias, and other parameters. Assuming a normal IMF but allowing for various other sources of uncertainty, we find $z_*=2.01^{+0.22}_{-0.17}$ for the luminosity-weighted mean star formation epoch. The main uncertainty is the slope of the IMF in the range $1-2\\,M_{\\odot}$: we find $z_* = 4.0$ for a top-heavy IMF with slope $x=0$.  The $M/L_B$ ratios of the cluster galaxies are compared to those of field early-type galaxies at $0.32\\leq z\\leq 1.14$.  Assuming that progenitor bias and the IMF do not depend on environment we find that the present-day age of stars in massive field galaxies is 4.1\\,\\%\\,$\\pm$\\,2.0\\,\\% ($\\approx 0.4$\\,Gyr) less than that of stars in massive cluster galaxies. This relatively small age difference is surprising in the context of expectations from ``standard'' hierarchical galaxy formation models, and provides a constraint on the physical processes that are responsible for halting star formation in the progenitors of today's most massive galaxies. ",
        "introduction": "Recent studies of near-infrared selected samples have uncovered a large population of red, massive galaxies at $z>2$ (e.g., {Labb{\\' e}} {et~al.} 2003; {Franx} {et~al.} 2003; {van Dokkum} {et~al.} 2003; {Daddi} {et~al.} 2004; {Yan} {et~al.} 2004; {Papovich} {et~al.} 2006). Red galaxies dominate the high-mass end of the mass function at $2<z<3$ ({van Dokkum} {et~al.} 2006) and are highly clustered (Daddi et al.\\ 2003; Grazian et al.\\ 2006; Quadri et al.\\ 2006). Most of these objects are too faint in the rest-frame ultra-violet to be selected as Lyman break galaxies (e.g., {F{\\\"o}rster Schreiber} {et~al.} 2004; {Reddy} {et~al.} 2005). The most straightforward interpretation of these newly-found objects is that they are progenitors of today's massive early-type galaxies.  Interestingly, a large fraction of these galaxies appear to have very high star formation rates, based on modeling of their spectral energy distributions, (stacked) X-ray emission, infrared emission, submm emission, and spectra (e.g., {F{\\\"o}rster Schreiber} {et~al.} 2004; {Rubin} {et~al.} 2004; {Daddi} {et~al.} 2004; {Labb{\\'e}} {et~al.} 2005; {Knudsen} {et~al.} 2005; {Papovich} {et~al.} 2006; {Webb} {et~al.} 2006). A similar or smaller fraction is best-fit by passively evolving models with little ongoing star formation (e.g., {Labb{\\'e}} {et~al.} 2005; {Papovich} {et~al.} 2006; {Kriek} {et~al.} 2006b). These results suggest that by $z\\sim 2.5$ we are entering the star formation epoch of massive early-type galaxies.  An important question is whether the properties of early-type galaxies at intermediate redshift are consistent with this interpretation. Furthermore, detailed studies of early-type galaxies in different environments may identify descendants of particular populations of massive galaxies at $z>2$. For example, it may be that the passively evolving galaxies at $z\\sim 2.5$ are progenitors of cluster early-type galaxies whereas the star-forming galaxies are progenitors of early-type galaxies in groups.  One of the most sensitive tools for determining the star formation epoch of early-type galaxies is the redshift evolution of the fundamental plane (FP) relation ({Djorgovski} \\& {Davis} 1987), as it reflects evolution in the $M/L$ ratios of galaxies and has very low scatter. The $z>0$ FP has been studied extensively in the past decade, both in clusters (e.g., {van Dokkum} \\& {Franx} 1996; {van Dokkum} {et~al.} 1998a; {Wuyts} {et~al.} 2004; {Fritz} {et~al.} 2005; {Holden} {et~al.} 2005; {J{\\o}rgensen} {et~al.} 2006) and in the field (e.g., {Treu} {et~al.} 1999, 2002, 2005a; {van Dokkum} {et~al.} 2001; Rusin et al.\\ 2003; van de Ven et al.\\ 2003; {van der Wel} {et~al.} 2004, 2005; {di Serego Alighieri} {et~al.} 2005). The results from these papers can be summarized as follows: 1) the $M/L$ ratios of the most massive cluster and field early-type galaxies evolve slowly and regularly, indicating early formation of their stars; and 2) there is evidence that low mass galaxies evolve faster than high mass galaxies, both in the field and in clusters, suggesting that they have younger stellar populations. The interpretation of the first result is complicated by the small number of clusters that have been studied so far, uncertainties in the initial mass function (IMF), and by selection effects due to progenitor bias: if the morphologies of galaxies change with time the sample of high-redshift early-type galaxies is only a subset of the sample of nearby early-type galaxies, leading to biased age estimates (see {van Dokkum} \\& {Franx} 2001). The interpretation of the second result is complicated by the fact that large corrections for selection effects need to be made (see {Treu} {et~al.} 2005a, 2005b; {van der Wel} {et~al.} 2005), and that galaxies are assumed to undergo no structural or dynamical changes with redshift.  The goal of the present paper is to better constrain the star formation epoch of massive field and cluster galaxies, using the FP. We present spatially resolved spectroscopic data from Keck and high-quality photometric data from the {\\em Hubble Space Telescope} (HST) for a sample of early-type galaxies in three clusters at $z\\approx 0.5$. The FPs of the three clusters are presented and discussed, and combined with literature data on the nearby Coma cluster, a sample of nearby galaxies drawn from the Sloan Digital Sky Survey (SDSS), and eleven clusters at $0.18\\leq z \\leq 1.28$. The evolution of the $M/L$ ratio derived from this large sample is used to constrain the star formation epoch of massive cluster galaxies. We combine the cluster sample with recently published samples of field galaxies, and determine the age difference between massive field and cluster galaxies using a self-consistent modeling approach which is different from previous studies. In two companion papers (van der Marel and van Dokkum 2006a,b; hereafter vdMvD06a,b) we utilize our spatially-resolved data to construct detailed dynamical models for the sample galaxies. The results are compared to those for local samples, providing an independent method of deriving $M/L$ ratio evolution, and allowing us to validate the many assumptions that enter into analyses based on the FP. The dynamical models also yield a normalized measure of the rotation rate [akin to $(v/\\sigma)^*$] of the sample galaxies, providing a method to identify S0 galaxies that have been visually misclassified as ellipticals, based on their kinematics.  The structure of the present paper is as follows. The analysis of the spectroscopic and photometric data for the three $z\\approx 0.5$ clusters is described in \\S\\,2 and \\S\\,3.  In \\S\\,4 the FPs of the three clusters are presented and discussed. In \\S\\,5 results for the three clusters are combined with literature data on 11 additional distant clusters and two local samples, and the $M/L_B$ evolution is determined from this large sample. Detailed information on the transformations of all the literature data to our system are given in an Appendix. The measured evolution of the $M/L_B$ ratio is interpreted in \\S\\,6. In this Section the implications for the star formation epoch of massive cluster galaxies are discussed, and the cluster data are compared to previously published data for field early-type galaxies. The main results are summarized and discussed in \\S\\,7. We assume $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$, and $H_0=71$\\,\\kms\\,Mpc$^{-1}$ where needed. ",
        "conclusions": "\\label{conclusions.sec}  In this paper we presented new measurements of kinematics and structural parameters in three distant galaxy clusters. In vdMvD06a,b the spatially-resolved photometric and kinematic profiles of the galaxies are used to measure the evolution of the $M/L$ ratio independent of many assumptions that enter the FP analysis. In the present study we focused on the FP, and the implications for the star formation epoch of the most massive galaxies in clusters and in the field.  We find that the stars in massive early-type galaxies in clusters have a mean luminosity weighted formation redshift $z_*= 2.01^{+0.22}_{-0.17}$. The implied ages are somewhat younger than most previous studies, which were based on smaller samples (e.g., {Holden} {et~al.} 2005). The quoted error reflects a combination of uncertainties in the data points, the stellar population synthesis models, and the significance of progenitor bias. In particular, varying the amount of progenitor bias from zero to the maximum allowed by the {van Dokkum} \\& {Franx} (2001) models [$\\sim 0.1 z$ in $\\Delta \\log (M/L_B$)] leads to a range in $z_*$ of $2.23-1.84$. A special form of progenitor bias is ``dry'' merging, i.e., (nearly) dissipationless mergers of early-type galaxies (e.g., {Tran} {et~al.} 2005b; {van Dokkum} 2005; {Bell} {et~al.} 2006; {De Lucia} {et~al.} 2006). Numerical simulations suggest that these mergers preserve the edge-on projection of the FP relation (e.g., {Gonz{\\' a}lez-Garc{\\'{\\i}}a} \\& {van Albada} 2003; {Boylan-Kolchin}, {Ma}, \\&  {Quataert} 2006; {Robertson} {et~al.} 2006), which would imply that our analysis is insensitive to this effect.  The main source of uncertainty that our analysis does not (fully) take into account is the form of the IMF. The rate of evolution, as parameterized by $\\kappa$ in our analysis, is very dependent on the logarithmic slope of the IMF in the range $1-2$\\,$M_{\\odot}$. {Tinsley} (1980) finds $\\Delta \\kappa \\sim -0.3 \\Delta x$, where $x$ is the slope of the IMF, and the more recent {Worthey} (1994) models give $\\Delta \\kappa \\sim -0.22 \\Delta x$. Although we explored the values of $\\kappa$ for most commonly used IMFs ({Salpeter} 1955; {Kroupa} 2001; {Chabrier} 2003) these are all very similar in the region around one Solar mass. These IMFs are probably appropriate for star formation in present-day disk galaxies, but may not be applicable to (the progenitors of) early-type galaxies. As is well known the large mass of heavy elements observed in the hot gas of massive clusters, the abundance ratios of early-type galaxies, numerical simulations of star formation at very high redshift, and other lines of evidence suggest that the IMF may have been top-heavy at early times (see, e.g., {Worthey}, {Faber}, \\&  {Gonzalez} 1992; {Larson} 1998; {Schneider} {et~al.} 2002, and references therein). The top-heavy IMF preferred by {Nagashima} {et~al.} (2005) has $x=0$ (where $x=1.35$ is the {Salpeter} [1955] slope), which implies $\\kappa_B = 1.4$. This value of $\\kappa$ gives a best fitting formation redshift of $z_* = 4.0$, higher than the $z \\sim 2$ range derived for Salpeter-like IMFs. This model is shown by the broken line in Fig.\\ \\ref{mlfit.plot}. We note that our observations probe a very different range of stellar masses than the abundance studies ($1-2$\\,\\msun\\ versus $\\sim 10$\\,\\msun), and that a top-heavy IMF does not necessarily imply an overall change in the slope but could instead be caused by a time-varying characteristic mass ({Larson} 1998).  Interestingly, there is some indirect evidence for a top-heavy IMF from the combination of our data with recent results for galaxies at higher redshifts. Although many massive galaxies at $z\\sim 2.5$ are still in the process of forming stars (see \\S\\,1), there is also a population with very low specific star formation rates and strong Balmer/4000\\,\\AA\\ breaks (e.g., Kriek et al.\\ 2006a,b). These objects formed most of their stars at redshifts $z\\sim 3$ or beyond, inconsistent with the star formation epoch that we derive for Salpeter-like IMFs but in agreement with our results for a top-heavy IMF. Alternative explanations are also possible; e.g., the apparently old $z\\sim 2.5$ galaxies may be ``rejuvenated'' at later times (see, e.g., Treu et al.\\ 2005b), or they could simply be the oldest objects in a wide distribution of progenitor galaxies.  The age difference between massive field and cluster galaxies is better constrained than their absolute ages, as it is less affected by the uncertainty in the IMF. Our self-consistent modeling approach gives different constraints than the standard approach, as it has less free parameters. Assuming that progenitor bias is similar for massive cluster and field galaxies, they have a Salpeter-like IMF, and they have the same metallicity and dust content at a given mass, we find that massive field galaxies are younger than massive cluster galaxies, but that the difference is small at $4.1 \\pm 2.0$\\,\\% (or $0.4 \\pm 0.2$ Gyr). If a top-heavy IMF with $x=0$ is assumed for both field and cluster galaxies the age difference is 3.0\\,\\%\\,$\\pm$\\,1.5\\,\\%.  We note that the quoted errors do not include systematic errors in the comparison of field and cluster galaxies, which are difficult to quantify. Possible sources of such errors are differences in dynamical structure that deviate from pure homology, selection effects induced by the magnitude limits of the field samples (although these effects are more relevant for low mass galaxies; see Treu et al.\\ 2005, van der Wel et al.\\ 2005), a difference in $\\kappa$ and/or $A$ (which would reflect differences in the IMF or the metallicity), and a difference in the amount of progenitor bias. We tested the effect of a difference in progenitor bias by assuming that cluster early-type galaxies are not affected by it whereas the $M/L$ ratios of field early-type galaxies need to be corrected by $-0.1 \\times z$. In that case the age difference increases by a factor of two, to $\\approx 9$\\,\\%. The importance of progenitor bias is bounded by the observed scatter in the color-magnitude relation of early-type galaxies (van Dokkum \\& Franx 2001). In clusters this scatter remains low all the way to $z\\sim 1.3$ (Mei et al.\\ 2006), but it is not yet clear whether this also applies to massive field early-type galaxies (see, e.g., Ferreras et al.\\ 2005).  The small age difference we find is consistent with most, but not all, previous studies. It is consistent with the high redshift field FP studies of {Treu} {et~al.} (2005b) and {van der Wel} {et~al.} (2005), the analysis of a sample of gravitational lenses by Rusin \\& Kochanek (2005), and with studies of the local Mg$_2$ -- $\\sigma$ relation and FP by {Bernardi} {et~al.} (1998, 2003a, 2006). However, other studies find larger age differences (always in the sense that field galaxies are younger than cluster galaxies). For nearby massive early-type galaxies, {Thomas} {et~al.} (2005) and {Clemens} {et~al.} (2006) find age differences of 1.5 -- 2 Gyr from complex models which include age, metallicity, $\\alpha$-enhancement, and (in the case of Clemens et al.) Carbon-enhancement as free parameters. Such large age differences are inconsistent with our analysis at $>3\\sigma$. Perhaps most importantly, our results are not in agreement with the recent study of {di Serego Alighieri} {et~al.} (2006), who compare the $M/L$ ratios of field galaxies at $z\\sim 1$ from {Treu} {et~al.} (2005b) and {di Serego Alighieri} {et~al.} (2005) to that of galaxies in two clusters with mean redshift $\\langle z\\rangle = 0.86$ from {J{\\o}rgensen} {et~al.} (2006). For galaxies with masses $>10^{11}$\\,\\msun\\ {di Serego Alighieri} {et~al.} (2006) find a very large age difference of 3.5 -- 4 Gyr (their Fig.\\ 4). The source of the discrepancy between their results and ours is, at least in part, due to the cluster data that are used in the comparison to field galaxies. As discussed in Appendix \\ref{lit.sec} the $M/L$ ratios of the galaxies in the {J{\\o}rgensen} {et~al.} (2006) study need to be corrected downward by a factor $1+z$. Applying this offset to the $\\langle z \\rangle = 0.86$ cluster galaxies in Figs.\\ 1 and 2 of di Serego Alighieri et al.\\ (2006) brings them in line with the $z\\sim 1$ field galaxies.\\footnote{See also the erratum published by di Serego Alighieri et al.\\ (2006) in response to our discovery of this discrepancy. Their revised conclusions are consistent with those presented here.}  The small age difference between massive field and cluster early-type galaxies that we find is remarkable in the context of ``standard'' hierarchical models, which had predicted that early-type galaxies in clusters have much older stellar populations than those in the general field (e.g., {Baugh}, {Cole}, \\& {Frenk} 1996; {Kauffmann} 1996; {Kauffmann} \\& {Charlot} 1998). Specifically, as shown in van Dokkum et al.\\ (2001), the models of Kauffmann et al.\\ (1999) and Diaferio et al.\\ (2001) predicted a systematic offset between field and cluster galaxies of $\\sim 0.25$ in $\\log (M/L_B)$ at all redshifts, clearly at odds with observations. The reason for this environmental dependence in these models is that gas cooling and subsequent star formation continues as long as a galaxy is the central galaxy in its halo. Star formation terminates when a galaxy becomes a satellite in a larger halo, which naturally leads to a large population of ``red and dead'' galaxies in clusters and blue star forming galaxies in the field. This inability of standard models to produce red field galaxies is well known, and among proposed solutions is to prevent cooling by postulating that the gas is heated by a central active galactic nucleus (AGN) (e.g., {Bower} {et~al.} 2006; {Kang}, {Jing}, \\& {Silk} 2006; Croton et al.\\ 2006). The recent study of {De Lucia} {et~al.} (2006) incorporates such AGN feedback in a very large cosmological simulation. In their simulation, the median star formation epoch $z_* \\sim 2.6$ for elliptical galaxies with masses $M>10^{11}$\\,\\msun, in good agreement with our results. Furthermore, judging from their Fig.\\ 7 the difference in age between ellipticals in halos of $10^{13}$\\,\\msun\\ and $10^{15}$\\,\\msun\\ is only $\\sim 0.7$ Gyr, consistent with our study. We note that this success does not necessarily imply that the models can reproduce the observed evolution of the FP. Building on the semi-analytical models of Baugh et al.\\ (2005) and Bower et al.\\ (2006), Almeida, Baugh, \\& Lacey (2006) find that evolution in the radii of early-type galaxies compensates for the evolution in their $M/L$ ratios, such that the zeropoint of the FP is approximately constant with redshift -- inconsistent with the observations.  Taken at face value, the age difference between field and cluster galaxies is $\\sim 0.4 \\pm 0.2$\\,Gyr, which should be fairly easily detectable at $z\\gtrsim 2$. Interestingly, this number agrees reasonably well with the age difference that {Steidel} {et~al.} (2005) find comparing Lyman break galaxies in the general field to those in a proto-cluster at $z=2.3$. It is also qualitatively consistent with the finding that red galaxies at $z\\sim 2.5$ cluster more strongly than blue galaxies (Quadri et al.\\ 2006)."
    },
    "0609/astro-ph0609314_arXiv.txt": {
        "abstract": "The universal baryonic mass fraction ($\\Omega_b/\\Omega_m$) can be sensitively constrained using X-ray observations of galaxy clusters. In this paper, we compare the baryonic mass fraction inferred from measurements of the cosmic microwave background with the gas mass fractions ($f_{\\rm gas}$) of a large sample of clusters taken from the recent literature.  In systems cooler than 4 keV, $f_{\\rm gas}$ declines as the system temperature decreases.  However, in higher temperature systems, $f_{\\rm gas}(r_{500})$ converges to $\\approx \\left(0.12\\pm0.02\\right)\\left(h/0.72\\right)^{-1.5}$, where the uncertainty reflects the systematic variations between clusters at $r_{500}$.  This is significantly lower than the maximum-likelihood value of the baryon fraction from the recently released {\\it WMAP} 3-year results.  We investigate possible reasons for this discrepancy, including the effects of radiative cooling and non-gravitational heating, and conclude that the most likely solution is that $\\Omega_m$ is higher than the best-fit {\\it WMAP} value (we find $\\Omega_m = 0.36^{+0.11}_{-0.08}$), but consistent at the $2\\sigma$ level. Degeneracies within the {\\it WMAP} data require that $\\sigma_8$ must also be greater than the maximum likelihood value for consistency between the data sets. ",
        "introduction": "For over a decade now cluster gas mass fractions as inferred from X-ray observations have been used as a probe of the universal ratio of baryon to total matter densities, $\\Omega_b/\\Omega_m$ (e.g., White et al.\\ 1993; David et al. 1995; Evrard 1997; Mohr et al.\\ 1999; Roussel et al.\\ 2000; Allen et al.\\ 2002; Lin et al.\\ 2003; Ettori 2003; Allen et al.\\ 2004).  Supplementing these gas mass fractions with constraints on $\\Omega_b$ from, e.g., cosmic microwave background (CMB) measurements or a combination of Big Bang Nucleosynthesis (BBN) predictions and D/H measurements from high redshift quasars, therefore allows one to measure the total matter density $\\Omega_m$.  The reliability of this test rests on the assumption that clusters have been able retain the original baryon inventory assigned to them in the early universe.  So-called ``non-radiative'' cosmological simulations, which include a hydrodynamic treatment of the baryons but neglect sources or sinks such as radiative cooling, star formation, and feedback, indeed indicate that clusters retain nearly all their baryons until the present day (e.g., Frenk et al.\\ 1999; Kay et al.\\ 2004; Crain et al.\\ 2006).  The same is generally true for simulations with cooling and feedback.  Although the fraction of baryons in the hot phase depends strongly on the model, most recent simulations predict a mild increase in the hot gas fraction with cluster mass, and little evolution with redshift (e.g., Tornatore et al.\\ 2003; Kravtsov et al.\\ 2005; Ettori et al.\\ 2006).  Although the cluster baryon fraction test has been examined previously in many studies, there are several good reasons for revisiting it.  First, new high-quality data obtained from the {\\it Chandra} and {\\it XMM-Newton} telescopes now allow us to probe both the surface brightness and temperature profiles of clusters out to large radii.  As a result, both the statistical and systematic observational uncertainties on the gas mass fraction are substantially improved.  Second, much improved (e.g., K-band) measurements of the stellar content of clusters are now available. Third, cosmological simulations can now robustly predict the baryon fractions within $r_{500}$, which is roughly the same radius the latest X-ray measurements reliably extend out to. Fourth, analyses of mock observations of realistic cosmological simulations allow us to, e.g., quantify the observational bias introduced by assuming strict hydrostatic equilibrium (HSE) in the derivation of X-ray gas mass fractions.  Finally, analysis of the recently released {\\it WMAP} 3-year cosmic microwave background (CMB) data has yielded very tight constraints on the universal baryon fraction (Spergel et al.\\ 2007).  As a result, there is now an excellent opportunity to check for consistency (or lack of) between these two very different tests.  We note that Ettori (2003) previously found evidence for a slight discrepancy between the universal baryon fraction inferred from the first year {\\it WMAP} data and cluster baryon fractions inferred from {\\it BeppoSAX} X-ray data.  It is interesting to see whether or not this discrepancy remains when one makes use of the better observational data and improved theoretical predictions that are now available.  In this paper we analyse the best literature X-ray data in a homogeneous way, and compare the results with the universal baryon fraction inferred from the recently released {\\it WMAP} 3-year data (Spergel et al.\\ 2007).  Unless otherwise stated, we assume a $\\Lambda$CDM cosmology with $h=0.72$. ",
        "conclusions": "Recent, good quality observations of massive clusters with {\\it Chandra} and {\\it XMM-Newton} put strong observational constraints on the gas mass fraction in massive clusters at large radius. In many cases the new data allow the fraction to be constrained out to $r_{500}$ (the radius at which the cluster density contrast is 500). Simulations of clusters suggest that the cluster gas fraction at this radius should closely reflect the average baryon mass fraction in the universe as a whole.  We find $f_{\\rm gas}(r_{500}) \\approx \\left(0.12\\pm0.02\\right)\\left(h/0.72\\right)^{-1.5}$, where the uncertainty reflects the systematic variations between clusters. This is lower than the best fit to the {\\it WMAP} 3-year result of $f_b=0.176\\pm0.02$.  We consider whether the discrepancy could be due to a large fraction of the cluster gas cooling to form stars and cold gas clouds, or whether it could be due to strong non-gravitational heating transporting $\\sim30$\\% of the cluster gas outside $r_{500}$. Observational limits on the stellar and cold gas content of clusters appear to rule out the first possibility.  In order to investigate the second, we compute the energy budget required to rearrange the cluster gas. The energy required significantly exceeds the plausible energy input from black holes, unless their mass accretion history is always associated with efficient jet production.  The most likely explanation is that $\\Omega_m$ lies in the range $0.36^{+0.11}_{-0.08}$.  Such a relatively high $\\Omega_m$ lies slightly above the 68 per cent confidence limits from {\\it WMAP}, but is consistent with the current data at the $2\\sigma$ level. Since the measurements of $\\Omega_m$ and $\\sigma_8$ are highly correlated in the {\\it WMAP} analysis, this means that $\\sigma_8$ is also likely to be higher than the formal best-fit value of $0.74$.  This has important implications for the abundance of collapsed objects prior to re-ionisation (Reed et al.\\ 2007)."
    },
    "0609/astro-ph0609408_arXiv.txt": {
        "abstract": "{We present the second part of an optical spectroscopic study of planetary nebulae in the \\object{LMC} and \\object{SMC}. The first paper, Leisy~\\&~Dennefeld~(1996), discussed the CNO cycle for those objects where C abundances were available. } {In this paper we concentrate more on other elemental abundances (such as O, Ne, S, Ar) and their implications for the evolution of the progenitor stars. } { We use a much larger sample of 183 objects, of which 65 are our own observations,  where the abundances have been re-derived in a homogeneous way. For 156 of them, the quality of data is considered to be satisfactory for further analysis. } {We confirm the difficulty of separating Type$\\,$I and non-type-I objects in the classical He-N/O diagram, as found in Paper I, a problem reinforced by the variety of initial compositions for the progenitor stars. We observed oxygen  variations, either depletion via the ON cycle in the more massive progenitor stars, or oxygen production  in other objects. Neon production also appears to be present. These enrichments are best explained by fresh  material from the core or from burning shells, brought to the surface by the $3^{rd}$ dredge-up, as reproduced in recent models (some including overshooting). All the effects appear stronger in the \\object{SMC}, suggesting a higher efficiency in a low metallicity environment. } {Neither oxygen nor neon can therefore be used to derive the initial composition of the progenitor star: other elements not affected by processing such as sulfur, argon or, if observed, chlorine, have to be preferred for this purpose. Some objects with very low initial abundances are detected, but on average, the spatial distribution of PNe abundances is consistent with the history of star formation (SF) as derived from field stars in both Clouds. } ",
        "introduction": "Planetary nebulae (PNe) are prime targets for studing the chemical evolution of nearby galaxies. This is particularly true for the two Magellanic Clouds, which have a different composition from our own Galaxy, thus allowing the study of metallicity effects. Their distances are known with reasonable accuracy, and thus the physical properties of the PNe can be determined easily; and they are close enough that a large number of PNe have already been identified, therefore giving statistical significance to the derived properties.  In our first paper (Leisy~\\&~Dennefeld~(\\cite{leis2}); hereafter Paper~I) we reported  our observations of PNe in the Magellanic Clouds, and discussed those objects (16 in \\object{LMC} and 15 in \\object{SMC}) for which carbon abundances were available. The present paper discusses the entire sample for which optical observations are available: 65 objects from our own observations, to which we add 118 other PNe already published in the literature. All the data were treated for the first time in a homogeneous way, allowing a global analysis of elemental abundances  in the Magellanic Clouds.  \\noindent At the end of their evolution, intermediate-mass stars pass through an  asymptotic giant branch (AGB) phase and experience helium shell flashes (Thermal-Pulse AGB). During these  phases, they  lose most of their envelopes. These envelopes contain both original, un-processed  material and  material synthesized in the deep layers of the star, and then  dredged-up during the Red Giant and AGB phases. Therefore substantial mass is ejected, which enriches the interstellar medium (ISM) with processed elements such as He, C, N and s-process elements (see for example Renzini~\\&~Voli~\\cite{ren1}, Iben~\\&~Renzini~(\\cite{iben3}), Dopita~\\&~Meatheringham~(\\cite{dop2}), Vassiliadis \\&~Wood~\\cite{vasi2}). But the ISM is also enriched in other, un-processed elements that were incorporated from the ISM at the time the progenitor star was formed, such as O, S, Ar, or Fe, which are elements believed to remain un-processed during the evolution of the PN progenitor star.   For those elements synthesized during the lifetime of their progenitor stars, the study of PNe allows one to determine how the production of chemical elements depends on initial mass or metallicity and how the products are {\\it dredged-up} during the late stages of stellar evolution. Although Kaler,~Iben~\\&~Becker~\\cite{kal1}, Becker \\&~Iben~(1979, 1980), Renzini~\\&~Voli~\\cite{ren1}, or Sweigart,~Greggio,~Renzini~(1989, 1990) have demonstrated reasonable agreement between observations and the dredge-up theory for giant stars, quantitative discrepancies still exist. We showed in Paper~I that initial metallicity  plays a crucial role in the efficiency of the various reactions and that additional processes, such as Hot Bottom Burning, predicted by theory (Gronewegen and de Jong,~\\cite{groe}) were effectively observed. Recent  evolution models of intermediate mass stars incorporating this are now available for a direct comparison with observations (Marigo et al.~(\\cite{mari1}), Marigo~(\\cite{mari2}) or Marigo et al.~\\cite{mari3}).  While PNe  play an important role in the chemical evolution of their host galaxies by furnishing helium, carbon, and nitrogen in great quantities, they can also be used to trace the evolution of those elements not affected by processing in their progenitor stars. With masses in the range $1~M_{\\sun}~\\leq~M_{i}~\\leq~8~M_{\\sun}$, evolution can be traced from PNe over the past 10 billion years, provided the mass of the central star and its initial metallicity are known. The latter is usually derived from oxygen, but in view of the metallicity effects shown in Paper~I and the involvement of O in the CNO cycle (as discussed later in this paper), other heavy elements, such as neon, sulfur and argon will also have to be used. Present day abundances in the ISM, used for comparison, are determined from H\\,{\\sc ii}~regions. \\\\ A large sample of objects, with various metallicities and an homogeneous data reduction and determination of abundances such as the one  provided here, is therefore crucial to understand the role of the various parameters involved. It will also allow a global analysis of the chemical evolution of the Magellanic Clouds to be compared with the results obtained  from H\\,{\\sc ii}~regions (e.g. by Dufour~\\cite{duf1}, 20 objects) and with those obtained from direct studies of a few stellar spectra (Spite et al.~\\cite{spit1}; Hill et al.~\\cite{hill1}). Such a program is a long-term observational enterprise, and first results have already been presented in various conferences (e.g. in Leisy \\& Dennefeld~\\cite{leis3}).  A short description of the observations and data reduction is given again in Sect.~2, as a reminder from Paper~I. The physical parameters for the objects, the abundances determinations, and the corresponding uncertainties are given in Sect.~3. Readers not interested in the details of the derivation of the PNe parameters  can go directly to Sect.~4, where the abundance patterns are discussed element by element. Implications for the chemical evolution of the Magellanic Clouds are presented in Sect.~5. ",
        "conclusions": ""
    },
    "0609/astro-ph0609122_arXiv.txt": {
        "abstract": "The nature of far-infrared dust emission toward main sequence stars, whether interstellar or circumstellar, can be deduced from submillimeter photometry. We present JCMT/SCUBA flux measurements at 850\\,\\micron\\ toward 8 stars with large photospheric excesses at 60--100\\,\\micron. 5 sources were detected at $3\\sigma$ or greater significance and one was marginally detected at $2.5\\sigma$. The inferred dust masses and temperatures range from 0.033 to $0.24\\,M_\\oplus$ and $43-65$\\,K respectively. The frequency behavior of the opacity, $\\tau_\\nu\\propto\\nu^\\beta$, is relatively shallow, $\\beta < 1$. These dust properties are characteristic of circumstellar material, most likely the debris from planetesimal collisions. The 2 non-detections have lower temperatures, $35-38$\\,K, and steeper opacity indices, $\\beta > 1.5$, that are more typical of interstellar cirrus. The confirmed disks all have inferred diameters $\\simgt 2''$ most lie near the upper envelope of the debris disk mass distribution, and 4 are bright enough to be feasible for high resolution imaging. ",
        "introduction": "Dust around main sequence stars with ages $\\simgt 10$\\,Myr must have been formed {\\it in situ} due to the short timescales for collisional destruction and radiative dispersal (Artymowicz 1988). The micron and sub-micron sized particles that are detected in infrared images of scattered light (e.g., Liu 2004) and emission (e.g., Su et al. 2005) are created from the collisions of larger particles, possibly from the late stages of planet formation. The study of such debris disks, therefore, constrain planet formation models (Kenyon \\& Bromley 2004) and new detections can signpost good targets for imaging of extrasolar planets (Zuckerman \\& Song 2004).  The faint emission from the cold dust is generally only apparent above the stellar photosphere at wavelengths longer than about 20\\,\\micron. The IRAS all sky survey has been the primary resource for debris disk searches, subsequently augmented with more targeted surveys by ISO (see de Muizon 2005) and most recently by Spitzer (Meyer et al. 2004; Rieke et al. 2005). These instruments are able to define the spectral energy distribution (SED) of the dust out to near its peak at about 100\\,\\micron, but photometry at longer wavelengths where the emission is optically thin is an essential complement by providing a measure of the dust mass. Moreover, by constraining the Rayleigh-Jeans side of the SED, the disk temperature and luminosity and the the power law index of the dust opacity can be determined. These are significantly different in debris disks and the interstellar medium and thus submillimeter observations can confirm or reject debris disk candidates that are based solely on unresolved far-infrared excesses.  Observations of debris disks at long wavelengths are difficult due to the intrinsic weakness of the emission and the poor transparency of the atmosphere. The Submillimeter Common-User Bolometer Array (SCUBA) on the James Clerk Maxwell Telescope (JCMT) has been the workhorse instrument to date (e.g., Holland et al. 1998) but has still had limited success: less than 20 systems have been detected at submillimeter wavelengths. Consequently, the evolution of mass and inner radius, derived from the dust temperature, is not well understood (Najita \\& Williams 2005).  Far-infrared surveys provide a good starting point to define a target list for submillimeter observations, although the coldest disks may be missed (Wyatt, Dent, \\& Greaves 2003; Najita \\& Williams 2005). A recent list of 58 nearby ($< 100$\\,pc) debris disk candidates was published by Zuckerman \\& Song (2004) based on an extensive search of the IRAS Faint Source Catalog by Silverstone (2000). Most of their candidates have not been observed beyond 100\\,\\micron\\ but many have large excesses and their SEDs extrapolate to detectable fluxes in the submillimeter regime.  A program to obtain 850\\,\\micron\\ photometry of stars in the Zuckerman \\& Song catalog began in Spring 2004 with the intent to survey as many sources with large predicted fluxes as possible. The unfortunate retirement of SCUBA in the summer of 2005 ended the project after only 8 sources had been observed. Nevertheless, the truncated survey yielded a very high (75\\%) detection rate and significant constraints on the nature of the dust toward the 2 other sources. This short paper describes the observations in \\S2, the spectral energy distributions of the 8 candidate debris disks in \\S3, discusses the results and implications in \\S4, and concludes in \\S5. ",
        "conclusions": "We have carried out 850\\,\\micron\\ photometry of 8 nearby stars to search for thermal emission from cold circumstellar dust. With one exception, the rms noise levels of these SCUBA observations lie between $\\sigma=1.5-2.3$\\,mJy. 5 sources were detected at $3\\sigma$ significance or higher and an additional one at $2.5\\sigma$. The high detection rate is attributable to their selection based on high IRAS 60 and 100\\,\\micron\\ excesses.  These 6 sources are a significant addition to the small number of debris disks detected at submillimeter wavelengths. The inferred dust masses range from $0.033-0.23\\,M_\\oplus$. Several are some of the most massive disks known for their age and show that even the upper end of the disk mass distribution may not have been fully explored. Dust temperatures range from $43-65$\\,K, and are similar to other debris disks. There is no correlation of either temperature of equivalent black-body radius with stellar luminosity or age.  Far-infrared and submillimeter observations are generally unable to resolve the emission from the cold dust. Flux measurements, or even sensitive upper limits, on the Rayleigh-Jeans side of the SED regime constrain the dust opacity index and are important for distinguishing between a disk or interstellar (most likely cirrus) origin. The former applies for the 6 detections which have $\\beta < 1$ and the latter for the 2 non-detections where the combination of high far-infrared and low submilllimeter fluxes constrain $\\beta > 1.5$.  The photospheric excesses are well fit by single temperature greybodys implying that the bulk of the dust, as measured by mass, lies far from the star, ranging from 42 to 84\\,AU for each source or about $1-2''$ in angular size. Two of the confirmed disks have high fractional luminosities and may be detectable in scattered light. An additional 2 are strong enough to image with the current generation of (sub)millimeter interferometers. Such high resolution observations may show clumping in the disk due to planetary induced resonances.  There are undoubtedly more sources in the Zuckerman \\& Song catalog, and perhaps others in the plethora of new Spitzer results (e.g., Bryden et al. 2006), that would be detectable if SCUBA were available. Its successor, SCUBA-2, will be about a factor of 3 times more sensitive for photometry and far quicker for mapping (Holland et al. 2003). It will distinguish between cirumstellar and interstellar emission and provide mass measurements and dust opacities for perhaps as many as an order of magnitude more debris disks."
    },
    "0609/astro-ph0609472_arXiv.txt": {
        "abstract": "We analyze the historical light curve of the symbiotic star YY Her, from 1890 up to December 2005. A secular declining trend is detected, at a rate of $\\sim $.01 magn in 1000 d, suggesting that the system could belong to the sub-class of symbiotic novae. Several outburst events are superposed on this slow decline. Three independent periodicities are identified in the light curve. A quasi-periodicity of 4650.7 d is detected for the outburst occurrence. We suggest that it is a signature of a solar-type magnetic dynamo cycle in the giant component. A period of 593.2 d modulates the quiescent light curve and it is identified as the binary period of the system. During outburst events the system shows a stable periodic oscillation of 551.4 d. We suggest that it is the rotation period of the giant.The secondary minima detected at some epochs of quiescence are probably due to dark spots on the surface of the rotating giant.  The difference between the frequencies of these two last periods is the frequency of a tidal wave in the outer layers of the giant. A period which is a beat between the magnetic cycle and the tidal wave period is also apparent in the light curve. YY Her is a third symbiotic system exhibiting these cycles in their light curve, suggesting that a magnetic dynamo process is prevalent in the giant components of symbiotic stars, playing an important role in the outburst mechanism of some of these systems. ",
        "introduction": "The long-term light curve (LC) of a symbiotic star (SS) is quite irregular, and shows phases of quiescence, with quasi regular brightness oscillations  and phases of activity. For a few  SS's, records of their photometric behaviour, going back about a century, are available. The inspection of these historical LC's reveals a complex photometric variability both during quiescence and during outburst epochs. Analysis of these  LC's can be used as a tool for gaining insight into the symbiotic nature  and may help revealing  properties of the cool component of these binary systems.  Formiggini \\& Leibowitz (1994) analyzed the historical LC of Z And and discovered  a $\\sim $8400 d period for the outburst activity, beside the 758.8 d binary one. During the outburst phases, a period of  $\\sim $656 d is also present. In BF Cyg, Leibowitz \\& Formiggini (2006, hereafter paper I) detected, in addition to the well known 757.3 d binary period of the system, a 6376 d cycle for the occurrence of the outbursts. Another significant periodicity of 798 d  was detected for this system, and was interpreted as the rotation period of the giant component of this binary system.  The discovery that outburst events occur with a constant time interval between them is relevant to the understanding of the origin of the outburst phenomenon and of the nature of the clock that regulates their appearance. In paper I we suggested that a magnetic dynamo process, similar to the well known solar cycle, can be the mechanism that regulates the activity events of BF Cyg.   In this paper we analyze the historical LC of  YY Her, which is among the prototype symbiotic stars. Its giant component is a M4 star (M\\H{u}rset \\& Schmid 1999). Its optical spectrum shows strong TiO bands, but no radial velocity data are available (Kenyon 1986). The nature of the hot component is not well established. Numerous strong emission lines are detected in the {\\em IUE ( International Ultraviolet Explorer)} spectra. The flat UV continuum  can be fitted by a hot main sequence accretor (Kenyon \\& Webbink 1984). The photometric history of YY Her since 1890 has already  been analyzed by Munari et al. (1997), and a periodic  590 d fluctuation was detected. Several outburst events are known for YY Her. After the 1997/98 event,  secondary minima appear in the B, V, R photometric data, but not in the U light curve (Hric, Petrik \\& Niarchos 2001). Ellipsoidal variations of the red giant have been suggested as the origin of these secondary minima (Mikolajewska et al. 2002).  In this work we analyze a 115 years historical LC of YY Her with a procedure that is similar to the one used for BF Cyg in paper I. Although  the photometric data available for YY Her are less regular, the results are so similar to those found for BF Cyg, that the two stars may be considered as nearly twin systems.    In Section 2 we describe the data sets used in our analysis and the strategy adopted for merging the data sets and constructing a time-series adequately scaled along the whole time interval of the available observations. In Section 3 we present the time-series analysis techniques used in the search for  periodicities and the periods detected in that analysis. We propose an interpretation of the detected periodicities and of the appearance of the secondary minima. In Section 4 we compare the characteristics of YY Her to those of BF Cyg and discuss some possible implications on the nature of the giant star in symbiotic systems.   \\section { The long-term light curve of YY Her}   The previous study of the photometric history of YY Her by Munari et al. (1997) is based on two data sets. The first one is of photographic magnitudes  from several sources and the second one is a merging of visual estimates retrieved from the  AAVSO and AFOEV data banks, after correcting for possible systematic errors. Several outburst events are evident in this LC but the different scale of the photographic and visual samples precludes  a whole overview of the behaviour of this system (see Fig. 3 of Munari et al. 1997).  The photographic magnitudes assembled in Table 6 of Munari et al. (1997) are sparse data mainly from the Harvard plates collection, and more frequent measurements from the Sonnenberg Sky Patrol archive, starting from JD 2434099 (1952) up to JD 2439059 (1965). We use these data such as published as representative of the ancient LC of YY Her. To these we added  the large AAVSO data set of confirmed visual estimates up to November 2001 and converted them into means of 10 d.   The last validated data from AAVSO correspond to November 2001. In order to obtain the LC  up to the present days, we retrieved the data from the AFOEV data-bank. We averaged these visual estimates over a time interval of 10 d, excluding the upper limits or the uncertain values. A systematic offset  of .028 magn between this set and the 10-d binned set from the of AAVSO was estimated comparing the data in the overlapping time interval. Scaling the data we were able to continue the  LC for YY Her up to December 2005.    \\begin{figure} \\includegraphics[width=85mm]{rfig1.eps} \\caption{A 115 yr LC of  YY Her, from the year 1890 up until Dec 2005. Dots refer to {\\em m}$_{pg}$, crosses indicate {\\em m}$_{v}$ data  from AAVSO and stars are scaled {\\em m}$_{v}$ from  AFOEV. Solid line is a linear regression of the quiescence data, in  each of the two data sets.} \\end{figure}  Fig. 1 is a plot of the entire LC obtained from all our data sources. Two distinct subsets are clearly seen in the figure. In order to establish the {\\em m}$_{pg}$ and the {\\em m}$_{v}$ LC's on a common zero-point level we considered the few measurements that populated the overlapping portions of the two light curves. Comparing the magnitudes of points that are close in their time of observations in the two LC's we derive the  value ({\\em m}$_{pg}$-{\\em m}$_{v}$)= .90.  A more accurate value has been obtained by also taking into account  the general trend of the system during quiescence state, in each of the two light curves separately. For this purpose, we eliminated  from both samples the measurements classified as uncertain and all the brightest data belonging to epochs of activity. Applying a linear regression to the quiescence state sections of each of the two light curves, a slow decline with time is detected in both of them. The slope of the two lines in Fig. 1 is .96 10$^{-05}$ and 4.3 10$^{- 06}$ for the {\\em m}$_{pg}$ and the {\\em m}$_{v}$ data respectively.  The fact that a negative slope is measured in the two independent LC's, at two different time intervals, is evidence of the reality of the trend. Considering the statistical uncertainty in the slope, due also to the different time-length of the two data samples, these result are consistent with a general decline at a constant rate throughout the entire time of the observations. The no trend result reported by Munari et al. (1997) is probably due to the relatively short time interval (1969-1993) considered by these authors. Our second subgroup of data points that does show the decline includes twelve additional years, based on AAVSO data.    When the decline trend is taken into account, the scaling factor between the two lines in Fig. 1 in the overlapping regions of the two data sets takes the value  .97 magn, and we applied it to the AAVSO data. We  then calculated a linear regression to the complete sample of quiescent data. The resulting rate of the steady decline of the system from 1890 up to November 2001 is  $\\sim $ .01 magn in 1000 d. The trend of decline detected here, suggests that the system is recovering from a major brightening event, which occurred sometime before  1890, the year of the first recorded magnitudes. A decreasing trend in the long-term LC is one of the characteristic of the sub-class of symbiotic novae such as HM Sge, RR Tel or  BF Cyg. This behaviour is associated with a single major nova-like event that occurred  in these systems, followed by a very slow fading that lasts for more than a century.   The dots in Fig. 2 show the measurements of YY Her from 1890 up to Dec 2005, de-trended for the secular decline as explained above. The total time covered by the observations is \\simi 115.6 yr, but the distribution of data within this time range is not homogeneous. Due to the scarcity of points of the very old observations, in the time series analysis that we applied to the data we have disregarded all points prior to JD 2420000. Therefore the analysis is applied to a 92.37 years LC. Several outbursts, with different amplitudes, are clearly seen in this LC, although only the last two  are well covered by frequent measurements. We shall address the bias due to the non homogeneous distribution of points in the LC that we have analyzed.     \\begin{figure} \\includegraphics[width=80mm]{rfig2.eps} \\caption{The de-trended LC of  YY Her. The line is the artificial LC calculated as explained in Section 3.5.} \\end{figure} ",
        "conclusions": "A 115 year light curve of the symbiotic star YY Her was analyzed. The mean optical brightness of the system was found to be declining at a constant rate of $\\sim $ .01 magn in 1000 d. This suggests that YY Her should be reclassified as a symbiotic nova. The optical LC of the star is modulated by three independent periods, {\\em P}1=4650 d, {\\em P}2=593.2 d and {\\em P}3=551.4 d. The first one is a mean time interval between successive outbursts of the star, of which 7.5 events have been recorded. We suggest that it is the period of a solar-type magnetic dynamo cycle operating in the outer envelope of the giant star of this system. {\\em P}2 is the orbital revolution period. It modulates the luminosity of the star at quiescence states of the system. {\\em P}3 is the rotation period of the giant star, modulating the emission during outbursts. At some quiescence epochs it also modulates the quiescence emission, probably through dark spots on the surface of the rotating giant. The difference between the spin and the orbital frequencies excites a tidal wave in the atmosphere of the giant. The LC of the system is also modulated at a frequency that is a beat of the magnetic dynamo frequency and the tidal wave frequency.   YY Her is the third symbiotic star exhibiting outburst events that occur at  nearly constant frequency, with a characteristic repetition time of a few thousands days. Similarly with the star BF Cyg, YY Her shows that tides in the atmosphere of the magnetically active giant star affect significantly the dynamo process operating in this star.  In contrast to BF Cyg, the deviation from locked synchronization in the YY Her system is in the sense that the giant star is spinning with a higher angular velocity than that of the orbital motion. In BF Cyg it is the other way around.   \\begin{tabular}{@{}lcrrrc@{}} \\\\ \\multicolumn {6} {|c|}{\\bf Table 1}{Comparison of the properties of BF Cyg and YY Her}\\\\ \\hline &   & BF Cyg &Ref.  & YY Her & Ref.   \\\\   \\hline \\\\ Giant Sp. Type    &     &  M5 III      & 1  & M4 III  & 1 \\\\ Luminosity L $\\odot$    &  Lo &  5200  & 2    &   1100  & 2 \\\\ Binary period     &  {\\em P}2 &  757.3      & 3    &   593.2 & 4\\\\ Giant Spin period&   {\\em P}3 &  798.8      &3    &   551.4 & 4 \\\\ Tidal wave period&   {\\em P}t & 14580       & 3   & -7825   &4 \\\\ Solar-type period &  {\\em P}1 &  5375       & 3   & 4650    & 4 \\\\ Beat period       & {\\em P}b  &  4436       & 3   & 2917    & 4 \\\\ \\\\ \\hline \\end{tabular} 1) M\\H{u}rset \\& Schmid (1999) 2) M\\H{u}rset et al. (1991) 3) Leibowitz \\& Formiggini (2006) 4) This paper"
    },
    "0609/astro-ph0609191_arXiv.txt": {
        "abstract": "We present a study of the galaxy population predicted by hydrodynamical simulations of galaxy clusters. These simulations, which are based on the {\\small GADGET-2} Tree+SPH code, include gas cooling, star formation, a detailed treatment of stellar evolution and chemical enrichment, as well as SN energy feedback in the form of galactic winds. As such, they can be used to extract the spectro--photometric properties of the simulated galaxies, which are identified as clumps in the distribution of star particles. Simulations have been carried out for a representative set of 19 cluster--sized halos, having mass $M_{200}$ in the range $5\\times 10^{13}$--$1.8\\times 10^{15}\\msun$. All simulations have been performed for two choices of the stellar initial mass function (IMF), namely using a standard Salpeter IMF with power--law index $x=1.35$, and a top--heavy IMF with $x=0.95$. In general, we find that several of the observational properties of the galaxy population in nearby clusters are reproduced fairly well by simulations. A Salpeter IMF is successful in accounting for the slope and the normalization of the color--magnitude relation for the bulk of the galaxy population. In contrast, the top--heavy IMF produces too red galaxies, as a consequence of their exceedingly large metallicity. Simulated clusters have a relation between mass and optical luminosity which generally agrees with observations, both in normalization and slope. Also in keeping with observational results, galaxies are generally bluer, younger and more star forming in the cluster outskirts. However, we find that our simulated clusters have a total number of galaxies which is significantly smaller than the observed one, falling short by about a factor 2--3. We have verified that this problem does not have an obvious numerical origin, such as lack of mass and force resolution. Finally, the brightest cluster galaxies are always predicted to be too massive and too blue, when compared to observations. This is due to gas overcooling, which takes place in the core regions of simulated clusters, even in the presence of the rather efficient supernova feedback used in our simulations. ",
        "introduction": "\\label{sec:intro}  A key question in the study of the formation and evolution of galaxies concerns the relationship between their observational properties and the large--scale cosmological environment. In recent years, a flourishing of observational campaign has provided a detailed description of the evolution of the galaxy population in clusters. Indeed, galaxy clusters play a key role in the characterization of the galaxy evolution. Each cluster provides a large sample of galaxies, all placed at the same redshift. Furthermore, clusters offer the possibility of sampling a variety of environments, from their dense core regions, to the outskirts where the properties of the cluster galaxy population tends to approach that of the field.  A diversity of the galaxy population in nearby clusters, with respect to that in the field, was noticed already by \\cite{1974ApJ...194....1O} and by \\cite{1980ApJ...236..351D}. Rich clusters were shown to contain a higher fraction of bulge--dominated (early type and S0) galaxies, and a correspondingly lower fraction of star forming galaxies, than poor systems. \\cite{1978ApJ...226..559B} noticed that moderately distant clusters ($z\\sim 0.3$) have a galaxy population which is generally bluer than nearby clusters. This result has been subsequently extended by a number of analyses and is currently interpreted as high redshift clusters having larger fractions of star--forming galaxies than local clusters \\citep[see ][for a review]{2004bdmh.confE.104P}. \\cite{1992MNRAS.254..601B} first noticed that early--type members of the Coma cluster lie on a tight relation in the color--magnitude diagram, a result that has been subsequently confirmed by a number of analyses for extended samples of nearby \\citep[e.g., ][and references therein]{1996A&A...309..749P,2003A&A...409...37A,2004ApJS..153..397B,2004ApJ...614..679L,2005ApJ...619..193M} and distant clusters, out to $z\\magcir 1$ \\citep[e.g., ][and references therein]{2000ApJ...541...95V,2003ApJ...596L.143B,2004MNRAS.353..353A,2006astro.ph..1165S}.  As for the luminosity function (LF hereafter) of cluster galaxies, much work has been done in recent years, with different groups reaching different conclusions as for its universality, shape and faint--end slope \\citep[e.g., ][for a review]{2004PASA...21..344D}. For instance, \\cite{2006A&A...445...29P} have recently analyzed the LF for sample of clusters identified in the X--ray band within the ROSAT All Sky Survey (RASS, \\citealt{VO92.1}) and covered by the Sloan Digital Sky Survey (SDSS \\citealt{2002AJ....123..485S}). They found no significant cluster to cluster variations of the LF, once calculated within the same physical radius ($r_{200}$ or $r_{500}$), with a shape fitted by a double Schechter function to account for the upturn at faint magnitudes. However, these results are at variance with respect to those from other analyses \\citep[e.g., ][]{2000A&A...353..930A}.  Within the widely accepted standard $\\Lambda$CDM cosmological scenario, galaxies arise from the hierarchical assembly of dark matter (DM) halos. The gravitational dynamics of these halos is relatively simple to describe to high precision with modern large supercomputer simulations \\citep[e.g.,][]{2005Natur.435..629S}. However, the observational properties of galaxies are determined by the combined action of the assembly of DM halos and by the physical processes which define the evolution of the cosmic baryons. A complex interplay between radiative gas cooling, star formation, chemical enrichment and release of energy feedback from supernovae (SN) and active galactic nuclei (AGN) is expected to determine the properties of the stellar population in galaxies. At the same time, the cluster environment is expected to play a significant role in altering the evolution of galaxies. For instance, ram pressure exerted by the hot intra-cluster medium (ICM) can lead to the removal of a substantial fraction of the interstellar medium (ISM; \\citealt{1972ApJ...176....1G}), thereby affecting galaxy morphology, star formation and luminosity \\citep[e.g., ][, and references therein]{1999MNRAS.308..947A,2004AJ....127.3361K}.  In this context, semi--analytical models of galaxy formation have been used since several years as a flexible tool to study galaxy formation within the cosmological hierarchical framework \\citep[e.g., ][and references therein]{1993MNRAS.264..201K,1999MNRAS.310.1087S,2000MNRAS.319..168C,2002ApJ...575...18M}. A powerful implementation of this method is that based on the so--called hybrid approach, which combines N--body simulations, to accurately trace the merging history of DM halos, and semi--analytic models to describe the physics of the baryons \\citep[e.g., ][]{1999MNRAS.303..188K}.  \\cite{SP01.2} applied this method to a DM simulation of a cluster, with high enough resolution to allow them resolving the population of dwarf galaxies. As a result, they found that several observational properties (e.g., luminosity function, mass-to-light ratio and morphological types) are rather well reproduced. \\cite{2001MNRAS.323..999D} applied a semi--analytical model to a DM simulation of a large cosmological box, with the aim of performing a combined study of kinematics, colors and morphologies for both cluster and field galaxies. They concluded that a good agreement with observations holds for cluster galaxies, while colors and star formation rates of field galaxies were shown to evolve more rapidly than observed. \\cite{2006astro.ph..6591C} applied the same semi--analytical model to a constrained simulation of the local universe and concluded that significant differences exist between the observed and the predicted properties of the large--scale distribution of galaxy groups. \\cite{2004MNRAS.349.1101D} incorporated in their model also a scheme of metal production to follow the enrichment of ICM and galaxies (see also \\citealt{2006MNRAS.tmp..453C}). Among their results, they found that the color--magnitude relation (CMR) is mainly driven by metallicity effects, the redder galaxies on the sequence being on average the more metal rich.  \\cite{2005MNRAS.361..369L} applied their semi--analytical model to a set of DM cluster simulations. They also included a prescription to account for the effect of ram--pressure stripping of the ISM as the galaxies move in the hot cluster atmosphere, and found it to have only a very little effect on the galaxy population.  A complementary approach to the semi--analytical models is represented by using full hydrodynamical simulations, which include the processes of gas cooling and star formation. The clear advantage of this approach, with respect to semi--analytical models, is that galaxy formation can be now described by following in detail the evolution of the cosmic baryons while they follow the formation of the cosmic web. However, the limitation of this approach lies in its high computational cost, which prevents it to cover wide dynamic ranges and to sample in detail the parameter space describing the processes of star formation and feedback. For these reasons, describing the process of galaxy formation with a self--consistent hydrodynamic approach within the typical cosmological environment of $\\sim 10\\,$Mpc, relevant for galaxy clusters, represents a challenging task for simulations of the present generation.  In a pioneering paper, \\cite{1994ApJ...437..564M} studied the effect of including galaxies for the energy feedback and chemical enrichment of the ICM. Since these simulations did not have enough resolution to identify galaxies, they have been placed by hand, identifying them with the peaks of the initial density field. \\cite{1996ApJ...472..460F} used for the first time a radiative simulation of a cluster and identified galaxies as concentrations of cooled gas. The aim of their study was to compare the dynamics of member galaxies to that of DM particles. They concluded that galaxies suffer for a substantial dynamical bias, a result which has not been confirmed by more recent hydrodynamical simulations \\citep[e.g., ][]{2005MNRAS.358..139F,2006astro.ph..5151B}.  Thanks to the ever improving supercomputing capabilities and efficiency of simulation codes, a number of groups have recently completed hydrodynamical simulations of galaxy clusters, which have good enough resolution to trace the galaxy population with better reliability. \\cite{2005ApJ...618..557N} used simulations of eight groups and clusters, performed with an adaptive mesh refinement code, including star formation, feedback from supernovae and chemical enrichment, to describe the spatial distribution of galaxies inside clusters. They found that galaxies are more centrally concentrated than DM sub--halos, with their number density profile described by a NFW shape \\citep{NA96.1}, although with a smaller concentration parameter than for the DM distribution. \\cite{2005MNRAS.361..983R} used SPH simulations of two clusters, including a similar physics, used spectrophotometric code to derive galaxy luminosities in different bands. They analyzed the resulting color--magnitude relations (CMR, hereafter) and luminosity functions, claiming for an overall general agreement with observations.  In this paper, we will present a detailed analysis of the galaxy population for a set of 19 simulated clusters, which span the mass range from $\\simeq 5\\times 10^{13}\\msun$ to $\\simeq 2\\times 10^{15}\\msun$. The simulations have been performed with the Tree-SPH code {\\small GADGET-2} \\citep{2005MNRAS.364.1105S}. They include the effect of radiative cooling, an effective model for star formation from a multiphase ISM \\citep{2003MNRAS.339..289S}, a phenomenological recipe for galactic winds, a detailed stellar evolution model, thereby accounting also for life--times and metal production from different stellar populations (\\citealt{2004MNRAS.349L..19T}; Tornatore et al. in preparation). These simulations have been carried out also by varying both the shape of the initial mass function (IMF, hereafter) and the feedback strength. The inclusion of a detailed model of chemical enrichment allows us to compute luminosities and colors for galaxies of different metallicities, by using the GALAXEV spectrophotometric code \\citep{2003MNRAS.344.1000B}. In our analysis we will concentrate on the properties of galaxy clusters at $z=0$. As we shall discuss through the paper, several observational trends are reproduced quite well by our simulations, although a number of significant discrepancies are found. For this reason, the aim of our analysis will be more that of understanding the directions to improve simulations, rather than seeking for a best fitting between model predictions and observations.  The plan of the paper is as follows. In Section 2 we provide the general characteristics of the simulated clusters and describe the relevant features of the {\\small GADGET-2} version used for this analysis. In Section 3 we will describe the method of galaxy identification and how luminosities in different bands are computed. Section 4 contains the description of the properties of the simulated galaxy population and their comparison with observational data. In particular, we will discuss the radial distribution of galaxies, the CMR, the mass--luminosity ratio, the luminosity function, the star formation rate and the color and age gradients. Our main results will be summarized and discussed in Section 5. We will discuss in an Appendix the effects of numerical resolution on the stability of the results of our analysis. ",
        "conclusions": "\\label{sec:conc} \\noindent In this paper we have presented an analysis of the galaxy population in cosmological hydrodynamical simulations of galaxy clusters at $z=0$. The inclusion of a detailed treatment of stellar evolution and chemical enrichment (\\citealt{2004MNRAS.349L..19T}; Tornatore et al., in preparation) in the {\\small GADGET-2} code \\citep{2005MNRAS.364.1105S} has allowed us to derive the properties of galaxies in the optical/near--IR bands. In our simulations each star particle is treated as a single stellar population (SSP) characterized by a formation time and a metallicity. Based on this, we apply the GALAXEV spectro-photometric code \\citep{2003MNRAS.344.1000B} to compute luminosities in different bands. Simulations have been carried out for a representative set of galaxy clusters, containing 19 objects with mass $M_{200}$ ranging from $5\\times 10^{13}\\msun$ to $1.8\\times 10^{15}\\msun$. All clusters have been simulated assuming both a standard Salpeter IMF \\citep{1955ApJ...121..161S} and a power--law top--heavier IMF with exponent $x=0.95$ \\citep{1987A&A...173...23A}. The main results of our analysis can be summarized as follows. \\begin{description} \\item[1.] Both the color--magnitude relation (CMR) and the $M/L$ ratio are in reasonable agreement with observational data for a Salpeter IMF. In contrast, using a top--heavy IMF provides too high a metallicity of galaxies, which turns into too red colors. This spoils the agreement with the CMR and with the $M/L$ ratio in the $K$ band. The CMR is confirmed to be a metallicity sequence, in the sense that more enriched galaxies systematically populate the brighter redder part of the sequence. \\item[2.] Galaxies are systematically older and redder in central cluster regions, thus in keeping with observational results. This trend extends at least out to $2r_{vir}$, thus showing that the galaxy population feels the presence of a cluster well beyond its virial region. Due to the overcooling occurring in the central cluster regions, BCGs are always much bluer and more massive than observed, and characterized by too high a recent star formation. Indeed, neglecting the contribution of stars formed at $z<1$ produces significantly redder BCGs, thus in better agreement with observational data. \\item[3.] The number density profile of galaxies is confirmed to steepen with the galaxy stellar mass, approaching the DM profile for the galaxies with $M_*>10^{10}\\msun$. However, when compared to the observations \\citep[e.g.,][]{2006astro.ph..6260P}, the simulation profiles are flatter that the observed ones inside $0.4r_{200}$. \\item[4.] Simulated clusters have about three times fewer galaxies above a given luminosity limit than real clusters. We have verified that this disagreement is directly related neither to lack of mass and force resolution, nor to numerical gas heating due to a non optimal choice of the softening parameter. This leaves as further possibilities more subtle numerical effects or a better suited implementation of energy feedback. \\item[5.] The luminosity function (LF) is shallower than the observed one in the bright end, thus confirming that feedback is not strong enough to suppress cooling in the most massive halos. In the faint end, the LF steepens and indicates the presence of an excess of small red galaxies. Although this result resembles that found in observational data by \\cite{2006A&A...445...29P}, we warn against its overinterpretation, in the light of the deficit of the overall number of galaxies found in the simulated clusters. \\end{description}  The results of our analysis support the capability of hydrodynamical simulations of galaxy clusters to reproduce the general trends characterizing their galaxy population. Therefore, simulations in which the properties of the stellar population are self--consistently predicted from the gas dynamics, can provide an alternative and complementary approach to semi--analytical methods. However, at present, our simulations have two main limitations in accounting for the observed properties of the cluster galaxy population.  First of all, the deficit of galaxies produced in our simulations suggests that they are not able to produce galaxies which are resistant enough to survive the tidal field of the cluster environment. Indeed, the shallow galaxy number density profile in central cluster regions shows that the problem is more apparent where effects of galaxy disruption are expected to be stronger. This result on the lack of galaxies is at variance with respect to the predictions of semi--analytical models (SAM) of galaxy formation, which produce roughly the correct number of cluster galaxies \\citep[e.g., ][]{2004MNRAS.349.1101D,2005MNRAS.361..369L}. Comparisons between the galaxy populations predicted by hydrodynamical simulations and by SAM have shown a reasonable level of agreement. However, these comparisons have been always performed by excluding the effect of energy feedback \\citep[e.g., ][]{2003MNRAS.338..913H,2006astro.ph..5750C} or explicit conversion of cooled gas particles into collisionless stars \\citep[e.g., ][]{YO02.1}.  Furthermore, we always find that the brightest cluster galaxies (BCGs) are much bluer and star forming that observed. While the adopted feedback scheme, based on galactic winds, is efficient in regulating star formation for the bulk of the galaxy population, it is not able to quench low--redshift star formation in the central cluster regions, to a level consistent with observations. Clearly, the required feedback mechanism should be such to leave the bulk of the galaxy population unaffected while acting only on the very high end of the galaxy mass distribution. Feedback from central AGN represents the natural solution and provides in principle a large enough energy budget \\citep[e.g., ][]{2006astro.ph..5323R}. Its detailed implementation in cosmological hydrodynamical simulations require understanding in detail the mechanisms for the thermalization of this energy in the diffuse medium \\citep[e.g., ][]{2005A&A...429..399Z,2006MNRAS.366..397S}.  An ambitious goal for hydrodynamical simulations of the next generation will be that of describing in detail the complex interplay between the history of star formation and the thermodynamical and chemical evolution of the diffuse cosmic baryons. For this to be reached, it is mandatory that simulations are able to produce a realistic population of galaxies, both in terms of color and of luminosity and mass distribution. Therefore, simulation codes will be required on the one hand to have numerical effects under exquisite control and, on the other hand, to include physically motivated schemes of energy and metal feedback."
    },
    "0609/astro-ph0609158_arXiv.txt": {
        "abstract": "We investigate the growth or decay rate of the fundamental mode of even symmetry in a viscous accretion disc.  This mode occurs in eccentric discs and is known to be potentially overstable.  We determine the vertical structure of the disc and its modes, treating radiative energy transport in the diffusion approximation. In the limit of very long radial wavelength, an analytical criterion for viscous overstability is obtained, which involves the effective shear and bulk viscosity, the adiabatic exponent and the opacity law of the disc. This differs from the prediction of a two-dimensional model. On shorter wavelengths (a few times the disc thickness), the criterion for overstability is more difficult to satisfy because of the different vertical structure of the mode.  In a low-viscosity disc a third regime of intermediate wavelengths appears, in which the overstability is suppressed as the horizontal velocity perturbations develop significant vertical shear.  We suggest that this effect determines the damping rate of eccentricity in protoplanetary discs, for which the long-wavelength analysis is inapplicable and overstability is unlikely to occur on any scale.  In thinner accretion discs and in decretion discs around Be~stars overstability may occur only on the longest wavelengths, leading to the preferential excitation of global eccentric modes. ",
        "introduction": "Accretion discs may support a wide variety of wave motions.  The basic epicyclic tendencies of orbiting particles, which result from gravitational and inertial forces, can combine with acoustic and buoyancy effects to allow many different modes of oscillation.  In some cases self-gravitational or magnetic forces may also be significant.  Waves in accretion discs are usually analysed without regard to the shear stress that is necessarily present in order to transport angular momentum and facilitate accretion.  This stress, probably of turbulent origin, is most often modelled as deriving from an effective viscosity, which might be expected always to damp the waves.  However, \\citet{K78} found that `pulsational instability' can occur under certain circumstances, by analogy with the `overstability' of stellar oscillations \\citep{E26}.  In particular, if the stress responds to wavelike perturbations in a certain way, a viscous overstability is possible, in which some of the energy being drawn by the stress from the differential rotation of the disc is diverted into waves of growing amplitude.  This effect has long been discussed as a means of exciting disturbances in accretion discs and planetary rings \\citep[e.g.][]{KF80,BGT85,PS86,KHM88,PL88,ST99}.  The stresses in accretion discs and planetary rings are probably not accurately described as arising from an isotropic viscosity in the sense of the Navier--Stokes equation.  Nevertheless, the general principle that overstability may occur as the stress responds to wavelike perturbations applies to all discs, even if the details are difficult to quantify.  Waves in thin discs are often studied in a two-dimensional approximation in which motion in the vertical ($z$) direction, perpendicular to the orbital plane, is absent and the disc remains in vertical hydrostatic balance.  This approximation is highly convenient but can only be justified under exceptional circumstances and indeed eliminates almost all of the wave modes.  The original analysis of \\citet{K78} was in fact based on a three-dimensional disc, but at the time of writing of his paper the behaviour of waves in such discs was not fully understood.  Subsequently it was shown that a thin disc acts as a waveguide that conducts in the radial direction a set of discrete modes having distinct vertical structures determined from an eigenvalue problem \\citep{OK85,L86,LP93,KP95}.  Some of these can be classified as f, p and g~modes by direct analogy with stellar oscillations \\citep{O98}. Generally the propagation of these waves is restricted to certain intervals of radius because they encounter turning points at Lindblad or other resonances.  Perhaps of greatest interest among these is the $\\fe$~mode, the fundamental mode of even symmetry about the midplane, in the classification of \\citet{O98}.  This mode is the closest equivalent of the inertial--acoustic wave (or `density wave') that appears in the analysis of a two-dimensional disc.  Its eigenfunction involves horizontal velocity perturbations which are symmetric about the midplane and have no nodes in their vertical structure, but which in general do vary significantly with $z$ and are accompanied by vertical velocity perturbations.  This is by far the dominant mode excited by tidal forcing at Lindblad resonances, and its excitation, propagation and dissipation have been investigated by \\citet{LO98}, \\citet{OL99} and \\citet{BOLP02}.  The $\\fe$~mode is also the obvious candidate for viscous overstability.  However, \\citet{K78} assumed that the horizontal velocities would be independent of $z$, which is not true of the $\\fe$~mode.  A self-consistent linear calculation of overstable modes in a three-dimensional disc has never been carried out, though the nonlinear behaviour of axisymmetric dics with full vertical structure has been simulated by Kley et al. (1993). They report that when $\\alpha$ is larger than a small value the disc exhibits global, axisymmetric, overstable modes confined to the outer edge of the computational domain.  The question of overstability also arises in the theory of eccentric discs.  A small eccentricity can be regarded as a special type of wave motion, corresponding to azimuthal wavenumber $m=1$, propagating in a circular disc.  Although the viscous overstability has been discussed mainly for axisymmetric waves, the local dynamics of waves of modest azimuthal mode number in a thin disc is equivalent to that of axisymmetric waves.  Studies of eccentric viscous discs in two and three dimensions have indeed found overstability \\citep{LPP94,O01}, meaning that the eccentricity often has a negative diffusion coefficient and would grow spontaneously and preferentially on small radial scales (comparable to the disc thickness).  This is contrary to the intuitive concept that viscosity should damp eccentricity. It is important to note that $m=1$ waves propagating in a (nearly) Keplerian disk are quite exceptional in that they can maintain a long radial wavelength over an extended radial distance.  Understanding the growth or decay rate of eccentricity in a disc is of great importance.  For example, in protoplanetary systems, eccentricity is a joint property of the planets and the disc since it is exchanged through secular and resonant interactions.  The behaviour of eccentricity within the disc therefore affects the eccentricity of all the planets as well.  The decretion discs around Be~stars appear to have global eccentric modes which may be excited by the viscous overstability \\citep{K83,O91}.  However, since the overstability is believed to grow preferentially on a wavelength comparable to the disc thickness it may be difficult to account for the appearance of coherent global modes.  It was shown by \\citet{O01} that the criterion for overstability is different in two- and three-dimensional discs, and that it can be suppressed either by introducing a sufficiently large bulk viscosity or by using a more sophisticated model of the turbulent stress that involves a sufficiently large relaxation time.  \\citet{LO06} have also shown that the viscous overstability is suppressed in dilute planetary rings when the non-Newtonian nature of the stress is taken into account.  The analysis of eccentric discs by \\citet{O01} assumes that the eccentricity is independent of $z$, which is correct if sufficiently large stresses are present to couple different strata in the disc.  In a low-viscosity disc, however, the eccentricity could vary with $z$ and might be better described by the structure of the $\\fe$~mode.  We are led, therefore, to investigate the behaviour of wave modes in a viscous disc.  In Section~\\ref{s:2d} we analyse a simple two-dimensional model by way of introduction, while the full three-dimensional model is presented in Section~\\ref{s:3d}.  We describe the methods of numerical solution in Section~\\ref{s:numerical} and the results follow in Section~\\ref{s:results}.  We conclude with a summary and discussion in Section~\\ref{s:conclusion}. ",
        "conclusions": "\\label{s:conclusion}  We have investigated the growth or decay rate of the fundamental mode of even symmetry in a three-dimensional viscous accretion disc.  Our numerical model resolves the vertical structure of the disc and its modes, treating radiative energy transport in the diffusion approximation.  Overstability can occur, in principle, for radial wavelengths longer than a few times the disc thickness.  It results almost entirely from the fact that the unperturbed disc necessarily has a shear stress in order to facilitate accretion, and the variation of this stress in the wave motion allows energy to be fed from the differential rotation into growing oscillations.  This tendency competes with the viscous damping of the mode.  Smaller contributions to the growth rate arise from non-adiabatic effects related to the $\\epsilon$- and $\\kappa$-mechanisms in stars, and thermal diffusion.  The pattern of behaviour in a low-viscosity disc ($\\alpha\\ll1$) is as follows.  For radial wavenumbers $kH\\sim1$ (i.e.~radial wavelengths a few times the disc thickness) the structure of the $\\fe$~mode resembles its counterpart in an inviscid disc.  It has a non-trivial vertical structure dictated by the stratification of the disc; the horizontal velocities involve vertical shear and are accompanied by vertical motion.  The mode can be overstable in this short-wavelength regime, depending on the parameters of the disc, but the presence of vertical shear and vertical motion acts against this tendency.  For very long radial wavelengths ($kH\\la\\alpha$) the mode is dominated by horizontal motion and the vertical shear is removed by the action of viscous stresses between different strata in the disc.  This type of motion is optimal for exciting the overstability but the (complicated) analytical criterion we obtain in this regime differs from the prediction of a two-dimensional model because the disc is not in vertical hydrostatic or thermal balance.  On intermediate wavelengths ($\\alpha\\ll kH\\ll1$) a transitional behaviour is found.  Here the vertical shear developing in the horizontal velocity perturbations leads to an enhanced viscous damping and the overstability is suppressed.  In order to bring out this subtle behaviour, we have restricted our attention in this paper to viscous models of accretion discs based on the Navier--Stokes equation and the alpha viscosity prescription.  We fully appreciate that stress in accretion discs is likely to arise from magnetohydrodynamic turbulence and to exhibit a more complicated rheology \\citep[e.g.][]{O01,O03}.  In particular, the non-zero relaxation time of the stress is known to act against the overstability \\citep{O01} as it does in planetary rings \\citep{LO06}. In this paper we have employed a bulk viscosity as a convenient surrogate for any such stabilizing tendencies.  A major application of this work is to the theory of eccentric accretion discs, since a small eccentricity can be regarded as a wave of azimuthal wavenumber $m=1$ and long radial wavelength propagating in a circular disc. Indeed, it is only for $m=1$ waves in (nearly) Keplerian discs that a long radial wavelength can be maintained over an extended radial distance. Overstability appears as a negative diffusion coefficient for eccentricity in the theory of \\citet{O01}, which uses approximations similar to the long-wavelength regime described above. The approximations break down for intermediate wavelengths in a very low-viscosity disc, because the viscous stresses are inadequate to couple the different strata and enforce uniformity of the eccentricity with height.  While the overstability may be suppressed by stress relaxation effects, as suggested by \\citet{O01}, an alternative possibility is that overstability may indeed occur on very long radial wavelengths, while shorter wavelengths require a treatment similar to that carried out here.  In protoplanetary discs the stress is of uncertain origin, but values in the range $10^{-4}\\la\\alpha\\la10^{-2}$ are usually quoted based on observed accretion rates and disc lifetimes, while the angular semithickness is in the range $0.05\\la H/r\\la0.1$.  In this case the long-wavelength analysis has no application, since the relevant wavelengths greatly exceed the size of the disc.  Instead, the regimes of short and intermediate wavelengths describe the behaviour of global eccentric modes.  This means that a correct determination of the the damping rate of eccentricity in a protoplanetary system requires a three-dimensional treatment of the disc, allowing for vertical shear in the horizontal velocities. We find that a typical decay rate of a global eccentric mode (with a radial wavelength of $20H$) is $10^{-5}\\,\\Omega^{-1}$ for $\\alpha=10^{-3}$ and $2\\times 10^{-6}\\,\\Omega^{-1}$ for $\\alpha=10^{-4}$. These are only rought estimates as we have not considered the opacity regimes relevant to protoplanetary discs. For $\\alpha=10^{-2}$ global modes may be overstable.  In other accretion discs $H/r$ may be much smaller and $\\alpha$ somewhat larger.  In this case the long-wavelength regime applies to global eccentric modes.  The fact that the overstability is suppressed for intermediate wavelengths means that only the largest-scale modes may be permitted to grow.  This would neatly account for the occurrence of global eccentric modes in decretion discs around Be~stars.  No modes other than the $\\fe$~mode were found to exhibit overstability.  In particular, a similar mechanism would not permit global warping of discs to be excited."
    },
    "0609/astro-ph0609702_arXiv.txt": {
        "abstract": "The chemical compositions of 26 metal-poor stars that exhibit strong CH and/or C$_{2}$ molecular bands are determined based on high-resolution spectroscopy. We define carbon-enhanced stars taking account of the carbon abundance ratio ([C/Fe]) and the evolutionary status, which is a slight modification over previous definitions. Twenty two stars in our sample satisfy our modified definition for Carbon-Enhanced Metal-Poor (CEMP) stars. In addition, we measure Na abundances for nine other carbon-enhanced stars for which abundances of other elements have been previously reported. Combining our new sample with the results of previous work, we investigate the abundance and evolutionary status of a total of 64 CEMP stars. The following results are obtained: (1) All but one of the 37 stars with [Fe/H] $\\geq -2.6$ exhibit large excesses of barium ([Ba/Fe] $> +0.5$), while the other 27 stars with lower metallicity exhibit a large scatter in their barium abundance ratios ($-1.2 <$ [Ba/Fe]$ < $+3.3). (2) A correlation between the carbon and barium abundance ratios ([C/Fe] and [Ba/Fe]) is found in Ba-enhanced objects (comprising 54 stars), suggesting that the origin of the observed carbon excess in Ba-enhanced stars is nucleosynthesis in asymptotic giant branch (AGB) stars, where the main $s$-process occurs. The correlation between the barium abundance ratio and that of carbon plus nitrogen ([(C+N)/Fe]) is relatively weak, because of the large excesses of nitrogen in some extremely metal-poor stars. (3) The majority of the Ba-enhanced stars have $-1.0 < $ [C/H] $ < 0.0$, and a clear cutoff exists at [C/H] $\\sim 0$, which we take as the limit of carbon-enrichment by metal-poor AGB stars. Within the above range, the [C/H] of the Ba-enhanced CEMP stars decreases, on average, by up to 0.6~dex from the main-sequence turnoff up the red-giant branch, suggesting some dilution of carbon enhancement during their evolution. The [C/H] values of Ba-normal stars are relatively low, with a wide distribution. (4) The difference in the distributions of evolutionary status between Ba-enhanced and Ba-normal CEMP stars suggested by our previous work is not statistically confirmed by the present, enlarged sample. (5) Excesses of Na are found in stars with extremely large enhancements of C, N and Ba, suggesting efficient production of this element by AGB nucleosynthesis. The implications of these results on the origins of carbon in CEMP stars, in particular for Ba-normal stars, are discussed. ",
        "introduction": "\\label{sec:intro}  Large modern surveys of metal-deficient stars in the Galaxy have revealed that the fraction of carbon-enhanced objects is significant at low metallicity \\citep{beers92, christlieb03}.  The fraction appears to increase with decreasing metallicity (Beers \\& Christlieb 2005; Lucatello et al. 2006), although the derived fraction depends somewhat on one's choice of the definition of the carbon-enhancement phenomenon. Many years ago, \\citet{keenan42} coined the term ``CH stars'' to refer to ``high-velocity'' carbon stars that exhibit very strong G-bands due to the CH molecule, and otherwise weak lines of other metals usually used to classify stellar spectra. The small sample of stars he had available at that time had spectra that suggested high luminosities, i.e., membership on the giant branch. The nomenclature was extended to objects of lower luminosity by \\cite{bond74}, who found subgiants with similar elemental compositions. More recently, surveys for carbon stars at high Galactic latitude have also been made \\citep{green94, totten98,totten00,christlieb01}. The Carbon-Enhanced Metal-Poor (CEMP) stars found by recent surveys of metal-poor stars may well be closely related to these objects.  An important question concerning the nature of CEMP stars is that of the astrophysical origin of the carbon excess that is observed in these objects. One established scenario is carbon production by nucleosynthesis in an asymptotic giant branch (AGB) star. Since most CEMP stars that are observed in the present Galaxy are too old to be intrinsic AGB stars\\footnote{There exists one recently-discovered exception, CS~30322--023, which \\citet{masseron06} argue is likely to be an extremely metal-poor AGB star that is presently undergoing thermal pulses. This star also appears in our present sample.}, it is usually assumed that the AGB star responsible for carbon production was once the primary of a binary system, and the elements produced during its lifetime and dredged to its outer atmosphere were transferred to its companion prior to the primary evolving to become a faint white dwarf. Observational support for this scenario is found from the observed excess of $s$-process elements, as well as carbon, in classical CH stars \\citep[e.g. ][]{wallerstein64, vanture92} and some CEMP stars \\citep[e.g., ][]{norris97a}. Other evidence comes from the binarity of the objects found in some CH and CEMP stars identified from radial-velocity monitoring \\citep{mcclure84, mcclure90,preston01}. The fraction of binary stars among CEMP stars has been recently studied by \\citet{lucatello05}, who concluded that all CEMP stars with s-process element overabundances (which Beers \\& Christlieb 2005 refer to as CEMP-s stars) could statistically belong to binary systems, given the limited observational constraints on the radial-velocity variations observed to date.\\footnote{Contributions of the s- and r- processes are estimated from the abundance ratios of neutron-capture elements such as [Ba/Fe], [Eu/Fe], and [Ba/Eu]. For instance, the term 's-process element-enhanced stars' indicates objects that have high [Ba/Fe]($ > +0.5$) and high [Ba/Eu] ($\\sim +1$).}  However, in some CEMP stars, no excess of neutron-capture elements is found \\citep[e.g., ][]{norris97b,aoki02b}; these stars are referred to as CEMP-no stars \\citep{beers05}. At least one CEMP star \\citep[CS~22892--052; ][]{sneden96} has been reported with an excess of neutron-capture elements associated with the r-process. Lucatello et al. (in preparation) has identified a few other such stars, hence the class CEMP-r is no longer based on a single object. There are a growing number of CEMP stars with apparent contributions from {\\it both} the r-process and the s-process, the CEMP-r/s stars (see Beers \\& Christlieb 2005 for further discussion of these classes). The existence of such a wide variety of elemental abundance patterns associated with CEMP stars strongly suggests that other mechanisms may be responsible for their nucleosynthetic history beyond those involving AGB stars.  \\citet{ryan05} compiled the elemental and isotopic abundances of 19 CEMP stars measured by previous work where the $^{13}$C/$^{12}$C ratio was studied, and considered possible differences between Ba-enhanced CEMP stars (CEMP-s) and Ba-normal ones (CEMP-no). This discussion was necessarily somewhat limited because of the small sample size.  Our long-running investigation of these interesting objects has led us to undertake a high-resolution spectroscopic study for an expanded sample of candidate CEMP stars. We have thus far obtained spectra for 26 stars with the Subaru Telescope High Dispersion Spectrograph \\citep{noguchi02}; we report on their chemical compositions in this paper. We also analyze redder spectra, in order to measure their Na abundances, for another nine stars for which abundances of other elements have been reported by previous work. Based on the results of our present abundance study and on previous work, we discuss the distribution of the carbon-enhancement, the fraction of Ba-enhanced stars, and the metallicity distribution among CEMP stars. The evolutionary status of our sample is taken into consideration in our discussion. The sample selection, observations, and data-reduction procedures are reported in \\S 2. Details of the abundance analyses are presented in \\S 3. Our modified definition of carbon-enhanced objects is given in \\S 4. We discuss the abundance distribution of barium and the differences between the Ba-enhanced and Ba-normal stars in \\S 5. In \\S 6 we present our conclusions. ",
        "conclusions": "In this paper we have determined chemical abundances for 26 metal-poor stars, and classified 22 stars among them as Carbon-Enhanced Metal-Poor (CEMP) stars, based on a new definition for the carbon-enhancement phenomenon including the effects of evolution along the giant branch on observed carbon abundances.  Sodium abundances were determined for another nine CEMP stars for which abundances of other elements were previously studied. Combining the results of previous work for 42 CEMP stars and the 22 other CEMP stars in our present sample, we investigated the abundance trends found in CEMP stars and discussed the origins of carbon excesses in these objects. We found that most CEMP stars (more than 80\\%) exhibit excesses of neutron-capture elements (e.g., Ba), suggesting significant contributions of AGB nucleosynthesis. However, the fraction of Ba-normal stars increases at lower metallicity. The difference in the metallicity distribution of Ba-enhanced and Ba-normal CEMP stars found in the present work should be a key to understanding the processes responsible for the Ba-normal CEMP stars.  Another difference between the Ba-enhanced and Ba-normal stars is found in their carbon abundance distributions. While the Ba-enhanced stars have quite high carbon abundances ($-1 \\lesssim $ [C/H] $ \\lesssim 0$), independent of their metallicity, the [C/H] values of Ba-normal CEMP stars distribute over a wide range, while their average value is relatively low. This difference also suggests that the processes responsible for the Ba-normal CEMP stars are quite different from the AGB nucleosynthesis that can explain the chemical abundances of Ba-enhanced CEMP stars.   Some possible interpretations are proposed for the origin of Ba-normal CEMP stars, in particular for the $\\alpha$-element-enhanced stars. Moreover, these stars are likely to be closely related to the two known hyper metal-poor stars with [Fe/H] below $-5$. However, the sample of such stars studied with high-resolution spectroscopy is still quite small (ten objects in the present sample).  Further observations for this class of objects are strongly recommended.  Measurements of abundances for Ba-enhanced CEMP stars based on high-resolution spectroscopy have been rapidly increasing in recent years. Although the abundances of these objects are usually interpreted within the paradigm of AGB nucleosynthesis and mass transfer in binary systems, their detailed abundance patterns are not well explained by current models. One example is the abundance pattern found for objects with very large excesses of neutron-capture elements, which cannot be explained by the present AGB models. Several CEMP stars are known to have enhancements of both r- and s-process elements, the so-called CEMP-r/s stars. More detailed abundance studies, as well as the monitoring of radial velocities to investigate their binarity, are strongly desired. We plan to investigate these problems in future papers in this series."
    },
    "0609/astro-ph0609228_arXiv.txt": {
        "abstract": "We present an application of the fast Independent Component Analysis (\\fstica) to the \\WMAP~3yr data with the goal of extracting the CMB signal. We evaluate the confidence of our results by means of Monte Carlo simulations including CMB, foreground contaminations and instrumental noise specific of each \\WMAP~frequency band. We perform a complete analysis involving all or a subset of the \\WMAP~channels in order to select the optimal combination for CMB extraction, using the frequency scaling of the reconstructed component as a figure of merit. We found that the combination KQVW provides the best CMB frequency scaling, indicating that the low frequency foreground contamination in Q, V and W bands is better traced by the emission in the K band. \\\\ The CMB angular power spectrum is recovered up to the degree scale, it is consistent within errors for all \\WMAP~channel combination considered, and in close agreement with the \\WMAP~3yr results. \\\\ We perform a statistical analysis of the recovered CMB pattern, and confirm the sky asymmetry reported in several previous works with independent techniques. ",
        "introduction": "\\label{intro}  The Cosmic Microwave Background (CMB) anisotropies are at present the main tracer of the physical processes occurred in the very early universe; a complete mapping of them may reveal crucial clues about the statistics of the primordial perturbations, the existence of gravity waves, the global geometry of the universe as well as the physical components responsible for the accelerated expansion eras in the very early and recent universe, known as inflation and dark energy, respectively. \\\\ For these reasons, a great experimental work is being carried out in order to measure the finest structure of the CMB anisotropies. In total intensity, right after successful sub-orbital observations, the Wilkinson Microwave Anisotropy Probe (\\WMAP) satellite is performing all sky observations of the CMB in five frequency bands ranging between 22 and 90~GHz (see \\citealt{spergel_etal_2006} and references therein). The first detections of the CMB polarization have been made from the ground \\citep{kovac_etal_2002}, balloons \\citep{montroy_etal_2005} as well as from the \\WMAP~satellite itself \\citep{page_etal_2006}. In a few years the {\\sc Planck} satellite\\footnote{http://www.rssd.esa.int/Planck} will be launched, performing all sky measurements of the CMB total intensity and polarization anisotropy in 9 frequency bands between 30 and 857~GHz, with an angular resolution reaching 5 arcminutes with a sensitivity of a few micro-Kelvin. Next generation sub-orbital and satellite observations are ongoing or planned\\footnote{http://lambda.gsfc.nasa.gov/}.  The observations are confirming that the Galactic and extra-Galactic foreground emissions represent one of the two most challenging obstacles to the complete knowledge of CMB anisotropies. The other one is represented by the control of instrumental systematics. In total intensity the sky at high Galactic latitudes appears dominated by the CMB emission at least in the frequency range between 60 and 90~GHz observed by \\WMAP, although residual foreground contamination might be responsible for the deviations of the observed signal from a Gaussian statistics (see \\citet{copi_etal_2006} and references therein). In polarization the sky remains foreground dominated even if the region containing the brightest Galactic emission is cut out \\citep{page_etal_2006}. In particular, foregrounds are likely to be comparable or higher to the the curl component in the CMB polarization at all frequencies, and in any region of the sky \\citep{baccigalupi_2003}. \\\\ For these reasons it is important to study and test data analysis algorithms which are able to remove foreground contamination from the data. The class of these tools is called component separation, and they all make use of the multi-frequency superposition of the background and foreground emissions in order to separate them out. The concepts and means they use, however, are very different. The category of non-blind methods uses prior knowledge of the foreground emission in order to recover the CMB component \\citep{brandt_etal_1994,tegmark_efstathiou_1996,hobson_etal_1998,bouchet_etal_1999, barreiro_etal_2004,stolyarov_etal_2005}. Blind methods exploit the statistical independence between background and foregrounds. Among them, the Independent Component Analysis \\citep[ICA, see][ and references therein]{amari_chichocki_1998} has provided interesting results on simulated data in total intensity \\citep{baccigalupi_etal_2000,maino_etal_2002} and polarization \\citep{baccigalupi_etal_2004,stivoli_etal_2006a}, as well as on the real data by \\COBE~\\citep{maino_etal_2003} and $BEAST$ \\citep{donzelli_etal_2006}; these results are based on an implementation of the ICA technique capable to rapidly achieve convergence to the solution if the hypothesis of statistical independence among the signals is verified \\citep[\\fstica][]{hyvarinen_1999}. Recently, hybrids techniques merging the blind and non-blind categories and largely based on a parametric description of the foreground emission have been proposed \\citep{delabrouille_etal_2003,patanchon_etal_2005,eriksen_etal_2006}.  As we already mentioned, the design and testing of the component separation algorithms has been successful on real data. This, together with the results presented here, indicates that these techniques may become standard data analysis techniques, complementary to the methods of CMB extraction exploited so far. In this work we apply the \\fstica~algorithm to the total intensity \\WMAP~3yr data. Our purpose is to identify the CMB among the separated components and study its pattern in comparison with the findings of other works exploiting different techniques. The paper is organized as follows. In Section \\ref{fastica} we outline the main features of the \\fstica~algorithm. In Section \\ref{sims} we evaluate its capabilities on simulated \\WMAP~data. In Section \\ref{wmap} we apply the \\fstica~to the \\WMAP~data and study the recovered CMB. Finally, in Section \\ref{conclusions} we draw our conclusions. ",
        "conclusions": "\\label{conclusions}  In this work we applied the fast Independent Component Analysis (\\fstica) component separation technique to the 3 years data of the Wilkinson Microwave Anisotropy Probe (\\WMAP). The algorithm retrieves the different components superposed in a multi-frequency observation as linear combinations of the input data at different frequencies, which maximize the mutual statistical independence. We first evaluate the expected performance by means of Monte Carlo chains on simulated \\WMAP~data varying the noise and CMB realization, and exploiting the existing foreground models. Among the recovered components, we identify one which is compatible with the CMB emission by checking the recovered frequency scaling. On simulations, the precision of the recovery on the frequency scaling is of the order of percent. We then apply the technique to the real data, again identifying the CMB component by means of the reconstructed frequency scaling. Different combinations of the \\WMAP~channels give consistent results; in terms of the reconstruction of the CMB frequency scaling, the best \\WMAP~configuration includes the three high frequency channels (Q, V and W, respectively at 41, 61, 94 GHz), plus the lowest one (K at 23 GHz). This indicates that the algorithm benefits from having a good tracer of the low frequency foregrounds in the \\WMAP~data in order to achieve a proper CMB cleaning. The recovered CMB power spectrum is in close agreement with the \\WMAP~one on all accessible scales. The CMB fluctuation asymmetry in the northern and southern hemisphere claimed by several authors exploiting different data analysis techniques is confirmed in this work, with the same amplitude.  The agreement of the present results on CMB with the \\WMAP~ones and those from other authors is remarkable and strengthens the confidence we have on the CMB pattern reconstructed from the existing data as a whole. At the same time, from the point of view of the development of component separation techniques based on the Independent Component Analysis (ICA) this work represents the achievement of a most important milestone in view of the application to the Planck and other experiments; it means that the algorithm proved itself to be stable against the \\WMAP~instrumental systematics, and realistic foregrounds. \\\\ On the other hand, we are aware of the limitations of the present analysis. First of all, \\WMAP~is not an ideal experiment for performing a map based component separation, since the different channels have markedly different resolutions, and one has to do a pre-processing step decreasing the angular resolution of the data to a common one, which is about 1 degree in the present case. Second, the precision of the separation is evaluated by means of Monte Carlo chains on simulated data, which rely on ingredients, e.g. foregrounds and systematics, which may be far from reality, and also incomplete in the case of \\WMAP, as the dust foreground is visible only in the W band at 94 GHz, and does not have a multi-frequency coverage required by most component separation algorithms.  Despite of all these oddities, the CMB pattern recovered by the \\fstica~is consistent with the results in earlier literature, and confirms effects like the asymmetric distribution of the CMB anisotropy power across the sky.  This performance is likely to be due to the high level of statistical independence between background and foregrounds in the data, which is able to drive the solution close to the right one even for data affected by instrument systematics and realistic foregrounds. Thus we believe the present technique might be proving itself useful for future CMB probes, either accessing the CMB statistics beyond the power spectrum, and to control the foreground emission in polarization, which is going to be less known than in total intensity, and substantially higher compared with the CMB."
    },
    "0609/gr-qc0609059_arXiv.txt": {
        "abstract": "We investigate cosmological solutions of Brans-Dicke theory with both the vacuum energy density and the gravitational constant decaying linearly with the Hubble parameter. A particular class of them, with constant deceleration factor, sheds light on the cosmological constant problems, leading to a presently small vacuum term, and to a constant ratio between the vacuum and matter energy densities. By fixing the only free parameter of these solutions, we obtain cosmological parameters in accordance with observations of both the relative matter density and the universe age. In addition, we have three other solutions, with Brans-Dicke parameter $\\omega=-1$ and negative cosmological term, two of them with a future singularity of big-rip type. Although interesting from the theoretical point of view, two of them are not in agreement with the observed universe. The third one leads, in the limit of large times, to a constant relative matter density, being also a possible solution to the cosmic coincidence problem. ",
        "introduction": "Recent observations suggest that the total energy density in the universe is greater than the (barionic + dark) matter density. On one hand, dynamical estimations lead to a ratio between the matter and critical densities around one third \\cite{calb}. On the other hand, measurements of anisotropies in the microwave background radiation indicate that our universe is spatially flat \\cite{bach}, suggesting the existence of another, unknown form of energy, usually called dark energy. Its presence is also corroborated by age estimations \\cite{age} and by the observed distance-redshift relation for supernovas Ia, which suggests that, in the present phase of universe evolution, the deceleration parameter $q$ is negative, and therefore the universe performs an accelerated expansion \\cite{SnIa}.  Several models have been proposed in order to explain those observational data. There are some candidates for dark energy, as, for example, the cosmological constant, the so called quintessence, or the generalized Chaplygin gas. Among them, the simplest and oldest one is the cosmological constant, also associated to the energy density of vacuum.  We can contextualize the study of dark energy in different gravitational theories. For instance, the theory of General Relativity, where the Einstein field equations are used, and where the gravitation constant $G$ is a universal constant. Another one is the Brans-Dicke theory \\cite{weinberg,brans}, a scalar-tensor theory in which the gravitational constant is a function of space-time, and where a new parameter, $\\omega$, is introduced. More recently the interest on this kind of theory was renewed, owing to its association with superstrings theories, extra-dimensional theories and models with inflation or accelerated expansion \\cite{pavon}-\\cite{barrow}.  It is generally assumed that General Relativity is recovered in the limit $\\omega\\rightarrow+\\infty$ (despite the existence of Brans-Dicke solutions for which this is not true \\cite{faraoni}). Astronomical observations in the realm of Solar System impose a very high inferior limit for $\\omega$. Nevertheless, such a result corresponds to the weak field limit, and applies only in the simplest case of constant $\\omega$. Therefore, it is possible that General Relativity is not adequate to describe the universe at early times, or needs corrections in the cosmological limit.  In this article we consider the Brans-Dicke theory and associate to dark energy the equation of state of vacuum. We investigate models in which the vacuum energy density decreases with the universe expansion, a hypothesis that has been considered as a possible solution to the cosmological constant problem, that is, to the question of why the presently observed value of $\\Lambda$ is about $120$ orders of magnitude below the value predicted by quantum field theories \\cite{sahni,peebles}.  Our goal is to find solutions of Brans-Dicke theory which satisfy a particular variation law for $G$. We shall use the Eddington-Dirac relation, based on the large number coincidence, $G\\approx H/m_\\pi^3$, where $H=\\dot{a}/a$ is the Hubble parameter and $m_\\pi$ is the pion mass \\cite{saulo2,saulo1}. We will then take $G=H/8\\pi\\lambda$, where $\\lambda$ has the order of $m_\\pi^3$. In addition, as usual, we will relate the Brans-Dicke scalar field to the gravitational constant through $\\phi=G_0/G$, where $G_0$ is a positive constant of the order of unity.  Together with that variation law for $G$, we shall consider two different ansatzen. The first one is given by $\\rho=3\\alpha H^2/8\\pi G$, where $\\rho=\\rho_m+\\rho_\\Lambda$ is the total density, and $\\alpha$ is an adimensional constant of the order of unity. This ansatz is suggested by observations, which show that $\\rho_m\\approx\\rho_c/3$, where $\\rho_c=3H^2/8\\pi G$ is the critical density. On the other hand, we know that $\\rho_\\Lambda$ has, at most, the same order of magnitude as $\\rho_m$, otherwise its presence would be more evident. This ansatz was already considered in \\cite{saulo3,agostinho}.  The second ansatz will be given by $\\Lambda=\\beta H^2$, where $\\beta$ is a constant of the order of unity. We are, in this case, inferring a variation law for the cosmological term, which has already been considered in the literature on the basis of different arguments \\cite{saulo1}-\\cite{aldro} (for other variation laws for the cosmological term, see, for instance, \\cite{wu}-\\cite{prd}). We will show that this ansatz leads to a set of solutions larger than the first one, containing its solutions as a particular case.  We will look for solutions for recent times, that is, we shall consider a spatially flat ($k$ = 0) Friedmann-Robertson-Walker space-time, filled with a perfect fluid whose matter component is pressureless ($p_m = 0$). For the cosmological term, we will take the equation of state of vacuum, $p_\\Lambda=-\\rho_\\Lambda$. ",
        "conclusions": "In this work we have found some exact solutions of Brans-Dicke cosmology, by using two different ansatzen. We have shown that the first ansatz is a particular case of the second one, when the deceleration parameter $q$ is constant.  In the first ansatz, the ratio between the energy densities of matter and vacuum is constant, characterizing a possible solution for the cosmic coincidence problem, that is, the approximated coincidence presently observed between $\\rho_m$ and $\\rho_c$. This possibility survives to a quantitative analysis, since a relative matter density around $1/3$, as indicated by observations, leads to an age parameter $Ht\\approx1$, corresponding to a universe age around $14$ Gyr, also in accordance with observational limits.  Nevertheless, this ansatz presents some problems as well. The most severe of them is the presence of a constant deceleration factor (equals to zero if $Ht = 1$). In spite of the claim of some authors (see, for example, \\cite{jain,sethi}) about the possibility of a uniform expansion along the whole universe evolution, a decelerated phase is usually considered necessary for large structure formation. For this reason, we should consider this ansatz valid only in the limit of late times, restricting in this way the predictive power of the model.  With the second ansatz, besides the case of constant $q$, we have found three other universes, with varying $q$, in which the dark energy density is negative and the Brans-Dicke parameter is $\\omega = -1$. In one of them the deceleration parameter is always highly positive. In a second one, it is always highly negative. Therefore, these two cases are interesting just from a theoretical viewpoint.  In the third case, on the other hand, the deceleration parameter is initially positive, becoming negative at later times, but always finite. In this case (as well as in one of the previous cases) one has a future {\\it big-rip}, with the scale factor, the matter density and the Hubble parameter diverging in a finite time, but with the relative matter density remaining finite. As one can see from equation (\\ref{137}), for an age parameter in the interval $0.8<Ht<1.3$, as defined by the observational limits, the deceleration parameter is $-2.0>q>-2.9$. Whence, with the help of (\\ref{139}), it is possible to see that we have, for the relative matter density, $1.1>\\Omega_m/G_0>0.7$. As we know, different observations restrict the matter density parameter to $0.2<\\Omega_m<0.4$. Therefore, this solution satisfies such observations, provided $G_0$ is in the interval $0.3<G_0<0.4$. Furthermore, for the whole evolution we have $\\Omega_m/G_0<7.5$, tending, for future times, to a constant value around $0.5$ (see Figure 3). This also characterizes a possible explanation for the cosmic coincidence.  It is interesting to observe that, in the three cases with varying $q$, the cosmological time varies linearly with $q$, which may, therefore, be used to define the time measurement. It is interesting to note as well that the total energy density can be negative, since the dark energy density is negative. In our last solution, for example, $\\rho$ is positive until $q\\simeq-2.3$, becoming negative since then.  There is a particular result which may seem a limitation of our solutions, namely the typical values found for the Brans-Dicke parameter $\\omega$. As we know, observations in the realm of Solar System impose very high inferior limits for it. Let us remember, however, that we are considering the simplest version of scalar-tensor theories, which plays just an effective role here. Corrections to General Relativity, if exist, may be scale-dependent, and, therefore, observations in the Solar System cannot, in principle, impose limits to corrections at the cosmological scale. Particularly, we should not expect any time dependence of the Brans-Dicke scalar field (and so of $G$) in the Solar System, where the metric is stationary. While no spatial dependence should exist in large scale, because of the cosmological principle.  Anyway, a generalization of the solutions studied here seems to be necessary, either by modifying our ansatzen in the case of early times, either by considering more general scalar-tensor theories, with $\\omega$ depending on the scale. The study performed here, though limited in its scope, shows the variability of solutions in these contexts.  We should also note that, for the solution of Brans-Dicke equations, it is enough to add, for instance, the Eddington-Dirac relation. The inclusion of additional constraints, even when empirically or theoretically justified, had the goal of limiting the set of solutions. Therefore, a possible line of investigation would be to relax our ansatzen, imposing to the Brans-Dicke equations, for example, just the constraint given by the Eddington-Dirac relation, enlarging in this way the class of possible solutions. Such a generalization may also include the radiation phase, although we do not have any indication about the validity of the Eddington-Dirac relation at early times.  To conclude, let us remind that, despite the analysis we have made about the observational limits for the age and matter density parameters, a more detailed analysis of the whole set of current observational data is still in order. In particular, a careful study of the distance-redshift relation for supernovas Ia constitutes the subject of a forthcoming publication."
    },
    "0609/astro-ph0609045_arXiv.txt": {
        "abstract": "Existing  determinations show that $n$(C)/$n$(Fe) is more than a  factor of 10  below solar in the atmospheres of three  white dwarfs  that appear to be externally-polluted.  These results are not easily explained if the stars have accreted interstellar matter, and we re-interpret these measurements as  evidence that these stars have accreted asteroids with a chondritic composition. ",
        "introduction": "In white dwarfs with effective temperatures lower than about 20,000 K, it is predicted that heavy elements  settle below the photosphere in the star's strong gravitational field so the atmosphere is either largely hydrogen or largely helium  (Paquette et al. 1986, Chayer et al. 1995).  This expectation is supported by the determination that ${\\sim}$75\\% of these stars display upper limits to their calcium abundance that are as much as a  factor of 10$^{6}$ below the solar value (Zuckerman et al. 2003).  Currently, the only viable hypothesis to explain the observational result that ${\\sim}$25\\% of cool white dwarfs have detected metals is that these stars are externally-polluted. Interstellar accretion (Paquette et al. 1986, Dupuis et al. 1993a,b, Koester \\& Wilken 2006) is often thought to account for the metals within white dwarf atmospheres, but accretion from a tidally-disrupted minor body such as an analog to an asteroid or Kuiper  Belt Object naturally explains the infrared excesses and metal pollutions of  the hydrogen-rich white dwarfs G29-38, GD 362 and GD 56 (Jura 2003, Becklin et al. 2005, Kilic et al. 2005, 2006, Reach et al. 2005).    Here we re-interpret existing abundance determinations of three   helium-rich white dwarfs  that  previously were thought to have accreted interstellar matter  to argue, instead, that their measured low values of $n$(C)/$n$(Fe) support the hypothesis  that these stars have accreted circumstellar  matter with a chondritic composition.  Since the diffusion times out of a white dwarf's atmosphere for different elements heavier than helium vary by less than a factor of 2.5 (Paquette et al. 1986, Dupuis et al. 1993a),  a  test of any accretion scenario is to compare the metals in the star's photosphere with the composition of the polluter.    Within meteorites of all kinds as well as the Earth's crust,  carbon is depleted relative to refractory species such as iron, magnesium and silicon  by at least a factor of 10 and often more.    Accretion of circumstellar matter with a  composition similar to rocky material in the inner solar system can lead to distinctively low values of $n$(C)/$n$(Fe) that cannot be naturally explained by accretion of interstellar matter.  On the basis of this argument,  a carbon deficiency in the photosphere of the main sequence star J37 in the open cluster NGC 6633 has been attributed to the accretion of material with a chondritic composition (Ashwell et al. 2005, Laws \\& Gonzalez 2003).    Carbon has been detected in some  helium-rich white dwarfs. For the stars with $T_{eff}$ $<$ 13000 K, the convective zones are sufficiently deep that carbon may be dredged up from the interior (Koester et al. 1982, Pelletier et al. 1986, Dupuis et al. 1993b,  MacDonald et al. 1998).      In three white dwarfs with $T_{eff}$ near 25000 K, Petitclerc et al. (2005) have reported an unexplained presence of enhanced carbon with, for example, $n$(C)/$n$(Si) $>$100.   Here we focus on stars with carbon deficiencies in the sense that $n(C)$ $<<$ $n(Fe)$.   The unique chemistry of carbon makes it a useful element for  distinguishing between circumstellar and interstellar models of atmospheric pollution.  In the early solar system,  carbon was probably largely concentrated in volatile molecules such as CH$_{4}$ and CO (see, for example, Lodders 2003, 2004) and not incorporated into rocky bodies that formed in the inner solar system such as the Earth and asteroids (see also Simonelli et al. 1997, Gail 2002). In environments where carbon is more abundant than oxygen, carbon can condense into refractory solids such as graphite, and this may have occurred during the formation of  some of the outer regions of  the solar system where water was depleted (see Krot et al. 2000) and also, perhaps, around ${\\beta}$ Pic (Roberge et al. 2006).   In any case, meteorites, which are the best available measure of the composition of asteroids, are deficient in carbon (Wasson \\& Kallemeyn 1988).  In contrast, as discussed  below, carbon is abundant in both the gaseous and solid phase  in the interstellar medium.      In $\\S2$ we discuss scenarios for interstellar accretion and describe why the alternative of accretion from circumstellar matter is preferable for many stars.  In $\\S3$ we describe the scenario for  white dwarf pollutions by minor bodies such as asteroids.   In $\\S4$ we discuss our results and in $\\S5$ we report our conclusions.  Below, we suggest that at least  7\\% of  white dwarfs possess asteroid belts analogous to the solar system's.  However, two recent surveys  of main-sequence solar-type stars find  that only 1 star out of 69  (Bryden et al. 2006) or 1 star out of 41  (Beichman et al. 2006)  have warm dust that might be associated with asteroids.     While these surveys could not have detected an analog to the current zodiacal cloud in the solar system,  our own asteroid belt was probably much  more massive in the past, and  associated warm dust   could have produced  a significant infrared excess at ${\\lambda}$ $<$ 30 ${\\mu}$m (Gaidos 1999). We show in an Appendix that when the effects of a stellar wind are included, the infrared excess from dust associated with asteroids is suppressed, and therefore,  both young and old solar-type main-sequence  stars may possess  undetected asteroid belts. ",
        "conclusions": "The notably low  carbon  to iron ratios  in the atmospheres of three white dwarfs are naturally  explained if these stars have accreted  carbon-poor asteroids with masses between  7 ${\\times}$ 10$^{-7}$ M$_{\\oplus}$ and  3 ${\\times}$ 10$^{-4}$ M$_{\\oplus}$.    The inferred existence of extrasolar carbon-deficient minor bodies  suggests that chemical and dynamical scenarios for the formation and evolution of  rocky planetesimals in the inner solar system can be generalized.   I thank B. Hansen, D. Koester, J. Wasson  and B. Zuckerman for valuable conversations.  This work has been partly supported by NASA."
    },
    "0609/astro-ph0609677.txt": {
        "abstract": "{} { We search for effects of metallicity on B and Be stars in the Small and Large Magellanic Clouds (SMC and LMC) and in the Milky Way (MW). We extend our previous analysis of B and Be stars populations in the LMC to the SMC. The rotational velocities of massive stars and the evolutionary status of Be stars are examined with respect to their environments.} { Spectroscopic observations of hot stars belonging to the young cluster SMC-NGC\\,330 and its surrounding region have been obtained with the VLT-GIRAFFE facilities in MEDUSA mode. We determine fundamental parameters for B and Be stars with the GIRFIT code, taking into account the effect of fast rotation, and the age of observed clusters. We compare the mean \\vsini~obtained by spectral type- and mass-selection for field and cluster B and Be stars in the SMC with the one in the LMC and MW.} { We find that (i) B and Be stars rotate faster in the SMC than in the LMC, and in the LMC than in the MW; (ii) at a given metallicity, Be stars begin their main sequence life with a higher initial rotational velocity than B stars. Consequently, only a fraction of B stars that reach the ZAMS with a sufficiently high initial rotational velocity can become Be stars; (iii) the distributions of initial rotational velocities at the ZAMS for Be stars in the SMC, LMC and MW are mass- and metallicity-dependent; (iv) the angular velocities of B and Be stars are higher in the SMC than in the LMC and MW; (v) in the SMC and LMC, massive Be stars appear in the second part of the main sequence, contrary to massive Be stars in the MW.} {} ",
        "introduction": "The origin of the Be phenomenon has given rise to long debates. Whether it is linked to stellar evolution or initial formation conditions remains a major issue. Thus, finding out  differences in the physical properties of B and Be stars populations belonging to environments with different metallicity could provide new clues to understand the Be phenomenon.  To investigate the influence of metallicity,  star-formation conditions and stellar evolution on the Be phenomenon, we have undertaken an exhaustive study of B and Be stars belonging to young clusters or field of the Small and Large Magellanic Clouds (SMC and LMC), because these galaxies have a lower metallicity  than the Milky Way (MW). For this purpose we made use of the new  FLAMES-GIRAFFE instrumentation installed at the VLT-UT2 at ESO, which is particularly well suited, in MEDUSA mode, to obtain high quality spectra of large samples needed for the study of stellar populations. In Martayan et al. (2006a, hereafter  M06), we reported on the identification of 177 B and Be stars  belonging to the young cluster LMC-NGC 2004 and its surrounding region. In Martayan et al. (2006b, hereafter Paper I),  we determined  fundamental parameters of a large fraction of the sample in the  LMC, taking into account rotational effects (stellar flattening, gravitational darkening) when appropriate. We then investigated the effects of metallicity on rotational velocities. We concluded that Be stars begin their life on the main sequence (MS) with a  higher initial velocity than B stars. Moreover, this initial velocity is sensitive to the metallicity. Consequently, only a fraction of the B stars that reach the ZAMS with a sufficiently high initial rotational velocity can become Be stars. However, no clear influence of metallicity on the rotational velocity of B stars was found.  The present paper deals with a large sample of B and Be stars in the SMC, which has a lower metallicity than the LMC. With the determination of fundamental parameters, and the study of the evolutionary status and abundances, we aim at confirming and enlarging our results derived from the study of B and Be stars in Martayan et al. (2005a) and from the LMC (Paper I).  % ",
        "conclusions": "With the VLT-GIRAFFE spectrograph, we obtained spectra of a large sample of B and Be stars in the SMC-NGC\\,330 and surrounding fields. We determined fundamental parameters for B stars in the sample, and apparent and parent non-rotating counterpart (pnrc) fundamental parameters for Be stars.  Using results from this study and those obtained for the LMC with the same instrumentation (Paper I), we made a statistical comparison of the behaviour of B and Be stars in the Magellanic Clouds with the MW.\\\\ \\begin{itemize} \\item there is no difference in rotational velocities between early-type stars in clusters and in fields in the LMC and in the SMC. \\item the lower the metallicity is, the higher the rotational velocities are. B and Be stars rotate faster in the SMC than in the LMC, and faster in the LMC than in the MW. \\item we have determined, for the first time, the distributions of the ZAMS rotational velocities of Be stars. The ZAMS rotational velocities are mass-dependent and metallicity-dependent. The gradients of the distributions seem to be similar whatever the metallicity. \\item Only a fraction of B stars, that reach the ZAMS with sufficiently high initial rotational velocities, can become Be stars. Wisniewski \\& Bjorkman (2006) seem to have found photometrically a similar result. \\item the angular velocities are similar for B stars in the LMC and in the MW and lower than those in the SMC. The same result is obtained for Be stars. \\item more massive Be stars in the SMC, which are evolved, are critical rotators. \\item in an evolutionary scheme, massive Be stars in the MC and more particularly in the SMC appear in the second part of the main sequence, to the contrary to massive Be stars in the MW, which appear in the  first part of the main sequence. \\item other categories of Be stars (intermediate and less massive Be stars) follow the same evolution whatever the metallicity of the environment. \\end{itemize}  Our results support Stepi\\'en's scenario (2002): massive stars with a weak or moderate magnetic field and with an accretion disk during at least 10\\% of their PMS lifetime would reach the ZAMS with sufficiently high initial rotational velocity to become Be stars.  Our observational results also support theoretical results by Meynet \\& Maeder (2000, 2002) and Maeder \\& Meynet (2001) for massive stars. In a more general way they illustrate the behaviour of massive stars and the importance of the processes linked to rotation and to metallicity.  In particular, our results show that rotation has a major role in the behaviour of massive stars. Moreover, this pioneer study of a large sample of B stars in low metallicity environments, such as the LMC (Paper I) and the SMC (this study), allows us to approach  the ``first stars''. We can reasonably expect that first massive stars are fast rotators.  In forthcoming papers, we will present results of CNO abundances determinations, discuss the proportion of Be stars and report on the discovery of binaries in the SMC.  %@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@@ \\footnotesize{"
    },
    "0609/astro-ph0609273_arXiv.txt": {
        "abstract": "We have conducted a survey of 61 southern white dwarfs searching for magnetic fields using Zeeman spectropolarimetry. Our objective is to obtain a magnetic field distribution for these objects and, in particular, to find white dwarfs with weak fields. We found one possible candidate (WD 0310$-$688) that may have a weak magnetic field of $-6.1\\pm2.2$ kG. Next, we determine the fraction and distribution of magnetic white dwarfs in the Solar neighborhood, and investigate the probability of finding more of these objects based on the current incidence of magnetism in white dwarfs within 20 pc of the Sun. We have also analyzed the spectra of the white dwarfs to obtain effective temperatures and surface gravities. ",
        "introduction": "White dwarfs are fossil remains of stellar evolution, and a study of the distribution of magnetic fields among white dwarfs may help elucidate the role played by magnetic fields in the evolution of low to intermediate mass stars.  Our current knowledge of the field distribution, i.e., the fraction of magnetic white dwarfs as a function of field strength (polar, surface average, or longitudinal field strength), is dictated by a compilation of surveys of various aims and field sensitivities. Spectroscopic surveys are useful in probing the Zeeman effect in the Balmer lines of hydrogen-rich DA white dwarfs, and may achieve field sensitivity of $\\ga 1$MG causing Zeeman splitting of $\\Delta\\lambda \\ga 20$\\AA\\ in low-dispersion spectra \\citep{ven1999,fer1998}, or $\\gtrsim 10$kG corresponding to $\\Delta\\lambda \\ga 0.2$\\AA\\ in high-dispersion spectra of bright candidates \\citep[e.g.,][]{koe1998}. On the other hand, low-dispersion spectropolarimetric surveys of DA white dwarfs may achieve sensitivities of $\\gtrsim 10$kG while offering increased accessibility to fainter white dwarfs \\citep{sch1995}.  Similar studies \\citep[e.g.,][]{put1997,jor1998,sch2001b} of the incidence of magnetic fields in non-DA white dwarfs (helium-line DB, or continuum-like DC) are more difficult because of the weakness of helium lines relative to hydrogen lines in stars with similar ages, although the presence of Zeeman-split metal lines may betray the presence of a magnetic field in such objects \\citep[e.g.,][]{rei2001}. Dichroism of continuous opacities induces optical circular polarization of the order of $V/I \\approx 0.1B$\\%, where $B$ is expressed in MG \\citep{kem1970,ang1981}. Therefore, polarimetry of DC white dwarfs may achieve field sensitivity $\\gtrsim 1$MG in high-quality surveys ($V/I\\gtrsim 0.1$\\%) or more generally $\\gtrsim 10$MG ($V/I\\gtrsim 1$\\%).  Spectroscopic and polarimetric surveys have accumulated sufficient data to describe the field distribution $\\gtrsim1$MG in all nearby white dwarfs (DA and non-DA), while only a fraction of the data is available to describe the distribution $\\gtrsim0.1$MG among brighter hydrogen-rich DA white dwarfs. The aim of spectropolarimetric surveys is to complement spectroscopic surveys and provide data for fainter DA candidates with even lower fields in the 10kG-1MG range. \\citet{kaw2003} derived the field distribution in the local white dwarf census of \\citet{hol2002}. The incidence of magnetic field appears constant per decade of field strength from 0.1MG to close to 1000 MG. Since the estimate of \\citet{kaw2003}, which included the low-field magnetic white dwarf 40 Eri B \\citep{fab2000}, \\citet{azn2004} added the low-field white dwarf LTT~9857 (3.1 kG) to the local census (i.e., within 20 pc of the Sun) thereby extending the field distribution well below 10 kG. The origin of fields in white dwarf stars is more difficult to establish quantitatively. \\citet{kawthesis} and \\citet{kaw2004b} revisited some assumptions about the magnetic white dwarf space density and the corresponding space density of their likely Ap/Bp progenitors. They concluded that our current knowledge of stellar formation rate and evolutionary time scales cannot account for the present day density of magnetic white dwarfs. Low field white dwarfs are also not accounted for, which prompted \\citet{kaw2004b} to suggest that additional progenitors are required.  In \\S 2 we present spectropolarimetry of 61 white dwarfs with $V \\lesssim 16.4$ and $\\delta < -30^\\circ$, among which 55 have $V \\lesssim 15.0$ and 5 are in close binaries. We also present spectropolarimetry of 4 subdwarf B (sdB) stars, where two of these were misclassified as white dwarfs. This complements a spectropolarimetric survey for magnetic fields among northern hemisphere white dwarfs carried out by \\citet{sch1995}. That survey sampled some 169 DA white dwarfs and resulted in the discovery of four new magnetic white dwarfs with fields between $\\sim$$10^5$ and $10^9$ G. We present our analyses of our observations in \\S 3 and in \\S 4 we discuss the properties of the local population of magnetic white dwarfs. Finally, in \\S 5 we summarize our results. Appendix A tabulates the properties of all magnetic white dwarfs known to date. ",
        "conclusions": "We have conducted a spectropolarimetric survey of southern white dwarfs, which resulted in no new detections of magnetic fields with the possible exception of WD~0310-688 for which we measured a longitudinal field of $-6.1\\pm2.2$ kG. We have also observed the known magnetic white dwarf EUVE~J0823-25.4 and measured a longitudinal field of $-600$ kG. However, this survey has helped constrain the incidence of magnetism in the Solar neighborhood. We reviewed the list of known white dwarfs in the Solar neighborhood and found that $21\\pm8 \\%$ are magnetic. We also investigated the probability of finding more magnetic white dwarfs within the Solar neighborhood based on the current magnetic white dwarf incidence. We found that a significant number of magnetic white dwarfs with fields $B < 100$ kG, remain to be detected within the Solar neighborhood."
    },
    "0609/astro-ph0609790_arXiv.txt": {
        "abstract": "We present basic observational data on the 6C** sample. This is a new sample of radio sources drawn from the 151\\,MHz 6C survey, which was filtered with radio criteria chosen to optimize the chances of finding radio galaxies at $z > 4$.  The filtering criteria are a steep-spectral index and a small angular size. The final sample consists of 68 sources from a region of sky covering 0.421\\,sr.  We present VLA radio maps, and the results of \\kband imaging and optical spectroscopy.  Near-infrared counterparts are identified for 66 of the 68 sources, down to a $3 \\sigma$ limiting magnitude of $K \\sim 22$~mag in a 3-arcsec aperture. Eight of these identifications are spatially compact, implying an unresolved nuclear source.  The $K$-magnitude distribution peaks at a median $K \\approx 18.7$~mag, and is found to be statistically indistinguishable from that of the similarly selected 6C* sample, implying that the redshift distribution could extend to $z \\,\\,\\gtsim\\,\\, 4$.  Redshifts determined from spectroscopy are available for 22 (32 per cent) of the sources, over the range of 0.2 $\\,\\,\\ltsim\\,\\, z \\,\\,\\ltsim\\,\\, 3.3$ . We measure 15 of these, whereas the other 7 were previously known. Six sources are at $z > 2.5$.  Four sources show broad emission lines in their spectra and are classified as quasars. Three of these show also an unresolved \\kband identification.  Eleven sources fail to show any distinctive emission and/or absorption features in their spectra. We suggest that these could be (i) in the so-called `redshift desert' region of $1.2 < z < 1.8$, or (ii) at a greater redshift, but feature weak emission line spectra. ",
        "introduction": "The chief purpose of the 6C** sample was to find radio galaxies at the highest redshifts ($z > 4$) and, based on its data, to directly constrain the co-moving space density of high-redshift radio sources. It is the second of the filtered 6C redshift surveys, being preceded by a pilot sample, namely 6C* (Blundell et al. 1998; Jarvis et al. 2001a,b).   \\begin{table*} \\begin{center} \\begin{tabular}{lrrccr} \\hline \\mc{1}{c}{Sample} & \\mc{1}{c}{Flux-density} & \\mc{1}{c}{Sky Area} & \\mc{1}{c}{Number of} & \\mc{1}{c}{Spectroscopic} & \\mc{1}{c}{Radio} \\\\ & \\mc{1}{c}{Limits (Jy)} & \\mc{1}{c}{(sr)} & \\mc{1}{c}{Sources} & \\mc{1}{c}{Completeness} & \\mc{1}{c}{Filtering}\\\\ \\hline 3CRR & $S_{178} \\geq 10.9$ & 4.24 & 173 & 100\\%& None\\\\ 6CE  & $ 2.00 \\leq S_{151} \\leq 3.93$ & 0.10 & 58 & 97\\% & None\\\\ 7CRS & $S_{151} \\geq 0.5$ & 0.022 & 128 & 90\\% & None\\\\ \\\\ 6C* & $ 0.96 \\leq S_{151} \\leq 2.0$ & 0.13 & 29 & 100\\% & $\\rm \\alpha^{\\scriptscriptstyle4.85{\\scriptscriptstyle\\rm G}}_{\\scriptscriptstyle151{\\scriptscriptstyle\\rm M}} \\geq 0.981$; $\\theta<15$ arcsec\\\\ 6C** & $S_{151} \\geq 0.5$ & 0.42 & 68 & 32\\% &  $\\rm \\alpha^{\\scriptscriptstyle1.4{\\scriptscriptstyle\\rm G}}_{\\scriptscriptstyle151{\\scriptscriptstyle\\rm M}} \\geq 1.0$; $\\theta<11$ arcsec\\\\ \\hline \\end{tabular} \\end{center} \\caption{Summary of the properties of the 3CRR, 6CE, 7CRS, 6C*  and 6C** redshift surveys. Note: spectroscopy has not been attempted yet on all members of the 6C** sample (see Section~\\ref{sec:spectroscopy}). } \\label{tab:samples} \\end{table*}   Previous to the 6C and 7C surveys, low-frequency studies of radio sources were based on the revised 3C sample (3CRR; Laing, Riley \\& Longair 1983).  The problem with this sample is that, as with any single flux-limited sample, there is a tight correlation between luminosity and redshift, in the sense that the highest redshift sources in the sample are necessarily the most luminous ones. One obvious implication is that based solely on the 3CRR sample, or in any single flux-density limited sample, one cannot distinguish between correlations of source parameters with redshift and with luminosity. The other implication is that the 3CRR sample, being very bright, does not find any objects at $z \\gtrsim 2$, due to the rarity of the extremely powerful (\\mbox{$L_{151} \\sim 10^{28}$ -- $10^{29}$~W~Hz$^{-1}$~sr$^{-1}$}) radio sources.  The 6CE (Rawlings, Eales \\& Lacy 2001) and 7CRS (Willott et al. 1998, 2002; Lacy et al. 1999a,b) complete samples, with fainter flux-density limits (see Table~\\ref{tab:samples}), were designed to address these problems and, in particular, to improve coverage of the 151\\,MHz luminosity ($L_{151}$), redshift plane. They were indeed very successful in extending the $L_{151} - z$ plane into lower radio luminosities at low-redshift and into $ 2 < z < 3$ at high radio luminosity (e.g. Willott et al. 2001).  However, the problem with these samples is that the requirement of low flux-density limits could only be achieved, in terms of complete spectroscopic follow-up being feasible, at the expense of the sky area covered.  This is particularly dramatic in the case of the 7CRS sample and it very much restricted the number of very high-redshift sources included in these surveys. The need to overcome this limitation, i.e. to have a redshift survey which covers a larger area of sky at flux levels comparable to that of the 7CRS, and thereby probe very high redshift radio sources ($z > 3$) in a statistically meaningful way, provided the main motivation for the construction of the filtered 6C redshift surveys\\footnote{A fainter and larger survey, containing $\\sim$ 1000 radio sources with $S_{151} > 0.1$~Jy, is also being put together by the groups at Oxford and Texas (Hill \\& Rawlings 2003).  This survey, named the TexOx-1000 (TOOT), aims to probe those radio-loud objects which are more typical at high redshift ($z > 3$), i.e. less luminous than the ones probed by 3CRR, 6CE and 7CRS at the same redshift, which are rare.}.   These surveys, 6C* and 6C**, were both filtered with more selective radio criteria than just a simple flux-density limit (see Table~\\ref{tab:samples}). Additional criteria were introduced in order to filter out a large fraction of the sources, mostly at low-redshift, thereby ensuring that optical follow-up would be limited to a manageable number of sources, biased towards high-redshift. The chosen criteria took into account, therefore, the characteristics often seen in very distant radio galaxies, namely steep radio spectral index $\\alpha$ and small angular size $\\theta$.  The first criterion has been the one most widely used in searches for the highest redshift radio galaxies (e.g. R\\\"ottgering et al. 1994; Chambers et al. 1996; Blundell et al. 1998; De Breuck et al. 2000; Cohen et al. 2004). It is based primarily on the convex shape of the radio spectra which flattens below rest-frame $\\sim$ 300~MHz. Since higher redshift objects are observed at higher rest-frame frequencies, their measured spectral indices are invariably steeper. The second criterion relies on strong statistical evidence for a negative evolutionary trend of linear size with redshift (e.g. Barthel \\& Miley 1988; Neeser et al. 1995; Blundell, Rawlings \\& Willott 1999) although it has been less widely used in targeting high-redshift radio sources.  Both criteria are imperfect in the sense that not all of the sources at low-redshift are eliminated and some fraction at high redshift are actually excluded. Jarvis et al. (2001b) quantified these fractions by comparing the redshift distribution of the 6C* sample with the radio luminosity function model of Willott et al. (2001).  They have found that employment of the filtering criteria has reduced the number of $z < 1.5$ sources by $\\approx$ 90 per cent, whereas in the redshift range $1.5 \\,\\,\\ltsim\\,\\, z \\, \\,\\ltsim \\,\\,3.0$ the fraction excluded varies between $\\sim 30 - 70$ per cent. The distribution, as a result, is skewed towards objects at $z \\gtrsim 2.0$, and the 6C* sample has a median redshift of $z \\sim 1.9$ (cf. 6CE and 7CRS with $\\sim 1.1$ median redshift).  Even though the effects of the filtering criteria introduce difficulties in interpreting the data from filtered samples, namely in assessing which population of sources is filtered out, they are very effective in biasing these samples to objects at high redshift. This has been demonstrated very successfully with the 6C* sample (Jarvis et al. 2001b), which led to the discovery of 6C0140+326 at $z = 4.41$ (Rawlings et al. 1996) -- the most distant radio galaxy known at the time of discovery -- in a sample of only 29 objects.  Nevertheless, the 6C* sample has only two objects at $z > 3$ and, therefore, the $L_{151}-z$ diagram still suffers from small number statistics at high redshift. The 6C** sample, deeper and larger than 6C* (Table~\\ref{tab:samples}), aims to ameliorate this problem and, ultimately, to extend the $L_{151}-z$ diagram to even higher redshifts ($z \\sim 5$).  The final spectroscopically complete 6C** sample will also potentially supply us with targets with which we can probe the most massive objects in the early Universe. Studies of the host galaxies of radio-loud AGN have shown that all powerful radio sources are associated with the most massive galaxies at every cosmological epoch (Jarvis et al. 2001a; De Breuck et al. 2002; Willott et al. 2003; Zirm et al. 2003).  Similarly, it has been shown that the most radio-loud quasars are those with the most massive black-holes (e.g. Dunlop et al. 2003; McLure \\& Jarvis 2004).  Moreover, recent studies have identified high redshift radio galaxies (HzRGs; $z > 2$) which are associated with the sites of forming proto-clusters (Kurk et al. 2004; Miley et al. 2004; Venemans et al. 2004, 2005), implying that HzRGs are the most likely progenitors of the nearby brightest clusters. Powerful radio sources are therefore key targets for studies of the formation and evolution of massive structures in the early Universe.  Also, the discovery of a radio galaxy within the epoch of reionization ($z > 6$) will allow for 21~cm absorption line studies of neutral hydrogen at that epoch (e.g. Carilli, Gnedin \\& Owen 2002).   Finally, one important characteristic of all the samples mentioned here is that they have been selected at low radio frequencies ($\\nu < 200$~MHz). This is a crucial factor in searches for high-redshift sources  (e.g. Blundell, Rawlings \\& Willott 1999, 2000) and strongly determines the content of a sample.  Low-frequency surveys typically contain a high proportion of steep-spectrum sources, where the observed emission is dominated by the isotropically emitting radio-lobe components. In high-frequency ($\\nu \\gtrsim 1$~GHz) surveys, in contrast, core-dominated flat-spectrum sources prevail. Consequently, surveys at high frequency tend to have more objects which are selected above the flux limit because their emission is boosted due to orientation of their jet axes, rather than their intrinsic power. Surveys at low-frequency ($\\nu < 200$~MHz), being dominated by isotropic radio emission, provide us, therefore, with much less biased samples of radio sources.    In this paper (hereafter Paper I) we present complete \\kband imaging of the 6C** sample, high-resolution radio maps for 42 of the sources, and optical spectroscopy for a sub-sample of the optically identified sources. In a companion paper (Paper II), we will present single-colour ($K$-band) photometric redshifts for all the sources and investigate the redshift distribution of the sample. The paper is set out as follows. In Section~\\ref{sec:sample} we describe the sample selection criteria.  In Section~\\ref{sec:observations} we outline the observing and data reduction procedures. The radio maps, \\kband images and photometry, and the redshifts and line parameters are presented in Section~\\ref{sec:analysis}. In Section~\\ref{sec:notes} we provide notes of each source in the sample with respect to both imaging and spectroscopy. The results are discussed in Section~\\ref{sec:discussion} and our main conclusions are summarized in Section~\\ref{sec:conclusions}.  Unless otherwise stated, we assume throughout that $H_{0}=70~ {\\rm km~s^{-1}Mpc^{-1}}$, $\\Omega_ {\\mathrm M} = 0.3$ and $\\Omega_ {\\Lambda} = 0.7$. The convention used for radio spectral index is $S_{\\nu} \\propto \\nu^{-\\alpha}$, where $S_{\\nu}$ is the flux-density at frequency $\\nu$.   \\begin{table*} \\scriptsize \\begin{center} \\begin{tabular}{lclrccrc} \\hline \\mc{1}{c}{} & \\mc{1}{c}{$K$-BAND}& \\mc{1}{c}{} & \\mc{1}{c}{} & \\mc{1}{c}{RADIO} & \\mc{1}{c}{} & \\mc{1}{c}{} & \\mc{1}{c}{} \\\\ \\hline \\mc{1}{c}{Source} & \\mc{1}{c}{Telescope/} & \\mc{1}{c}{Dates Observed} & \\mc{1}{c}{Total} & \\mc{1}{c}{Map} & \\mc{1}{c}{Dates} & \\mc{1}{c}{Peak Flux} & \\mc{1}{c}{RMS}\\\\ \\mc{1}{c}{Name} & \\mc{1}{c}{Instrument} & \\mc{1}{c}{} & \\mc{1}{c}{Exp. (s)} & \\mc{1}{c}{Freq. (GHz)} & \\mc{1}{c}{Obs.} & \\mc{1}{c}{(mJy~beam$^{-1}$)} & \\mc{1}{c}{(mJy~beam$^{-1}$)}\\\\ \\hline 6C**0714+4616 & \\uk/\\uf & {\\bf 2002-01-09} & 540 & 4.9  & 1996-12-06 & 22.16  & 0.36\\\\ 6C**0717+5121 & \\uk/\\uf & {\\bf 1999-03-12} & 3240& 8.4  & 1996-12-18 &~~1.30 & 0.10\\\\ & \\uk/\\ui & 2003-01-28 & 540 \\\\ 6C**0726+4938 & \\uk/\\ui & {\\bf 2003-01-28} & 540 & 4.9  & 1996-12-18 &~~1.29 & 0.13\\\\ 6C**0737+5618 & \\uk/\\uf & 1999-03-12 1999-03-13 & 4320\\\\ & \\uk/\\ui & 2003-01-28 & 540\\\\ & \\gem   & {\\bf 2004-02-03} & 2700& 8.4    & 1996-12-18 &~~0.35 & 0.09\\\\ \\\\ 6C**0744+3702 & \\uk/\\uf & {\\bf 1999-03-12} & 2160& 8.4  & 1996-12-18 &~~0.85 & 0.09\\\\ 6C**0746+5445 & \\uk/\\uf & {\\bf 1999-03-12} & 1620& 4.9  & 1996-12-06 &~~0.55 & 0.12\\\\ 6C**0754+4640 & \\uk/\\uf & 1999-03-13 & 3000\\\\ & \\gem    & {\\bf 2004-02-10} & 2460& 4.9    & 1996-12-06 &~~0.48 & 0.10\\\\ 6C**0754+5019 & \\uk/\\ui & 2003-01-28 & 540\\\\ & \\gem    & {\\bf 2004-02-10} & 2460& 4.9    & 1996-12-06 &~~8.75 & 0.21\\\\ \\\\ 6C**0801+4903 & \\uk/\\uf & 1999-03-11$^{\\dagger}$ {\\bf 2003-12-18} & 1080 & 4.9 & 1996-12-06 &~~0.52 & 0.12\\\\ & \\uk/\\ui & 2003-01-28 & 540\\\\ 6C**0810+4605 & \\uk/\\uf & 2002-01-09 {\\bf 2002-01-13} & 2160& 4.9 & 1996-12-06 &~~0.59 & 0.10\\\\ 6C**0813+3725 & \\uk/\\uf & 2002-01-09 {\\bf 2003-12-23} & 1080& 4.9 & 1996-12-06 & ~~0.34 & 0.08\\\\ 6C**0824+5344 & \\uk/\\uf & 2002-01-09 {\\bf 2002-01-13} & 2160& 4.9 & 1996-12-06 &~~5.56 & 0.19\\\\ 6C**0829+3902 & \\uk/\\uf & {\\bf 1999-03-13} & 2640 & 4.9 & 1996-12-06 &~~5.96 & 0.15\\\\ \\\\ 6C**0832+4420 & \\uk/\\uf & 1999-03-11$^{\\dagger}$ {\\bf 2003-12-15} & 1080 & 4.9 & 1996-12-06 & ~~0.51 & 0.12\\\\ 6C**0832+5443 & \\uk/\\uf & {\\bf 2002-01-09} & 540& 4.9 & 1996-12-06 & ~~0.37 & 0.10\\\\ 6C**0834+4129 & \\uk/\\uf & {\\bf 2002-01-12} & 540& 8.4 & 1996-12-18 & ~~3.02 & 0.13\\\\ 6C**0848+4803 & \\uk/\\uf & 1999-03-11$^{\\dagger}$ {\\bf 2003-12-15} & 1620& 4.9 & 1996-12-06 &~~3.68 & 0.12\\\\ 6C**0848+4927 & \\uk/\\uf & {\\bf 1999-03-13} & 1080& 4.9 & 1996-12-06 & 19.77 & 0.32\\\\ 6C**0849+4658 & \\uk/\\uf & 1999-03-11$^{\\dagger}$ {\\bf 2003-12-15} & 1080& 8.4 & 1996-12-18 & 19.66 & 0.37\\\\ \\\\ 6C**0854+3500 & \\uk/\\uf & {\\bf 1999-03-12} & 540 & 8.4 & 1996-12-18 & ~~0.49 & 0.12\\\\ 6C**0855+4428 & \\uk/\\uf & {\\bf 2002-01-12} & 540 & 4.9 & 1996-12-06 & ~~3.33 & 0.15\\\\ 6C**0856+4313 & \\uk/\\uf & {\\bf 2003-12-15} & 540 & 4.9 & 1996-12-06 & ~~5.87 & 0.17\\\\ 6C**0902+3827 & \\uk/\\uf & 2000-04-12 {\\bf 2003-12-18} & 1080 & 8.4 & 1996-12-18 & ~~8.47 & 0.18\\\\ 6C**0903+4251 & \\uk/\\uf & {\\bf 2000-04-13} & 540 & 8.4 & 1996-12-18 & ~~8.23 & 0.18\\\\ 6C**0909+4317 & \\uk/\\uf & {\\bf 2000-04-12} & 3240& 4.9 & 1996-12-06 & ~~0.88 & 0.20\\\\ \\\\ 6C**0912+3913 & \\uk/\\uf & {\\bf 2000-04-13} & 540& 1.4 & & 26.03  & 0.37\\\\ 6C**0920+5308 & \\uk/\\uf & {\\bf 2002-11-22} & 540 & 4.9 & 1996-12-06 & ~~4.60 & 0.15\\\\ 6C**0922+4216 & \\uk/\\uf & {\\bf 2002-11-22} & 540 & 1.4 & & 199.65 & 0.13\\\\ 6C**0924+4933 & \\uk/\\uf & {\\bf 2002-11-22} & 540 & 4.9 & 1996-12-06 & ~~8.47 & 0.21\\\\ 6C**0925+4155 & \\uk/\\uf & 2000-04-14 2003-12-18 2003-12-23 & 1620 \\\\ & \\gem    & 2004-02-12$^{\\dagger}$ & 840\\\\ & \\keck   & {\\bf 2004-04-30}  & 1920& 1.4 &  & 76.41 & 0.15\\\\ \\\\ 6C**0928+4203 & \\uk/\\uf & {\\bf 2002-11-22} & 540 & 4.9 & 1996-12-06 & ~~8.52 & 0.24\\\\ 6C**0928+5557 & \\uk/\\uf & {\\bf 2000-04-14} & 540 & 4.9 & 1996-12-06 & ~~0.92 & 0.11\\\\ 6C**0930+4856 & \\uk/\\uf & {\\bf 2002-11-22} & 540 & 4.9 & 1996-12-18 & ~~0.79 & 0.19\\\\ 6C**0935+4348 & \\uk/\\uf & 2000-04-12 2003-12-18 2003-12-23 & 540 \\\\ & \\gem    & {\\bf 2004-02-12} & 2280& 1.4 &   & 29.02 & 0.14\\\\ 6C**0935+5548 & \\uk/\\uf & {\\bf 2000-04-13} & 540 & 4.9 & 1996-12-06 & ~~0.72 & 0.14\\\\ \\\\ 6C**0938+3801 & \\uk/\\uf & 2000-04-13 {\\bf 2004-02-09} & 1080 & 8.4 & 1996-12-18 & ~~2.89 & 0.13\\\\ & \\uk/\\ui & 2003-01-28 & 540\\\\ 6C**0943+4034 & \\uk/\\uf & {\\bf 2000-04-14} & 540 & 1.4 & & 56.70 & 0.13\\\\ 6C**0944+3946 & \\uk/\\uf & 2000-04-14 {\\bf 2003-12-18} & 1080& 1.4 &  & 53.96 & 0.13\\\\ & \\uk/\\ui & 2003-01-28 & 540\\\\ 6C**0956+4735 & \\uk/\\uf & {\\bf 2002-11-22} & 540 & 4.9 & 1996-12-06 & ~~0.75 & 0.18\\\\ \\\\ 6C**0957+3955 & \\uk/\\uf & {\\bf 2002-11-22} & 540& 8.4 & 1996-12-18 & ~~0.77 & 0.10\\\\ 6C**1003+4827 & \\uk/\\uf & {\\bf 2002-11-22} & 540& 4.9 & 1996-12-18 & ~~2.89 & 0.21\\\\ 6C**1004+4531 & \\uk/\\uf & {\\bf 2002-11-22} & 540& 4.9 & 1996-12-06 & ~~0.60 & 0.13\\\\ 6C**1006+4135 & \\uk/\\uf & {\\bf 2002-11-22} & 540& 8.4 & 1996-12-18 & ~~0.43 & 0.11\\\\ 6C**1009+4327 & \\uk/\\uf & {\\bf 2002-11-22} & 540& 1.4 &  & 133.03 & 0.14\\\\ 6C**1015+5334 & \\uk/\\uf & {\\bf 2000-04-14} & 540& 4.9 & 1996-12-06 & ~~0.58 & 0.12\\\\ \\hline \\end{tabular} {\\caption{\\label{tab:im_journalk1} Log of the \\kband and radio imaging observations of the sources present in the 6C$^{**}$ sample.  The bold lettering indicates the \\kband images shown in Fig.~\\ref{fig:kband_images1}. Columns 5, 6, 7 and 8 list the frequency, observation date, peak flux density and rms noise of the radio maps shown in Fig.~\\ref{fig:kband_images1}. The maps at 4.9~GHz and 8.4~GHz were obtained with the VLA in its A-configuration, as described in Section~\\ref{sec:radio}, with integrations times varying from 2 to 5 minutes. The maps at 1.4~GHz are from the FIRST survey (http://sundog.stsci.edu/first/catalogs.html). Note: the Keck--NIRC image of 6C**0925+4155 was observed through a K$_{\\rm S}$ filter. $\\dagger$ signifies that the observation was made under non-photometric conditions.  }} \\end{center} \\end{table*}  \\addtocounter{table}{-1} \\begin{table*} \\scriptsize \\begin{center} \\begin{tabular}{lclrccrc} \\hline \\mc{1}{c}{} & \\mc{1}{c}{$K$-BAND}& \\mc{1}{c}{} & \\mc{1}{c}{} & \\mc{1}{c}{RADIO} & \\mc{1}{c}{} & \\mc{1}{c}{} & \\mc{1}{c}{} \\\\ \\hline \\mc{1}{c}{Source} & \\mc{1}{c}{Telescope/} & \\mc{1}{c}{Dates Observed} & \\mc{1}{c}{Total} & \\mc{1}{c}{Map} & \\mc{1}{c}{Dates} & \\mc{1}{c}{Peak Flux} & \\mc{1}{c}{RMS}\\\\ \\mc{1}{c}{Name} & \\mc{1}{c}{Instrument} & \\mc{1}{c}{} & \\mc{1}{c}{Exp. (s)} & \\mc{1}{c}{Freq. (GHz)} & \\mc{1}{c}{Obs.} & \\mc{1}{c}{(mJy~beam$^{-1}$)} & \\mc{1}{c}{(mJy~beam$^{-1}$)}\\\\ \\hline 6C**1017+3436 & \\uk/\\uf & {\\bf 2000-04-14} & 540& 8.4 & 1996-12-18 & ~~0.48 & 0.12\\\\ 6C**1018+4000 & \\uk/\\uf & {\\bf 2000-04-14} & 540& 8.4 & 1996-12-18 & ~~0.44 & 0.11\\\\ 6C**1035+4245 & \\uk/\\uf & {\\bf 2000-04-12} & 3240& 8.4& 1996-12-18 & ~~0.44 & 0.11\\\\ 6C**1036+4721 & \\uk/\\uf & {\\bf 2002-11-22} & 540& 1.4 &  & 358.25 & 0.15\\\\ 6C**1043+3714 & \\uk/\\uf & {\\bf 2000-04-14} & 540& 1.4 &  & 191.04 & 0.30\\\\ 6C**1044+4938 & \\uk/\\uf & 2000-04-13 {\\bf 2000-04-14} & 3780& 8.4 & 1996-12-18 & ~~1.26 & 0.17\\\\ \\\\ 6C**1045+4459 & \\uk/\\ui & {\\bf 2003-01-28} & 540& 1.4 &  & 45.63 & 0.14\\\\ 6C**1048+4434 & \\uk/\\ui & {\\bf 2003-01-28} & 540& 1.4 &  & 86.02 & 0.14\\\\ 6C**1050+5440 & \\uk/\\uf & {\\bf 1999-03-06} & 3240& 1.4&  & 56.99 & 0.17\\\\ 6C**1052+4349 & \\uk/\\uf & {\\bf 2002-12-20} & 540 & 1.4&  & 45.77  & 0.13\\\\ 6C**1056+3303 & \\uk/\\uf & {\\bf 2002-12-20} & 540 & 1.4&  & 149.92 & 0.11\\\\ 6C**1100+4417 & \\uk/\\uf & {\\bf 2002-01-13} & 540 & 1.4&  & 59.79  & 0.14\\\\ \\\\ 6C**1102+4329 & \\uk/\\uf & 2002-01-13 2003-12-18 {\\bf 2003-12-23} & 1620& 1.4&  & 81.15 & 0.13\\\\ 6C**1103+5352 & \\uk/\\uf & 2002-01-13 {\\bf 2003-12-23} & 1080& 1.4&  & 256.49 & 0.14\\\\ 6C**1105+4454 & \\uk/\\uf & {\\bf 2002-01-13} & 540& 1.4&  & 36.60 & 0.14\\\\ 6C**1106+5301 & \\uk/\\uf & {\\bf 2002-05-05} & 540& 1.4&  & 36.27 & 0.14\\\\ 6C**1112+4133 & \\uk/\\uf & {\\bf 2002-12-20} & 540 & 1.4&  & 11.77 & 0.13\\\\ 6C**1125+5548 & \\uk/\\uf & {\\bf 2002-01-13} 2002-05-05 & 1080& 1.4&  & 24.44 & 0.14\\\\ \\\\ 6C**1132+3209 & \\uk/\\uf & {\\bf 2002-01-13} & 540& 1.4&   & 27.52 & 0.14\\\\ 6C**1135+5122 & \\uk/\\uf & {\\bf 2000-04-12} & 540 & 1.4&  & 50.91 & 0.14\\\\ 6C**1138+3309 & \\uk/\\uf & 2002-01-13 2002-12-20 {\\bf 2003-12-18}  & 1620& 1.4&  & 36.05 & 0.11\\\\ 6C**1138+3803 & \\uk/\\uf & {\\bf 2002-01-13} & 540& 1.4&  & 35.34 & 0.15\\\\ 6C**1149+3509 & \\uk/\\uf & {\\bf 2002-12-20} & 540 & 1.4&  & 54.16 & 0.13\\\\ \\hline \\end{tabular} {\\caption{\\label{tab:im_journal2} {\\em continued} }} \\end{center} \\end{table*} ",
        "conclusions": "\\label{sec:discussion}  \\subsection{Number of chance coincidences expected among the near-infrared identifications}  We have identified all but two of the sources in the 6C** sample, down to a limiting magnitude of $K \\sim 21$~mag (8-arcsec diameter aperture). Note, however, the caveat that one of the non-detected sources (6C**0935+4348) has an uncertain identification (see notes on this source in Section~\\ref{sec:notes-imaging}).  Using K-band number counts from Gardner, Cowie \\& Wainscoat (1993) and assuming an average search radius of 2.5 arcsec, we now estimate the number of chance coincidences expected inside the matching area. This is a function of K-band magnitude reached.  Deep infrared surveys (e.g. Gardner, Cowie \\& Wainscoat 1993) have shown that there are about $7.8 \\times 10^{4}$ galaxies per square degree brighter than $K = 21$\\,mag. Therefore, we expect 0.12 chance coincidences within a 2.5\\,arcsec matching radius in the fields imaged with NIRI and NIRC. With UFTI we reach $K \\sim 19.5$\\,mag (8-arcsec diameter aperture) with the shortest exposures (540 sec.) and $K \\sim 20$\\,mag (8-arcsec diameter aperture) with the deepest ones (3240 sec.). The $K$-band number counts are down to $\\sim 2.9 \\times 10^{4}$ and $\\sim 4.0 \\times 10^{4}$ per square degree, respectively, and the chance coincidences expected within the same matching radius are 0.06 and 0.04. Finally, with UIST we reach $K \\sim 19$\\,mag (8-arcsec diameter aperture), to which the number counts are $\\sim 2.0 \\times 10^{4}$ per square degree. Therefore, we expect 0.03 chance coincidences between a radio and near-infrared source within a 2.5\\,arcsec matching radius.    \\subsection{\\kband magnitude distribution}   \\begin{figure} \\begin{center} \\caption{Histogram of the \\kband magnitudes, measured in a 3-arcsec diameter aperture, of the 6C** sample (solid line) and the 6C* sample (dashed line). The dotted line represents the distribution of $K$-band magnitudes, measured in a 5-arcsec aperture, for the sources in the 7C-I and 7C-II regions of the 7C Redshift Survey. The bin width in each histogram is \\mbox{$\\Delta K = 0.5$~mag}. The non-detected sources (6C**0737+5618 and 6C**0935+4348) are represented by arrows.} \\label{fig:histmag} \\end{center} \\end{figure}  In Fig.~\\ref{fig:histmag} we compare the distribution of $K$-magnitudes of the 6C** sample with those of the similarly selected 6C* sample (Jarvis et al. 2001a) and of the 7C-I and 7C-II regions of 7CRS (Willott et al. 1998, 2002, 2003). The later were selected at the same flux level as 6C** but without any filtering.\\footnote{The 7C-I region covers an area of sky of 0.0061\\,sr and contains 37 sources with $S_{151} > 0.5$\\,Jy. The 7C-II region cover an area of sky 0.0069\\,sr and contains 38 sources with $S_{151} > 0.5$\\,Jy.}.  The median \\kband magnitude of the 6C** sample is $\\approx 18.7$~mag (3-arcsec diameter aperture), with a maximum magnitude of $\\gtrapprox 21.7$~mag, while the 6C* sample has a median \\kband magnitude of $\\approx 18.8$~mag in the same aperture, with the faintest magnitude being $\\approx 20$~mag. It can be seen that although the distributions of both samples peak at the same median magnitude, that of the 6C** sample extends both to brighter and fainter magnitudes, with a slightly asymmetrical tail extending towards bright magnitudes ($K \\simeq 14.5$~mag). A Kolmogorov-Smirnov test shows that the two datasets are consistent with being drawn from the same underlying distribution with probability $p = 0.51$.  Therefore, we expect the 6C** sample to have a similar redshift distribution to that of the 6C* sample.  The 6C* sample has a median redshift of $z \\simeq 1.9$, with a minimum redshift of 0.513 and maximum redshift of 4.410 (Jarvis et al. 2001b).  The median \\kband magnitude of the combined 7-CI and 7C-II regions is $\\approx 17.4$~mag (5-arcsec aperture). This shows that the filtering criteria employed to select the 6C** sample is being effective in biasing the sample towards fainter sources, and possibly higher redshifts. A full derivation of redshift estimates for the members of the 6C** sample, based on their \\kband magnitudes, will be presented in Paper II, along with a detailed analysis of their space density.    \\begin{table} \\small \\begin{center} \\begin{tabular}{lcl} \\hline \\mc{1}{c}{Name} & \\mc{1}{c}{$K$} & \\mc{1}{c}{Comments}\\\\ \\mc{1}{c}{} & \\mc{1}{c}{(mag)} & \\mc{1}{c}{}\\\\ \\hline 6C**0717+5121 & 17.79(5)   &  Faint red continuum, no lines\\\\ 6C**0737+5618 & $>21.8$(3) &  Blank\\\\ 6C**0744+3702 & 19.44(8)   &  Blank\\\\ 6C**0754+4640 & 20.09(5)   &  Blank\\\\ 6C**0813+3725 & 18.50(8)   &  Blank\\\\ 6C**0829+3902 & 19.41(5)   &  Blank\\\\ 6C**0848+4927 & 18.22(8)   &  Blank\\\\ 6C**0902+3827 & 19.29(5)   &  Faint red continuum, no lines\\\\ 6C**0925+4155 & 20.30(5)   &  Blank\\\\ 6C**0938+3801 & 18.13(8)   &  Blank\\\\ 6C**1050+5440 & 19.71(8)   &  Blank\\\\ \\hline \\end{tabular} \\end{center} {\\caption{\\label{tab:no-redshift} Attempted spectroscopic observations not yielding a redshift.}} \\end{table}   \\subsection{Sources without a redshift}\\label{sec:non-redshift2} We now discuss the sources for which we could not determine a redshift from optical spectroscopy (listed in Table~\\ref{tab:no-redshift}).  It is likely that some of the these sources lie in the redshift range $1.2 < z < 1.8$, where it is difficult to measure redshifts, due to the absence of the strongest lines of the radio galaxy spectrum in the optical wavelength region (e.g. Lacy et al. 1999b). Moreover, some may have SEDs which are similar to those of the $1 < z < 2$ EROs, found in the 7CRS redshift survey (Willott et al. 1999), which may be counterparts to radio galaxies which lack strong emission lines in their spectra. Other sources will be at $z > 1.8$, but have a weak emission line spectrum, which falls below our detection limit ($\\sim 5 \\times 10^{-20}$\\,W\\,m$^{-2}$ in $\\sim$ 30~min. exposures).  \\begin{table} \\small \\begin{center} \\begin{tabular}{lcc} \\hline \\mc{1}{c}{Name} & \\mc{1}{c}{$K$} & \\mc{1}{c}{$z$}\\\\ \\mc{1}{c}{} & \\mc{1}{c}{(mag)} & \\mc{1}{c}{}\\\\ \\hline 6C**0714+4616 & 16.3(8) & 1.466\\\\ 6C**0849+4658 & 17.3(8) & -- \\\\ 6C**0922+4216 & 15.9(8) & 1.750 \\\\ 6C**1003+4827 & 16.9(8) & -- \\\\ 6C**1036+4721 & 17.1(8) & 1.758 \\\\ 6C**1052+4349 & 17.1(8) & -- \\\\ 6C**1056+5730 & 17.3(8) & --\\\\ 6C**1138+3803 & 17.3(3) & -- \\\\ \\hline \\end{tabular} \\end{center} {\\caption{\\label{tab:unresolved} Summary of unresolved \\kband sources. We assume that these sources are quasars based on their unresolved continuum emission. Redshifts are given for sources where spectroscopy is available. We note that in all these cases spectroscopy has confirmed the source as a quasar. }} \\end{table}   For all the sources with no measured redshifts it is possible that deeper spectroscopy will enable redshift determination. Indeed, this is the case of 6C**0744+3702, which is not detected in our one hour exposure spectrum (on WHT), but for which Keck spectroscopy (De Breuck et al. 2001) finds a weak-emission-line spectrum at $z = 2.992$ (see also Table~\\ref{tab:literature}). However, as found by De Breuck et al. (2001) some steep-spectrum radio galaxies are not detected even in long exposures ($\\simeq 1-2$~hrs) with 8-10 metre class telescopes. These sources seem to be preferentially characterised by compact radio morphologies ($\\theta \\,\\, \\ltsim \\,\\, 2-3$~arcsec).  One of the objects we targeted (6C**0737+5618) has these characteristics, and is also one of the faintest (\\mbox{$K > 21.8$\\,mag} in a 3-arcsec aperture) objects in our sample. Several attempts in different observing runs were made at obtaining optical spectroscopy for this object. Different position angles and increasingly higher exposure times were used each time. None the less this source remained undetected in our WHT spectra.  Its \\kband magnitude suggests a very high redshift, which could well be above the optical detection limit at $z \\,\\, \\gtsim \\,\\, 6.5$, if the $K-z$ relation holds to such high redshifts. Moreover, if a source is at $z \\,\\gtsim\\, 6$, i.e. within the epoch of reionization, it may have its Ly$\\alpha$ emission severely suppressed by the intervening intergalactic medium (IGM, Gunn \\& Peterson 1965, Becker et al. 2001). However, it is possible that Ly$\\alpha$ remains visible even when it is embedded in a neutral IGM (e.g. Haiman 2002; Santos 2004; Furlanetto, Zaldarriaga \\& Hernquist 2006).  Another possibility is that 6C**0737+5628 is at a lower redshift but heavily obscured (e.g. like WN J0305+3525 in Reuland et al. 2003), possibly due to it being observed shortly after the jet-triggering event.  Given the nature of our selection criteria, i.e. a small radio angular size, the galaxies which are at high-redshift ($z \\gtrsim 2 $) in our sample are preferentially young and seen shortly after the jet-triggering event (Blundell \\& Rawlings 1999).  The projected linear size of the radio emission, which may be used as an age indicator for the radio source, (Blundell, Rawlings \\& Willott 1999) is also found to anti-correlate with the 850\\,$\\mu$m flux-density (Willott et al. 2002), which in turn is an indicator of the dust content of a galaxy. We speculate that the spectra for some of the sources for which we could not obtain a redshift may be similar to those of 6C**0754+5019 and 6C**0935+4348, with higher or lower degrees of Ly$\\alpha$ suppression and/or chemical enrichment indicators as dictated by their evolutionary stage (Vernet et al. 2001)."
    },
    "0609/astro-ph0609759_arXiv.txt": {
        "abstract": "Polaris, the nearest and brightest  classical Cepheid, is a member of at least a triple system.  It has a wide ($18''$) physical companion, the F-type dwarf Polaris B. Polaris itself is a single-lined  spectroscopic binary with an orbital period of 30 years (Kamper, 1996, JRASC, 90, 140). By combining {\\it Hipparcos} measurements of the instantaneous proper motion with long-term measurements and the Kamper radial-velocity  orbit,  Wielen et al. (2000, A\\&A, 360, 399) have  predicted the  astrometric  orbit of the close  companion.  Using the {\\it Hubble  Space Telescope} and the Advanced Camera for Surveys' High-Resolution Channel with an ultraviolet (F220W) filter, we have now directly detected the close companion.  Based on the Wielen et al. orbit, the {\\it Hipparcos} parallax, and our measurement of the separation ($0.176''$ $\\pm$ $0.002''$), we find a preliminary mass of 5.0 $\\pm$ 1.5 M$_{\\odot}$ for the Cepheid and 1.38 $\\pm$ 0.61 M$_{\\odot}$ for the close companion. These  values will be refined by additional {\\it HST} observations scheduled for the next 3 years.  We have also  obtained a {\\it Chandra}  ACIS-I image of the Polaris field. Two distant companions C and D are not X-rays sources  and hence are not young  enough  to be  physical  companions  of  the  Cepheid.  There  is one additional stellar X-ray source in the field, located $253''$ from Polaris A, which is a possible companion. Further investigation of such a distant companion is valuable to confirm the full extent of the system. ",
        "introduction": "Polaris, like most massive stars, is a member of a multiple system.  It is also a supergiant (F5 Ib) which is the nearest and brightest classical Cepheid.  It is a low amplitude somewhat quirky Cepheid with a variable amplitude, which is pulsating in the first overtone mode, as shown by the {\\it Hipparcos} parallax (Feast and Catchpole, 1997)  The goals of this discussion are to present preliminary results on two topics. The first is a direct dynamical measurement of the mass, which is of particular interest because it is a Cepheid.  The second is to explore how many physical companions belong to the system, which is of interest as the ``footprints\" of star formation. ",
        "conclusions": ""
    },
    "0609/astro-ph0609103_arXiv.txt": {
        "abstract": "% Star clusters are subject to density irregularities in their host galaxy, such as giant molecular clouds (GMCs), the galactic disc and spiral arms, which are largely ignored in present day ($N$-body) simulations of cluster evolution. Time dependent external potentials give rise to tidal forces that accelerate stars leading to an expansion and more rapid dissolution of the cluster. I explain the basic principles of this tidal heating in the {\\it impulse approximation} and show how related disruption time-scales depend on properties of the cluster. ",
        "introduction": "Assume a cluster and a moving perturbing object (GMC, spiral arm, disc, etc.) with relative separation and velocity $X$ and $V$, respectively. The acceleration due to the perturber is referred to as $g(X)$. The {\\it tidal} acceleration ($\\at$) on a star at a distance $x$ from the cluster centre can be obtained by linearising $g(X)$: \\begin{equation} \\at=\\frac{\\dr v_{\\rm x}}{\\dr t} = g(X+x)-g(X)\\simeq x\\,\\frac{\\dr g}{\\dr X}. \\label{eq:1} \\end{equation} The total velocity increase ($\\dv$) of the star after the perturbation can be obtained by substituting $\\dr X = V \\dr t$ and integrating Eq.~\\ref{eq:1}. If the time it takes the perturber to pass the cluster is much shorter than the crossing time of stars in the cluster, $x$ can be assumed constant. This is referred to as the {\\it impulsive approximation}. The resulting \\dv\\ then simply scales with $x$, $V$ and the maximum acceleration ($\\gm$) as: $\\dv \\propto x\\,\\gm/V$.  The total kinetic energy gain per unit of cluster mass ($\\Delta E$) can then be written as \\begin{equation} \\Delta E = \\frac{1}{2}|\\dv^2| \\propto \\frac{\\gm^2\\,\\rrms}{V^2}, \\label{eq:2} \\end{equation} where we used that the mean-square position of stars (\\rrms) is $\\rrms=3\\bar{x^2}$. {\\it Eq.~\\ref{eq:2} shows that $\\Delta E$ is large for slow perturbations on large clusters}.  For passing point masses $\\gm=GM/p^2$, with $G$ the gravitational constant, $M$ the mass of the perturber and $p$ the distance of closest approach (see \\citealt{2006MNRAS.tmp..808G} for the exact definition of Eq.~\\ref{eq:2}, taking into account an extended density profile for a GMC). For a one-dimensional density perturbation, such as the perpendicular crossing of a galactic disc and, to lesser extend, the crossing of a spiral arm, $\\gm \\propto GH\\rho_0$, with $H$ the scale height/length and $\\rho_0$ the central density. For more details see \\citet{1987degc.book.....S} and references therein. ",
        "conclusions": ""
    },
    "0609/astro-ph0609335_arXiv.txt": {
        "abstract": "We announce the discovery of the second transiting hot Jupiter discovered by the Trans-atlantic Exoplanet Survey. The planet, which we dub \\trp, orbits the nearby star \\gsc\\ every 2.47063 days.  From high-resolution spectra, we determine that the star has $T_{\\rm eff}=5960\\pm100\\,\\mathrm{K}$ and $\\log{g}=4.4\\pm0.2$, implying a spectral type of G0V and a mass of $1.08^{+0.11}_{-0.05}\\,M_{\\sun}$. High-precision radial-velocity measurements confirm a sinusoidal variation with the period and phase predicted by the photometry, and rule out the presence of line-bisector variations that would indicate that the spectroscopic orbit is spurious.  We estimate a planetary mass of $1.28^{+0.09}_{-0.04}\\,M_{\\rm Jup}$.  We model $B$, $r$, $R$, and $I$ photometric timeseries of the 1.4\\%-deep transits and find a planetary radius of $1.24^{+0.09}_{-0.06}\\,R_{\\rm Jup}$.  This planet lies within the field of view of the NASA \\textit{Kepler} mission, ensuring that hundreds of upcoming transits will be monitored with exquisite precision and permitting a host of unprecedented investigations. ",
        "introduction": "\\label{sec:intro}  Observations of the ten known transiting hot Jupiters have provided precise planetary radii and masses, and tested formation and structure models for extrasolar planets \\citep[see][]{Laughlin_Wolf_Vanmunster:apj:2005a, Charbonneau_Brown_Burrows:PPV:2006a}. More detailed studies of the nearby planets have probed their atmospheres and led to the direct detection of their thermal emission \\citep[e.g.][]{Charbonneau_Brown_Noyes:apj:2002a, Charbonneau_Allen_Megeath:apj:2005a, Deming_Brown_Charbonneau:apj:2005a, Deming_Seager_Richardson:nat:2005a}.  Three of these planets were known from radial-velocity surveys of the solar neighborhood, and were subsequently observed to transit.  The remaining seven were discovered from photometric observations. The radial-velocity confirmation of transiting planet candidates involves extensive use of large-aperture telescopes. With the goal of maximizing the yield of transiting planets around bright stars and minimizing the time required of large observatories, several teams are undertaking wide-field photometric surveys using small telescopes \\citep[for a review, see][]{Charbonneau_Brown_Burrows:PPV:2006a}. Our collaboration is conducting the \\anchor{http://www.astro.caltech.edu/~ftod/tres/}{Trans-atlantic Exoplanet Survey}% \\footnote[14]{\\url{http://www.astro.caltech.edu/$\\sim$ftod/tres/}}% \\ (TrES): TrES--1 was the first nearby transiting planet to be discovered photometrically \\citep{Alonso_Brown_Torres:apjl:2004a}.  Such photometric surveys yield numerous transit candidates, of which the majority are astrophysical false positives that are discarded by follow-up photometric and spectroscopic observations \\citep[e.g.,][]{ODonovan_Charbonneau_Alonso:ApJ-submitted:2006a}. However, eliminating a blend, wherein a bright star forms a chance superposition or a hierarchical triple with a faint eclipsing binary, can require a careful analysis \\citep[][]{Torres_Konacki_Sasselov:apj:2004b, Mandushev_Torres_Latham:apj:2005a, ODonovan_Charbonneau_Torres:apj:2006a}.  We present here the discovery of the planet \\trp, and describe the process by which we confirmed its planetary nature and deduced its bulk properties. ",
        "conclusions": "\\label{sec:discuss}  In order to determine $M_{\\star}$ and its uncertainties, we compared our estimates of $T_{\\rm eff}$ and $\\log{g}$ with evolutionary models from \\cite{Yi_Demarque_Kim:apjs:2001a}, assuming solar metallicity. For each isochrone, we identified the range of $M_{\\star}$ for which the $T_{\\rm eff}$ and $\\log{g}$ lay within our 1\\,$\\sigma$ errors.  We took the best-fit model as our estimate of $M_{\\star}$, and the span of permitted models (over all ages greater than 500~Myr) to be our uncertainty.  We then used the resulting value, $M_{\\star} = 1.08^{+0.11}_{-0.05}\\, M_{\\Sun}$, and the spectroscopic orbit (\\S2) to estimate $M_{p} = 1.28^{+0.09}_{-0.04}\\, M_{\\rm Jup}$.  We also evaluated the stellar radius $R_{\\star}$ in a similar fashion, and found results that were consistent with, but less tightly constrained than that from the light-curve modeling (below). The uncertainty contributed by that in the evolutionary models is less than 0.02 solar masses. Based on the absolute visual magnitude ($M_{V}$=4.5) predicted by the best-fit model, we estimate the distance to be approximately 230\\,pc. We estimate the reddening in the direction of \\trs\\ to be $E(B-V) \\sim 0.05$ and the extinction to be $\\sim 0.15$\\,mag from comparison of its observed colors with the intrinsic colors predicted by the model.  To estimate $R_{\\star}$, $i$, and the planetary radius $R_{p}$, we simultaneously fit our light curves using the analytical transit curves of \\cite{Mandel_Agol:apjl:2002a} and the color-dependent quadratic limb-darkening parameters from \\cite{Claret:aa:2000a}, which were matched to the spectroscopically-estimated properties of the star.  We identified the best-fit solution by fixing the value of $M_{\\star}$ at its best estimate, $1.08\\ M_{\\Sun}$, and minimizing the ${\\chi}^2$ to all the photometry.  We note that the available time series are well-described by the model, with the exception of the IAC80 $B$, for which the in-transit data fall below the model.  We speculate that those data, which were gathered at high airmass, may have been imperfectly corrected for extinction, which is a larger effect at $B$ than the other band passes.  The best-fit solution obtains ${\\chi}_{\\nu}^2=1.15$ ($\\nu=2065$), and its values for $\\{R_{\\star}, R_{p},i\\}$ are listed in Tables~\\ref{tab:tres2} and \\ref{tab:tres2b}. The uncertainties in these quantities are dominated by our uncertainty in $M_{\\star}$. To derive 1\\,$\\sigma$ errors for each of $\\{R_{\\star}, R_{p}, i\\}$, we change the value of that parameter and fix it at a new value, and then allow the other two parameters to float, as well as allow for a value of $M_{\\star}$ within our uncertainty. (The uncertainties in $P$ and $T_{c}$ are sufficiently small so as not to contribute significantly to the errors in $R_{\\star}$, $R_{p}$, and $i$.) We repeat this procedure until the best-fit solution produces an increase in the ${\\chi}^2$ corresponding to a 1\\,$\\sigma$ change.  Our estimate of the planetary radius, $R_{p} = 1.24^{+0.09}_{-0.06}\\ R_{\\rm Jup}$, implies a mean density of $0.83^{+0.12}_{-0.09} \\ {\\rm g\\, cm^{-3}}$, indistinguishable from that of TrES-1 \\citep[using the values from][]{Sozzetti_Yong_Torres:apjl:2004a}, despite the fact that \\trs\\ is nearly twice as massive. We also note that the impact parameter, $b = a\\, \\cos{i} / R_{\\star} = 0.84\\pm0.02$, is the largest of any known transiting exoplanet.  We intend to improve our estimates of the planetary and stellar parameters by undertaking a more detailed analysis of the stellar spectrum as we did for TrES--1 \\citep{Sozzetti_Yong_Torres:apjl:2004a}, and by gathering very high-precision $z$-band photometry \\citep[e.g.][]{Holman_Winn_Latham:preprint:2006a}.  Such data will permit us to look for transit timing variations indicative of additional planets in the \\trs\\ system \\citep{Agol_Steffen_Sari:mnras:2005a, Holman_Murray:Science:2005a, Steffen_Agol:mnras:2005a}. \\trs\\ lies within the FOV of the NASA \\textit{Kepler} mission. During the four year mission, \\textit{Kepler} will observe nearly 600 transits of \\trs. The precision with which \\textit{Kepler} will observe these transits will enable an extremely sensitive search for additional planets in the \\trs\\ system through their dynamical perturbations.  Moreover, the large impact parameter means that very subtle changes in its value could be detected. Such variations are predicted \\citep{Miralda-Escude:apj:2002a} to occur as a result of either additional planets, or the stellar quadrupole moment. \\textit{Kepler} may also detect the reflected light from \\trs\\ \\citep{Jenkins_Doyle:apj:2003a} and hence determine the long-sought geometric albedo and phase function of a hot Jupiter.  The large impact parameter also makes \\trs\\ particularly favorable for determining the angle between the stellar spin-axis and the orbital axis via the Rossiter--McLaughlin effect \\citep{Gaudi_Winn:preprint:2006a}. \\cite{Williams_Charbonneau_Cooper:preprint:2006a} discuss the use of {\\it Spitzer} IRAC observations spanning the time of secondary eclipse to resolve the surfaces of extrasolar planets.  The large impact parameter of the \\trs\\ orbit is ideal for this application, since it grants access to both longitudinal and latitudinal flux variations across the dayside hemisphere of the planet."
    },
    "0609/astro-ph0609453_arXiv.txt": {
        "abstract": "The High Resolution Fly's Eye (HiRes) experiment has observed the GZK cutoff.  HiRes observes two features in the ultra-high energy cosmic ray (UHECR) flux spectrum: the Ankle at an energy of $4\\times10^{18}$ eV and a high energy suppression at $6\\times10^{19}$ eV.  The later feature is at exactly the right energy for the GZK cutoff according to the $E_{1/2}$ criterion. HiRes cannot claim to observe a third feature at lower energies, the Second Knee.  The HiRes monocular spectra are presented, along with data demonstrating our control and understanding of systematic uncertainties affecting the energy and flux measurements. \\vspace{1pc} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609665_arXiv.txt": {
        "abstract": "We present $IZJHKL'$ photometry of the core of the cluster NGC\\,6611 in the Eagle Nebula. This photometry is used to constrain the Initial Mass Function (IMF) and the circumstellar disk frequency of the young stellar objects. Optical spectroscopy of 258 objects is used to confirm membership and constrain contamination as well as individual reddening estimates. Our overall aim is to assess the influence of the ionizing radiation from the massive stars on the formation and evolution of young low-mass stars and their disks. The disk frequency determined from the $JHKL'$ colour-colour diagram suggests that the ionizing radiation from the massive stars has little effect on disk evolution (\\cite[Oliveira et al. 2005]{oliveira05}). The cluster IMF seems indistinguishable from those of quieter environments; however towards lower masses the tell-tale signs of an environmental influence are expected to become more noticeable, a question we are currently addressing with our recently acquired ultra-deep (ACS and NICMOS) HST images. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609386_arXiv.txt": {
        "abstract": "Observations of the inner radian of the Galactic disk at very high energy (VHE) gamma-rays have revealed at least 16 new sources. Besides shell type super-nova remnants, pulsar wind nebulae (PWN) appear to be a dominant source population in the catalogue of VHE gamma-ray sources.  Except for the Crab nebula, the newly discovered PWN  are resolved at VHE gamma-rays to be spatially extended (5-20 pc). Currently, at least  3 middle aged ($t>10$ kyrs) PWN (Vela X, G18.0-0.7, and G313.3+0.6 in the ``Kookaburra'' region) and  1 young PWN MSH 15-5{\\it2} ($t=1.55$~kyrs) have been identified to be VHE emitting PWN (sometimes called ``TeV Plerions'').  Two more candidate ``TeV Plerions'' have been identifed and have been reported at this conference \\cite{Svenja}. In this contribution, the gamma-ray emission from Vela X is explained by a nucleonic component in the pulsar wind. The measured broad band spectral energy distribution is compared with the expected X-ray emission from primary and secondary electrons. The observed X-ray emission and TeV emission from the three middle aged PWN are compared with each other. ",
        "introduction": "\\label{intro} The discovery of extended gamma-ray emission from pulsar wind nebulae (PWN) opens an exciting possibility to study the acceleration of particles in ultra-relativistic shocks. Acceleration at relativistic shocks is of relevance for the understanding of the mostly non-thermal emission seen from relativistic jets in active galactic nuclei (AGN) and gamma-ray bursts (GRBs). While the Lorentz factor assigned to the relativistic flow in AGN and GRBs is currently believed to be in the range of $\\Gamma\\approx 10\\ldots100$ based upon arguments of opacity against pair production, the Lorentz factor of pulsar winds $\\gamma$ is largely undetermined. Using the potential drop of the last open field lines as a constraint: $\\gamma\\le e\\Phi_\\mathrm{open}/(m_e c^2)=2\\times 10^8 (\\dot P_{-13}/P)^{1/2}$ with $P$ indicating the rotational period and $dP/dt=\\dot P_{-13}\\times 10^{-13}$~s/s the rate of slowing down. The Lorentz factor has been determined in model dependent ways to be in the range of $10^6-10^7$ for the Crab nebula (see e.g. \\cite{Arons96} and \\cite{BogoAhar00}). Even if this value is possibly smaller in other PWN, we are witnessing particle acceleration in ultrarelativistic shocks driven by the energetic pulsar wind.  The total energy available by slowing down of the pulsar (assuming a moment of inertia $I=I_{45} 10^{45}$~g~cm$^2$): $\\dot E=10^{34}~I_{45} \\dot P_{-13}/P^3$~erg/s.\\\\ Theoretical calculations of Fermi-type acceleration in relativistic shocks using widely different approaches (see e.g.  \\cite{BedOstro98}, \\cite{Achter01}) lead to the consistent conclusion, that (largely independent of the detailed parameters of the shock)  particles in the downstream plasma are accelerated to follow a universal power-law for large Lorentz factors of $s= 2.2\\ldots 2.3$ which is compatible with e.g. the observed X-ray synchrotron spectrum of e.g. the Crab nebula.\\\\ However, recent progress in the understanding of the downstream turbulence spectrum, a revision of this ``universal acceleration'' picture may be necessary (see e.g. \\cite{LePeRe06}, \\cite{NiemOs06}). In the more recent calculations, the particle spectrum is found to be softer  than in previous calculations (see above), to deviate from a universal power-law, and  to show a stronger dependence on the Lorentz factor of the upstream medium and on the particular implementation of downstream turbulence. Similar conclusions have been found independently with particle-in-cell (PIC) simulations (\\cite{Hoshi92}, \\cite{Spit05}, \\cite{AmAr06}).\\\\ The presence of ions in the pulsar wind obviously complicate the structure of the shock and have been described e.g. in \\cite{HoshiAr91}, \\cite{Gall92}, \\cite{Hoshi92}. Hoshino et al. (1992) \\cite{Hoshi92}  found that an admixture of ions in the wind can lead to acceleration of positrons in the downstream region by resonant absorption of magnetosonic waves emitted by the gyrating ions. In a recent PIC simulation \\cite{AmAr06} a larger ratio of ion mass $m_i$ to electron mass $m_e$ in the simulation has been used (previously, PIC simulations had been limited to values of $m_i/m_e\\approx 20$) and various energy fraction of ions have been considered. The overall efficiency of acceleration has been found to increase with the relative energy fraction carried by ions in the wind. The authors also show that the spectrum of non-thermal particles varies. \\\\ The presence of the so-called  wisps in the Crab nebula has been used to argue for the presence of ions in the wind (see e.g. \\cite{GallAr96}): the compression caused by the ions leads to increased magnetic field and correspondingly more intense synchrotron radiation. Wisp-like features have been observed from PSR B1509-58 \\cite{Gaens02}, however, the predicted time variation for the northern arc was not found  in later observations \\cite{Del06}. The Vela PWN shows variable features along its jet which have been interpreted to be the result of a kink instability \\cite{Pav03}. The ring like features in the Vela PWN have not been seen to vary with time even though the expected time scale should be similar to the one observed from the Crab nebula.   \\\\ Besides the existence of wisp-like structures near the wind shock, ions are expected to leave other observable signatures. The downstream energy distribution  of the ions follows a relativistic Maxwellian distribution with some modifications due to the energy loss of ion energy transfered to the non-thermal tail of accelerated pairs \\cite{Hoshi92}. The temperature of the distribution is close to $\\gamma m_i c^2$. At some distance to the shock, the ions will move diffusively outwards, loosing energy by adiabatic expansion. Gamma-rays and neutrinos will be produced predominantly in inelastic scattering of the nucleons on the ambient medium. Depending on the diffusion coefficient in the PWN, a large fraction of the particle energy can be converted into gamma-rays and neutrinos. \\\\ Here, we consider specifically the Vela X PWN, which has recently been detected to emit VHE gamma-rays \\cite{HESS06}. In the final section, we will discuss similarities and differences to the other ``TeV Plerions''. For a more general review of gamma-ray production in PWN see \\cite{Bed06} in these proceedings. \\begin{figure} \\centering \\includegraphics[width=\\linewidth]{2-10kev_total_hsv.eps} \\caption{In grey scale (see the online proceedings for a color version): ASCA 2-10 keV mosaic showing the flux in bins of 0.25 arc~min$^2$ size. The contours represent the HESS significance map \\cite{HESS06} with the contours starting at $5~\\sigma$ above the background noise and incrementing by one standard deviation.} \\label{fig:1}       % \\end{figure} ",
        "conclusions": ""
    },
    "0609/astro-ph0609209_arXiv.txt": {
        "abstract": "{All galaxies that have been adequately examined so far have shown an extended stellar halo.} {To search for such a halo in the LMC we have obtained low-resolution spectra for 100 LMC RR Lyrae stars, of which 87 are in the field and 13 in the clusters NGC1835 and NGC2019.} {We measured radial velocities for 87 LMC RR Lyrae stars, and metallicities for 78 RR Lyrae stars, nearly tripling the previous sample. These targets are located in 10 fields covering a wide range of distances, out to 2.5 degrees from the center of the LMC. } {Our main result is that the mean velocity dispersion for the LMC RR Lyrae stars is $\\sigma_{RV} = 50 \\pm 2$ km/s. This quantity does not appear to vary with distance from the LMC center. The metallicity shows a Gaussian distribution, with mean [Fe/H]= -1.53 $\\pm 0.02$ dex, and dispersion $\\sigma_{[Fe/H]} = 0.20\\pm 0.02$ dex in the Harris  metallicity scale, confirming that they represent a very homogeneous metal-poor population. There is no dependence between the kinematics and metallicity of the field RR Lyrae star population.} {Using good quality low-resolution spectra from FORS1, FORS2 and GEMINI-GMOS we have found that field RR Lyrae stars in the LMC show a large velocity dispersion and that this indicate the presence of old and metal-poor stellar halo. All the evidence so far for the halo, however, is from the spectroscopy of the inner LMC regions, similar to the inner flattened halo in our Galaxy. Further study is necessary to confirm this important result.} ",
        "introduction": "The LMC is an irregular, disk dominated, bulge-less, nearly face-on galaxy. It is of fundamental importance for studies of stellar populations and of the interstellar medium. It is being used to study the presence of dark objects in the Galactic Halo through microlensing (e.g., Alcock et al. 2000a,b), it is an important probe of the Galaxy's formation history, and it plays a key role in determinations of the cosmological distance scale (e.g., Freedman et al. 2001).  All galaxies that have been adequately examined so far have shown an extended star halo. Furthermore the Lambda CDM cosmology predicts a dark matter halo for galaxies like the LMC. Unfortunately, outlining a stellar halo in the LMC (and its metal and age distribution), a key observation for galaxy formation scenarios, is made difficult by the large extent of this galaxy on the sky and inclination, causing the halo to be projected against the LMC disk.  Recently, using low resolution spectroscopy with the FORS1 spectrograph at the Very Large Telescope (VLT), Minniti et al. (2003) and Borissova et al. (2004; hereafter Paper~I) on the basis of  the radial velocities of 43 RR Lyrae stars obtained a true velocity dispersion of 53 km/s, which is higher than the  velocity dispersion of any other LMC population previously measured. This high  velocity dispersion indicates that a kinematically hot metal-poor old halo exists in the LMC.  Cole at al. (2005) reported metallicities and radial velocities derived from spectra at the near-infrared calcium triplet for 373 red giants in a 200 arcmin$^2$ area at the optical center of the LMC bar. They found that the velocity dispersion of the whole sample is $24.7\\pm0.4$ km~s$^{-1}$. When cut by metallicity, the most metal-poor 5$\\%$ of stars ($\\mathrm{[Fe/H]} < -1.15$) show $40.8\\pm1.7$ km~s$^{-1}$, more than twice the value for the most metal-rich suggesting that an old, thicker disk or halo population is present, in agreement with our results from RR Lyrae stars.  The counter-argument against the evidence of a spherical halo in the outer regions of LMC was presented by Alves (2004a,b), who re-examined the content of the 2MASS database at its periphery. Based on the dereddened colors of the LMC's red giant branch stars, he found an outward radial gradient of decreasing metallicity. The isodensity contours of the RGB stars rule out the existence of a spherical halo at the periphery, and reveal instead an inclined disk. Similar evidence, again from photometric data, was presented by Gallart et al.\\ (2004) from a high-quality color-magnitude diagram (CMD) for a $36 \\times 36$ arcmin field located at 7 kpc from the LMC center. The surface brightness profile of the LMC remains exponential to this large galactocentric radius and shows no evidence of disk truncation. Combining the information on surface brightness and stellar population, they concluded that the LMC disk extends (and dominates over a possible stellar halo) out to a distance of at least 7 kpc. Clementini et al. (2003) and Gratton et al. (2004) presented new photometry and spectroscopy for more than a hundred RR Lyrae stars in two fields located close to the bar of the LMC. Their average magnitude, the local reddening, individual metallicities, the luminosity-metallicity relation, and the distance to the LMC were derived. They calculated the dispersion of the mean radial velocity of the 48 stars with multiple observations of 40 km/s.  Subramaniam (2006) investigated the distribution of the RR Lyrae star population in the inner LMC and found a halo-like location and disc-like density distribution. He also found that RR Lyrae stars in the inner LMC are elongated more along the line of sight than along the major axis of the disk. Thus there are many indications that the LMC has a halo, which is similar to the inner halo of our Galaxy, but the results are puzzling. The outer halo of the LMC, which is characterized by a spatially extended old population with large velocity dispersion, has still not been detected.  In this paper we present new spectroscopic data extending our sample to 137 field RR Lyrae stars and 13 RR Lyrae stars in the LMC globular clusters NGC1835 and NGC2019. The next section describes the data and the third section discusses the radial velocities of the extended sample. The fourth section includes the metallicities and the last section is a summary of the results. ",
        "conclusions": "We have presented new low-resolution spectra for 100 RR Lyrae field stars in the Large Magellanic Cloud. This sample more than triples that presented in Paper~I and allows us to strengthen our previous results regarding the kinematics and metallicity properties of this old and metal-poor population in our nearest neighboring galaxy.   The average metallicity of the whole sample of RR Lyrae stars in the LMC is [Fe/H]= -1.53$ \\pm 0.02$ dex. The metallicity distribution is well fitted by a single Gaussian. There is no significant variation of the velocity dispersion with metallicity. This is at odds with the results of Cole et al.\\ (2005) who found an increasing velocity dispersion with decreasing metallicity. However, these authors looked at different stellar populations in the LMC, while our sample comprises only RR Lyrae stars. Thus the trend between the metallicity and velocity dispersion may be due to age rather than to metallicity.  Our main result is that the mean velocity dispersion of the LMC field RR Lyrae stars is $\\sigma_{RV} = 50 \\pm 2$ km/s. This quantity does not appear to vary with distance from the LMC center up to 2 degrees.  The mean velocities show some field to field variation, and our conclusions are weaker for this quantity. They appear to differ slightly from the velocities of the HI gas, but our data do not allow us to measure the bulk rotation."
    },
    "0609/nucl-th0609035_arXiv.txt": {
        "abstract": "Within a self-consistent thermal model using an isospin and momentum dependent interaction (MDI) constrained by the isospin diffusion data in heavy-ion collisions, we investigate the temperature dependence of the symmetry energy $E_{sym}(\\rho ,T)$ and symmetry free energy $F_{sym}(\\rho ,T) $ for hot, isospin asymmetric nuclear matter. It is shown that the symmetry energy $E_{sym}(\\rho ,T)$ generally decreases with increasing temperature while the symmetry free energy $F_{sym}(\\rho ,T)$ exhibits opposite temperature dependence. The decrement of the symmetry energy with temperature is essentially due to the decrement of the potential energy part of the symmetry energy with temperature. The difference between the symmetry energy and symmetry free energy is found to be quite small around the saturation density of nuclear matter. While at very low densities, they differ significantly from each other. In comparison with the experimental data of temperature dependent symmetry energy extracted from the isotopic scaling analysis of intermediate mass fragments (IMF's) in heavy-ion collisions, the resulting density and temperature dependent symmetry energy $% E_{sym}(\\rho ,T) $ is then used to estimate the average freeze-out density of the IMF's. ",
        "introduction": "The equation of state (EOS) of isospin asymmetric nuclear matter, especially the nuclear symmetry energy, is essential in understanding not only many aspects of nuclear physics, but also a number of important issues in astrophysics \\cite% {ireview98,ibook,bom,diep03,pawel02,lat01,baran05,steiner05,listeiner}. Information about the symmetry energy at zero temperature is important for determining ground state properties of exotic nuclei and properties of cold neutron stars at $\\beta $-equilibrium, while the symmetry energy or symmetry free energy of hot neutron-rich matter is important for understanding the liquid-gas phase transition of asymmetric nuclear matter, the dynamical evolution of massive stars and the supernova explosion mechanisms. Heavy-ion reactions induced by neutron-rich nuclei provide a unique means to investigate the symmetry energy \\cite{ireview98,ibook,baran05}. In particular, recent analyses of the isospin diffusion data in heavy-ion reactions \\cite{betty04,chen05,lichen05} have already put a stringent constraint on the symmetry energy of cold neutron-rich matter at sub-normal densities. On the other hand, the temperature dependence of the symmetry energy or symmetry free energy for hot neutron-rich matter has received so far little theoretical attention \\cite{chen01,zuo03,lichen06EsymT}.  For finite nuclei at temperatures below about $3$ MeV, the shell structure and pairing as well as vibrations of nuclear surfaces are important and the symmetry energy was predicted to increase slightly with the increasing temperature \\cite{don94,dean}. Interestingly, an increase by only about $8\\%$ in the symmetry energy in the range of $T$ from $0$ to $1$ MeV was found to affect appreciably the physics of stellar collapse, especially the neutralization processes \\cite{don94}. At higher temperatures, one expects the symmetry energy to decrease as the Pauli blocking becomes less important when the nucleon Fermi surfaces become more diffused at increasingly higher temperatures \\cite{chen01,zuo03,lichen06EsymT}. Based on a simplified degenerate Fermi gas model at finite temperatures, two of present authors \\cite{lichen06EsymT} have recently studied the temperature dependence of the symmetry energy and it was shown that the experimentally observed decrease of the nuclear symmetry energy with the increasing centrality or the excitation energy in isotopic scaling analyses of heavy-ion reactions can be well understood analytically within the degenerate Fermi gas model. In particular, it was argued that the symmetry energy extracted from isotopic scaling analyses of heavy-ion reactions reflects the symmetry energy of \\emph{bulk nuclear matter} for the emission source. Furthermore, it was found that the evolution of the symmetry energy extracted from the isotopic scaling analysis is mainly due to the variation in the freeze-out density rather than temperature when the fragments are emitted in the reactions carried out under different conditions.  In the present work, within a self-consistent thermal model using an isospin and momentum dependent interaction (MDI) constrained by the isospin diffusion data in heavy-ion collisions, we study systematically the temperature dependence of the nuclear matter symmetry energy $E_{sym}(\\rho ,T)$ and symmetry free energy $F_{sym}(\\rho ,T)$. It is shown that the nuclear matter symmetry energy $E_{sym}(\\rho ,T)$ generally decreases with increasing temperature while the symmetry free energy $F_{sym}(\\rho ,T)$ exhibits opposite temperature dependence. The decrement of the symmetry energy with temperature is essentially due to the decrement of the potential energy part of the symmetry energy with temperature. Furthermore, the difference between the nuclear matter symmetry energy $E_{sym}(\\rho ,T)$ and symmetry free energy $F_{sym}(\\rho ,T)$ is found to be quite small around nuclear saturation density. Using the resulting density and temperature dependent symmetry energy $E_{sym}(\\rho ,T)$, we estimate the average freeze-out density of the fragment emission source based on the measured temperature dependent symmetry energy from the isotopic scaling analysis in heavy-ion collisions.  The paper is organized as follows. In Section \\ref{EOS}, we introduce the isospin and momentum dependent MDI interaction and the detailed numerical method to obtain the EOS of the symmetric nuclear matter and pure neutron matter at finite temperatures. Results on the temperature dependence of the symmetry energy and symmetry free energy are presented in Section \\ref% {symmetry energy}. In Section \\ref{isoscaling}, we discuss the experimental data of the isotopic scaling in heavy-ion collisions by means of the obtained density and temperature dependent symmetry energy. A summary is given in Section \\ref{summary}. ",
        "conclusions": "\\label{summary}  Within a self-consistent thermal model using the isospin and momentum dependent MDI interaction with $x=0$ and $-1$ constrained by the isospin diffusion data in heavy-ion collisions, we have investigated the temperature dependence of the nuclear matter symmetry energy $E_{sym}(\\rho ,T)$ and symmetry free energy $F_{sym}(\\rho ,T)$. It is shown that the nuclear matter symmetry energy $E_{sym}(\\rho ,T)$ generally decreases with increasing temperature while the symmetry free energy $F_{sym}(\\rho ,T)$ exhibits opposite temperature dependence. The decrement of the symmetry energy with temperature is essentially due to the decrement of the potential energy part of the symmetry energy with temperature. The temperature effects on the nuclear matter symmetry energy and symmetry free energy are found to be stronger at lower densities while become much weaker at higher densities. Furthermore, the difference between the nuclear matter symmetry energy $% E_{sym}(\\rho ,T)$ and symmetry free energy $F_{sym}(\\rho ,T)$ is found to be quite small around nuclear saturation density, although significantly large at very low densities.  Comparing the theoretical density and temperature dependent symmetry energy $% E_{sym}(\\rho ,T)$ with the $C_{sym}(\\rho ,T)$ parameter extracted from the isotopic scaling data from TAMU and the INDRA-ALADIN collaboration at GSI, we found that the experimentally observed evolution of the symmetry energy is mainly due to the change in density rather than temperature, as shown in the previous work \\cite{lichen06EsymT}. Meanwhile, we have estimated the average freeze-out density of the fragment emission source formed in these reactions by comparing the calculated $E_{sym}(\\rho ,T)$ or $F_{sym}(\\rho ,T) $ with the measured $C_{sym}(\\rho ,T)$. Our results indicate that the extracted average freeze-out densities are not sensitive to whether the experimentally measured $C_{sym}(\\rho ,T)$ parameter is the symmetry energy $% E_{sym}(\\rho ,T)$ or the symmetry free energy $F_{sym}(\\rho ,T)$ in the temperature and density ranges reached in the TAMU and INDRA/GSI experiments. They are, however, sensitive to the $x$ parameter used in the MDI interaction, namely, the density dependence of the symmetry energy. Therefore the isotopic scaling in heavy-ion collisions provides a potentially good probe for the density dependence of the nuclear matter symmetry energy provided the average density of the emission source can be determined simultaneously in the isotopic scaling measurements."
    },
    "0609/astro-ph0609501_arXiv.txt": {
        "abstract": "We present the first $K^\\prime$-band, long-baseline interferometric observations of the northern Be stars $\\gamma$~Cas, $\\phi$~Per, $\\zeta$~Tau, and $\\kappa$~Dra.  The measurements were made with multiple telescope pairs of the CHARA Array interferometer, and in every case the observations indicate that the circumstellar disks of the targets are resolved. We fit the interferometric visibilities with predictions from a simple disk model that assumes an isothermal gas in Keplerian rotation.  We derive fits of the four model parameters (disk base density, radial density exponent, disk normal inclination, and position angle) for each of the targets.  The resulting densities are in broad agreement with prior studies of the IR excess flux and the resulting orientations generally agree with those from interferometric H$\\alpha$ and continuum polarimetric observations. We find that the angular size of the $K^\\prime$ disk emission is smaller than that determined for the H$\\alpha$ emission, and we argue that the difference is the result of a larger H$\\alpha$ opacity and the relatively larger neutral hydrogen fraction with increasing disk radius.  All the targets are known binaries with faint companions, and we find that companions appear to influence the interferometric visibilities in the cases of $\\phi$~Per and $\\kappa$~Dra.  We also present contemporaneous observations of the H$\\alpha$, H$\\gamma$, and Br$\\gamma$ emission lines. Synthetic model profiles of these lines that are based on the same disk inclination and radial density exponent as derived from the CHARA Array observations match the observed emission line strength if the disk base density is reduced by $\\approx 1.7$ dex. ",
        "introduction": "%  Be stars are rapidly rotating B-type stars that manage to eject gas into a circumstellar disk (observed in H emission lines, an IR excess flux, and linear polarization; \\citealt{por03}).  Be star disks are ephemeral and vary on timescales from days to decades (in some cases disappearing altogether for extended periods).   This inherent time variability suggests that gas injection into the disk is only partially the result of the fast spin and equatorial extension of the Be star \\citep*{tow04} and that magnetic, pulsational, wind driving, and/or other processes are required for mass loss into the disk \\citep{owo05}.  The evolutionary status of Be stars is still a subject of considerable debate \\citep*{mcs05,zor05}, but there is growing evidence that many Be stars were spun up through mass transfer from a binary companion that has since become a neutron star, white dwarf, or hot subdwarf \\citep{pol91,gie00}.  The direct resolution of Be star disks has come through high resolution radio observations ($\\psi$~Per; \\citealt{dou92}) and long baseline interferometry in a narrow band around the H$\\alpha$ emission feature \\citep{ste05}. The first H$\\alpha$ studies focused on the bright, northern sky Be stars $\\gamma$~Cas \\citep*{tho86,mou89,ste95,ber99} and $\\zeta$~Tau \\citep{vak98}.  In a seminal paper, \\citet{qui97} presented results for seven Be stars based upon narrow-band H$\\alpha$ observations made with the MkIII interferometer at Mount Wilson Observatory.  They were able to resolve the disks around all seven stars, and they found diameters in the range $2.6 - 4.5$~mas (Gaussian FWHM).  In most cases the visibilities indicated an elongated shape; for example, in the case of $\\zeta$~Tau, the major axis was resolved\u00a0(4.53~mas diameter) but the orthogonal minor axis remained unresolved, consistent with the suggestion that we are viewing the disk almost edge-on.\u00a0 The disk orientation position angle on the sky inferred from polarimetry matched the interferometric results in every case.\u00a0 These results have been extended and improved through new interferometric observations with the Navy Prototype Optical Interferometer (NPOI) in an important series of papers by Tycner and collaborators \\citep{tyc03,tyc04,tyc05,tyc06}.  The relative flux contribution from the disk becomes larger in the infrared spectral range due to bound-free and free-free emission from the dense and ionized gas in the disk. \\citet{wat86} developed a disk model for the IR emission that he subsequently applied to determine the disk properties of 101 Be stars based upon their far-IR emission observed with the {\\it IRAS} satellite \\citep*{cot87,wat87}, and \\citet{dou94} further extended this analysis to study the near-IR excess in 144 Be stars. \\citet{how01} demonstrated that the near-IR flux excess is correlated with the H$\\alpha$ emission strength in Be stars as expected for a common origin in the disk.  Thus, the near- and mid-IR ranges \\citep*{rin99} offer a new opportunity for resolving the disks of Be stars through long baseline interferometry.  \\citet{ste01} presented a wind model for the infrared emission from Be stars and the corresponding interferometric visibilities for the models.\u00a0 Their model for $\\gamma$~Cas, for example, predicts that the disk should appear approximately two times larger in the $K$-band continuum and in the Br$\\gamma$ emission line than found by narrow-band H$\\alpha$ interferometry.  \\citet{ste03} made additional predictions for 16 Be stars of the expected $K$-band interferometric visibilities.   The first attempt to resolve the infrared disk flux was recently made by \\citet{che05} who made $N$-band observations with the VLTI/MIDI instrument of the bright southern Be star $\\alpha$~Ara, and their results indicate that the disk is smaller than that predicted by \\citet{ste03}, perhaps due to disk truncation by a faint binary companion.   $K$-band interferometric observations from VLTI/VINCI of $\\alpha$~Ara were reported by \\citet{dom03} who interpreted the visibilities using a rotationally distorted photospheric model.  Here we present the first near-IR interferometric observations of four bright, northern sky Be stars that we obtained with the Georgia State University Center for High Angular Resolution Astronomy (CHARA) Array at Mount Wilson Observatory \\citep{mca05,ten05}. We describe the interferometric and complementary spectroscopic observations in \\S2 and then we present a simple disk model in \\S3 that we use to predict the interferometric visibilities and spectral line profiles.   We discuss the specific results for each of the four targets in \\S4.  Our results are summarized in \\S5. ",
        "conclusions": "%  The CHARA Array observations represent the first $K^\\prime$ interferometric measurements of northern Be stars, and we find that the disks of all four targets are at least partially resolved.  The derived disk inclination angle and position angle in the sky generally agree with similar estimates from narrow-band H$\\alpha$ interferometry where available. We find, for example, that the disk of $\\zeta$~Tau appears with a nearly edge-on orientation, which is consistent with the prediction for Be stars with shell spectral features \\citep{han96}. All four of the targets are close binaries with faint companions, and these companions appear to influence the $K^\\prime$ visibilities of $\\phi$~Per and $\\kappa$~Dra, an interesting but complicating factor for the interpretation of the disk properties. We find, for example, that the radial gradient in disk density (density exponent $n$ in Table~6) is smaller in binaries with smaller semimajor axes, which suggests that binary companions do influence the disk properties.  We developed a simple Keplerian disk model to create synthetic angular images of the disks, and we used these models to create model interferometric visibility patterns and spectral emission line profiles for comparison with our observations.   These models assume isothermal disks and ignore scattering in the radiative transfer process, two significant limitations that should be addressed in a more complete model.  The disks are characterized by two physical parameters, the base density and density exponent, and two observational parameters, the disk normal inclination and position angle.  We fit the CHARA Array visibilities for all four parameters for both single and binary star models, and our results are in good agreement with prior studies of the IR-excess in these stars.  The parameters derived from the interferometry lead to model emission lines that match the observed ones if the base densities are reduced by a factor of about 1.7~dex.  This mismatch is not unexpected considering the many simplifying assumptions made in the model.  We find that the angular diameters of the Be disks in the $K^\\prime$-band are consistently smaller than those found from H$\\alpha$ interferometry. We think that this difference results from the larger H$\\alpha$ opacity and the hydrogen ionization structure in the disk. The $K^\\prime$ disk emission comes primarily from free-free processes in the mainly ionized gas, and the free-free optical depth will vary with disk radius according to the square of the density.  In the thin disk approximation, the optical depth at radial distance $R$ will vary as $\\tau \\propto \\rho_0^2 ~R^{-2n+3/2} / \\cos i$ (see eq.\\ [2] and [3]).  On the other hand, the H$\\alpha$ emission depends on the relatively low neutral hydrogen number density, and we assumed above that this varies inversely with the stellar dilution factor (eq.\\ [9]; $W(r)\\propto R^{-2}$ at large radius and $Z=0$). Thus, we expect the H$\\alpha$ optical depth to vary approximately as $\\tau \\propto \\rho_0^2 ~R^{-2n+3/2+2} / \\cos i$, i.e., with a shallower effective density exponent of $n-1$. Consequently, the H$\\alpha$ optical depth reduction with increasing disk radius will be less pronounced than in the case of the IR optical depth, resulting in a larger spatial extension in the H$\\alpha$ emission compared to that for the $K^\\prime$ emission.  This expectation is borne out in the predicted spatial images. We show an example of the derived spatial structure in Figure~14 for our single star model of $\\gamma$~Cas.  The solid line shows the summed $K^\\prime$-band intensity of the image projected onto the major axis of the disk and measured in stellar radii from the center of the star. We used our model to create an H$\\alpha$ image adopting the base density from the fit of the observed emission line (Table~8).  The intensity was formed by summing across a 2.8~nm wavelength band centered on H$\\alpha$ in order to match the H$\\alpha$ interferometric observations of \\citet{tyc06}. Note that changing the adopted bandwidth will simply result in a rescaling of the contrast between the central star and the surrounding disk emission.  The disk image appears darker along the line of nodes and along the line of sight because of the minimal shear broadening that occurs at those disk azimuthal angles and the subsequent reduction in the range of wavelength over which H$\\alpha$ is optically thick.  The dashed line in Figure~14 shows the summed H$\\alpha$ image intensity projected onto the major axis. We see that H$\\alpha$ emission is indeed more extended along the major axis and that the half maximum intensity radii are located at $2.6 R_\\star$ and $7.4 R_\\star$ for the $K^\\prime$ band and H$\\alpha$ summed intensities, respectively.  The ratio of these half maximum intensity radii is 0.35, which is comparable to the observed ratio of 0.55 of the FWHM of the major axes of the $K^\\prime$-band and H$\\alpha$ emission in the Gaussian elliptical fits (Table~7).  Thus, the differences in the radial exponent dependence of density squared and the neutral hydrogen fraction may partially explain the larger appearance of the H$\\alpha$ disks.  \\placefigure{fig14}     %  The CHARA Array observations demonstrate the promise of these kinds of measurements for our understanding of Be stars. Many Be stars were probably spun up by mass transfer in a past close interaction, and future observations over a range in $(u,v)$ and orbital phase coverage will reveal the nature of their companions.   The disks of Be stars are driven by dynamic processes and evolve over timescales of months to years, so long-term observations will help elucidate the physics of disk formation and dispersal.  There is ample evidence that the disks are structured and prone to the formation of spiral arm features \\citep{oka02}.  Furthermore, in Be binaries like $\\phi$~Per, the portion of the disk closest to the hot companion may appear brighter due to local heating by the companion's radiation \\citep{gie98,hum01}.  The techniques of phase closure using multiple telescopes in long baseline interferometry \\citep{kra05} will soon give us the means to study these disk processes directly among the nearby Be stars."
    },
    "0609/astro-ph0609392_arXiv.txt": {
        "abstract": "{} {{It is commonly assumed that the massive winds of AGB stars are dust-driven and pulsation-enhanced. However, detailed frequency-dependent dynamical models which can explain the observed magnitudes of mass loss rates and outflow velocities have been published so far only for C-stars. This letter reports on first results of similar models for oxygen-rich AGB stars. The aim is to provide a better understanding of the wind driving mechanism, the dust condensation sequence, and the role of pulsations.}} {{New dynamical models for dust-driven winds of oxygen-rich AGB stars are presented which include frequency-dependent Monte Carlo radiative transfer by means of a sparse opacity distribution technique and a time-dependent treatment of the nucleation, growth and evaporation of inhomogeneous dust grains composed of a mixture of} Mg$_2$SiO$_4$, SiO$_2$, Al$_2$O$_3$, TiO$_2$, and solid Fe.} {{The frequency-dependent treatment of radiative transfer reveals that the gas is cold close to the star ($700-900\\,$K at $1.5-2\\,R_\\star$) which facilitates the nucleation process. The dust temperatures are strongly material-dependent, with differences as large as 1000\\,K for different pure materials, which has an important influence on the dust formation sequence. Two dust layers are formed in the dynamical models: almost pure glassy Al$_2$O$_3$ close to the star ($r\\!\\ga\\!1.5\\,R_\\star$) and the more opaque Fe-poor Mg-Fe-silicates further out. Solid Fe or Fe-rich silicates are found to be the only condensates that can efficiently absorb the stellar light in the near IR. Consequently, they play a crucial role for the wind driving mechanism and act as thermostat. Only small amounts of Fe can be incorporated into the grains, because otherwise the grains get too hot. Thus, the models reveal almost no mass loss, and no dust shells.}} {The observed dust sequence Al$_2$O$_3$ $\\to$ Fe-poor Mg-Fe-silicates for oxygen-rich AGB stars having low $\\to$ high mass loss rates is in agreement with the presented model and can be understood as follows: Al$_2$O$_3$ is present in the extended atmosphere of the star below the wind acceleration region, also without mass loss. The Mg-Fe-silicates form further out and, therefore, their amount depends on the mass loss rate. The driving mechanism of oxygen-rich AGB stars is still an unsolved puzzle.} ",
        "introduction": "The mass loss mechanism of AGB stars and red supergiants is a long-standing astrophysical problem.  In the carbon-rich case, an extraordinary condensate exists (amorphous carbon) which is very stable, \\ie it can exist already close the star, and is very opaque in the optical and near IR spectral region.  Detailed dynamical models with time-dependent dust formation theory (Winters\\etal 2000)\\nocite{wlj2000} show that the formation of amorphous carbon can provide sufficient radiation pressure to drive massive outflows, consistent with the basic characteristics of C-star winds. This result has been confirmed by dynamical models with frequency-dependent radiative transfer by H{\\\"o}fner\\etal(2003)\\nocite{hga2003}.  However, in the oxygen-rich case, no such condensate exists. The most stable metal-oxides like Al$_2$O$_3$ are too rare. The abundant pure silicates like Mg$_2$SiO$_4$ are already less stable and almost completely transparent around $1\\,\\mu$m where most of the stellar flux escapes.  Solid Fe and Mg-Fe-silicates are opaque but even less stable. Stationary models of dust-driven O-rich AGB star winds with grey radiative transfer (Ferrarotti\\plus Gail 2006) and dynamical models with pulsation and grey radiative transfer (Jeong\\etal 2003) nevertheless came to the conclusion that the winds of O-rich AGB stars are dust-driven, where the stellar pulsation helps to provide the necessary density conditions to form the dust close to the star (``pulsation-enhanced'').  In contrast, the a posteori frequency-dependent radiative transfer analysis of non-linear pulsation models with simplified dust formation theory by Ireland\\plus Scholz (2006) did not find much radiation pressure on dust (Al$_2$O$_3$ and Mg$_{2x}$Fe$_{2-2x}$SiO$_4$) in O-rich Mira variables, with radiative accelerations as small as 0.08 to 0.29 times the gravitational deceleration.  Recent mid-IR observations of O-rich AGB stars in globular clusters with {\\sc Spitzer} (Lebzelter\\etal 2006) and of galactic bulge AGB stars with ISO (Blommaert et al.\\ 2006) show a clear correlation between the kind of condensate and the mass loss rate $\\dot{M}$, called the ``observational dust condensation sequence'': stars with low $\\dot{M}$ show mainly Al$_2$O$_3$, whereas stars with higher $\\dot{M}$ show increasing amounts of Mg-Fe-silicates.  From {\\sc Midi} interferometry of the red supergiant $\\alpha$\\,Ori, Verhoelst et al.\\ (2006)\\nocite{vdmhc2006} concluded that Al$_2$O$_3$ grains are already present at radial distances as small as $r\\!=\\!1.5\\,R_\\star$. ",
        "conclusions": "\\begin{enumerate}  \\item This letter reports on a {\\it negative result}. Detailed dynamical models with frequency-dependent Monte Carlo radiative transfer and time-dependent formation of dirty dust grains cannot explain the observed magnitude of mass loss rates from oxygen-rich AGB stars, even in case of extreme stellar parameters (\\ie high $L_\\star/M_\\star$ ratios).  \\item The {\\it role of solid iron and Fe-rich silicates} is crucial for the wind driving mechanism. These condensates are the only ones that are opaque around $1\\,\\mu$m and, thus, only these condensates can efficiently absorb the stellar light. Since the Fe containing condensates are not particularly stable, they form at too large distances from the star in order to provide an efficient mass loss mechanism.  \\item Previous {\\it grey models} (Jeong\\etal 2003, Ferrarotti\\plus Gail 2006) have calculated the radiation pressure on dust with Rosseland mean opacities which leads to a severe overestimation in the O-rich case, because the mid-IR dust absorption resonances are strongly peaked in a spectral region where the stellar flux is low.  The local flux $F_{\\!\\lambda}$ is not $B_\\lambda(\\Td)$-like, because the dust features are optically thick which makes the radiation field more isotrop and reduce $F_\\lambda$ just where $\\kappa_{\\lambda,\\rm\\,ext}^{\\rm\\;dust}$ is large.   \\item The {\\it dust condensation sequence} is strongly affected by radiative transfer effects. Pure, glassy condensates like Al$_2$O$_3$ have lower dust temperatures than solid Fe or Fe-rich silicates. The differences are as large as 1000\\,K, which favours the formation of the glassy condensates and prevents the formation of Fe-inclusions close to the star. The results in this paper are consistent with the observational finding of Al$_2$O$_3$ at radial distances as small as $1.5\\,R_\\star$ around $\\alpha$\\,Ori (Verhoelst\\etal 2006) as well as with the observed dust condensation sequence in O-rich AGB stars (Blommaert\\etal 2006, Lebzelter\\etal 2006), because Al$_2$O$_3$ can already exist in an extended atmosphere without mass loss, whereas the Mg-Fe-silicates form in the more distant wind regions which require mass loss.  \\item The {\\it mass loss mechanism} of oxygen-rich AGB stars and red supergiants is still a puzzle. Pulsations alone cannot drive an outflow because the radiative cooling of the gas is too efficient, even in non-LTE (Woitke 2003, Schirrmacher  et al.\\ 2003)\\nocite{woi2003,sws2003}. According to the results of this letter, even a combination of stellar pulsation and radiation pressure on dust is insufficient to drive the mass loss.  Do we have to re-visit Alfv{\\'e}n-waves (\\eg Vidotto\\plus Jatenco-Pereira 2006)? \\nocite{vp2006}  \\end{enumerate}"
    },
    "0609/astro-ph0609671_arXiv.txt": {
        "abstract": "We carefully consider observational and theoretical constraints on the global properties of secondary stars in cataclysmic variable stars (CVs). We then use these constraints to construct and test a complete, semi-empirical donor sequence for CVs with orbital periods $P_{orb} \\leq 6$~hrs. All key physical and photometric parameters of CV secondaries (along with their spectral types) are given as a function of $P_{orb}$ along this sequence. This provides a benchmark for observational and theoretical studies of CV donors and evolution.  The main observational basis for our donor sequence is an empirical mass-radius relationship for CV secondaries. Patterson and co-workers have recently shown that this can be derived from superhumping and/or eclipsing CVs. We independently revisit all of the key steps in this derivation, including the calibration of the period excess-mass ratio relation for superhumpers and the use of a single representative primary mass for most CVs. We also present an optimal techique for estimating the parameters of the mass-radius relation that simultaneously ensures consistency with the observed locations of the period gap and the period minimum. We present new determinations of these periods, finding $P_{gap,+} = 3.18 \\pm 0.04$~hrs (upper edge), $P_{gap,-} = 2.15 \\pm 0.03$~hrs (lower edge) and $P_{min} = 76.2 \\pm 1.0$~min (period minimum).  We test the donor sequence by comparing observed and predicted spectral types ($SpTs$) as a function of orbital period. To this end, we update the $SpT$ compilation of Beuermann and  co-workers and show explicitly that CV donors have later $SpTs$ than main sequence (MS) stars at all orbital periods. This extends the conclusion of the earlier study to the short-period regime ($P_{orb} < 3$~hrs). We then compare our donor sequence to the CV data, and find that it does an excellent job of matching the observed $SpTs$. Thus the empirical mass-radius relation yields just the right amount of radius expansion to account for the later-than-MS spectral types of CV donors. There is remarkably little intrinsic scatter in both the mass-radius and $SpT-P_{orb}$  relations, which confirms that most CVs follow a unique evolution track.  The donor sequence exhibits a fairly sharp drop in temperature, luminosity, and optical/infrared flux well before the minimum period. This may help to explain why the detection of brown dwarf secondaries in CVs has proven to be extremely difficult.  We finally apply the donor sequence to the problem of distance estimation. Based on a sample of 22 CVs with trigonometric parallaxes and reliable 2MASS data, we show that the donor sequence correctly traces the envelope of the observed $M_{JHK}-P_{orb}$ distribution. Thus robust lower limits on distances can be obtained from single-epoch infrared observations. However, for our sample, these limits are typically about a factor of two below the true distances. ",
        "introduction": "\\label{intro}  Cataclysmic variable stars (CVs) are compact, interacting binary systems in which a white dwarf primary accretes from a low-mass, roughly main-sequence donor star. The mass transfer and secular evolution of CVs is driven by angular momentum losses. In systems with long orbital periods ($P_{orb} \\gtappeq 3$~hrs), the dominant angular momentum loss mechanism is thought to be magnetic braking (MB) due to a stellar wind from the donor star. In the canonical ``disrupted magnetic braking'' evolution scenario for CVs, MB stops when the secondary becomes fully convective, at $P_{orb} \\simeq 3$~hrs. At this point, the donor detaches from the Roche lobe, and gravitational radiation (GR) becomes the only remaining angular momentum loss mechanism. The GR-driven shrinkage of the orbit ultimately brings the secondary back into contact at $P_{orb} \\simeq 2$~hrs, at which point mass transfer resumes. The motivation for this scenario is the dearth of mass-transferring CVs in the so-called period gap between $P_{orb} \\simeq 2$~hrs and $P_{orb} \\simeq 3$~hrs.  As a CV evolves, its donor star is continuously losing mass. As long as the mass-loss time scale ($\\tau_{\\dot{M}_2} \\sim M_2/\\dot{M}_2$) is much longer than the donor's thermal time  scale ($\\tau_{th} \\sim \\frac{GM^2_2}{R_2 L_2}$), the secondary is able to maintain thermal equilibrium and should closely follow the standard main sequence track defined by single stars. Here and throughout, we use $M_2$, $R_2$ and $L_2$ to denote a donor's mass, radius and luminosity, and $\\dot{M}_2$ to denote the rate at which it is losing mass to the primary. If the condition $\\tau_{\\dot{M}_2} >> \\tau_{\\dot{M}_{th}}$ is not met, the donor will be driven out of thermal equilibrium and become oversized compared to an isolated main sequence (MS) star of identical mass.  So how do thermal and mass-loss time scales compare for the donor stars in CVs? Or, to put it another way, are the donors main sequence stars? This question has been addressed twice in recent years, in very different ways. First, Beuermann et al. (1998; hereafter B98) showed that, at least for periods $P_{orb} \\gtappeq 3$~hrs (i.e. above the period gap), the spectral types ($SpT$s) of CV donor stars are significantly later than those of isolated MS stars. \\footnote{For the record, earlier studies along these lines were carried out by  Echevarr\\'{i}a (1983), Patterson (1984), Friend et al (1990ab) and Smith \\& Dhillon (1998).} At the very longest periods ($P_{orb} \\gtappeq 5-6$~hrs), this is probably due to the donors being somewhat evolved (B98; also see Podsiadlowski, Han \\& Rappaport 2003). However, for all other systems, the late $SpT$s of the donors are simply a sign of their losing battle to maintain thermal equilibrium.  Second, Patterson et al. (2005; hereafter P05) produced an empirical mass-radius sequence for CV donors based on a sample of masses and radii derived mainly from superhumping CVs. One of their key results was that CV  donors are indeed oversized compared to MS stars, for all periods $P_{orb} \\ltappeq 6$~hrs. They also detected a clear discontinuity in $R_2$ for systems with similar $M_2$ above and below the period gap. This is just what is expected in the disrupted magnetic braking picture: the MB-driven systems just above the gap should be losing mass faster than than the GR-driven systems just below. They should therefore be further out of thermal equilibrium.  In the present paper, we will build on these important studies by constructing a complete, semi-empirical donor sequence for CVs. More specifically, our goal is to derive a benchmark evolution track for CV secondaries that incorporates all of the best existing observational and theoretical constraints on the global donor properties. \\footnote{Here and below, we use the term ``donor sequence'' to describe the dependence of the secondary star parameters on the orbital period of a CV. We call a sequence ``complete'' if all important physical ($P_{orb}$, $M_2$,  $R_2$, $T_{eff}$, $SpT$) and photometric (absolute magnitudes in UBVRIJHKLM) donor properties are fully specified along the entire evolutionary track.} Along the way, we will update both of the earlier studies and show that they are mutually consistent: the mass-radius relation obtained by P05 is just what is needed to account for the late $SpT$s observed by B98.  Our donor sequence should be useful for many practical applications. Perhaps most fundamentally, it provides a benchmark for what we mean by a ``normal'' CV and can thus be used to test CV evolution scenarios that predict the properties of CV donors. The sequence can also be used to simplify and improve on ``Bailey's method'' for estimating the distances to CVs (Bailey 1981). The new version of the method requires only knowledge of the orbital period and an infrared magnitude measurement. Conversely, for systems with known distances, the sequence can be used to estimate the donor contribution to the system's flux in any desired bandpass. Finally, the present work is also a stepping stone towards another goal: the construction of new, semi-empirical mass-transfer rate and angular momentum loss laws for CVs. This should be possible, since the degree of departure of a donor from the MS track is a direct measure of the mass loss from it and hence of the angular momentum loss that drives $\\dot{M}_2$. ",
        "conclusions": "The main goal of this paper has been the construction of a complete, semi-empirical donor sequence for CVs that is based on our best understanding of donor properties and is consistent with all key observational constraints. This donor sequence is provided in Table~\\ref{tab:donorseq} and is based primarily on an updated version of the mass-radius relationship for CV secondaries determined by P05. It also relies on a MS-based mass-effective temperature relation that has been shown to be appropriate for CV donors by theoretical work, and on up-to-date stellar atmosphere models. By design, the donor track also reproduces the observed locations of the period gap and the period minimum, and extends beyond this to the period bouncer regime. We have shown that this sequence provides an excellent match to the observed $SpT$-$P_{orb}$ relation for unevolved CV donors (i.e. $P_{orb} \\ltappeq 6$~hrs), and that it correctly traces the lower envelope of the $M_{JHK}-P_{orb}$ distributions of CVs.  Along the way, we have revisited and updated two important results on CV donors by P05 and B98. Regarding the former study, we have carried out a full, independent analysis of the superhumper and eclipser data that was used by P05 to construct their empirical $M_2-R_2$ relation. We constructed and used our own $q-\\epsilon$ calibration for superhumpers, verified the validity of the constant $M_1$ assumption on which the method rests, imposed the locations of the period gap and the period minimum as constraints on the mass-radius relation and used a self-consistent fitting technique to extract optimal parameters for the $M_2-R_2$ relations for long-period CVs, short-period CVs and period-bouncers.  Regarding the study by B98, we have updated their $SpT$ data base for CV secondaries. Our new sample contains 91 CV donors with spectroscopically-determined $SpT$ estimates (compared to 54 donors in the B98 sample). Comparing the $SpT$-$P_{orb}$ distribution to that of MS stars, we find that CV donors have later spectral types at all orbital periods. This extends the finding of B98 to the short-period regime ($P_{orb} \\ltappeq 3$~hrs).  We find that there is remarkably little intrinsic scatter about both the mean $M_2-R_2$ and $SpT-P_{orb}$ relations for CVs with $P_{orb} < 6$~hrs: probably no more than a few percent in $R_2$ and less than 1 spectral sub-type. This suggests that, on long time-scales, most CVs do indeed follow a unique evolution track, just as is theoretically expected (e.g. Kolb 1993; Stehle, Ritter \\& Kolb 1996). The large scatter in luminosity-based accretion rate estimates at fixed $P_{orb}$ (e.g. Patterson 1984) is therefore probably caused by fluctuations around the mean mass transfer rate on time-scales that are short compared to the time-scale on which the donor loses mass and the binary evolves (see B{\\\"u}ning \\& Ritter 2004 and references therein). In principle, the empirical $M_2-R_2$ relationship for CV donors can be inverted to infer the long-term average mass-transfer rate as a function of orbital period; this will be the subject of a separate paper.  An important feature of our donor sequence is a sharp decline in effective temperature, luminosity and optical/IR brightness, well before the period minimum is reached. In fact, the L/T transition essentially coincides with $P_{min}$ on the evolution track (although this statement depends somewhat on the validity of our spectral type calibration in a sparsely populated regime). Beyond $P_{min}$, even the accretion-heated WD alone is expected to outshine the donor at optical wavelengths and it is comparably bright to the donor even in the infrared. All of this helps to explain why brown dwarf donors have proven so difficult to detect in CVs, especially in suspected period bouncers.  Finally, we have looked at one obvious application of our donor sequence, the determination of distance estimates from infrared photometry. We have shown that robust lower limits can indeed be obtained, but that, in our calibration sample, these are typically a factor of two smaller than the true distances. Thus, even in the infrared, the donor typically contributes only $\\simeq$30\\% of the total flux. Distance estimates that explicitly allow for this average donor contribution are typically good to about a factor of 2 in our calibration sample."
    },
    "0609/astro-ph0609447_arXiv.txt": {
        "abstract": "{% In this contribution, I briefly review recent progress in detecting and measuring the properties of relativistic iron lines observed in stellar-mass black hole systems, and the aspects of these lines that are most relevant to studies of similar lines in Seyfert-1 AGN.  In particular, the lines observed in stellar-mass black holes are not complicated by complex low-energy absorption or partial-covering of the central engine, and strong lines are largely independent of the model used to fit the underlying broad-band continuum flux.  Indeed, relativistic iron lines are the most robust diagnostic of black hole spin that is presently available to observers, with specific advantages over the systematics--plagued disk continuum.  If accretion onto stellar-mass black holes simply scales with mass, then the widespread nature of lines in stellar-mass black holes may indicate that lines should be common in Seyfert-1 AGN, though perhaps harder to detect.} ",
        "introduction": "Broad Fe~K emission lines plausibly arising in the inner disk around black holes, and shaped by the strong Doppler shifts and gravitational redshifts expected there, were first reported in stellar-mass black holes (Barr, White, \\& Page 1985; van der Woerd, White, \\& Kahn 1989). Though these detections were achieved with low--resolution gas spectrometers, their discovery was met with enthusiasm, and theoretical papers showing their reported widths to be consistent with extreme dynamical broadening soon followed (Fabian et al.\\ 1989, Laor 1991).  It was also quickly realized that Fe~K emission lines are merely the most prominent feature of the overall response of an accretion disk to external irradiation by hard X-ray flux (George \\& Fabian 1991).  Although strong Fe~K lines in Seyfert AGN were well-known at the time broad lines were found in stellar-mass black holes, the lower flux in AGN lines made it difficult to measure their shape.  In the {\\it ASCA} era, however, studies of broad iron lines in Seyfert-1 AGN advanced rapidly.  The SIS (solid-state imaging spectrometer, an X-ray CCD spectrometer) aboard {\\it ASCA} revealed that the lines in some Seyfert-1 AGN were not only broad, but asymmetric --- consistent with strong Doppler shifts and gravitational redshifts expected in the inner accretion disk (see, e.g. Tanaka et al.\\ 1995; Iwasawa et al.\\ 1996; Nandra et al.\\ 1997a, Nandra et al.\\ 1996b).  Although the SIS was well-suited to studies of the line profile in Seyfert-1 AGN, ruinous photon pile-up resulted when stellar-mass black holes were observed, greatly complicating efforts to recover robust spectra for detailed line measurements (see, e.g., Ebisawa et al.\\ 1996).  In the {\\it XMM-Newton} and {\\it Chandra} era, considerable effort has again been devoted to the study of relativistic iron lines.  While relativistic lines appear to be confirmed in some cases, complex low-energy absorption or alternative continuum spectral models initially cast some doubt on relativistic lines in other cases (see, however, the contributions by Fabian et al.\\, Nandra et al.\\, and Reeves et al.\\ in these proceedings).  In the stellar-mass black hole regime, the fast read-out modes and advanced spectrometers aboard {\\it XMM-Newton} and {\\it Chandra} have made it possible to measure the properties of broad iron lines without photon pile-up.  A number of asymmetric lines have been revealed, showing that they are also shaped by relativistic blurring at the inner accretion disk.  In this sense, {\\it XMM-Newton} and {\\it Chandra} have done for stellar-mass black holes what {\\it ASCA} did for Seyfert-1 AGN.  In Section 2, recent observations of relativistic Fe~K lines in stellar-mass black holes are reviewed.  In Section 3, the robustness of relativistic Fe~K lines and their utility as spin diagnostics is critically examined.  In Section 4, different methods of constraining black hole spin are discussed.  Finally, the relevance of relativistic lines in stellar-mass black holes to studies of Seyfert-1 AGN and objectives for the future studies of relativistic lines are discussed in Section 5. ",
        "conclusions": ""
    },
    "0609/astro-ph0609321_arXiv.txt": {
        "abstract": "The ASTRA Cassegrain Spectrophotometer and its automated 0.5-m telescope at Fairborn Observatory in Arizona will produce a large quantity of high-precision stellar flux distributions. A separate paper \\citep{ADE+07} presented a review of the design criteria for the system and an overview of its operation. This paper discusses the techniques used in the data reduction to final flux calibrations.  Extraction of 1-d spectra from the 2-d images will be performed by a highly automated version of CCDSPEC \\citep{GH02}. The characteristics of the CCD are automatically applied to the images, including the location of dead rows and hot pixels. In order to achieve the goal of better than 1\\% precision, large numbers of bias and flat field frames will be used in the reduction process. There will be a continual programme to monitor the image quality. Finally, optimally extracted spectra will be obtained, including the removal of scattered light and cosmic rays.  The Earth's atmosphere has a considerable effect on the stellar flux as measured from the surface. The principal sources of extinction, Rayleigh and aerosol scattering, ozone and telluric line absorption, are discussed, along with methods used to determine their effects on the observed spectra. Correction for telluric lines is the most problematical, due to their non-linear variation with airmass. By using a large network of constant stars to monitor atmospheric extinction it is possible to determine the extinction coefficients to generally better than 1\\% and to assess their temporal variability.  The spectrophotometric observations are placed on an absolute flux scale by reference to stars with known values of true flux at top of Earth's atmosphere. These standard stars have been calibrated against terrestrial sources of known properties. Unfortunately, very few stars have been calibrated directly. The ASTRA fluxes will be calibrated against the best available Vega flux distribution. The constant stars used in the extinction determinations will provide the internal calibration network of secondary flux standards.  The available absolute calibrations are accurate to typically 1--2\\%. Ultimately this uncertainty will limit the accuracy of the final fluxes of other stars. However, the internal precision will be significantly higher, and should more-accurate absolute calibrations become available the fluxes can be re-calibrated to higher accuracy. ",
        "introduction": "The ASTRA Cassegrain Spectrophotometer and its automated 0.5-m telescope at Fairborn Observatory in Arizona will produce a large quantity of high-precision stellar flux distributions. A separate paper, \\citep{ADE+07}, presented a review the design criteria for the system and an overview of its operation. In this paper we discuss the techniques used in the data reduction to final flux calibrations.  ASTRA will produce a vast quantity of high quality data. The reduction procedures will be automated as far as possible, using software sentinels to watch for unusual events. ",
        "conclusions": "The reduction of ASTRA spectrophotometry requires careful extraction and calibration in order to achieve stellar flux measurements with internal (star-to-star) precision better than 1\\% for stars brighter than $\\sim$10th magnitude. In the regions heavily affected by telluric lines we will not necessarily reach the 1\\% precision level. For all observations, including faint spectrophotometric standards, we will use full error propagation to include both internal and external errors.  The available absolute calibrations are accurate to typically 1--2\\%. Ultimately this uncertainty will limit the accuracy of the final fluxes of other stars. However, the internal precision will be significantly higher, and should more-accurate absolute calibrations become available the fluxes can be re-calibrated to higher accuracy."
    },
    "0609/astro-ph0609117_arXiv.txt": {
        "abstract": "\\noindent The possible existence of additional long-period planetary-mass objects in the extrasolar planetary systems 47~UMa and 14~Her is investigated.  We combine all available radial-velocity data on these stars, spanning up to 18 years.  For the 47~UMa system, we show that while a second planet improves the fit to all available data, there is still substantial ambiguity as to the orbital parameters of the proposed planetary companion 47~UMa~c.  We also present new observations which clearly support a long-period companion in the 14~Her system.  With a period of 6906$\\pm$70 days, 14~Her~c may be in a 4:1 resonance with the inner planet.  We also present revised orbital solutions for 7 previously known planets incorporating recent additional data obtained with the 2.7m Harlan J.~Smith Telescope at McDonald Observatory. ",
        "introduction": "The longest-running radial-velocity surveys are now approaching time baselines of 16-18 years \\citep{butleretal96, cochran97}.  These surveys now achieve internal measurement precisions such that the signals from long-period giant planets ($P\\gtsimeq 10$yr) are now entering the realm of detectability \\citep{limitspaper}.  For example, the High Resolution Spectrograph (HRS) on the Hobby-Eberly Telescope (HET) now has a velocity precision of 3-4 \\ms \\citep{cochran04}, the Keck HIRES is achieving 1-2 \\ms\\ since its 2004 CCD upgrade \\citep{butler06}, and the HARPS instrument has demonstrated precision better than 1 \\ms \\citep{lovis06}.  Of particular interest are putative multi-planet systems, wherein the residuals of a known-planet's orbit show Keplerian periodicity indicative of a distant giant planet companion in the system.  Finding systems which contain long-period giant planets as well as planets in closer orbits will address important questions about the uniqueness of our own Solar System's architecture.  Long-period planet candidates have been reported for 47~UMa ($P=7.1$ yr, Fischer et al.~2002), 55~Cancri ($P=12.4$ yr, McArthur et al.~2004), HD~217107 ($P>7$ yr, Vogt et al.~2005), and HD~72659 ($P=9.9$ yr, Butler et al.~2006).  Of particular interest are systems which contain both short-period and long-period jovian planets.  Such systems could provide clues to address the question of how the processes of planet formation and migration can result in both ``hot'' and ``cold'' Jupiters in the same planetary system.  In this paper, we present improved fits to the known planets in nine systems using all available data, and we investigate the possibility of additional long-period objects.  In \\S 2, we briefly describe the data acquisition and fitting procedure.  The results are given in \\S 3, including our solutions for seven additional planetary systems in the McDonald Observatory planet search program.  All of our radial-velocity measurements for these objects are presented in Tables 5-14.  In this paper, we use the terms ``McDonald'' to refer to data taken with the McDonald Observatory 2.7~m Harlan J.~Smith (HJS) Telescope, and ``HET'' to refer to data taken with the 9.2~m Hobby-Eberly Telescope. ",
        "conclusions": "The Keplerian orbital fits are shown in Figures~1-4, and the orbital parameters implied by those fits are given in Table 2.  In computing the planetary minimum mass M~sin $i$ and semimajor axis $a$, the stellar masses derived by \\citet{santos04} listed in Table 1 were used, with adopted uncertainties of 0.05 \\Msun.  \\subsection{47 UMa (=HD~95128): Ambiguities concerning the second planet}  \\citet{butler96} first reported the 1090-day companion to 47~UMa using data from Lick Observatory.  With additional velocity measurements over 13 years, \\citet{fischer02} announced a long-period second planet, 47~UMa~c, with a period of 2594$\\pm$90 days and a mass of 0.76 \\Mjup. \\citet{naef04} presented ELODIE observations of 47~UMa, and noted that the second planet was not evident in their data, which were fit well with a single Keplerian model.  We now fit four data sets for 47~UMa: Lick (Fischer et al~2002, N=91), ELODIE (Naef et al.~2004, N=44), 2.7m HJS telescope (N=35), and the Hobby-Eberly Telescope (HET, N=201).  The HET data, which consist of multiple exposures per visit, were binned using the weighted mean value of the velocities in each visit.  The quadrature sum of the rms about the mean and the mean internal error bar was adopted as the error bar of each binned point (N=63).  The offset between the overlapping HJS and HET datasets was used to merge them into one, which was used in all fits. The one-planet fit (Model~1) and the residuals to that fit are shown in Figure~1.  We emphasize that this fit includes all available published data, over a time span of more than 18 years, and includes 195 high-precision measurements obtained with the HET High-Resolution Spectrograph at 61 independent epochs, which are given in Table 6.  The total rms about the combined one-planet fit is 10.4 \\ms.  The rms scatter about the one-planet fit for each of the four datasets is:  Lick--10.7 \\ms, ELODIE--13.0 \\ms, HJS--13.5 \\ms, HET--4.9 \\ms.  A periodogram of the residuals of all of the data to the 1-planet fit is shown in Figure 2. No clear peak rises above the noise level at any period between 2 and 10000 days; the total duration of the available data is now about 6900 days (18.3 years).  While a peak is present at about 2212 days, close to the period reported for 47~UMa~c by \\citet{fischer02}, its false-alarm probability (FAP) is 27.2\\%.  To further explore the possible presence of 47~UMa~c, we fit all of the datasets with a two-planet model fixed at the parameters of \\citet{fischer02} (Model~2), and then repeated the fit allowing all parameters to be free except for $e$ and $\\omega$ of the second planet (Model 3), which were fixed at 0.005 and 127\\degrees, respectively, after \\citet{fischer02}.  No models achieved convergence with those two parameters free.  The fit obtained by Model 3 is shown in the right panel of Figure~2.  Model 2 had a reduced chi-square ($\\chi^{2}_{\\nu}$) of 5.81 and an rms of 12.7 \\ms\\ about the fit, whereas Model 3 had a $\\chi^{2}_{\\nu}$ of 2.11 and an rms of 8.4 \\ms.  For comparison, the one-planet fit (Model 1) had a $\\chi^{2}_{\\nu}$ of 3.23 and an rms of 10.4 \\ms.  Noting that the poor fit of Model 2 was largely due to errors in the period of the inner planet, we re-did the fits allowing the parameters for the inner planet to be free while fixing those of the outer planet at the values reported by \\citet{fischer02} (Model 4).  The $\\chi^{2}_{\\nu}$ of this fit was 2.68, with an rms of 10.2 \\ms.  These tests are summarized in Table 3.  The free two-planet (Model 3) fit was the best of the four, in terms of both the goodness-of-fit criterion ($\\chi^{2}_{\\nu}$) and the total rms scatter about the fits.  The rms about the individual datasets for this fit is the following: Lick--8.0 \\ms, ELODIE--11.1 \\ms, HJS--10.1 \\ms, HET--5.3 \\ms.  The parameters for the 47~UMa planetary system given in Table~2 are those obtained by Model~3.  Although the best-fit set of parameters obtained a period of 7586 days for the outer planet, there is no corresponding peak on the periodogram shown in Figure~2.  However, the periodogram method is not as reliable when the periodic signal approaches or exceeds the total duration of observations, as is the case here, where the total time baseline is 6942 days.  We note that we have been able to reproduce the result of \\citet{fischer02} by this method; i.e.~a periodogram analysis of the Lick data alone after removing 47~UMa~b revealed a strong peak at 2083 days, but still with a FAP of 0.15\\%.  We are also able to recover the \\citet{fischer02} parameters of 47~UMa~c from the Lick data alone.  Since the total time coverage of the Lick data presented in \\citet{fischer02} is 5114 days (14~yr), an analysis of those data alone is not adequate to fully constrain the 7586-day period obtained by our best-fit model which adds four years to the total duration of observations.  It is possible that the shorter period for 47~UMa~c reported by \\citet{fischer02} is an alias of the true period; at present, our fits indicate that period to be about three times longer, but with substantial uncertainty.  To further test the methods by which we conclude that the parameters of 47~UMa~c reported by \\citet{fischer02} are dubious, we performed some Monte Carlo simulations.  From each of the three data sets considered in the fits described above, we generated 1000 simulated sets of velocities consisting of a Keplerian signal plus 7 \\ms\\ of Gaussian noise for each of the two planets.  The parameters for the inner planet were those listed in Table 2, and those of the outer planet were those of \\citet{fischer02}.  These simulated datasets retained the times of observation and the error bars of the originals.  The simulated data were then fit with a one-planet model exactly as described above, then the residuals of the one-planet fit were examined by the periodogram method, to determine whether the signal of the second planet was recovered.  The criteria for recovery were that the period of the second planet had to be detected correctly and with a FAP of less than 0.1\\%.  This FAP was computed using the analytic FAP formula of \\citet{hb86}.  Of the 1000 trials, only 6 did not result in a successful recovery of the signal of the second planet.  The correct period was recovered 999 times, and the FAP exceeded 0.1\\% only 5 times; the worst FAP was 0.24\\%.  For comparison, the analytic FAP of the 2212-day peak in the residuals of the 1-planet fit is 1.7\\%, a factor of 7 higher.  These results indicate that our method should have been able to detect the signal of 47~UMa~c, had it been present with the parameters given by \\citet{fischer02} and \\citet{butler06}.  We conclude that while an additional long-period object may be present, the data currently available do not provide sufficient evidence for an orbital solution.   \\subsection{14 Her (=HD~145675): Evidence for an outer companion}  Thirty-five radial-velocity measurements of 14~Her (=HD~145675)  obtained at McDonald Observatory were combined with published data from Keck HIRES \\citep{butler06} and ELODIE \\citep{naef04}.  The fit to the combination of these three data sets is shown in Figure~3, and the system parameters implied by that fit are given in Table 2.  It is evident from Figure~3 that the single Keplerian fit is inadequate (rms=13.0 \\ms); indeed, the residuals to the fit show a clear curvature.  \\citet{naef04} noted an upward linear velocity trend of 3.6$\\pm$0.3 m s$^{-1}$ yr$^{-1}$ in their observations which covered the time interval 1994 to 2003, but the Keck data of \\citet{butler03} indicated no such trend.  When the complete 12~yr duration of observations is examined, we see that those Keck data were serendipitously obtained during a period of relatively constant residual velocity, and the more recent McDonald data are now indicating a downward trend.  Fitting a double-Keplerian model (Figure~4) with $T_0$ for the outer body fixed at 2449100.0 yields an eccentricity consistent with zero: $e=0.02\\pm 0.06$.  Fixing the eccentricity of 14~Her~c at 0.0 gives a minimal orbit solution with period $P=6906\\pm 70$ days, and a velocity semi-amplitude $K=24.5\\pm 1.4$ \\ms. These orbital elements imply a minimum mass M~sin~$i$=2.1 \\Mjup\\ at a semimajor axis of 6.9~AU.  This represents the lowest possible mass and the shortest period for the outer companion.  The rms scatter about the three data sets is as follows: ELODIE -- 10.4 \\ms, Keck -- 2.9 \\ms, McDonald -- 5.8 \\ms.  The total rms about this two-planet fit is 8.4 \\ms, with a $\\chi^2_{\\nu}$ of 1.67.  The minimal orbit proposed above for 14~Her~c has a periastron distance of 6.2~AU, and 14~Her~b has an apastron distance of 3.76~AU.  The orbit of 14~Her~c, at a large semimajor axis with a small eccentricity, is similar to that of HD~72659b \\citep{butler03} and HD~50499b \\citep{vogt05}.  We emphasize that such fits are preliminary, and are only given in order to place a lower limit on the period and mass of this object.  \\citet{polishguys} performed dynamical simulations of the 14~Her system, and proposed the existence of 14~Her~c based on fits to data from \\citet{butler03} and \\citet{naef04}. The two most-favored models had 14~Her~c near the 3:1 or 6:1 mean-motion resonance (MMR), which coincided with highly stable orbital configurations as determined by their dynamical simulations.  The preliminary fit obtained in this work suggests a 4:1 resonance for 14~Her~c.  We attempted to fit a double-Keplerian model with 14~Her~c fixed at the 3:1 MMR parameters indicated in \\citet{polishguys}, but obtained a substantially worse fit ($\\chi^2_{\\nu}=2.2$).  An upper limit on the mass of 14~Her~c can be estimated from the results of adaptive-optics (AO) imaging by \\citet{luhman02}, who used the Keck~II AO system to search for companions to 25 extrasolar planet-host stars. No candidate companions were found around 14~Her, and the detection limits derived from their study exclude stellar-mass objects objects at orbital separations $\\gtsimeq$ 9~AU (see Luhman \\& Jayawardhana 2002, their Fig.~10).  A similar study by \\citet{patience02} using the Lick AO system also excluded stellar-mass objects beyond about 12.7~AU.  We can therefore use the upper limit provided by \\citet{luhman02} and our lower bound to constrain the mass of the outer companion between about 2.1 and 80 \\Mjup.  A more definitive statement on its nature, of course, requires many more years of observations, until a second velocity turnaround is confirmed.  The large separation implied by these data make 14~Her an attractive target for future direct imaging attempts and for astrometric follow-up.  The astrometric perturbation due to 14~Her~c would be $\\alpha =$0.9~mas, using the minimal orbit solution given above.  For sin~$i$=0.5, this signal would be $\\alpha =$1.8~mas, equivalent to that of $\\epsilon$~Eri~b \\citep{fritz}.  Since the astrometric perturbation increases with semimajor axis and planet mass, the signal could be much larger \\citep{sozzetti05}.   \\subsection{Revised orbital parameters for 7 additional systems}  The addition of new data from the HJS Telescope presented in this paper, and the use of multiple independent data sets in fitting Keplerian orbits, have generally improved the precision of the derived planetary parameters by 25-50\\%.  In particular, the precision of the orbital periods have been improved by the addition of new data, due to the increased number of orbits now observed.  In this section, we briefly describe the results of our combined fits for seven additional systems for which the revised orbital solutions are in agreement with previously published results. Our parameters are within 2$\\sigma$ of previous values except for $T_0$ and $\\omega$ of $\\upsilon$~And~c, and $P$ and $\\omega$ of $\\upsilon$~And~d, which differ from \\citet{butler06} by between 2.7 and 3.8 sigma.  The inner planet of HD~217107, with a period of 7.1 days, was first reported in \\citet{fischer99}, and additional CORALIE data supporting this discovery were published by \\citet{naef01}.  A shallow parabolic trend in the residuals was noted by \\citet{fischer01}.  \\citet{vogt05} used more recent Keck data to postulate an outer companion with $P>7$ years. However, this object has not completed a full orbit, and hence there is a wide range of possible solutions.  \\citet{wright05} revised the period of HD~217107c to 20,140 days (55 years), and \\citet{vogt05} give a period of 3150 days; our fits with the short period improved $\\chi^2_{\\nu}$ by only 0.07 over the long period.  We fit four datasets for HD~217107: Lick \\citep{fischer99}, CORALIE \\citep{naef01}, Keck \\citep{vogt05}, and 20 observations from McDonald.  We fit a double-Keplerian model using the parameters for planet~c from \\citet{vogt05} as a starting point for the least-squares fitting procedure.  The rms of this double-Keplerian fit is 9.0 \\ms, and the results are given in Table 2.  The resulting parameters are in agreement with \\citet{vogt05} to within 2$\\sigma$.  \\citet{vogt05} performed separate fits to the Lick and Keck data and commented on the high degree of uncertainty in the parameters of HD~217107c.  The combined fits given in this work support the $\\sim$3150-day period reported by \\citet{vogt05}, but the uncertainties remain large.  At present, no definitive statements can be made until this object completes a significant fraction of an orbit, or until such time as it can be detected via direct imaging or astrometry, techniques which are well-suited to very long-period objects.  The three planets around $\\upsilon$ And (=HD~9826) were fit with a triple-Keplerian model, combining the Lick data of \\citet{butler06} and the Advanced Fiber Optic Echelle (AFOE) data of \\citet{butler99} with 41 observations from McDonald.  Lick data preceding the Hamilton spectrometer upgrade (1995 February) were excluded following \\citet{butler97}.  The total rms about the fit is 14.6 \\ms.  As shown in Table 4, analysis of the residuals to a triple-Keplerian fit resulted in only a marginally significant periodicity at 2000 days (FAP=0.4\\%).  \\citet{marcy99} first detected the inner companion to HD~168443 using the HIRES spectrograph on Keck~I.  The outer 4.8-yr planet was then reported in \\citet{marcy01} and confirmed by \\citet{udry02}.  A two-planet fit to the McDonald data combined with that of \\citet{butler06} and \\citet{udry02} yields an rms of 9.5 \\ms.  The parameters given in Table~2 are in agreement with those of \\citet{butler06} within 2$\\sigma$. Periodogram analysis of the residuals to the double-Keplerian fit revealed a 63-day periodicity with a FAP of 0.12\\% (Table 4).  Performing the same analysis on the three datasets separately, however, showed no enhanced power at that period.  Additionally, the period of the inner planet is 58 days, preventing any object in a 63-day orbit.  The following five planets showed no residual periodicities of interest. The hot Jupiter orbiting HD~179949 was found by the Anglo-Australian Planet Search program using the 3.9m Anglo-Australian Telescope (AAT) \\citep{tinney01}.  The rms about the combined fit is 11.5 \\ms, and the fitted parameters agree with those of \\citet{butler06} to within 1$\\sigma$.  For 16~Cyg~B (=HD~186427), data from the discovery paper \\citep{cochran97} and Lick data from \\citet{butler06} were combined with 37 additional measurements from McDonald Observatory.  The rms about the combined fit is 10.6 \\ms, and the orbital elements are within 2$\\sigma$ of those given in \\citet{butler06}.  For HD 195019b, Lick data from \\citet{butler06} were combined with 19 McDonald observations.  The rms about the combined fit is 15.8 \\ms, and the parameters agree with those of \\citet{butler06} within 2$\\sigma$.  For HD 210277b, data from \\citet{butler06} and \\citet{naef01} were combined with 21 measurements from McDonald.  The rms about the combined fit is 6.8 \\ms, and the planetary parameters agree with those of \\citet{butler06} to within 2$\\sigma$."
    },
    "0609/astro-ph0609784_arXiv.txt": {
        "abstract": "It has long been realized that dark matter halos formed in cosmological N-body simulations are characterized by density profiles $\\rho(r)$ that, when suitably scaled, have similar shapes. Additionally, combining the density and velocity dispersion profiles $\\sigma(r)$, each of which have decidedly nonpower-law shapes, leads to quantity \\psd\\ that is a power-law in radius over 3 orders of magnitude in radius.  Halos' velocity anisotropy profiles $\\beta(r)$ vary from isotropic near the centers of halos to quite radially anisotropic near the virial radius.  Finally, there appears to be a nearly linear correlation between $\\beta$ and the logarithmic density slope $\\gamma$ for a wide variety of halos.  This work is part of a continuing investigation of the above interrelationships and their origins using analytical and semi-analytical techniques.  Our finding suggest that the nearly linear $\\beta$--$\\gamma$ relationship is not just another expression of scale-free \\psd\\ behavior.  We also note that simultaneously reproducing density and anisotropy profiles like those found in simulations requires $\\beta(r)$ and $\\gamma(r)$ to have similar shapes, leading to nearly linear $\\beta$--$\\gamma$ correlations.  This work suggests that the $\\beta$--$\\gamma$ and power-law \\psd\\ relations have distinct physical origins. ",
        "introduction": "Cosmological N-body simulations are important tools that allow us to reconstruct how gravitationally bound objects like galaxies and clusters form and evolve in the universe.  The base-level simulations involve Newtonian gravitational interactions between particles that represent masses of dark matter exclusively.  These types of simulations suppress the short-range interactions between masses, making the dynamics collisionless, \\ie\\ particles move under the influence of a global potential.  The consensus view of the results of such simulations is that all bound structures have common density structures determined by a few parameters.  The most often discussed density distribution in the literature is the empircal Navarro-Frenk-White \\citep[][NFW]{nfw96,nfw97} function.  However, many different functions have been invoked to describe the density profiles of dark matter structures resulting from individual simulations \\citep[\\eg][]{s68,m98,n04}.  While there is a considerable amount of literature attempting to explain the universality of N-body density profiles \\citep{sw98,l00,n01,b05,lu06}, no widely accepted explanation exists.  We are looking to gain a fundamental understanding of the physics driving the universality of dark matter halos in collisionless simulations.  Such insight may eventually provide a way of reconciling the differences between simulated halos and the inferred properties of halos of observed galaxies, such as the cusp--core controversy \\citep[\\eg][]{bmr01,sgh05}. However, the results of the present paper cannot be directly applied to observed galaxies and clusters because the latter are affected by dissipational baryonic processes, which we ignore in the present work; we concentrate exclusively on the effects of gravity in collisionless systems.  In the past few years, two interesting relationships have been uncovered that link various dynamical properties of halos formed in simulations like those discussed above.  \\citet{tn01} have found that there is a scale-free relationship between a phase-space proxy and radius; $\\psd \\propto r^{-\\alpha}$, where $\\rho$ is density, $\\sigma$ is velocity dispersion, $r$ is radial distance, and $\\alpha$ is a constant for any given halo.  Interestingly, this relationship is not unique to N-body results.  \\citet{a05} have shown that halos created semi-analytically also have power-law \\psd\\ behavior and speculate that this feature is a result of violent relaxation.  The other relationship relates the density and velocity anisotropy profiles of N-body halos \\citep{hm04}.  Specifically, the curve describing the logarithmic radial derivative of the logarithm of the density has the same basic shape as the radial profile of a halo's anisotropy.  As with the power-law behavior of \\psd, there is no obvious basis for this relationship, but together they signal that there are unifying physical processes at work in halo formation.  This paper and its companion \\citep[][hereafter Paper I]{b06} take steps towards uncovering what these processes are.  In that work, we found that the mechanical equilibrium of halos with velocity isotropy is not the cause of power-law \\psd\\ profiles, as at least two well-known equilibrium density profiles (Hernquist and King) do not produce \\psd\\ power-laws.  However, \\citet{n04} and \\sersic density distributions in mechanical equilibrium produce nearly scale-free \\psd\\ profiles.  We have also investigated density profiles that result from solving the isotropic Jeans equation with the constraint that $\\psd \\propto r^{-\\alpha}$.  In this restricted case, density profiles are similar to those presented in \\citet{n04} but are reproduced better by \\sersic models.  We continue and expand this work by again considering the density profiles that result from solving the Jeans equation under the assumption of scale-free \\psd\\ (motivated by N-body and semi-analytical results), but now the velocity distributions will not be restricted to be isotropic.  The details of solving the ``constrained'' Jeans equation are given in \\S\\ref{constrain}.  The anisotropic Jeans equation admits a wide variety of solutions. We use generic, but sensible, forms for our adopted anisotropy profiles. These profiles and the solutions that result are discussed in detail in \\S\\ref{constb} and \\ref{varyb}.  Our main conclusion follows from these sections; equilibrium density profiles replicate those seen in simulations if the radial density and anisotropy profiles of any given halo have similar shapes.  The degree of connectedness between the \\psd\\ and density slope-anisotropy relations is considered in Section~\\ref{fundy}.  We end with a summary of our findings and our conclusions. ",
        "conclusions": "It is well-established empirically that \\psd\\ is a power-law in radius for virialized, collisionless halos \\citep{tn01,a05,dm05}.  We exploit this relationship and transform the usual Jeans equation into the constrained Jeans equation, solely in terms of the logarithmic density slope $\\gamma$ and the anisotropy $\\beta$.  Doing this allows us to solve directly for $\\gamma$ once an anisotropy distribution has been chosen.  In an earlier paper, \\citet{b06}, we discussed the solutions of this constrained Jeans equation when the anisotropy is assumed to be zero everywhere.  Here, we have extended our earlier work to include constant, but nonzero, anisotropies as well as anisotropy distributions.  In particular, we have looked at two flexible, but distinct, types of anisotropy distributions; a power-law and a hyperbolic tangent.  We find that, in general, the solutions of the constrained Jeans equation that most resemble the empirical NFW and N04 density profiles have $\\blo \\approx 0$ and $\\bhi \\ga 0.5$. Anisotropy must be an increasing function of radius in halos like the ones produced in N-body simulations.  A central conclusion of our present work is that there is a distinct similarity between the shapes of radial density and anisotropy profiles.  Density profiles that best match the NFW form are found when the anisotropy profile is of the tanh form (right panels of Figure~\\ref{tbnfw}), which contains the Osipkov-Merritt and Mamon-Lokas forms as special cases.  N04-like solutions are found with either profile (right panels of Figures~\\ref{plbn04} and \\ref{tbn04}), but with the tanh $\\beta$ profile, the anisotropy must be roughly constant over the range $10^{-3} r_{\\rm vir} < r < r_{\\rm vir}$ (Figure~\\ref{tbn04}).  Another aspect of this profile similarity is given by the $\\beta$ vs. $\\gamma$ plots (Figures~\\ref{plhmnfw}, \\ref{plhmn04}, \\ref{thmnfw}, and \\ref{thmn04}).  The combination of $\\beta$ and $\\gamma$ profiles that produce the most linear curves \\citep[as suggested by the results of N-body simulations;][]{hm04,hm06}, are, of course, those profiles that resemble each other most closely.  The inflection point in the tanh anisotropy distribution also appears in the NFW density slope distribution (right panels of Figure~\\ref{tbnfw} and Figure~\\ref{thmnfw}).  In other words, when this second-derivative symmetry is lost, as with NFW-like density and power-law anisotropy profiles (right panels of Figure~\\ref{plbnfw}), so is the near-linear correlation between $\\beta$ and $\\gamma$.  However, we note that the most linear $\\beta$--$\\gamma$ relations found here are also those that are the most quantitatively different from the results of simulations, in that the solid red lines in Figures~\\ref{plhmn04}a and \\ref{thmnfw}a extend outside the simulation uncertainty range.  This suggests that the relationship between $\\beta$ and $\\gamma$ for any single halo is not strictly linear.  Utilizing the results of semi-analytic models of halo formation, we argue that the $\\beta$--$\\gamma$ relationship is not just a manifestation of a scale-free \\psd.  Rather, a near-linear relation between $\\beta$ and $\\gamma$ supports an earlier argument regarding the nature of halo density profiles \\citep{b05}.  Velocity distributions with nearly isotropic cores and more radially anisotropic envelopes result from radial collapses.  This form of anistropy distribution necessarily creates a density profile that is less steeply rising in the central regions, since there are fewer radial orbits that plunge close to the center and can support a strong cusp.  Likewise, more tangential orbits in the outer regions would increase the density since the average time an orbit spends far from the center would increase.  At least in a qualitative sense, this picture supports the observed similarities between profile shapes; more isotropy (smaller $\\beta$) implies less steeply varying density (smaller $\\gamma$) and vice versa.  The outcome of this work is that there appear to be at least two physical processes at work in the formation of collisionless dark matter halos in N-body simulations.  The evidence presented suggests that one process creates the scale-free \\psd\\ distribution and must be a fairly generic collapse process as the \\psd\\ behavior is seen in the more restrictive ESIM halos.  As a consequence, another process links the anisotropy and density profiles in a unique, nearly linear way."
    },
    "0609/astro-ph0609267_arXiv.txt": {
        "abstract": "We calculate the thermal and dynamical evolution of the surface layers of an accreting neutron star during the rise of a superburst. For the first few hours following unstable $^{12}$C ignition, the nuclear energy release is transported by convection. However, as the base temperature rises, the heating time becomes shorter than the eddy turnover time and convection becomes inefficient. This results in a hydrodynamic nuclear runaway, in which the heating time becomes shorter than the local dynamical time. Such hydrodynamic burning can drive shock waves into the surrounding layers and may be the trigger for the normal X-ray burst found to immediately precede the onset of the superburst in both cases where the {\\sl Rossi X-Ray Timing Explorer} was observing. ",
        "introduction": "\\label{sec:intro}  Superbursts are hour-long X-ray bursts powered by unstable thermonuclear burning on the surfaces of accreting neutron stars in low mass X-ray binaries. Discovered with long term monitoring campaigns by BeppoSAX and the {\\sl Rossi X-Ray Timing Explorer} (RXTE), their recurrence time ($\\approx 1 \\trm{ yr}$), duration ($\\approx \\trm{ hours}$), and energies ($\\approx 10^{42} \\trm{ ergs}$), are $\\sim1000$ times greater than that of normal type I X-ray bursts  (see \\citealt{Kuulkers:04, Cumming:04a, Strohmayer:06} for reviews).  The current view is that superbursts are fueled by the unstable burning of a thick layer of $^{12}$C \\citep{Cumming:01, Strohmayer:02}, which calculations show is present in the ashes of H/He burning in normal bursts \\citep{Brown:98, Schatz:03, Cooper:06}. Fits to superburst lightcurves suggest ignition column depths of $0.5-3\\times10^{12} \\trm{ g cm}^{-2}$ and energy releases of $2 \\times 10^{17} \\trm{ ergs g}^{-1}$, implying $^{12}$C mass fractions $X_{12} > 10\\%$ \\citep{Cumming:06}. For accretion rates $\\dot{M} < 0.3 \\dot{M}_{\\rm Edd}$ (where $\\dot{M}_{\\rm Edd}$ is the global Eddington accretion rate), a larger $^{12}$C fraction is needed  ($X_{12} \\ga 20\\%$) if conditions for the thermal instability are reached before the $^{12}$C burns stably away \\citep{Cumming:01, Cumming:06}.  Previous theoretical studies have focused on two aspects of the $^{12}$C burning scenario: the thermal structure of the surface layers just before ignition \\citep{Cumming:01, Strohmayer:02, Brown:04, Cooper:05, Cumming:06, Cooper:06} and the thermal evolution of the fully burned layers as they cool following the burst \\citep{Cumming:01,  Strohmayer:02, Cumming:04b, Cooper:05, Cumming:06}. In this paper, we investigate the rise of the superburst starting from $^{12}$C ignition.  Unlike the cooling studies, which assume the entire layer of fuel burns instantly, we resolve the thermal and dynamical evolution during the burst rise. Understanding the details of the rise is important because it sets the initial conditions for the cooling models.  The principle result of our study is that burning is so explosive during a superburst rise that the heating time $t_h \\equiv (d\\ln T_b/ dt)^{-1}$ becomes shorter than the dynamical time $t_d = h / c_s \\approx 10^{-6} \\trm{ s}$, where $T_b$, $h$, and $c_s$ are the temperature, pressure scale height, and sound speed at the base of the burning layer.  Such a small $t_h$ is possible because the burning takes place in a strongly degenerate layer (cf. \\citealt{Woosley:76, Taam:78}). During a normal burst, H/He ignites in a weakly degenerate layer and the burning terminates once radiation pressure dominates at $T_b \\approx 2\\times10^9 \\trm{ K}$; throughout the burn $t_h \\ga 10^{-3} \\trm{ s}$ (e.g., \\citealt{Weinberg:06}).  Once $t_h \\la t_d$, convection cannot transport the nuclear energy release and the base undergoes a local thermonuclear runaway, rapidly burning all the $^{12}$C within a thin shell. A hydrodynamic combustion wave forms and propagates into the surrounding fuel, likely driving shock waves into the overlying layers. By contrast, in normal X-ray bursts the energy is transported exclusively by convection and heat diffusion (see \\citealt{Bildsten:98}).  Our primary aim here is to describe those ignition conditions that result in a hydrodynamic nuclear runaway. That is, we examine the range of parameters for which the minimum heating time satisfies $t_{h, {\\rm min}} < t_d$. We therefore limit our present analysis to the sequence of events leading up to the formation of a combustion wave. In \\S~\\ref{sec:initial} we describe the initial conditions of our rise models and in \\S~\\ref{sec:rise} we present our main results. We conclude in \\S~\\ref{sec:summary} and briefly discuss how the hydrodynamic nature of the burning can influence the observed rise. ",
        "conclusions": "\\label{sec:summary}  We have shown that the superburst rise evolves through three nuclear burning stages: an hour-long convective stage, a runaway stage, and a hydrodynamic stage. As $T_b$ rises during the convective stage, the heating rate increases and eventually burning is so rapid that $t_h \\approx t_e$, resulting in a local nuclear runaway. During the runaway, $t_h$ becomes even shorter than the dynamical time $t_d$, and a hydrodynamic combustion wave forms.  The combustion wave propagates from the site of the runaway either subsonically as a deflagration or supersonically as a detonation. The issues are similar to those found in theoretical studies of type Ia supernova, in which a C/O white dwarf ignites $^{12}$C at (or near) its center, where the initial density is $\\sim 10^9 \\trm{ g cm}^{-3}$. In a superburst, the combustion wave is subject to much larger gravitational accelerations. One can show that if the superburst combustion wave is a deflagration, it must always be subject to strong Rayleigh-Taylor instabilities. The correct solution might therefore be a detonation rather than a deflagration.  A normal X-ray burst is found to immediately precede the onset of the superburst in both cases where the RXTE Proportional Counter Array was observing \\citep{Strohmayer:02, Strohmayer:02b}. These precursors may be triggered by hydrodynamic perturbations originating beneath the H/He layer.  In particular, if the combustion wave propagates as a detonation, it will drive shocks into the surrounding layers (see \\citealt{Zingale:01}). A deflagration will also generate shocks if the burning front is sufficiently distorted by Rayleigh-Taylor instabilities. The shocks will initially be weak, with overpressures at $y_b$ of $\\Delta p/ p \\sim E_{\\rm nuc} / E_{\\rm init} \\approx 0.2$, where $E_{\\rm init}$ is the initial internal energy of the gas at the ignition depth. However, the shocks steepen as they propagate upwards, becoming quite strong by the time they reach the H/He layer at $y \\approx 3\\times10^8 \\trm{ g cm}^{-2}$. To estimate the strength, approximate the initially weak, upward-propagating, shock as a plane parallel acoustic wave for which $(\\Delta p)^2/\\rho c_s$ is constant by flux conservation. For $y \\ga 10^{10} \\trm{ g cm}^{-2}$, relativistic degeneracy pressure dominates and $\\rho \\propto y^{3/4}$, $c_s \\propto y^{1/8}$. Thus, while the shock is weak, the overpressure increases with decreasing column depth as $\\Delta p / p \\propto y^{-9/16}$. When the shock is strong $\\Delta p / p \\propto y^{-(1+6\\beta)/4}$ where $\\beta \\simeq 1/(2+[2\\gamma/(\\gamma-1)]^{1/2})$ (\\citealt{Whitham:74}; see also \\citealt{Matzner:99}). Here $\\gamma =4/3$ which gives $(1+6\\beta)/4 = 0.56 \\simeq 9/16$ so that over the entire range in $y$ we have \\beq \\Delta p / p \\approx 0.2 (y / y_{12})^{-9/16}, \\eeq and at the H/He layer $\\Delta p / p \\approx 20 y_{12}^{9/16}$. In an upcoming paper we show that a shock of this strength deposits sufficient entropy into the H/He layer that it can trigger He burning and thereby engender a precursor to the main burst."
    },
    "0609/astro-ph0609051_arXiv.txt": {
        "abstract": "The solar magnetic activity cycle is responsible for periodic episodes of severe space weather, which can perturb satellite orbits, interfere with communications systems, and bring down power grids. Much progress has recently been made in forecasting the strength and timing of this 11-year cycle, using a predictive flux-transport dynamo model \\cite{dik05,dik06}. We can strengthen the foundation of this model by extending it to match observations of similar magnetic activity cycles in other Sun-like stars, which exhibit variations in their Ca~{\\sc ii} H and K emission on time scales from 2.5 to 25 years \\cite{bal95}. This broad range of cycle periods is thought to reflect differences in the rotational properties and the depth of the surface convection zone for stars with various masses and ages. Asteroseismology is now yielding direct measurements of these quantities for individual stars, but the most promising asteroseismic targets are in the southern sky ($\\alpha$~Cen~A, $\\alpha$~Cen~B, $\\beta$~Hyi), while the existing activity cycle survey is confined to the north. We are initiating a long-term survey of Ca~{\\sc ii} H and K emission for a sample of 92 southern Sun-like stars to measure their magnetic activity cycles and rotational properties, which will ultimately provide independent tests of solar dynamo models. ",
        "introduction": "Astronomers have been making telescopic observations of sunspots since the time of Galileo, gradually building a historical record showing a periodic rise and fall in the number of sunspots every 11 years. We now know that sunspots are regions with an enhanced local magnetic field, so this 11-year cycle actually traces a variation in surface magnetism. Attempts to understand this behavior theoretically often invoke a combination of differential rotation, convection, and meridional flow to modulate the field through a magnetic dynamo. Significant progress in this area quickly followed after helioseismology allowed the detailed characterization of the Sun's interior structure and dynamics over the past decade (e.g., the convection zone depth and meridional flow).  \\begin{figure} \\epsfxsize 3.0in \\epsffile{fig1.eps} \\caption{A flux-transport dynamo model correctly reproduces the relative amplitudes of the past 8 solar cycles, and then predicts the strength and timing of the next cycle (from \\cite{dik06}, reproduced with permission).} \\end{figure}  This has gradually led to the development of a predictive flux-transport dynamo model, which has recently been used to forecast the timing \\cite{dik05} and strength \\cite{dik06} of the upcoming solar magnetic cycle. By calibrating the relationship between sunspot area and surface magnetic flux using modern measurements, historical sunspot data going back more than 125 years could be used to test the model. After initializing the dynamo with the first several cycles, the model correctly reproduced the relative amplitudes of the subsequent 8 cycles, and was then used to predict the characteristics of future cycles (see Fig.~1). This is not merely of academic interest, since the magnetic maximum of the Sun is accompanied by severe space weather events. Advance knowledge of the timing and overall strength of the coming cycle allows us to prepare for the adverse consequences, which include the destabilization of orbiting satellites, interference with communications systems, and the disruption of power grids.  While the success of this dynamo model in reproducing the solar cycle is very encouraging, there will soon be new opportunities for independent tests using the magnetic activity cycles of other Sun-like stars. The properties of the dynamo---most notably the cycle period---are sensitive to specific physical circumstances, such as the meridional flow, the degree of latitudinal differential rotation and the depth of the convection zone. By matching the dynamo model to the single set of circumstances represented by the Sun, we risk the possibility of fine tuning, which could ultimately limit our comprehension of the underlying physical mechanisms. If we exploit the magnetic activity cycles observed in a wide variety of Sun-like stars, we can extend the calibration to dozens of unique circumstances and ensure that the physical basis of the model is more broadly applicable. Specifically, observations of stellar activity cycle periods and differential rotation profiles (from this proposed survey, and its existing northern counterpart) as well as stellar convection zone depths (from ground- and space-based asteroseismology) will allow us to test the dynamo model by inverting the implied stellar meridional flow, and comparing it to the solar case.  Although we can rarely observe spots on other Sun-like stars directly, these areas of concentrated magnetic field produce strong emission in the Ca~{\\sc ii}~H (396.8~nm) and K (393.4~nm) spectral lines. The intensity of the emission scales with the amount of non-thermal heating in the chromosphere, making these lines a useful spectroscopic proxy for the strength of, and fractional area covered by, magnetic fields \\cite{lei59}. Disk-integrated time-series measurements of Ca~{\\sc ii}~H and K emission in the Sun easily reveal the 11-year cycle (see Fig.~2). Similar measurements for a large sample of northern Sun-like stars began at the Mount Wilson Observatory in 1966 \\cite{wil78}, and continue to this day \\cite{bal06}.  \\begin{figure}[t] \\epsfxsize 3.0in \\epsffile{fig2.eps} \\caption{The Ca~{\\sc ii}~HK index of the Sun observed by the Mount Wilson survey (adapted from \\cite{bal95}).} \\end{figure}  The Mount Wilson Ca~HK survey, as it is known, has revealed magnetic activity variations in Sun-like stars with cycle periods ranging from 2.5 to more than 25 years. The frequent time sampling for 111 stars since 1980 also allows the measurement of rotation periods, from the repeated signature of individual active regions rotating into view and then out again. In some cases the rotation periods appear to drift over longer timescales, due to the migration of active regions to lower latitudes which are rotating faster (the same mechanism that is responsible for the ``Butterfly Diagram'' on the Sun). Thus, it is sometimes possible to measure or place limits on the degree of latitudinal differential rotation in other Sun-like stars. Early analysis of the Mount Wilson data revealed an empirical correlation between the mean level of magnetic activity and the rotation period normalized by the convective timescale \\cite{noy84a}. The data also established a relation between the rotation rate and the period of the observed activity cycle, which generally supports a dynamo interpretation \\cite{noy84b}.  It was several decades before similar observations were conducted for a large sample of Sun-like stars in the southern sky \\cite{hen96}. These data confirmed the bimodal distribution in mean stellar activity {\\it levels} first seen in the northern sample \\cite{vp80}, but they could say nothing about activity {\\it cycles} because the data were confined to a single epoch of observation. Another decade has passed, and there is still no time-domain survey of Ca~HK emission in southern Sun-like stars. This is particularly disappointing in light of the fact that the most promising candidates for asteroseismology, which will soon allow the direct measurement of stellar convection zone depths (see Fig.~3), are stars in the southern sky---most notably the brightest solar-type stars $\\alpha$~Cen~A, $\\alpha$~Cen~B, and $\\beta$~Hyi which has long been studied as a ``future Sun''.  \\begin{figure*}[t] \\epsfxsize 5.5in \\hskip 0.5in \\epsffile{fig3.eps} \\caption{The expected oscillation signal from the edge of the convective envelope of a 1.1 M$_{\\odot}$ star with (dotted) and without (solid) overshoot (adapted from \\cite{mon00}). Data with this level of precision will soon emerge from space-based asteroseismology missions such as CoRoT and Kepler.} \\end{figure*}  Solar-like oscillations have already been detected in these stars \\cite{bed04,kje05,bed01}, but tests of solar dynamo models will not be possible until their magnetic activity cycle periods are established. There is some evidence from sparse x-ray observations of $\\alpha$~Cen~A that it may exhibit an activity cycle as short as 3.5~years \\cite{rob05}, despite the fact that it is only 10\\% more massive than the Sun and just slightly older. A firm measurement of the cycle period in $\\alpha$~Cen~A, particularly when combined with forthcoming asteroseismic data, could provide a very stringent second test of the flux-transport dynamo model. ",
        "conclusions": ""
    },
    "0609/astro-ph0609098_arXiv.txt": {
        "abstract": "I assess the similarities and differences between the star-formation modes in quiescent spiral galaxies versus those in violent starburst regions, including galactic nuclei. As opposed to the quiescent star-formation mode, current empirical evidence on the star-formation processes in the extreme, high-pressure environments induced by galaxy encounters strongly suggests that star {\\it cluster} formation is an important and perhaps even the dominant mode of star formation in such starburst events. This implies that by using star clusters as unique diagnostic probes, we can trace a galaxy's most violent star formation history very well, at least for the past few Gyr. The sizes, luminosities, and mass estimates of the young massive star clusters are entirely consistent with what is expected for young Milky Way-type globular clusters (GCs). Recent evidence lends support to the scenario that GCs, which were once thought to be the oldest building blocks of galaxies, are still forming today. ",
        "introduction": "Luminous, massive yet compact star clusters (YMCs; often with masses $m_{\\rm cl} \\ge 10^5 {\\rm M}_\\odot$) are the hallmarks of the most intense starbursts. YMCs are therefore important as benchmarks of cluster formation and evolution, and also as important tracers of the (stellar) initial mass function (IMF) and other physical properties of starbursts. This is so, because star clusters can extremely well be represented by single-age, single-metallicity ``simple'' stellar populations (SSPs). This sets the YMCs formed in violent starburst events apart from the apparently similar nuclear star clusters, since the latter are often characterised by stellar population mixes (e.g., Walcher et al. 2006). The defining characteristics of YMCs have been explored in starburst regions in several dozen galaxies, both in normal spirals and in gravitationally interacting galaxies.  The crucial question remains, however, whether or not at least some of the YMCs observed in extragalactic starbursts might survive to become (possibly somewhat more metal-rich) counterparts of the Galactic globular clusters (GCs) when they reach a similar age. If we could resolve this issue convincingly, one way or the other, the implications would be far-reaching for a wide range of astrophysical questions, including our understanding of how galaxy formation, assembly and evolution proceeds, and what the process and requisite conditions are for star (cluster) formation.  The evolution of young clusters depends crucially on their stellar IMF: if its slope is too shallow, i.e., if the clusters are significantly deficient in low-mass stars compared to, e.g., the solar neighbourhood, they will likely disperse within about a Gyr of their formation (e.g., Gnedin \\& Ostriker 1997; Goodwin 1997; Smith \\& Gallagher 2001; Mengel et al. 2002; de Grijs \\& Goodwin, in prep.). At present, there are two principal approaches which one can use to address the underlying IMFs of extragalactic YMCs.  \\subsection{The Cluster Luminosity Function: the case of M82}  In de Grijs et al. (2003a,b) we reported the discovery of an approximately log-normal cluster luminosity and mass function (CLF, CMF) for the subpopulation of approximately coeval star clusters at the intermediate age of $\\sim 1$ Gyr in M82's fossil starburst region ``B''.  This provided the first deep CLF (CMF) for a star cluster population at intermediate age, which thus serves as an important benchmark for theories of the evolution of star cluster systems [see also Goudfrooij et al. (2004) for a related important result for NGC 1316, at $\\sim 3$ Gyr]. The fact that we considered an approximately coeval subset of the M82 B cluster population, combined with our use of the 100 per cent completeness limit as our base line ensures the robustness of the CMF peak detection. Additional arguments in favour of this robustness are provided in de Grijs et al. (2005).  Most analyses of the young Large Magellanic Cloud cluster system (with ages $\\le 2 \\times 10^9$ yr), starting with the seminal work by Elson \\& Fall (1985), seem to imply that the CLF of YMCs is well described by a power law. However, in de Grijs \\& Anders (2006) we showed that this power-law behaviour breaks down below $m_{\\rm cl} \\simeq 10^3 {\\rm M}_\\odot$, as well as for our youngest LMC cluster age bins. Although the latter could partially be due to the so-called ``infant mortality'' effect, by which some 60 -- 90 per cent of newly formed clusters are disrupted in the first $\\sim 10$ Myr of their lifetime (e.g., Bastian et al. 2005), the deviation of the LMC CLFs at low cluster masses is significant for ages up to at least 1 Gyr, beyond which our observational completeness limit prevents us from making definite statements on the CLF shape at low mass (de Grijs \\& Anders 2006).  Nevertheless, for old GC systems with ages $\\ge 10^{10}$ yr, the CLF shape is -- on the other hand -- well established to be roughly log-normal, and almost universal among local galaxies (barring a slight metallicity dependence). This type of observational evidence has led to the popular, but thus far mostly speculative theoretical prediction that not only a power-law, but {\\it any} initial CLF (CMF) will be rapidly transformed into a log-normal distribution because of {\\it (i)} stellar evolutionary fading of the lowest-luminosity (mass) objects to below the detection limit; and {\\it (ii)} disruption of the low-mass clusters due both to interactions with the gravitational field of the host galaxy, and to internal two-body relaxation effects leading to enhanced cluster evaporation (but see Parmentier \\& Gilmore 2006).  From our detailed analysis of the expected evolution of CMFs starting from initial log-normal and initial power-law distributions (de Grijs et al. 2005), we conclude that the observed turnover in the M82 B CMF is inconsistent with a scenario in which the 1 Gyr-old cluster population originated from an initial power-law mass distribution. This applies to a large range of ``characteristic'' cluster disruption time-scales, and is supported by arguments related to the initial density in M82 B, which would be unphysically high if the present cluster population were the remains of an initial power-law distribution (particularly in view of the effects of the cluster ``infant mortality'', which requires large excesses of low-mass unbound clusters to be present at the earliest times).  In de Grijs et al. (2003c) we showed that the CMFs of YMCs in many different environments are well approximated by power laws with slopes $\\alpha \\simeq -2$. However, except for the intermediate-age cluster systems in M82 B and NGC 1316 (Goudfrooij et al. 2004), the {\\it expected} turn-over (or peak) mass (based on comparisons with present-day GC systems and taking evolutionary fading into account) in most YMC systems observed to date occurs close to or below the observational detection limit, simply because of their greater distances and shallower observations. As such, these results are not necessarily at odds with each other, but merely hindered by observational selection effects.  \\subsection{High-resolution spectroscopy: individual cluster analysis}  With the ever increasing number of large-aperture ground-based telescopes equipped with state-of-the-art high-resolution spectrographs and the wealth of observational data provided by the {\\sl Hubble Space Telescope}, we may now finally be getting close to resolving the potentially far-reaching issue of YMC-to-GC evolution conclusively. To do so, one needs to obtain {\\it (i)} high-resolution spectroscopy, in order to obtain dynamical mass estimates, and {\\it (ii)} high-resolution imaging to measure their sizes (and luminosities). As a simple first approach, one could then construct diagnostic diagrams of YMC mass-to-light ratio vs. age, and compare the YMC locations in this diagram with SSP models using a variety of IMF descriptions (cf. Smith \\& Gallagher 2001; Mengel et al. 2002; Bastian et al. 2006). However, such an approach, while instructive, has serious shortcomings:  {\\it (i)} In this simple approach, the data can be described by {\\it both} variations in the IMF slope {\\it and} variations in a possible low-mass cut-off; the models are fundamentally degenerate for these parameters.  {\\it (ii)} While the assumption that these objects are approximately in virial equilibrium is probably justified at ages greater than a few $\\times 10^7$ yr (at least for the stars dominating the light), the {\\it central} velocity dispersion (as derived from luminosity-weighted high-resolution spectroscopy) does not necessarily represent a YMC's total mass. It is now well-established that almost every YMC exhibits significant mass segregation from very early on, so that the effects of mass segregation must be taken into account when converting central velocity dispersions into dynamical mass estimates (see also Lamers et al. 2006; Fleck et al. 2006; Moll et al. 2006 and in prep.).  {\\it (iii)} With the exception of a few studies (e.g., M82-F; Smith \\& Gallagher 2001), the majority of YMCs thus far analysed in this way have ages around 10 Myr. Around this age, however, red supergiants (RSGs) appear in realistic stellar populations. Unfortunately, the model descriptions of the RSG phase differ significantly among the various leading groups producing theoretical stellar population synthesis codes (Padova vs. Geneva vs. Yale), and therefore the uncertainties in the evolutionary tracks are substantial.  {\\it (iv)} For nuclear star clusters, one also needs to relax the SSP approximation, and allow multiple stellar populations to be present in the mix; this has the potential of matching the observed high-resolution spectra very well, however (Walcher et al. 2006). ",
        "conclusions": ""
    },
    "0609/astro-ph0609344_arXiv.txt": {
        "abstract": "We have performed several simulations of black hole systems (non-rotating, black hole spin parameter $a=0.0$ and rapidly rotating, $a=0.95$) with a geometrically thin Keplerian disk using the newly developed RAISHIN code. The simulation results show the formation of jets driven by the Lorentz force and the gas pressure gradient. The jets have mildly relativistic speed ($\\gtrsim 0.4~c$). The matter is continuously supplied from the accretion disk and the jet propagates outward until each applicable terminal simulation time (non-rotating: $t/\\tau_{\\rm S} = 275$ and rotating: $t/\\tau_{\\rm S} = 200$, $\\tau_{\\rm S} \\equiv r_{\\rm S}/c$). It appears that a rotating black hole creates an additional, faster, and more collimated matter-dominated inner outflow ($\\gtrsim 0.5~c$) formed and accelerated by the twisted magnetic field resulting from frame-dragging in the black hole ergosphere. This is the first known simulation confirming the formation of an inner magnetically-driven, matter-dominated jet by the frame-dragging effect from a black hole co-rotating with a thin Keplerian disk threaded by a vertical magnetic field. This result indicates that jet kinematic structure depends on black hole rotation and on the initial magnetic field configuration and strength. ",
        "introduction": "Both magnetic and gravitational fields play an important role in the dynamics of matter in many astrophysical systems. In a highly conducting plasma, the magnetic field can be amplified by gas contraction or shear motion. Even when the magnetic field is weak initially, the magnetic field can grow on short time scales and influence the gas dynamics of the system. The magnetic influence on gas dynamics is particulary important for a compact object such as a black hole or a neutron star. Relativistic jets have been observed or postulated in various astrophysical objects, including active galactic nuclei (AGNs) (e.g., Urry \\& Pavovani 1995; Ferrari 1998), microquasars in our galaxy (e.g., Mirabel \\& Rodiriguez 1999), and gamma-ray bursts (GRBs) (e.g., Zhang \\& M\\'{e}sz\\'{a}ros 2004; Piran 2005; M\\'{e}sz\\'{a}ros 2006). The most promising mechanisms for producing the relativistic jets involve magnetohydrodynamic centrifugal acceleration and/or magnetic pressure driven acceleration from an accretion disk around the compact objects (e.g., Blandford \\& Payne 1982; Fukue 1990), or involve the extraction of rotating energy from a rotating black hole (Penrose 1969; Blandford \\& Znajek 1977).  General relativistic magnetohydrodynamics (GRMHD) codes with fixed spacetimes have been developed to investigate relativistic magnetorotators (RMRs) (e.g., Koide et al. 1998; De Villiers \\& Hawley 2003; Gammie et al.\\ 2003; Komissarov 2004; Ant\\'{o}n et al.\\  2005; Anninos et al.\\ 2006). These codes have been used to study the Blandford-Znajek effect near a rotating black hole (Koide 2003; Komissarov 2005; McKinney 2005), and the formation of GRB jets in collapsars (Mizuno et al.\\ 2004a, 2004b; De Villiers et al.\\ 2005b).  Recently, in order to investigate the properties of accretion flows onto a black hole associated with the magneto-rotational instability (MRI) (Balbus \\& Hawley 1991), many simulations have been performed using a thick torus-like disk (the disk thickness $H/r > 0.1$, where $H$ is height of the disk and $r$ is radius from a black hole) with weak poloidal magnetic fields in a torus (the plasma beta, $\\beta = p_{\\rm gas}/p_{\\rm mag} > 100$, where $p_{\\rm gas}$ is gas pressure and $p_{\\rm mag}$ is magnetic pressure) (De Villiers et al.\\ 2003, 2005a; Krolik et al.\\ 2005; Hawley \\& Krolik 2006; Beckwith et al.\\ 2006; McKinney \\& Gammie 2004; McKinney 2006). The initial ``poloidal-loop'' magnetic fields in the torus contribute to the generation of MRI, diffusion of matter and magnetic field, and jet generation. However, in their simulations the structure of magnetic fields that piled up and twisted near the black hole is different from the magnetic fields which are twisted by a thin Keplerian disk and/or the frame-dragging effect of the rotating black hole ($H/r \\sim 0.06$) with a stronger initial vertical magnetic field ($\\beta < 40$) (Koide, Shibata, \\& Kudoh 1998, 1999; Koide et al. 2000; Nishikawa et al. 2005). In the thin disk simulations MRI does not grow since the wavelength at the maximum growth rate is larger than the height of the thin disk. Koide et al.\\ (2000) have found that jets are formed from thin Keplerian accretion disks for both counter- and co-rotating black holes. In the co-rotating disk case, the jet had a two-layered structure: an inner gas pressure-driven jet, and an outer magnetically driven jet (Koide et al. 2000).  In this Letter we report on new simulations showing jet formation from the black hole magnetosphere co-rotating with a thin Keplerian disk threaded by a vertical magnetic field using the recently developed three-dimensional, GRMHD code RAISHIN (Mizuno et al. 2006).  Our new result shows that an inner matter-dominated jet is generated by the magnetic field twisted in the ergosphere of the rotating black hole in addition to the outer jet formed by the co-rotating accretion disk seen by Koide et al.\\ (2000). While our inner jet is similar the inner jet created by a counter-rotating disk found by Koide et al. (2000), in this case the inner jet is located closer to the rotation axis. ",
        "conclusions": "We have performed simulations of jet formation from a geometrically thin accretion disk near both non-rotating and rotating black holes using the newly developed RAISHIN code (Mizuno et al. 2006). Simulation results show the formation of magnetically-driven jets accelerated by the Lorentz force and the gas pressure. In general, results are similar to those found in previous GRMHD simulations with a thin Keplerian disk (Koide et al.\\ 1998, 1999, 2000; Nishikawa et al. 2005). However, in this study the rotating black hole creates a second, faster, and more collimated inner matter-dominated jet formed by the twisted magnetic field resulting from frame-dragging in the black hole ergosphere. Thus, kinematic jet structure depends on black hole rotation. While an inner jet is created by a counter-rotating thin Keplerian disk (Koide et al. 2000), the inner jet seen in our simulations is a new result not found in a previous co-rotating thin Keplerian disk case (Koide et al.\\ 2000). This difference appears to be the result of the weaker initial magnetic field strength used in the present simulations. With stronger magnetic field (Koide et al.\\ 2000), transport of angular momentum from the disk is more rapid, accreting matter falls more quickly, is decelerated by the centrifugal barrier near the black hole strongly, and generates a stronger shock. As a result, jets are accelerated by the gas-pressure rather than the Lorentz force of magnetic field twisted by frame-dragging. The jet seen in the co-rotating disk simulation in Koide et al. (2000) corresponds to the outer jet seen in our present co-rotating disk simulation. Therefore this is a first simulation confirming the formation of an inner matter-dominated jet driven by the magnetic field twisted by the frame-dragging effect from a black hole co-rotating with a thin Keplerian disk threaded by a vertical magnetic field.  A two component jet structure has also been seen in GRMHD simulations of a black hole co-rotating with a thick torus (e.g., Hawley \\& Krolik 2006; McKinney 2006). One component is a matter-dominated outflow (funnel wall jet) with a mildly relativistic velocity ($\\sim 0.3c$) along the centrifugal barrier accelerated and collimated by magnetic and gas pressure forces in the inner torus and the surrounding corona. This formation mechanism and jet are the same as the outer jet seen in our simulations. Therefore the funnel wall jet in Hawley \\& Krolik (2006) is the same as the outer jet in our GRMHD simulations. Their other component is a highly-relativistic Poynting flux dominated jet that is produced from the formation of a large scale radial magnetic field within the funnel. In our simulation, such a highly-relativistic Poynting flux dominated jet is not seen. This is likely caused by the difference in the initial magnetic field configuration. In Hawley \\& Krolik (2006) and McKinney (2006), there are initial ``poloidal-loop'' magnetic fields inside the torus. In the simulations, the magnetic field is twisted and expands from the torus as a magnetic tower and fills the funnel region (Hirose et al. 2004). In the late stage of the simulations a highly-relativistic Poynting flux dominated jet is formed in the low density regions in the funnel. On the other hand, our magnetic field vertically threads the disk and ergosphere initially. The vertical magnetic field near the black hole region is twisted by the frame-dragging effect (see Fig. 2) and forms the inner jet. The inner jet is much denser than the highly-relativistic Poynting flux dominated jet of Hawley \\& Krolik (2006) and McKinney (2006) because in the inner jet the matter is supplied from the accretion disk and the free-falling corona. While our jets are formed by the same basic mechanisms as in Hawley \\& Krolik (2006) and McKinney (2006), the different initial magnetic configuration has led to different jet properties and a slower matter dominated jet spine. Previous results with the co-rotating disk case of Koide et al. (2000) found that a stronger magnetic field suppressed the formation of the presently observed matter-dominated jet spine."
    },
    "0609/astro-ph0609458_arXiv.txt": {
        "abstract": "We describe further results of a program aimed to yield $\\sim$10$^4$ fully characterized optical identifications of ROSAT X-ray sources. Our program employs X-ray data from the ROSAT All-Sky Survey (RASS), and both optical imaging and spectroscopic data from the Sloan Digital Sky Survey (SDSS). RASS/SDSS data from 5740~deg$^2$ of sky spectroscopically covered in SDSS Data Release 5 (DR5) provide an expanded catalog of 7000 confirmed quasars and other AGN that are probable RASS identifications. Again in our expanded catalog, the identifications as X-ray sources are statistically secure, with only a few percent of the SDSS AGN likely to be randomly superposed on unrelated RASS X-ray sources. Most identifications continue to be quasars and Seyfert 1s with $15<m<21$ and $0.01<z<4$; but the total sample size has grown to include very substantial numbers of even quite rare AGN, e.g., now including several hundreds of candidate X-ray emitting BL Lacs and narrow-line Seyfert 1 galaxies. In addition to exploring rare subpopulations, such a large total sample may be useful when considering correlations between the X-ray and the optical, and may also serve as a resource list from which to select the ``best\" object (e.g., X-ray brightest AGN of a certain subclass, at a preferred redshift or luminosity) for follow-on X-ray spectral or alternate detailed studies. ",
        "introduction": "The premier X-ray imaging survey of the astronomical sky is the ROSAT All-Sky Survey (hereafter RASS; Voges et al. 1999, 2000). In the 1990s, RASS surveyed most of the celestial sphere in the $0.1-2.4$~keV range with the Position Sensitive Proportional Counter (PSPC; Pfeffermann et al. 1988) to a typical limiting sensitivity of $\\sim 10^{-13}$~erg~cm$^{-2}$~s$^{-1}$, cataloging more than $10^5$ X-ray sources in the RASS Bright and Faint Source Catalogs.  With RASS's typical X-ray positional uncertainty of order $\\sim10-30''$, the effort involved in optically identifying a large fraction of the cataloged abundant RASS sources may appear daunting; and yet, several large-scale identification efforts are currently making substantial progress. Thousands of optical identifications for RASS sources are now suggested, albeit with various levels of identification confidence on the nature of the counterparts, ranging from: statistical cross-correlations with Schmidt photographic plate images (e.g., see radio/optical/X-ray matches cataloged in Flesch and Hardcastle 2004); to optical spectral identification estimates from low spectral-resolution photographic Schmidt objective prism plates (e.g., see Zickgraf et al. 2003; Mickaelian et al. 2006); to digital optical photometry and multi-object spectrophotometry (e.g., see RASS Bright and Faint Source Catalog identifications in Anderson et al. 2003).  For comparison to such recent RASS efforts involving thousands of suggested optical identifications each, it may be recalled that the most complete optical identification effort at similar depth from {\\it Einstein} X-ray data was the Extended Medium Sensitivity Survey (EMSS; e.g., Gioia et al. 1984, Stocke et al. 1991). Quasars/AGN were the predominant class in the EMSS, accounting for about half of the $\\sim800$ {\\it Einstein} identifications. Although the EMSS total of $\\sim400$ AGN identifications is quite large in aggregate,  certain subsets are critically small when subdivided; for example, although the EMSS provided fundamental insights into BL Lac evolution (Morris et al. 1991), the entire EMSS sample included only $\\sim$40 BL~Lacs.  In order to substantively expand on such earlier {\\it Einstein} programs, current efforts should provide samples significantly exceeding $\\sim10^3$ spectroscopically secure optical identifications. Our approach, detailed previously in Anderson et al. 2003 (hereafter, Paper 1), relies on the very good sensitivity match between AGN detectable in X-rays with RASS and AGN accessible to good quality optical spectra with the Sloan Digital Sky Survey (SDSS; York et al. 2000). Our ultimate aim has been the cataloging and characterization of $\\sim10^4$ X-ray identifications; the AGN in this sample are not only statistically secure in their identifications as X-ray sources, but are also accompanied by high quality and uniform data in the X-ray from RASS and in the optical (both photometry and spectroscopy) from SDSS.    In section 2, we briefly review relevant aspects of the SDSS data, and the selection and confirmation of candidate AGN optical counterparts from SDSS; detailed descriptions of these aspects, and a much more complete set of references, may be found in Paper 1. In section 3, we present results from our RASS/SDSS program for 5740~deg$^2$ of sky (about $4\\times$ the area coverage considered in Paper 1), providing an updated and expanded catalog of X-ray and optical properties of 7000 likely SDSS AGN identifications for RASS X-ray sources. In sections 3-4, we discuss some AGN subclasses of special interest, including more than 260 X-ray emitting BL Lac candidates, and 770 candidate X-ray emitting narrow-line Seyfert 1s and related quasars. Section 5  discusses the reliability of the identifications and some ensemble properties of the expanded sample. Section~6 provides an updated example optical/X-ray correlation study. A short summary is provided in Section~7. ",
        "conclusions": "The SDSS optical and RASS X-ray surveys are well-matched to each other, allowing efficient large-scale identification of X-ray source optical counterparts. Application of our approach to SDSS DR5 data has provided homogeneous identification and RASS/SDSS flux and spectroscopic data for a large sample of X-ray emitting quasars and other kinds of AGN. The combination of SDSS multicolor selection and RASS data---and in some cases FIRST radio information---is highly efficient for the selection of X-ray emitting quasars/AGN. In our updated analysis encompassing 5740~deg$^2$ of sky, 7000 plausible X-ray emitting quasars/AGN have been optically identified, including hundreds of rare cases such as BL~Lacs and NLS1s. The RASS/SDSS survey is rapidly approaching $\\sim10^4$ fully and homogeneously characterized optical counterpart identifications. The large sample will allow for a variety of more detailed studies of various AGN subclasses and individual objects of special interest, as well as for studies of ensemble correlations between optical and X-ray wavebands. In closing, we also note that many of the cataloged objects have X-ray fluxes accessible to at least the next generation of envisioned high-quality X-ray spectroscopy experiments, such as Constellation-X.  \\bigskip"
    },
    "0609/astro-ph0609491_arXiv.txt": {
        "abstract": "In this paper we present the results of the analysis of the INTEGRAL data for the black hole transient GRO J1655-40. The observations consist of four ToO (100 ks each) observations's AO-3 spread from February to April of 2005. We present here the preliminary spectral analysis of this source between 5 and 200 keV. Also, some comments of the light curves obtained during this period are shown. ",
        "introduction": "The X-ray transient GRO J1655-40 (also called X-ray Nova Scorpii 1994) was discovered with the Burst and Transient Source Experiment (BATSE) on board the Compton Gamma-Ray Observatory (CGRO) on 1994 July 27 \\citep{zhang94}. The optical counterpart was discovered soon after by \\citet{bailyn95a}. Subsequent optical studies showed that the system is an LMXB composed of a blue subgiant (spectral type F4 IV) as the secondary and a black hole as the primary ($m_{BH}=7.02{\\pm}0.22$$M_{\\odot}$) \\citep{orosz97a}, located at a distance of 3.2 kpc \\citep{tingay95} (see also \\citet{foellmi06}). \\citet{bailyn95b} (see also \\citet{orosz97b} and \\citet{hooft98}) established the orbital inclination of the system to be ${\\simeq}70^{\\circ}$.   Galactic black hole binaries (hereafter BHB) show interesting luminosity/spectral properties which appears to be strongly related with the accretion behaviour onto the black hole (hereafter BH). Matter accretes onto the BH through an accretion disk, which has nearly Keplerian orbits; the last stable orbit is called ``Innermost Stable Circular Orbit'' (i.e. ISCO) \\citep{shapiro83} which is located at $R_S=6R_{g}$ (for a Schwarzschild BH; where $R_{g}=GM/c^{2}$ is the gravitational radius). Note that not always the matter arrives at this radius; this depends on the particular state.  There are five states known in BHB: quiescent, low-hard, intermediate, high-soft and very high states in the unification scheme (e.g. \\citet{esin97}); the inner radius of the accretion disk is thought to decrease from quiescence to VHS. At lower mass accretion rate, corresponding to several percent of the Eddington luminosity, a BHB usually enters the low-hard (LH) state and at very low accretion rates it reaches the quiescent state, which may be just an extreme of the LH state. In both of these states, the spectrum of a BHB is dominated by a hard, nonthermal power-law component (photon index ${\\sim}1.7$); it is most plausibly explained as due to Comptonization of soft photons (from a hot optically thick disk) by a hot optically thin plasma. \\citet{narayan96} postulated that in these states the disk does not extend down to the ISCO, but is truncated at some larger radius and the interior volume is filled by a hot ($T_{e}{\\sim}100$ keV), radiative inefficient, advection-dominated accretion flow or ADAF.  The Multicolor Disk Model (MCD) (\\citet{mitsuda84}, which is an application of the \\citet{shakura73} disk) is used to describe the thermal component that is dominant in the Very High (VH) and High-Soft (HS) states. In Figure 4.1 \\citep{mcclintock04} the unification model from \\citet{esin97} is pictured, using both MCD and ADAF models. It represents how the geometry of the accretion flow changes as the mass accretion rate from the donor star varies. When the accretion is low, the flow consists of two zones (disk and ADAF) and for the two states in which the accretion is high, the disk extends down to the ISCO. In all five states the disk is bathed in a corona which resembles a continuation to the ADAF. Although this model had success in the explanation of evolution of the states, it has important limitations not discussed in this paper.  In this paper, we briefly report on observations of the BHB GRO J1655-40 made by INTEGRAL. After describing the observations and our analysis techniques, we present the preliminary results of the spectral fits. We then briefly discuss the results and end with some concluding remarks.  \\begin{figure} \\centering \\vspace{4cm} \\includegraphics[width=1.0\\linewidth]{groj1655_sig2.ps} \\caption{Mosaic image (obtained in revolution 295) of the GRO J1655-40 region as seen by ISGRI in the 20-40 keV energy range. Besides the target source, several other high energy sources are visible.\\label{mosaic}} \\end{figure} ",
        "conclusions": "As we noted above, there was a transition from an LH state to an HS state \\citep{brocksopp06}. This can be explained in the common accepted view of black hole X-ray transients, in which the low-mass companion star undergoes Roche-lobe overflow, but at a sufficiently low rate that the gas is accumulated in the (cool) disk until a critical level is reached, at which point the outburst occurs (\\citet{lasota01}; \\citet{meyer04}). In between outbursts, the accretion rate is low and is conjectured that an ADAF model region fills the inner accretion disk. With this scenario, in the quiescent and low states the accretion is low and the spectrum is hard as is shown in our spectrum from revolution 290. During the outburst, the spectrum changes to a ${\\sim}1$keV thermal spectrum, often interpreted as the radiation from a geometrically thin, optically thick accretion disk, superposed on a softer power law (photon indices ${\\Gamma}{\\sim}2-3$) extending to ${\\sim}200$ keV, as shown in our spectra from revolutions 295 and 296. When the accretion disk becomes thick and very near to the ISCO, then the reflection features are manifested in the shape of the spectra (see \\citet{done01} and references therein), as can be seen in our spectra from revolutions 299 and 304. It is worth noticing that, when matter is close to the BH, jet production is an alternative mechanism for the origin of the hard X-rays.  We can conclude that our spectra is consistent with major models for outburst evolution, including model of \\citet{esin97}."
    },
    "0609/astro-ph0609172_arXiv.txt": {
        "abstract": "A remarkable series of breakthroughs in numerical relativity modeling of black hole binary mergers has occurred over the past few years.  This paper provides a general overview of these exciting developments, focusing on recent progress in merger simulations and calculations of the resulting gravitational waveforms. ",
        "introduction": "The coalescence of a comparable mass black hole binary (BHB) is a powerful source of gravitational waves and proceeds in three phases: inspiral, merger, and ringdown \\cite{Flanagan97a}. The inspiral stage, when the BHs are widely separated and following quasi-circular trajectories, and the ringdown stage, during which the final merged BH settles into a quiescent Kerr state, can both be treated analytically. However the merger phase, in which the two BHs plunge together and merge to form a highly distorted remnant BH, occurs in the arena of very strong, dynamical gravitational fields and can only be calculated using numerical relativity.  This final merger will produce an intense burst of gravitational radiation with a luminosity $\\sim 10^{23} L_{\\odot}$, briefly emitting more energy than the combined light from all the stars in the visible universe.  Such bursts are expected to be among the strongest sources for LISA, which will observe mergers of massive BHBs.  Mergers of stellar mass and intermediate mass BHBs are likely to be the strongest sources for ground-based gravitational wave detectors such as LIGO and VIRGO.  Observing the gravitational waves from the final merger will allow unprecedented tests of general relativity in the dynamical, strong field regime -- provided we know the waveforms that general relativity predicts.  The final merger of BHBs also has compelling astrophysical implications.  In particular, when the BHs have unequal masses, the resulting gravitational wave emission is asymmetric; since the gravitational waves carry momentum, the merged remnant BH suffers a recoil kick \\cite{Favata:2004wz}. If this kick is large enough, it could eject the merged remnant from its host structure, thereby affecting the overall rate of merger events \\cite{Merritt:2004xa}.  In addition, since BHBs are generally expected to be spinning, their mergers could produce interesting spin dynamics and couplings \\cite{Merritt:2002hc}.  For more than 30 years, numerical relativists have attempted to calculate the merger of comparable mass BHs and the resulting gravitational waveforms.  This has proved to be a very difficult undertaking indeed. In particular, the simulation codes have been plagued by a host of difficulties, typically resulting in various instabilities that caused them to crash before any sizeable fraction of a binary orbit could be evolved.  Recently, however, a series of dramatic breakthroughs has occurred in numerical relativity, resulting in accurate and robust simulations of BHB mergers and their gravitational waveforms. This paper provides a general overview of these exciting developments. We begin with a brief overview of numerical relativity and BHB calculations.  The heart of the paper then focuses on recent progress in computing BHB orbits and mergers, and the resulting waveforms. The paper concludes with a summary and outlook for the future.  With the goal of reaching as wide an audience as possible, technical details are deliberately kept to a minimum; interested readers can find more detailed information in the references.\\footnote{Since this paper is not a full review of the subject, the reference list is representative rather than comprehensive.  We attempted to cite key papers from the major numerical relativity efforts that were steppingstones to the current breakthroughs.} We follow conventional practice by setting $G = 1$ and $c = 1$, which allows us to measure both time and distance in terms of mass $M$.  In particular, $1 M \\sim (5 \\times 10^{-6})(M/M_{\\odot}){\\rm sec} \\sim 1.5 (M/M_{\\odot})  {\\rm km}$. Spatial indices are taken to have the range $i = 1,2,3$. Note that the simulation results scale with the masses of the BHs, and are thus applicable to LISA as well as to ground-based detectors. ",
        "conclusions": "The past few years have seen a remarkable series of breakthroughs in numerical relativity modeling of BHB mergers.  The first BHB orbit was achieved by Br\\\"ugmann, Tichy, and Jansen~\\cite{Bruegmann:2003aw} using special comoving gauge conditions in a traditional numerical relativity code in late 2003.  Roughly a year and a half later, the first plunge, merger, and ringdown calculation by Pretorius appeared, using nonstandard techniques and including the extraction of the gravitational waveform \\cite{Pretorius:2005gq}.  Less than six months later, in November 2005, the moving puncture method was introduced by the UTB \\cite{Campanelli:2005dd} and Goddard \\cite{Baker:2005vv} groups, enabling merger calculations with simple but novel gauge conditions in traditional numerical relativity codes.  This was rapidly exploited by its developers and, by the early spring of 2006, the Goddard group had obtained a consistent solution for the gravitational wave burst from the merger of two equal mass Schwarzschild BHs, independent of the (quasi-astrophysical) initial conditions \\cite{Baker:2006yw}. Very shortly thereafter, the UTB group produced the first evolutions of equal mass, spinning BHB mergers, demonstrating the orbital hangup when the BH spins are aligned with the orbital angular momentum \\cite{Campanelli:2006uy}.  The first attempt to model unequal mass mergers using the moving puncture method was carried out by the Penn State group \\cite{Herrmann:2006ks} and first appeared in January 2006. This was soon followed by similar work at higher resolution and larger initial separations by the Goddard group \\cite{Baker:2006vn}. Overall, by the summer of 2006, a large segment of the numerical relativity community was actively adopting, adapting, and exploiting the moving puncture method.  At the present time there is broad consensus that the merger of two equal mass Schwarzschild BHs produces a final remnant BH with spin $a \\sim 0.7M$, and that the amount of energy radiated in the form of gravitatational waves, starting with the final few orbits and proceeding through the plunge, merger and ringdown, is $\\sim 0.04M$. The UTB and Goddard groups are working with Pretorius to compare their waveforms; preliminary results using longer runs by the UTB group and Pretorius show good agreement with the waveforms obtained by the Goddard group shown in Figure \\ref{fig:waves}.  Plans are underway to include other groups in this comparison effort.  The outlook for continued progress in BHB merger simulations is very bright.  New work with moving punctures continues to appear \\cite{Sperhake:2006cy,Hannam:2006vv,Campanelli:2006fg}. Pretorius is carrying out new and longer runs with his generalized harmonic code.  The Caltech-Cornell collaboration has made important progress in carrying out orbits with their spectral code based on a fully first-order formulation of the Einstein equations \\cite{Scheel:2006gg} and hopes to achieve mergers soon.  The impressive progress in this field, occurring across a broad front, is very encouraging.  Stay tuned!   \\begin{theacknowledgments} It is a pleasure to acknowledge John Baker, Jim van Meter, and Richard Matzner for valuable comments on the manuscript, and to thank Frans Pretorius and Manuela Campanelli for the use of figures from their papers. This work was supported in part by NASA grant O5-BEFS-05-0044. \\end{theacknowledgments}"
    },
    "0609/astro-ph0609728_arXiv.txt": {
        "abstract": "We discuss the region of transition between galactic and extragalactic cosmic rays. The exact shapes and compositions of these two components contains information about important parameters of powerful astrophysical sources and the conditions in extragalactic space as well as for the cosmological evolution of the sources of high energy cosmic rays. Several types of experimental data, including the exact shape of the ultrahigh energy cosmic rays, their chemical composition and their anisotropy, and the fluxes of cosmogenic neutrinos have to be included in the solution of this problem. \\vspace{1pc} ",
        "introduction": "In view of the expected high experimental statistics from the Auger Observatory~\\cite{Auger} the interest in the transition from galactic to extragalactic cosmic rays has increased significantly. It is now a standard belief that above some high energy of order 10$^{19}$ eV all observed cosmic rays should come from extragalactic sources because they cannot be contained in the Galaxy long enough to be accelerated~\\cite{Cocconi}. For the rest of this paper we shall assume that the Ultra-High Energy Cosmic Rays (UHECR) are accelerated at powerful astrophysical objects.  Extragalactic cosmic rays would lose energy in propagation from their sources to us if their sources are isotropically and homogeneously distributed in the Universe. The main energy loss process is the photoproduction interaction in the microwave background (MBR) that causes the GZK effect, the steepening of the cosmic ray spectrum above 6$\\times$10$^{19}$ eV~\\cite{GZK}. Several fits of the existing experimental data have been published in recent years that derived different values of the most important astrophysical and cosmological parameters: the acceleration (injection in terms of cosmic ray propagation) energy spectrum of these particles and the maximum acceleration energy, their chemical composition and the cosmological evolution of their astrophysical sources. In the assumption that most of UHECR are protons injection spectra as different as E$^{-2.7}$~\\cite{BGG05} and E$^{-2.0}$~\\cite{BW03} and cosmological evolution of the type $(1+z)^m$ with $m$ values from 0 to 3 have been obtained.  Other attempts~\\cite{APO05a,APO05b,HST06} have assumed that extragalactic cosmic rays have at their sources the same mixed chemical composition as the low energy galactic cosmic rays. In such a case the main energy loss process is the disintegration of the heavy nuclei mostly in interactions in MBR. Hadronic interactions become important only after the energy per nucleon exceeds the photoproduction threshold. The fits of the observed cosmic rays spectrum under this assumption gives an intermediate $E^{-2.2-2.4}$ injection spectrum.  In all these attempts the fits of the observations show the end of the galactic cosmic rays spectrum which is obtained by subtraction of the propagated extragalactic spectrum from the experimentally observed one. This process gives some limits of the astrophysical parameters~\\cite{ddmts05} when the subtraction gives unphysical negative values.  An important and interesting question is what such models predict  for the chemical composition of UHECR, which for our purposes we shall define as cosmic rays of energy above 10$^{18}$~eV. The assumption that extragalactic cosmic rays are protons obviously lead to a composition that becomes very light at GZK and super-GZK energies. Under the assumption that extragalactic cosmic rays have a mixed composition the conclusions are more complicated and the question of the highest acceleration energy is more important.  In terms of high energy astrophysics the location of the transition region can be used to estimate the strength of the cosmic ray sources in different types of astrophysical objects, from spiral galaxies like our own to Active Galactic Nuclei (AGN). In our current understanding the acceleration spectra depend on the strength of astrophysical shocks and the maximum acceleration energy is a function of the magnetic field strength $B$ and fluctuations $\\delta B/B$. Other important parameters are the cosmological evolution of the UHECR sources and the strength of the extragalactic magnetic fields around the sources and the average magnetic field in the cosmologically nearby Universe.  We shall discuss several types of data and their predictions in different models: UHECR spectra and chemical composition, the anisotropy in this energy range, and the production of {\\em cosmogenic}~\\cite{BZ69} neutrinos by extragalactic cosmic rays. ",
        "conclusions": "The transition region between galactic and extragalactic cosmic rays is currently not very well defined and studied. It has a significant astrophysical importance, because the energy content of the extragalactic cosmic rays, as well as their injection spectrum and composition, will define much better the type of their sources. In addition, a derivation of the cosmological evolution of the sources will not only restrict the number of available scenarios, but also provide an additional measure of the cosmological evolution of powerful astrophysical objects.  This will not be easy to do even with the expected much higher statistics and detection precision of the Auger Observatory. We should employ all available types of measurements to solve this problem as well as it is necessary. These measurements include the cosmic ray composition in the transition region and above it, the cosmic ray anisotropies, and the possible detection of cosmogenic neutrinos.  The first experimental result on the shape of this transition presented by the Fly's Eye group~\\cite{Birdetal93} used the simultaneous change of the cosmic ray composition and energy spectrum shape. We have to do the same now - to combine all observational data, look for consistency between different sets, and with models of the extragalactic cosmic rays origin.  The possible detection of cosmogenic neutrinos would be a powerful test if we succeed in collecting a reasonable statistics of such events. It is unlikely this will happen very soon, but the expanding efforts for designing and building detectors for ultrahigh energy neutrinos are very encouraging.  {\\bf Acknowledgments} This talk is based on work performed in collaboration with D.~Seckel, D.~DeMarco, J.~Alvarez-Mu\\~{n}iz, R.~Engel and others. Partial support of my work comes from US DOE Contract DE-FG02 91ER 40626 and NASA Grant ATP03-0000-0080."
    },
    "0609/astro-ph0609202_arXiv.txt": {
        "abstract": "{  Multiplet structures are a common feature in pulsating stars, and can be the consequence of rotational splitting, mode interaction or sinusoidal amplitude variations.  In this paper we examine the phenomenon of (nearly) equidistant triplets, which are unlikely to be caused by rotational splitting, in different types of pulsating stars: a $\\delta$~Scuti variable (1 Mon), an RR Lyrae variable (RR Lyr) and a short-period Cepheid (V743 Lyr). We examine the hypothesis that one of the modes forming the triplet results from a combination of the other two modes. The analyses were carried out on recent data sets by using multiple-frequency analyses and statistics with the package PERIOD04. In particular, the small departures from equidistance were calculated for the three selected stars.  For the $\\delta$~Scuti variable 1 Mon, the departure from equidistance  is only 0.000079 $\\pm$ 0.000001 cd$^{-1}$ (or 0.91 $\\pm$ 0.01 nHz). For 1~Mon the Combination Mode Hypothesis with a mode excited by resonance is the most probable explanation. For the star RR Lyr, the hypothesis of resonance through a combination of modes should be considered. The results for the best-studied cepheid with a Blazhko period (V743~Lyr) are inconclusive because of an unfavorable period of 1.49d and insufficient data. } ",
        "introduction": " ",
        "conclusions": "We have examined the phenomenon of (nearly) equidistant frequency triplets in a $\\delta$~Scuti star, an RR~Lyrae star, and a cepheid variable and tested whether the smallest amplitude component could be a combination of the other two modes. For the $\\delta$~Scuti variable 1 Mon, a departure from equidistance  of only 0.000079 $\\pm$ 0.000001 cd$^{-1}$ (or 0.91 $\\pm$ 0.01 nHz). Here the Combination Mode Hypothesis with a mode excited by resonance is the most probable explanation. For the star RR~Lyr, the hypothesis of resonance through a combination of modes was also found to be plausible. Finally, the results for the best-studied cepheid with a Blazhko period (V743 Lyr) were inconclusive because of an unfavorable period of 1.49d and insufficient data."
    },
    "0609/astro-ph0609805_arXiv.txt": {
        "abstract": "We have used the Submillimeter Array for the first interferometric CO imaging toward the starburst-Seyfert nucleus of the southern barred spiral galaxy NGC 1365, which is one of the four galaxies within 30 Mpc that have $L_{8-1000 \\; \\mu {\\rm m}} \\geq 10^{11} \\Lsol$. Our  mosaic maps of \\twelveCO, \\thirteenCO, and \\CeighteenO (J=2--1) emission at up to 2\\arcsec\\ (200 pc) resolutions have revealed a circumnuclear gas ring and several CO clumps in the central 3 kpc. The molecular ring shows morphological and kinematical signs of bar-driven gas dynamics, and the region as a whole is found to follow the star formation laws of Kennicutt. We have found that some of the gas clumps and peaks in CO brightness temperature, which we collectively call CO hotspots, coincide with the radio and mid-infrared sources previously identified as dust-enshrouded super star clusters. This hotspot-cluster association suggests that either the formation of the most massive clusters took place in large molecular gas concentrations (of $\\Sigma_{\\rm mol} \\sim 10^{3} $ \\Msol\\ \\persquarepc\\ in 200 pc scales) or  the clusters have heated their ambient gas to cause or enhance the CO hotspots. The active nucleus is in the region of weak CO emission and is not associated with distinctive molecular gas properties. ",
        "introduction": "} Central regions of spiral galaxies sometimes host vigorous star formation called (circum)nuclear starbursts. This phenomenon forms stars at a higher rate and more preferentially in compact massive clusters than a non-starburst environment \\citep[e.g.,][]{Telesco93,Barth95}. Since stars form from molecular gas, properties of the molecular gas are expected to control the starbursts. The study of such properties in nearby starbursts tells us about physical processes involved in the phenomenon, and the information serves as a basis for modeling starbursts throughout galaxy evolution.   The typical extent of a (circum)nuclear starburst in our vicinity is $\\lesssim$1 kpc, which is only 20\\arcsec\\ for a distance of 10 Mpc.  High angular resolution is therefore essential for starburst studies. In particular, high resolution imaging of molecular gas with millimeter interferometers has been a powerful tool in the field \\citep[e.g.,][]{Jogee05}. Until very recently, however, such arrays had been operated only at the mid-latitudes in the northern hemisphere, leaving many southern galaxies unattainable.   We have now observed \\object{NGC 1365}, a starburst-Seyfert composite galaxy in the southern sky, using the Submillimeter Array (SMA)\\footnote{ The Submillimeter Array is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy and Astrophysics, and is funded by the Smithsonian Institution and the Academia Sinica. } . This is part of our SMA survey of southern starbursts, aiming to explore molecular gas and dust in some of the most important but undeservedly less observed starburst galaxies in the southern sky. NGC 1365 is one of our targets because it has the circumnuclear starburst typical of barred spirals, and because the infrared luminosity of the starburst is among the largest in non-merging galaxies, bridging more luminous starbursts in distant mergers and less luminous ones in undisturbed spirals. Other galaxies reported so far from our survey are \\object{M83}, \\object{NGC 253}, and \\object{NGC 3256} \\citep{Sakamoto04,Sakamoto06a,Sakamoto06b}.   We have observed CO emission to trace the molecular gas. Our resolution of 2\\arcsec\\ is a significant improvement over the best previous ones of 5\\arcsec--10\\arcsec. In addition, the three isotopic lines that were simultaneously imaged, i.e., the J=2--1 transitions of \\twelveCO, \\thirteenCO, and \\CeighteenO, provide information about the physical properties of the molecular gas. Among other things, we have discovered bright CO spots, or CO hotspots, that are associated with young massive star clusters in the starburst. We discuss their properties and relation to the starburst.   The goals of this paper are (1) to report the detailed structure, kinematics, and properties of molecular gas in the central region of NGC 1365 on the basis of the new high-resolution data, (2) to search for the relation between circumnuclear starburst and the gas properties, and (3) to see if there is anything special in the molecular gas in the vicinity of the Seyfert nucleus. The next section further introduces NGC 1365, and \\S \\ref{s.observations} describes our SMA observations and data reduction. The observational results are presented in \\S \\ref{s.results}, and are discussed in \\S \\ref{s.bardynamics} --  \\S \\ref{s.starformation}. A summary is given in \\S \\ref{s.summary}. ",
        "conclusions": "\\label{s.summary} We have observed 1.3 mm CO lines and continuum in the center of the barred spiral galaxy NGC 1365 at high spatial resolutions. The circumnuclear starburst in the galaxy is among the most luminous and intense in non-merging galaxies. Our main findings regarding the properties of molecular gas and their relation to the starburst and to the Seyfert nucleus are as follows.   \\begin{enumerate} \\item There is an oval ring of molecular gas around the nucleus, having a \\tilde2 kpc extent and twin-peak enhancements at its apocenters, and connected to gas lanes on the leading side of the stellar bar. The gas morphology and kinematics in the region are consistent with the model that the ring is at the transition region between \\xone- and \\xtwo-like gas stremlines in the bar. The structure is typical of barred galaxies.   \\item The molecular ring coincides with the starburst ring seen in radio emission, suggesting that the gas distribution largely determines the extent and the shape of the starburst. Moreover, the intensity of the starburst averaged over the circumnuclear region is what is expected from the molecular gas surface density of the region and the star formation laws in \\citet{Kennicutt98}. A plausible synopsis of the galaxy's starburst is as follows: The long (\\tilde18 kpc) stellar bar funneled a large amount of molecular gas (\\tilde$10^{9}$ \\Msol) toward the galactic center, and the resulting high gas surface density (\\tilde$10^{3}$ \\Msol\\ \\persquarepc) in the relatively large circumnuclear ring ($r\\tilde1$ kpc) resulted in the starburst (\\tilde3 \\Msol\\ \\peryr\\ \\persquarekpc) of large total luminosity (\\tilde$10^{11}$ \\Lsol) following the star formation laws. --- A remaining question is why the transient starburst (gas consumption timescale \\tilde$10^{8}$ yr) is seen now.   \\item There are compact ($\\lesssim 200$ pc) peaks of CO integrated intensity and brightness temperature, which we call CO hotspots, on the circumnuclear ring. Some of the hotspots coincide with the radio and mid-IR sources that have been identified as dust-enshrouded super star clusters. The association of the CO hotspots and super clusters allows two interpretations. The first one is that the most luminous and presumably massive super clusters were formed in regions with the highest gas surface densities (\\tilde$10^{3}$ \\Msol\\ \\persquarepc\\ in 200 pc scale). The second one is that the luminous embedded clusters heated their ambient gas to form the hotspot. The former, which is more conventional in terms of the interpretation of CO intensities at $\\gtrsim 100$ pc scales, implies a relation between the cluster formation and gas complexes already identifiable at a 200 pc resolution.   \\item The active nucleus coincides with the dynamical center of the galaxy within 30 pc, according to the CO velocity field. Neither integrated intensity nor brightness temperature of CO peaks at the nucleus. We found no anomaly in molecular gas that may be related to the AGN.  \\end{enumerate}"
    },
    "0609/astro-ph0609034_arXiv.txt": {
        "abstract": "We report the results of 1D particle-in-cell simulations of ultrarelativistic shock waves in proton-electron-positron plasmas. We consider magnetized shock waves, in which the upstream medium carries a large scale magnetic field, directed transverse to the flow. Relativistic cyclotron instability of each species as the incoming particles encounter the increasing magnetic field within the shock front provides the basic plasma heating mechanism. The most significant new results come from simulations with mass ratio $m_p/m_\\pm = 100$. We show that if the protons provide a sufficiently large fraction of the upstream flow energy density (including particle kinetic energy and Poynting flux), a substantial fraction of the shock heating goes into the formation of suprathermal power-law spectra of $e^--e^+$. Cyclotron absorption by the pairs of the high harmonic ion cyclotron waves, emitted by the protons, provides the non-thermal acceleration mechanism. As the proton fraction increases, the non-thermal efficiency increases and the $e^--e^+$ power-law spectra harden.  When the proton fraction is small (pair plasma almost charge symmetric), the $e^-$ and $e^+$ have approximately equal amounts of non-thermal heating. At the lower range of our simulations with mass ratio 100, when the ions contribute 56\\% of the upstream flow energy flux, the pairs' non-thermal acceleration efficiency by energy is about 1\\%, increasing to 5\\% as the ions' fraction of upstream flow energy increases to 72\\%. When the fraction of upstream flow energy in the ions rises to 84\\%, the efficiency of non-thermal acceleration of the pairs reaches 30\\%, with the $e^+$ receiving most of the non-thermal power.  We suggest that the varying power law spectra observed in synchrotron sources that may be powered by magnetized winds and jets might reflect the correlation of the proton to pair content enforced by the underlying electrodynamics of these sources' outflows, and that the observed correlation between the X-ray spectra of rotation powered pulsars with the X-ray spectra of their nebulae might reflect the same correlation. ",
        "introduction": "}  Collisionless shock waves in relativistic astrophysical flows have been implicated as the energizing mechanism for emission in a variety of non-thermal sources - termination shocks of pulsar winds (\\cite{slane05,kaspi04}), hot-spot shocks terminating jets from radio galaxies (\\cite{heavens87, wilson00, balcus05}) and internal shocks in AGN and microquasar jets (\\cite{marscher85, turler00, turler04}), internal and external shocks in GRBs (\\cite{piran04}) and outflows from other high energy transients (e.g., \\cite{gaensler05}).  While shock jump conditions applied to models of magnetohydrodynamic flow set the constraints on what fraction of flow energy goes into heat of all sorts,  the partition of that energy between ions and electrons (and positrons, when these are present), and the distributions of particle momenta (thermal or non-thermal), are kinetic questions not addressed by the macroscopic conservation laws. The kinetic questions are those of collisionless plasma physics - two body encounters cannot generate the entropy implied by the jump conditions on the observed scales of transition from non-radiative flow (upstream, in the shock excitation model) to the observed regions of synchrotron emission (downstream, in the shock model), nor can collisional processes lead to non-thermal particle spectra.  A fundamental approach in which the flow and kinetic structure of relativistic shocks are resolved on all scales offers the opportunity to evaluate the nature of the electromagnetic turbulence generated, as a function of the upstream parameters of the flow, simultaneously with a determination of the heated (possibly non-thermally accelerated) particle spectra both up and downstream. For this purpose, fully kinetic particle-in-cell (PIC) simulations (\\cite{birdsall91,dawson95}) provide a powerful investigatory tool, since these model the plasma from first principles, with the only approximations being those that go into the method - primarily the ``cloud-in-cell'' algorithm, which in effect softens the interactions to the grid scale.  PIC simulations of relativistic transverse magnetosonic shocks in a symmetric pair plasma were first reported by \\cite{langdon88} and by \\cite{gallant92}, using spatially 1D models. Magnetized shocks are characterized by  coherent gyration (\\cite{alsop88}) and cyclotron instability (\\cite{hoshino91}) of the plasma species as they encounter the shock front - magnetic reflection mediates the density and velocity transition, while the instabilty of the induced gyration mediates the entropy generation. Simulations of the Weibel instability when $B_0 =0$ with applications to unmagnetized relativistic shocks have been reported recently by a number of authors (\\cite{silva03,fred04,hed04,jaro05,hed05}). Spitkovsky and Arons (in preparation) have found comparable results, and show how the shock structures in pair plasmas depend on the magnetic field strength and orientation with respect to the flow.  All these results indicate that non-thermal particle acceleration is either weak or absent in shocks in pair plasmas.  \\cite{hoshino92} found that when the upstream flow includes heavy ions which carry a significant fraction of the upstream flow energy, non-thermal acceleration {\\it does} occur, as a consequence of the absorption of ion waves emitted at harmonics of the ion cyclotron frequency reabsorbed by the pairs, even though the shock structure is one dimensional and no cross field diffusion occurred. The interaction is resonant, in contrast to the non-resonant scattering underlying Diffusive Shock Acceleration (DSA). Computational limitations required the use of a low rest mass ratio, $m_i /m_\\pm = 10, 20$.  Those inhomogeneous shock calculations showed no electron acceleration. While most of the characteristics of the instability scale with just the energy density ratio between the protons and the pairs, the efficacy of the waves as accelerators of electrons and positrons also depends on the polarization properties of the proton cyclotron waves. The lack of electron acceleration at low mass ratio stems from the polarization of the waves excited. In a quasi-neutral pair plasma the waves are linearly polarized. If the heavy ions are a small fraction by number, the waves excited by the ion maser remain almost linearly polarized. Since linearly polarized waves are an equal mixture of left and right circularly polarized modes, equal electron and positron acceleration might have been expected. However, the simulations also revealed that the mechanism is a significant accelerator only if the energy density in the ions $U_{1i} = m_i N_{i1} ( \\gamma_1 -1)  c^2$, where $N_{i1}$ is the upstream ion density and $\\gamma_1$ is the upstream flow Lorentz factor, is comparable to the upstream energy density in the pairs, $U_{1\\pm} = m_\\pm ( \\gamma_1 -1) c^2 (N_{1+} + N_{1-})$. Then $N_{1i}/(N_{1+} + N_{1-}) = (m_\\pm /m_i) (U_{1i}/U_{1\\pm}) \\sim m_\\pm /m_i$. For the small mass ratios employed, the dispersion relations for small amplitude waves show that at comparable energy densities in the species, the waves are still strongly left handed, with consequent preference for positron acceleration.  The earlier work also entirely neglected upstream thermal spread in the momenta of the species. The acceleration mechanism relies upon cyclotron absorption of waves in the pairs at high harmonics of the ion cyclotron frequency. Since the shock in the pairs alone heats the pairs to a temperature $T_{2\\pm} \\approx m_\\pm c^2 \\gamma_1$ (if $\\sigma_1 \\equiv B_1^2/4\\pi \\rho_1 \\gamma_1 c^2 \\ll 1 $, with $B_1$ the upstream magnetic field strength and $\\rho_1$ the upstream rest mass density, both measured in the downstream frame), efficient non-thermal acceleration of the pairs requires cyclotron absorption of ion waves at frequency $\\omega \\approx \\Omega_{c\\pm} = eB /m_\\pm c \\gamma_1 = (m_i /m_\\pm ) eB /m_i c \\gamma_1 = (m_i /m_\\pm ) \\Omega_{ci}$; that is, the relativistic cyclotron instability of the ion ring in the shock front's leading edge must generate significant power at harmonic numbers $ n \\sim m_i /m_\\pm \\gg 1$. When the upstream plasma is completely cold, this is not a problem - the growth rate of the instability in a homogeneous medium scales as $\\Omega_{ci} n^{-1/9}$ (\\cite{hoshino91}). We shall see that if there is thermal dispersion in the upstream ion momenta $\\Delta p $, growth is inhibited, with no substantial wave power at harmonic numbers above $m_i c \\gamma_1 /\\Delta p$.  In this paper we reinvestigate the one dimensional shock structure of a transverse magnetosonic shock, with particular attention to the effect of more realistic mass ratios and the incorporation of thermal dispersion in the upstream medium.  In \\S \\ref{sec:shocks} we undertake the study of the structure and acceleration properties of transverse relativistic shocks.  \\S \\ref{sec:pairshocks} contains a study of shocks in pure pair plasma. Consistent with previous results, we find no signs of non-thermal particle acceleration. We introduce a proton component in \\S \\ref{sec:dens} and find, for the first time in the literature, evidence of non-thermal acceleration not only of positrons but also of electrons. In \\S \\ref{sec:high-mass} we show that at high mass ratio the waves become close to linearly polarized when the ion and pair energy densities are comparable. This leads to comparable heating and acceleration rates for positrons and electrons. In \\S \\ref{sec:temp} we include in the shock simulations a thermal spread of the ions' upstream distribution and study how this affects the results. Finally, we present a summary of our results and our conclusions in \\S \\ref{sec:discuss}. Aspects of the linear stability of relativistic protons gyrating in a relativstically hot pair plasma are described in the Appendix, \\S \\ref{sec:linear}. ",
        "conclusions": "}   In accord with previous results (\\cite{langdon88,gallant92}), we find full thermalization in pure $e^- - e^+$ shocks.  The reason is the rapid reabsorption of waves at and above the upper hybrid frequency $\\Omega_{UH\\pm}$, which causes the shock to behave as if it were collisional - rapid emission of the EM waves by the relativistic cyclotron instability at the particle reflection front, with rapid reabsorption of the waves, and exchange of real photons, is a mechanism just as capable as the exchange of virtual photons in Coulomb collisions in producing local thermodynamic equilibrium.  The conclusion, that pure pair shocks yield relativistic Maxwellians in the downstream medium, is not limited to the one-dimensionality of our model, as is shown elsewhere in 3D simulations by Spitkovsky and Arons (in preparation). Those simulations show that at sufficiently low $\\sigma$ ($\\sigma < 10^{-3}$ is low enough), the relativistic Weibel instability ({\\it e.g.}, \\cite{yoon87}) takes over from cyclotron instability in the pairs as the deceleration and thermalization mechanism of the shock wave, even when the background magnetic field is fully transverse to the flow.  The physical reasons for this result are discussed in detail elsewhere, in the context of the 3D results\\footnote{Other recent simulations of Weibel mediated shocks in $e^\\pm$ and $e-p$ plasmas (\\cite{silva03, hed04, fred04, hed05, jaro05}) have had insufficient size and temporal length (or have been constrained by periodic boundary conditions in the flow direction) to show the development of the shock wave; as a result, they have been, in effect, studies of the foreshock.}. However, the essence is that the magnetic fields generated by the Weibel instability in an unmagnetized shock saturate in a tangled network of tangled magnetic filaments (\\cite{milos06}), taking a distance $l_{therm} \\sim 25 c/\\omega_{p1}$. These filaments scatter and thermalize the incoming particles.  Since $l_{therm} /r_{L1\\pm} \\sim 25 \\sqrt{\\sigma_{1\\pm}}$, Weibel based thermalization (a multi-diemensional effect) controls the shock structure for $\\sigma_{1\\pm} < 0.002$.  For larger magnetizations, the cyclotron effects studied here dominate, somewhat modified by mild crinkling of the shock front.  In the shock wave, the incoming flow continuously injects cold ($M_s \\gg 1$) ions into the pair plasma, the latter being already heated by the leading shock in the pairs alone. These freshly injected cold ions efficiently produce high harmonic waves which are cyclotron resonant with the pairs, even though the heated ions that are found further downstream from the pair shock do not efficiently radiate high harmonic waves. Magnetic reflection at the shock front produces a continuous supply of cold gyrating ions, thus ensuring energy transfer from the ions to the non-thermal pairs in the downstream medium, different from what is the case in the analogous initial value problem in a uniform medium.  When the ion density is a small fraction of the pair density, non-thermal acceleration becomes symmetric between electrons and positrons. It also becomes less efficient, leading to steeper power-law slopes in the non-thermal part of the spectrum. As shown in Fig.~\\ref{fig:allhist100}, the spectra take a Maxwellian form up to energies around $E_\\pm \\approx \\gamma_1 m_\\pm c^2$, then have a power law form at higher energies. The energy above which the power law distribution sets in, and the slope, depend on $N_i / N_\\pm$.  Thus, apparent universality (if it exists) of power law spectra in synchrotron sources where {\\it relativistic} shocks are inferred to play a role in non-thermal accleration would be a consequence of universality in the composition of the upstream flow, expressed as the ratio $N_{1i}/N_{1\\pm} $, if the cyclotron resonance mechanism described here really is the dominant non-thermal heating process and if all relativistic shocks are sufficiently close to the transverse geometry studied here.  At present, nothing has been published on the full 3D structure of these shocks\\footnote{\\cite{hed04} in essence study the foreshock of an $e-p$ relativistic shock in 3D with a mass ratio $m_p /m_e = 15$.}. Our expectation is that the 3D structure of electron-positron-ion transverse shocks will show a transition to unmagnetized physics at a low value of $\\sigma_{1,total}$.  The value of this transition magnetization is quite uncertain, since current neutralization of the ion filaments formed by the Weibel instability by electron return currents (\\cite{lyub06}) may greatly increase the thermalization length, thus giving more room in $\\sigma$ space for cyclotron dynamics to ultimately mediate the shock structure. These issues are under investigation.  Relativistic flows that get their energization from Poynting losses from compact objects (possibly all in our list of pulsar winds, microquasar, AGN and GRB jets) may well be in this category, since the pair and ion content may both be set by the underlying electrodynamics of the compact objects (rotating neutron star, magnetized disk.) Quantification of how close shocks have to be to transverse requires consideration of a wider range of shock geometries. Investigation of other non-thermal heating mechanisms, such as Diffusive Shock Acceleration and magnetic pumping, require higher dimensional simulations. Spitkovsky and Arons (in preparation) address both of these issues.  In the shocks, the relative proportion of electron versus positron acceleration follows the ion to pair density, as expected from the changes in wave polarization inferred from the linear theory.  Our study of the shock acceleration process confirms, therefore, the results found by \\cite{hoshino92}, extending them to a higher mass ratio. Also confirmed is the expectation that signs of electron acceleration would be found if one could consider a plasma that contained the same amount of upstream flow energy given to a smaller fraction of ions. The overall evolution of the particles is consistent with resonant cyclotron absorption of magnetosonic waves by the pairs, whose frequencies are harmonics of the ion cyclotron frequency, being the main process providing non-thermal acceleration. Magnetic reflection of the upstream ions in the shock front sets up the unstable ion cyclotron motion which causes the wave emission. The non-thermal acceleration efficiencies and power law indices of the particle spectra are shown, as a function of the basic governing parameter $N_i/N_{-}$, in Table \\ref{tab:epp100}, when the upstream flow is cold. The reduction of the acceleration efficiency for a single value of the Mach number $M_s = 10 \\ll m_i/m_\\pm = 100$, with results summarized in Table \\ref{tab:epp100T}, shows the strong effect of large upstream momentum dispersion in the upstream flow. Equivalently, the absorption of ion cyclotron waves in the shock structure by the pairs is a strong non-thermal acceleration mechanism when the upstream flow is quite cold ($M_s > m_i/m_\\pm $), as results from adiabatic expansion in relativistic outflows from compact objects.  We conclude by pointing out an interesting feature of this acceleration mechanism. The acceleration efficiency and the spectral form of the accelerated particles depend on the composition of the flow upstream of the shock. That composition depends on the source of the relativistic flow within which the shock wave occurs. Relativistic flows of the type considered here are thought to be the consequence of magnetized wind and jet acceleration by compact objects, and pair creation within and near such objects. These properties are the result of the compact objects' electrodynamics, which constrains some aspects of the outflowing plasma composition to particular values - for example, in pulsars and magnetars, the ion flux is constrained to be the Goldreich-Julian value. Thus, one might expect a correlation between the radiative emission properties of the shocked flow and the emission properties of the underlying compact object, as appears to be the case in the X-ray emission from  pulsar wind nebulae and the spectra of the pulsed X-ray emission from the pulsars driving those nebulae (\\cite{gotthelf03})."
    },
    "0609/astro-ph0609752_arXiv.txt": {
        "abstract": "Mean-field theory describes magnetohydrodynamic processes leading to large-scale magnetic fields in various cosmic objects. In this study magnetoconvection and dynamo processes in a rotating spherical shell are considered. Mean fields are defined by azimuthal averaging. In the framework of mean-field theory, the coefficients which determine the traditional representation of the mean electromotive force, including derivatives of the mean magnetic field up to the first order, are crucial for analyzing and simulating dynamo action. Two methods are developed to extract mean-field coefficients from direct numerical simulations of the mentioned processes. While the first method does not use intrinsic approximations, the second one is based on the second-order correlation approximation. There is satisfying agreement of the results of both methods for sufficiently slow fluid motions. Both methods are applied to simulations of rotating magnetoconvection and a quasi-stationary geodynamo. The mean-field induction effects described by these coefficients, e.g. the $\\alpha$-effect, are highly anisotropic in both examples. An $\\alpha^2$-mechanism is suggested along with a strong $\\gamma$-effect operating outside the inner core tangent cylinder. The turbulent diffusivity exceeds the molecular one by at least one order of magnitude in the geodynamo example. With the aim to compare mean-field simulations with corresponding direct numerical simulations, a two-dimensional mean-field model involving all previously determined mean-field coefficients was constructed. Various tests with different sets of mean-field coefficients reveal their action and significance. In the magnetoconvection and geodynamo examples considered here, the match between direct numerical simulations and mean-field simulations is only satisfying if a large number of mean-field coefficients are involved. In the magnetoconvection example, the azimuthally averaged magnetic field resulting from the numerical simulation is in good agreement with its counterpart in the mean-field model. However, this match is not completely satisfactory in the geodynamo case anymore. Here the traditional representation of the mean electromotive force ignoring higher than first-order spatial derivatives of the mean magnetic field is no longer a good approximation. \\\\[2mm] \\textit{Keywords:} Magnetohydrodynamics; Mean-Field Theory; Dynamo Coefficients; Magnetoconvection; Geodynamo ",
        "introduction": "\\label{sec:1}  Large-scale cosmic magnetic fields as the Earth's, the solar and the galactic magnetic fields are maintained by hydromagnetic dynamos \\cite[e.g.][]{weiss_02}. However, global computational dynamo models simulate only the geodynamo resonably well, while it remains difficult to tackle the stellar or galactic dynamo problem in this way. Despite the increasing computational power, a direct numerical treatment of the governing equations is not yet feasible in the latter cases, because of the huge range of spatial and temporal scales needed to be resolved there \\citep{tobias_02,weiss_00,shukurov_02}. First attempts of global MHD simulations have been carried out by Gilman \\citep{gilman_81,gilman_83} adopting the Boussinesq approximation and later by Glatzmaier \\citep{glatzmaier_84,glatzmaier_85} using an anelastic model. Both aimed at modelling the solar dynamo. Though cyclic dynamo solutions were obtained in some cases, the models exhibited a wrong poleward migration of the magnetic field \\citep{glatzmaier_85}. In a recent attempt, \\citet{brun_04} succeeded in simulating a solar-like differential rotation, but could not reproduce basic features of the solar cycle yet.  An alternative approach is provided by mean-field electrodynamics \\citep{steenbeck_66,moffatt_78,raedler_80}, which is a theory focussing only on large scale, i.e. averaged fields. Highly complex small-scale or residual parts need not to be known in detail, only the averaged cross product of the residual velocity and magnetic field, in the following called the mean electromotive force, is relevant and accounts for the evolution of the mean field. Advantageously, the difficulties in resolving the small-scale structures can be avoided. Usually, the action of the small-scale velocity on the mean magnetic field as expressed by the mean electromotive force is parametrised, and the parameters are the mean-field coefficients defined below. Most prominent among them are those which describe the $\\alpha$-effect. They are closely related to the fundamental induction effect of cyclonic convection \\citep{parker_55}. Other mean-field coefficients describe the advection and diffusion of the mean magnetic field. In other words, mean-field theory supplies theoretical insight as well as formalised physical concepts in order to interpret and, in principle, also to quantify dynamo action.  Despite their relative simplicity, mean-field models have successfully reproduced basic features of the solar cycle \\citep[see e.g.][]{steenbeck_69,ossen_03} and are moreover unique in coherently simulating many features of the magnetic field in spiral galaxies \\citep{beck_96,shukurov_02}. But, whether or not mean-field models show dynamo action depends strongly on the mean-field coefficients. They are not known a priori but instead in general determined in a heuristic way.  Concerning the geodynamo, the situation is different. Many features of the Earth's magnetic field are successfully reproduced by nonlinear three-dimensional simulations of the magnetohydrodynamics in the Earth's core.  Although some model parameters still do not reach realistic values and in particular viscous effects and therefore the  size of viscous boundary layers are by orders of magnitude overestimated, the simulations exhibit a magnetic field dominated by an axial dipole at the Earth's surface that is maintained over several magnetic diffusion times \\citep{glatz_95b,kuang_97,christensen_98}. In addition, the time-dependence of the dipole moment, including secular variations, excursions and reversals, resembles the observed Earth's magnetic field \\citep{glatz_95a,kutzner_02,wicht_04}. The success in simulating the geodynamo can be attributed to the rather moderate vigour of turbulence in the Earth's outer core compared to the much more turbulent dynamics in the solar convection zone. The difference is formally expressed in terms of very different magnetic Reynolds numbers: While $\\Rm\\approx 10^3$ for the Earth's outer core \\citep{fearn_98}, a representative value of $\\Rm$ at the base of the solar convection zone is $10^{10}$ \\citep{ossen_03}.  Even though present geodynamo models are fully self-consistent, the interpretation of the mechanism by which they generate magnetic field relies frequently in a heuristic manner on mean-field concepts \\citep{glatz_95b,kageyama_97,olson_99}. Thus mean-field theory is very useful and indispensable at the present time. However, the applicability of mean-field concepts as tools for analysing dynamo processes in numerical simulations suffers from the poor knowledge of the mean-field coefficients and reliable methods to derive them.  There are two seminal approaches which have been followed in order to determine mean-field coefficients. The first one aims at deriving (quasi-)analytical expressions of the mean electromotive force and the mean-field coefficients \\citep[see e.g.][]{moffatt_78,raedler_80}. This requires the integration of the governing equation for the residual magnetic field with the help of closure methods. Most commonly used is the second-order correlation approximation (SOCA) in which only statistical moments up to the second order are taken into account, while moments of higher order are neglected. Early investigations with SOCA and simple assumptions on the turbulence, starting with the seminal paper by \\citet{steenbeck_66}, are summarized in \\citet{raedler_80b,raedler95,raedler00c}. More sophisticated assumptions on the turbulent motion let \\citet{kichatinov_92} and \\citet{ruediger_93} provide $\\alpha$-coefficients in the high conductivity limit for various rotation rates. In the context of the galactic dynamo, \\citet{ferriere_92,ferriere_93a,ferriere_93b} considered supernova explosions and derived expressions for the components of the $\\boldsymbol{\\alpha}$- and $\\boldsymbol{\\beta}$-tensors in cylindrical geometry. Steps beyond SOCA for specific flows were undertaken by \\citet{nicklausetal88}, \\citet{carvalho92,carvalho94} and \\citet{schmitt84,schmitt03}. In the context of the Karlsruhe dynamo experiment the velocity field possesses a number of symmetries, which allow a rather direct computation of mean-field coefficients \\citep{raedleretal97,raedleretal98,raedleretal02,raedleretal02b}. A particular approach beyond SOCA is the so-called $\\tau$-approximation, first introduced by \\citet{vainshteinetal83}, recently revived by \\citet{rogachevskiietal03}, \\citet{raedleretal03} and \\citet{brandenburgetal05,brandenburgetal05b}, and critically reviewed by \\citet{raedleretal06d}.  The second approach to the mean-field coefficients makes use of numerical modelling, in which the mean electromotive force, $\\emf$, as well as the mean-field, $\\bm$, are determined numerically as output of a MHD-simulation. Further on, a linear relation between $\\emf$ and $\\bm$ is assumed, which has to be inverted in order to solve for the unknown mean-field coefficients. A fundamental problem related to this approach is that the number of unknown variables is in general much higher than the number of equations resulting from the linear relation between $\\emf$ and $\\bm$. Therefore, all work presented so far refers either to specific situations in which certain constraints reduce the number of mean-field coefficients beforehand, or some of the mean-field coefficients are considered as small and are neglected. \\citet{ziegler_96} calculated the $\\boldsymbol{\\alpha}$-tensor for the galactic dynamo on the basis of numerical simulations of supernova explosions and confirmed the results given by \\citet{ferriere_93a}. Following a similar approach, \\citet{ossendrijver_01,ossendrijver_02} and \\citet{kaepylaeetal06} considered magnetoconvection in the solar convection zone and used box simulations to determine the $\\alpha$-tensor in Cartesian geometry. Similarly, \\citet{gieseckeetal05} derived a geodynamo $\\alpha$-effect from box simulations of rotating magnetoconvection. In a further attempt to determine not only the $\\boldsymbol{\\alpha}$- but the $\\boldsymbol{\\beta}$-tensor as well for accretion and galactic disks, \\citet{brandenburg_02} and \\citet{kowaletal06} applied box simulations of turbulence. \\citet{schrinner05} calculated all mean-field coefficients in a spherical shell with a convection pattern relevant for the geodynamo. The method and results are described in detail in this paper. For shorter contributions see also \\citet{schrinneretal05,schrinneretal06}.  The idea of this work is to take advantage of global numerical simulations of rotating magnetoconvection and the geodynamo and to compare them with respective mean-field calculations, where mean fields are defined by azimuthal averaging. This will lead to an estimation of the reliability of mean-field theory and its often used approximations. Furthermore, such a comparison will help to improve the conceptual understanding of dynamo mechanisms which are observed in the numerical simulations.  As already pointed out, both aims are intimately associated with the derivation of the corresponding mean-field coefficients. Hence, emphasis is placed on the developement of two methods which contribute to each of the principal approches mentioned above. Both methods have been applied to a simulation of rotating magnetoconvection and a quasi-stationary geodynamo. They are consistent with each other in a parameter regime in which the second-order correlation approximation is justified and serve as powerful tools to determine a number of relevant mean-field coefficients. While most of the quoted earlier work refers to a Cartesian  geometry, global mean-field coefficients for the astrophysically more relevant domain of a rotating spherical shell are presented here, and specific problems related to the spherical geometry are discussed.  The plan of the paper is as follows: In section 2 the equations, the boundary conditions and the parameters of the considered numerical models are given. Section 3 summarizes the mean-field concept and introduces the mean electromotive force. In section 4 two approaches to determine the mean-field coefficients are developed, a general numerical one and a semi-analytical approach using SOCA. In section 5 the mean-field coefficients for two models, rotating magnetoconvection and a quasi-steady geodynamo, are derived and their quenching as well as limitations of the validity of SOCA are discussed. A comparison of the mean fields for both models, derived from numerical simulations and from mean-field theory, respectively, is made in section 6, and the range of validity of the representation of the mean electromotive force are discussed. Section 7 summarizes our conclusions. In the appendices details of the derivation of the mean-field coefficients are given and the mean-field energy balance is considered. ",
        "conclusions": "\\label{sec:7}  The knowledge of the mean-field coefficients is decisive in order to analyse and to model dynamo action in many astrophysical bodies. In this paper, two approaches to determine mean-field coefficients have been developed. While the numerical approach does not use intrinsic approximations, the analytical approach is based on the second-order correlation approximation.  The mean-field view is applied to two examples: a simulation of rotating magnetoconvection and of a quasi-stationary geodynamo. In both examples similar processes take place: the mutual generation of poloidal and toroidal magnetic fields by an $\\alpha$-effect, flux expulsion from the dynamo region due to a $\\gamma$-effect, and a strong turbulent diffusion, which might be moderated by a $\\delta$-effect.  The calculation of mean-field coefficients provides insight into the reliability of frequently applied approximations in the framework of mean-field theory. Most dubious among them is the reduction of the $\\alpha$-tensor to an isotropic tensor, which leaves dominating non-diagonal components unconsidered. A further, important simplification is the second-order correlation approximation. It typically leads to overestimated amplitudes of mean-field coefficients, whereas their profiles are  rather unaffected.  The mean-field picture of geodynamo models is completed by the simulation of axisymmetric fields by means of a mean-field model, involving all mean-field coefficients determined. Test calculations with different sets of mean-field coefficients confirm their relation to the above mentioned dynamo processes. In  addition, a comparison with azimuthally averaged fields resulting from direct numerical simulations reveals their significance. In the magnetoconvection and the geodynamo example considered here, the match between direct numerical simulations and mean-field simulations is good only if a large number of mean-field coefficients are involved which contribute to $\\boldsymbol{\\alpha}$, $\\boldsymbol{\\gamma}$, $\\boldsymbol{\\beta}$ and $\\boldsymbol{\\delta}$. The application of corresponding mean-field coefficients derived in the second-order correlation approximation leads to similar results.  The reliability of mean-field models relies on a proper parametrisation of $\\emf$ in terms of the mean magnetic field. In the magnetoconvection example the traditional representation of the mean electromotive force considering no higher than first order derivatives of the mean magnetic field is valid. However, already in the geodynamo example it seems no longer to be justified and the assumption of scale separation is not fulfilled. This limits the applicability of some approximations commonly used within the mean-field theory. Nonetheless, even in the geodynamo example the spatial structure of the axisymmetric fields obtained by mean-field modelling corresponds roughly with that of the azimuthal averages extracted from the corresponding direct numerical simulation.  \\appendix"
    },
    "0609/astro-ph0609087_arXiv.txt": {
        "abstract": "We report detailed chemical abundance analysis of 27 RGB stars towards the Galactic bulge in Baade's Window for elements produced by massive stars: O, Na, Mg, Al, Si, Ca and Ti.  All of these elements are overabundant in the bulge relative to the disk, indicating that the bulge is enhanced in Type~II supernova ejecta and most likely formed more rapidly than the disk.  Our [Mg/Fe] ratios, confirmed by [Al/Mg], declines much more slowly with [Fe/H] than O, Si, Ca and Ti.  The [Mg/Fe] ratio stays above +0.25 dex up to well above solar metallicity.  We attribute the rapid decline of [O/Fe] to a metallicity-dependent modulation of the oxygen yield from massive stars, perhaps connected to the Wolf-Reyet phenomenon.  The explosive nucleosynthesis alphas, Si, Ca and Ti, relative to Fe, possess identical trends with [Fe/H], consistent with their putative common origin. We note that different behaviors of hydrostatic and explosive alpha elements can be seen in the stellar abundances of stars in Local Group dwarf galaxies. We also attribute the decline of Si,Ca and Ti, relative to Mg, to metallicity-dependent yields for these explosive alpha elements from Type~II supernovae.  The production of explosive alphas in Type~Ia supernovae likely explains the absence of obvious difference between Mg and Si, Ca, Ti trends in the Galactic thin disk.  An alternative explanation for the [Mg/$<$SiCaTi$>$] increase with metallicity is an excess population of 30--35M$_{\\odot}$ stars in the bulge that grew in importance with [Fe/H].  In this model the bulge formation timescale would have been significantly longer than our favored scenario of declining metallicity-dependent yields from Type~II supernovae.  The starkly smaller scatter of [$<$SiCaTi$>$/Fe] with [Fe/H] in the bulge, as compared to the halo, is consistent with expected efficient mixing for the bulge.  Since the metal-poor bulge [$<$SiCaTi$>$/Fe] ratios are higher than $\\sim$80\\% of the halo, the early bulge could not have formed from gas with the present-day halo composition.  If the bulge formed from halo gas, it occurred before $\\sim$80\\% of today's stellar halo; and the halo subsequently reduced its alpha/Fe ratios significantly.  The lack of overlap between the thick and thin disk composition with the bulge does not support the idea that the bulge was built by a thickening of the disk driven by the bar.  The trend of [Al/Fe] is very sensitive to the chemical evolution environment: a comparison of the bulge, disk and Sagittarius dwarf spheroidal galaxy shows a range of $\\sim$0.7 dex in [Al/Fe] at a given [Fe/H]; presumably due to a range of Type~II/Type~Ia supernova ratios in these systems. ",
        "introduction": "We present the second in a series of papers on the abundances of a sample of Galactic bulge K giants.  While our long term aim is to undertake a large scale survey of the Galactic bulge, the first step is to determine the correct stellar abundance scale.  This is straightforward for most stellar populations, but has been a subject of debate for bulge and metal rich stars.  Paper I \\citep{f06} reports our new iron abundance scale that supersedes \\citep{mr94}. Here we employ our new iron abundance scale to derive the abundances of alpha elements in the Galactic bulge stars.  In the third paper, we turn to the iron peak, r-, and s-process elements.  In future papers we will use the methods developed here to extend our studies to larger samples using multiobject spectroscopy.  The initial aim is to settle the iron abundance scale and then to follow with well determined abundances for a large number of additional elements.  When issues concerning the abundance determination of these additional elements have been settled, it will be possible to investigate large samples with confidence.  The understanding that it is possible for stellar populations to be both old and metal rich was gained surprisingly early.  In referring to the bulge of M31, \\citet{b63} states ``And the process of enrichment has taken very little time. ... So the CN giants that contribute most of the light in the nuclear region of the Nebula must also be called old stars; they are not young''.  The journey from insight to demonstration took nearly four decades, however.  Some significant steps along the way include the first integrated light spectroscopy of the bulge \\citep{w78}, surveys of bulge M giants \\citep{bb84} the first survey of bulge K giants \\citep{r88}, the first detailed abundances in the Galactic bulge (MR94), and the first demonstration that the bulge is as old as metal rich globular clusters \\citep{ort95}. Recent studies that confirm the old, metal rich nature of the bulge population \\citep{kr02,z03}, and most recently, the revised bulge abundance scale (Paper~I).  The events leading to the development of the new iron abundance scale are reviewed in Paper~I and in \\citet{r06}. Our new abundance scale is based on the use of weak iron lines in Arcturus (HR3905); we can use the very well-constrained parameters for Arcturus to perform a differential analysis, which effectively cancels out systematic problems with $gf$-values and stellar atmospheres.  The adopted lines are weak enough to remain on the linear portion of the curve of growth even for our most metal rich Galactic bulge giants.  The second advance is the use of multiple approaches to assess bulge membership and stellar parameters; this effort overcomes earlier problems with self-consistent abundance analysis that did not rely on photometry.  A significant thread in the abundance studies has been the discovery that alpha elements are enhanced in bulge giants.  This was demonstrated convincingly for the case of Mg and Ti (and to some extent Na and Al) in MR94.  Preliminary studies of high resolution spectra obtained at Keck, based on the MR94 iron abundance scale and line list, confirmed the Mg enhancement and found all alpha elements enhanced \\citep{rm00,mr04}.  It is widely accepted that at least Mg is enhanced in extragalactic spheroidal populations \\citep{wfg92}.  The modern effort to revisit Whitford's (1978) spectroscopy of the integrated light of a Galactic bulge field compared to external galaxies \\citep{puzia2002} supports moderate alpha enhancement in the bulge and shows that the bulge falls on the lower end of the elliptical galaxy sequence.  The long standing operating paradigm is that alpha elements (e.g. O, Mg, Si, Ca) are produced in the supernovae of short lived massive stars while iron is predominantly produced in core collapse SNe that occur on a timescale that is 1-2 orders of magnitude longer \\citep{t79,w89,t95,mcw97}. \\citet{m90} combined the SN yields and sophisticated chemical evolution models.  \\citet{m99} employ the SN yields of \\citet{tnh96} for Type Ia and \\citet{ww95} for Type II SNe. The yields from numerical models of SNe, along with corroborating abundance determinations in very metal poor stars, are the theoretical underpinning that support this paradigm.  The models of \\citet{m90}, \\citet{m99}, and \\citet{m06} incorporate star formation predictions and the requirement that the bulge form rapidly ($<1$ Gyr).  The alpha enhancement has been known since MR94 and is the constraint (along with the observations of an old turnoff age) that underpin the empirical arguments for a very old Galactic bulge.   The growing evidence that the Galactic bulge formed early and rapidly finds support in numerous observations at high redshift.  The discovery of the Lyman break selected high redshift galaxy population by \\citet{st96} was the first direct identification of a population of galaxies at $z>3$ with high star formation rates and the potential (in terms of mass) to evolve into galaxies like the Milky Way.  Since that discovery there have been numerous other high redshift populations discovered that are equally plausible bulge progenitors; all share the property of having high rates of star formation and large stellar mass.  \\citet{st96} and \\citet{adel03} have also detected and studied metal enriched wind outflows from Lyman Break galaxies.  \\citet{pett02} find winds and derive alpha enhancements in a gravitationally lensed Lyman break galaxy observed at high spectral resolution.  These offer additional evidence of high star formation and also indicate that winds should eventually be incorporated into the chemical evolution models.  However, the evolutionary path of these high redshift galaxies cannot unequivocally link them to a Milky Way like result in the present day Universe.  The union of these ideas--that there is overwhelming evidence for a rapid formation timescale for the bulge both from the chemical enrichment, turnoff age, and high redshift perspectives, is described in \\citet{r06}. If the bulge has undergone secular evolution \\citep{kk04} then it has done so (it would seem) mostly in the dynamical sense, with its population having been built early and rapidly.  However, in the case of the bulge, as opposed to giant ellipticals, there are good reasons to suspect that the starburst formation scenario \\citep{elm99} is not the complete picture. \\citet{imm04} explore models in which the bulge forms from the accumulation of smaller units (as might be favored in the LCDM cosmology scenario). It may be possible to detect fossil evidence of these subclumps by correlations between abundance and kinematics, or differing element trends (e.g. [Mg/Fe] vs [Fe/H] at a different level).  A large scale survey program to study these trends will require significant samples of stars with abundances determined to very high precision.  Contamination by the thick disk and foreground disk could mimic the expected signal from a subclump.  Further, it should be noted that the bar, a rapid formation timescale, and high metallicity all point toward a formation model in which the mass is in place early; nonetheless, scenarios of inhomogeneous evolution should be explored with larger surveys.  Evidence for the presence of a bar structure in the bulge is now indisputable.  Observations by the {\\sl COBE} satellite and subsequent modeling \\citep{dw95} argue for the bar based on surface photometry. \\citet{mao} find evidence for the bar from the stellar population and from stellar dynamics \\citep{zhao94}.  A recent study \\citep{soto06} finds that kinematics consistent with the bar set in at [Fe/H]$>-0.2$, with more metal poor stars being more consistent with an isotropic velocity ellipsoid.  Bars are also thought to be formed by secular evolution \\citep{kk04}. There has been extensive theoretical discussion about the origin of the bar based on dynamical processes that are secular in nature \\citep{c01}. There is prima facie evidence for ongoing star formation in the nuclear region and modeling of the luminosity function in fields within 50~pc of the nucleus requires a constant star formation rate, with models dominated by early bursts of formation are ruled out by models of the luminosity function \\cite{figer04}.  A study of infrared abundances in old stars in the bulge by \\citet{ram00} finds no abundance gradient from latitude 0$^{\\rm o}$ to -$6^{\\rm o}$, an observation expected if secular evolution is an important process in establishing the vertical thickness of the bar. At present, there is strong evidence from the turnoff age and chemical evolution that the bulge formed early and rapidly.  On the other hand, the bulge has been shown to have a clear bar structure and other properties (no vertical abundance gradient) more consistent with secular evolution.  There is at present a conflicted situation with regards to age and formation constraints from the turnoff and chemistry, and possible secular evolution from the presence of the bar structure. We believe that it is vital to not only measure abundances of alpha elements in a larger sample, but to also measure the abundances of more species, so as to provide better constraints for the chemical evolution models. We have the opportunity in the bulge to track the abundance trends of many elements in detail.  This is something that we will be unlikely to achieve for any other spheroidal population in the near future.  This approach can be used to ask more detailed questions, such as constraints on the initial mass function as well as setting limits on the timescale of the formation epoch to greater precision.  The presence or absence of vertical gradients in the alphas and [Fe/H] gives an additional important constraint on whether secular processes have played a role in the formation of the bulge.  We also undertake a detailed comparison between the bulge and thick disk/halo.  These well studied populations serve both as test samples to confirm the soundness of our abundance analysis and as a comparison sample against which the bulge abundance patterns can be compared.  This will enable us to relate our findings in the bulge to the formation of well studied populations in the Milky Way.  The paper is organized as follows.  In Section~2, we discuss the spectral data and the creation of a new unblended line list for the elements under study.  We review the stellar parameters adopted for these stars and their uncertainties in Section 3.  In Section~4, we describe the uncertainties of our abundance determinations, including a full quantitative analysis method. In Section~5, we apply a simple correction for the varying alpha-enhancements of our target stars.  In Section~6, we check our abundance results, showing that our adopted parameters return abundances in ionization equilibrium and in agreement with previous literature studies.  Our primary bulge abundance results are presented in Section~7, which includes the evidence for enhanced alpha element ratios at all metallicities in the bulge.  Further discussion of the abundances of Na and Al are in Section~8, including evidence that the Type II SNe that enriched the bulge included some metal-rich stars. In Section~9, we find that the metal-poor stars in the bulge show the same Na-O anticorrelation as seen in metal-poor globular clusters and discuss what that may mean for the early formation of the bulge. ",
        "conclusions": "We have performed a detailed chemical abundance analysis of 27 RGB stars towards the Galactic bulge in Baade's Window, based on Keck/HIRES echelle spectra.  In this paper we focus on light elements thought to be produced by massive stars:  the alpha elements O, Mg, Si, Ca and Ti, and also Al and Na.  We employed a differential analysis, line by line, relative to the high resolution, high S/N, spectral atlas of the red giant Arcturus by \\citet{hw00}.  We used only unblended lines in the equivalent width abundance analysis; this list should be useful for future chemical abundance studies of red giants.  The advantage of differential analysis is that systematic errors, due to various inadequacies of the model atmospheres, cancel out when performed relative to similar stars; certainly the oscillator strengths of the atomic lines do not enter into the abundance determinations. In order to relate differential abundances relative to Arcturus to the solar differential scale we have determined the composition of Arcturus, relative to the Sun, using the \\citet{kfb84} solar atlas and the \\citet{hw00} Arcturus atlas, which have superior wavelength coverage, S/N, and resolution than our science spectra.  Our abundance results for Arcturus show that its chemical abundance ratios are completely consistent with the composition of thick disk stars; the membership of Arcturus in the thick disk population is also supported by its kinematic parameters.  In addition the bulge spectra, we analyse 17 disk giants that have been observed at slightly lower resolution using other telescopes. These local objects enable a literature comparison of the abundance results with large disk dwarf surveys.  The comparison confirms that the abundance results for our giants are in good general agreement with published dwarf surveys; but it uncovered possible small ($< 0.1$ dex) systematic zero-point differences for Al and Si.  The stellar model atmosphere parameters we use are derived from the analysis method performed in Paper~I of this series.  The analysis relies on V$-$K photometry, differential Fe excitation temperatures and differential Fe ionization temperatures.  Plots showing the ionization equilibrium of Fe and Ti indicate a small, $\\sim$0.05 dex difference between Fe~I and Fe~II, and no obvious difference between Ti~I and Ti~II; although the titanium results were affected by more scatter than those for iron.  The ionization plots show that any non-LTE effects are the same, to within $\\sim$0.05 dex, across our entire sample of disk and bulge giant stars, which cover a range in [Fe/H] of 2 dex and T$_{\\rm eff}$ by 1000K.  Our main result is that all five alpha elements, plus Na and Al, are enhanced in the bulge stars relative to the Galactic thin and thick disks; oxygen, however, may be enhanced only marginally relative to the thin disk. These elements are thought to be produced in mostly massive stars, either during hydrostatic burning phases or during explosive nucleosynthesis in Type II SNe events.  Consequently, our abundance results indicate that massive stars contributed more to the chemical enrichment of the bulge than to the disk.  The enhancement of these massive star products indicates that the ratio of Type~II/Type~Ia supernova was higher in the bulge than in the disk. \\citet{t79} propsed what is now the established paradigm, that the declining trend of [O/Fe] with [Fe/H] observed from the halo to the solar neighborhood was due to the late addition of Fe, but not O, from Type~Ia SNe, due to the longer progenitor lifetimes of Type~Ia SNe. These enhancements of O and other alpha elements relative to iron are often associated with shorter formation timescales, and, thus, we conclude that the bulge formed faster than the disk. Chemical evolution models of the abundances of MR94 (whose results are similar to those found here) by \\citet{m99}, \\citet{fws03}% and \\citet{m06} find a $\\sim 500$~Myr (rapid) timescale for the chemical enrichment of the bulge as a central conclusion. While we believe that a short formation timescale for the bulge is the most attractive explanation for our findings, we cannot ignore other possible scenarios (IMF skewed to high mass, lower binary fraction in the bulge reducing the relative contribution Type~Ia SN ejecta, etc.).  For this reason the search for abundance patterns in the bulge that are characteristic of long-lived stars (e.g. low-mass AGB stars) will provide very useful constraints on the bulge formation timescale.  Overall, we find that Mg is strikingly enhanced relative the thick and thin disk populations, and that the trend of [Mg/Fe] with repsect to [Fe/H] to well above solar metallicity.  The strong enhancement of Mg with respect to the other alpha extralactic Mg enahncements; Mg has long been considered a proxy for the measurement of alphas in galaxies, and this view may have to change. Indications of relatively low Ca abundances \\citep{tmb03} in external galaxies may reflect a general dichotomy between Mg and the rest of the alphas.  Future studies should explore these findings.  Al and to some extent Na follow the same trend as Mg, but O and the other alphas are only marginally enhanced relative to the thick disk, and follow declining trends with [Fe/H].   This result was hinted at in \\citet{rm00} and in \\citet{mr04}.  But we see these trends clearly in our dataset now.  Our abundance trends indicate that the bulge is chemically distinct from the thick and thin disk stars in the Solar vicinity.  This does not rule out some kind of secular evolution \\citep{pn90} from a primordial massive thick disk. Because they have similar chemical compositions below [Fe/H]$\\sim -0.9$, it is possible that the early bulge gas could have come from a primordial thick disk. However, the bulge and thin disk chemistry is so disjoint that no relationship seems possible between those populations.  When we intercompare indivdiual alpha element raios we find our second major result, that the explosive nucleosynthesis alphas (Si, Ca, and Ti) have very similar trends with [Fe/H], but O and Mg (produced during hydrostatic burning) have unique trends.  We note that Venn et al.'s (2004) review of dwarf galaxy abundances also reports differing trends for the explosive and hydrostatic alphas (in that case, Mg vs Ca and Ti).  We find that [O/Fe] declines significantly more steeply than the other alphas for [Fe/H]$\\geq -0.05$ dex.  We suggest that this is obsevational evidence of the reduced oxygen yields due to stellar winds and the Wolf-Rayet phenomenon that was predicted by Maeder (1980, 1991).  While MR04 proposed this idea based on analysis of 8 bulge stars from this sample, we believe that the oxygen trend in the present data requires that this idea be considered seriously.  Maeder's predictions suggest that the decline in oxygen yield should e accompanied by an increase int he yield of carbon; that would provide a test of this idea. Another possible test comes from the prediction of \\citet{ma00} who propose that Wolf-Rayet winds are significant source of flourine (although this is debated; see Placis et al. 2005).   Abundances of C and F in the Galactic bulge stars could provide a crucial test of the wind hypothesis.  While the bulge trend of [O/Fe] with [Fe/H] is more steep than that of the exposive alphas, the rend of [Mg/Fe] with metallicity is quite shallow, with [Mg/Fe]$=+0.3$ even for the the most metal-rich bulge stars.  We consider that the bulge Mg trend poses a significant challenge for understanding nucleosynthesis by Type II SNe and the chemical evolution of the Galactic bulge.  Given the importance of this question, we have been very careful to check this finding (even though it is well known that Mg is enhanced in E and S0 galaxies). However, the agreement of the locak disk giant [Mg/Fe] values with the dwarf studies and the internal consistency of our abundances from several Mg lines makes it very unlikely that our results are in error.  Furthermore, the bulge [Al/Mg] trend, which is in excellent agreement with that seen for the thick and think disks, removes any lingering doubt about our measured [Mg/Fe] ratios.  To explain the unusual [Mg/$<$SiCaTi$>$] trend in the bulge compared with the disk, either more Mg was produced by sources present in the bulge and not in the disk, or there are sources of Si, Ca, and Ti in the disk that are less common or absent in the bulge.  If excess Mg is produced in the bulge, there must be an accompanying source of Al that exactly matches the amount needed to maintain the [Al/Mg] trend with [Fe/H] seen in the thick and thin disks. If there is an additional source of Mg and Al, it is most unlikely to arise in Type~Ia SNe (see Iwamoto et al. 1999); thus the unusual composition of the bulge must rest with the Type~II SNe.  Considering the dominance of the products of massive star SNe in the bulge relative to the disk, we believe that the bulge preserves the true metallicity dependent yield of hydrostatic (e.g. O, Mg, Al) versus explosive (Si Ca, Ti) elements.  We suggest that the [Mg/SiCaTi] trend in the bulge is due to Mg from Type~II SNe, combined with a metallicity dependent decline in the yields of Si, Ca, and Ti.  Unfortunately, the predicted yields from WW95 are contrary to the required enhancement of Mg relative to Si, Ca, and Ti at high metallicity.  It is significant that WW95 did not include metal dependent winds in their models; these winds result in significant mass loss for metal-rich massive stars and deplete the star of mass that would contribute to the hydrostatic alphas.  If mass loss in metal-rich massive stars explains the Mg vs O disparity, then we may ask why differences between Mg, Si, Ca, and Ti are not so evident in the disk.   The solution is to resort to the idea that in the disk the Type~Ia SNe produce significant amounts of Si, Ca, and Ti; this is expected based on the predictions of \\citet{iwa99}.  We do not favor the possibility that massive Type~II SNe progenitors eject Mg more efficiently at high metallicity, making them the source of metal-dependent Mg enhancements in the bulge.  This would require that Type~Ia SNe in the disk produce large amounts of Mg (and Al) contrary to theory.  An alternative scenario that might reproduce our unusual [Mg/SiCaTi] trend would posit a massive star IMF skewed to more massive stars as [Fe/H] increases.  A related alternative requires a sub-population of massive stars in the bulge that grows in importance with metallicity.  The difficulties presented by this scenario are the non-universal IMF, the unknown reason for the metal-dpeendent IMF, and the question of why the bulge abundance patterns are seen nowhere else.  Given the difficulties with the second scenario we favor the former idea: that the bulge [Mg/$<$SiCaTiI$>$] trend is due to a metallicity-dependent decline in the yield of Si, Ca and Ti from very massive Type~II SN progenitors, most likely driven by stellar winds.  The mechanism has the advantage that it presents fewer discrepancies with the current understanding of chemical evolution.  Another major finding is that the trend of [$<$SiCaTiI$>$/Fe] with [Fe/H] in the Galactic bulge has a remarkably small dispersion, $\\sim$0.05 dex, in stark contrast to the scatter of this ratio in Galactic halo stars. The result indicates that the bulge evolved with a high degree of homogeneity, suggesting that either efficient mixing processes occurred during bulge evolution, or that large numbers of SN events completely sampled the nucleosynthesis yield function in all bulge locations; this is not surprising, given the current bulge infall time of $\\sim$10$^6$ years. Our abundance results underscore the inhomogeneous evolution of the Galactic halo as evidence by its abundance pattern, consistent with ideas of accretion and low density chemical evolution.  The bulge evolution with efficient mixing, indicated by the high degree of homogeneity in [$<$SiCaTiI$>$/Fe] ratios suggests the presence of energetic processes; combined with the enhancement of nucleosynthesis products from massive stars, this might reasonably be expected to favor the so-called ELS scenario \\citep{els62}, in which the bulge formed early and rapidly through violent relaxation.  We also found that the [$<$SiCaTiI$>$/Fe] trend in the metal-poor bulge stars lies at a value equal to the maximum seen in the halo, exceeding the mean in the halo by $\\sim$0.13 dex; well in excess of the estimated maximum systematic uncertainty of 0.04 dex. Indeed the metal-poor bulge overlaps with the top $\\sim$20\\% of halo [$<$SiCaTiI$>$/Fe] ratios. This indicates that the metal-poor bulge could not have formed out of Galactic halo gas with the present day halo composition.  However, it is possible for the metal-poor bulge to have formed from halo gas, as suggested by \\citet{wg92}, if the halo composition at the onset of bulge formation was different than today, characterized by high [$<$SiCaTiI$>$/Fe] ratios.  For this to have occurred the mass in halo stars must be greater today than when the metal-poor bulge was formed, and the onset of bulge formation must have occurred very early in Galactic history, at roughly the same time as the halo.  Our plot comparing [Al/Fe] versus [Fe/H] in the Sgr dSph galaxy with the Galactic bulge and thin disk shows that [Al/Fe] depends very sensitively on environmental parameters.  Therefore, it should provide a useful probe for chemical evolution.  Further studies of the Al abundances in many different populations and environments should be undertaken, both to exploit this finding and to understand its underlying origin.  Compared to the original study of McWilliam \\& Rich (1994) we have increased our resolution by roughly a factor of four, and our S/N by a factor of 3. Thanks to the Keck telescope and the HIRES spectrograph, abundance analysis of Galactic bulge giants can be done on data that is as good as those avaiable for the solar vicinity and globular cluster stars.  We have developed new approaches that place the iron abundance scale on a far more firm footing than before, and undertake a differential analysis relative to Arcturus that to a great extent, mitigates systematic errors associated with the imperfect nature of the stellar atmosphere modeling.  With these improvements, we have confirmed largely the results of MR94, finding that the bulge has strong enanchements of Mg, Al, and Na and lesser, but significant enhancements of the other alphas. However, we find important new results.  Oxygen and Mg have disjoint abundance trends, perhaps requiring that winds in the most metal-rich massive stars deplete the hydrostatic burning zones of mass.  The explosive alphas follow a different trend from Mg, Na, and Al.  Consideration of the explosive alphas shows that the bulge has a far more homogeneous composition than the halo, extending even to the most metal-poor bulge members.  The notion of a bulge that formed early, and rapidly, remains attractive.  This conclusion is consistent with observations of galaxy formation at high redshift.  However, it remains vital to study additional elements and larger samples of stars, because the fossil record locked in the bulge's composition has the potential to provide a detailed record of its history that is unmatched by anything that can be inferred from the observations of distant galaxies."
    },
    "0609/astro-ph0609564_arXiv.txt": {
        "abstract": "We study the hydrodynamic evolution of cosmological background neutrinos. By using a spherically symmetric Newtonian hydrodynamic code, we calculate the time evolution of the density profiles of neutrino matter in cluster and galactic scales. We discuss the possible observational consequences of such evolution and the resulting density profiles of the degenerate neutrino `stars' in galaxies and clusters. ",
        "introduction": "The observations of neutrino oscillations indicate that at least one flavour of neutrinos have non-zero rest mass, and at least one type of sterile neutrinos with mass of eV order must be taken into account for the three scales of mass squared differences if the LSND experiment is confirmed \\cite{Oscillation}. A huge number of primordial neutrinos were produced in the early universe \\cite{Peebles}, \\cite{Dodelson}, and they were cooled down as the universe expanded, decoupling from baryons much earlier than the photon decoupling epoch. The temperature of the neutrinos dropped to below 0.1~meV after several billions of years, and so if their rest mass $m_{\\nu}$ is an eV or so, the neutrinos have become non-relativistic. Sterile neutrinos may be produced in the early universe by the oscillations of the active neutrinos \\cite{Dolgov}. These sterile neutrinos may have rest mass greater than 10 eV and may play a role in structure formation. Also, a neutrino rest mass $m_{\\nu}$ of the eV to keV order can account for most of the dark matter, which dominates the mass in galaxies and clusters \\cite{Top}. In this article, we study the gravitational collapse of such non-relativisitc cosmic background active and sterile neutrinos, and we discuss the formation of degenerate neutrino dark halos in cluster and galactic scales. It has been proposed that such degenerate neutrino `stars' can be found in centers of galaxies \\cite{Viollier}, \\cite{Apparao}. ",
        "conclusions": "We have demonstrated the possibility that the neutrinos can form structures in galaxies and clusters. Including the formation time, our model is consistent with the observation but not the top-down scenario. Our hydrodynamical simulation results (see Fig.~4) agree with recent observed data indicating that the central density of dark matter profile is constant \\cite{Tyson}.  Light neutrinos (1 eV) cannot form structure by their own gravitational attraction. However they can still form structures under the gravitational field of a heavier component (see Fig.~2 and 3). It is well known that light primordial neutrinos (eV order) cannot account for all the dark matter as $\\Omega_{\\nu}$ is much less than 0.3 for $m_{\\nu}=1$ eV \\cite{Hu}. There must exist heavier particles such as sterile neutrinos which dominate the mass of the dark matter and provide such a gravitational field in the clusters. These sterile neutrinos (10 eV order) may also form smaller structures in galaxies (about 10-20 kpc), and all these smaller structures link up together to form a large structure in clusters. Therefore, neutrinos with rest mass 10 eV or above are important in structure formation, and we should not neglect their contribution. The range of possible masses of active neutrinos and sterile neutrinos can both be constrained by more data from the distribution of hot gas in clusters and the rotation curves of galaxies.  From Fig.~3 and 7, we can see that the central density of a neutrino star oscillates, which may affect the hot gas in clusters. It may force oscillate the hot gas, which may be observable. Also, the oscillations in the hot gas may convert the kinetic energy into internal energy of the hot gas and thus heat them up, which may at least partially account for the heat source of the hot gas."
    },
    "0609/astro-ph0609422_arXiv.txt": {
        "abstract": "%  I discuss the role of short-duration eruptive mass loss in the evolution of very massive stars.  Giant eruptions of Luminous Blue Variables (LBVs) like the 19th century event of $\\eta$ Carinae can remove large quantities of mass almost instantaneously, making them significant in stellar evolution.  They can potentially remove more mass from the star than line-driven winds, especially if winds are clumped such that O star mass-loss rates need to be revised downward. When seen in other galaxies as ``supernova impostors'', these eruptions typically last for less than a decade, and they can remove of order 10 M$_{\\odot}$ as indicated by massive nebulae around LBVs. Such extreme mass-loss rates cannot be driven by radiation pressure on spectral lines.  Instead, these outbursts must either be continuum-driven super-Eddington winds or outright hydrodynamic explosions, both of which are insensitive to metallicity.  As such, this eruptive mode of mass loss could play a pivotal role for massive metal-poor stars in the early universe. ",
        "introduction": "The main question I wish to address here is whether the majority of mass lost during the lifetime of the most massive stars occurs primarily via steady line-driven stellar winds, or instead through violent, short-duration eruptions.  The two extremes are shown graphically in Figure 1.  This is critical for understanding how $\\dot{M}$ scales with metallicity.  In this contribution I draw attention to the role of LBV eruptions, advocating for their importance.  The essential points of the argument, already made by Smith \\& Owocki (2006), are the following:  \\begin{itemize}  \\item Recent studies of hot star winds indicate that MS mass-loss rates are lower than previously thought due to the effects of clumping (Fullerton et al.\\ 2006; Bouret et al.\\ 2005; Puls et al.\\ 2006; Crowther et al.\\ 2002; Hillier et al.\\ 2003; Massa et al.\\ 2003; Evans et al.\\ 2004).  Revised $\\dot{M}$ values are inadequate to reduce the star's mass enough to reach the WR phase.  \\item Observations of nebulae around LBVs and LBV candidates have revealed very high ejecta masses -- of order 10 M$_{\\odot}$ (Fig.\\ 2). Some objects show evidence for multiple shell ejections on timescales of 10$^3$ years.  These eruptions could remove a large fraction of the total mass of the star.  \\item The extreme mass-loss rates of these LBV bursts imply that line opacity is too saturated to drive them, so they must instead be either continuum-driven super-Eddington winds (see Owocki et al.\\ 2004) or outright hydrodynamic explosions.  Unlike steady winds driven by lines, the driving in these eruptions may be largely independent of metallicity, and might play a role in the mass loss of massive metal-poor stars.  \\end{itemize}  For reasons that may be obvious, my talk at this meeting and the original paper (Smith \\& Owocki 2006) were deliberately provocative. Rather than repeat that discussion, I will briefly elaborate on a few of these issues, and will consider alternatives and further implications. ",
        "conclusions": ""
    },
    "0609/astro-ph0609108_arXiv.txt": {
        "abstract": "The relation between the Evershed flow and moving magnetic features (MMFs) is studied using high-cadence, simultaneous spectropolarimetric measurements of a sunspot in visible (630.2~nm) and near-infrared (1565~nm) lines. Doppler velocities, magnetograms, and total linear polarization maps are calculated from the observed Stokes profiles. We follow the temporal evolution of two Evershed clouds that move radially outward along the same penumbral filament. Eventually, the clouds cross the visible border of the spot and enter the moat region, where they become MMFs. The flux patch farther from the sunspot has the same polarity of the spot, while the MMF closer to it has opposite polarity and exhibits abnormal circular polarization profiles. Our results provide strong evidence that at least some MMFs are the continuation of the penumbral Evershed flow into the moat. This, in turn, suggests that MMFs are magnetically connected to sunspots. ",
        "introduction": "\\label{sec:intro}  Moving magnetic features (MMFs) are small flux concentrations observed in the moat surrounding sunspots. They were discovered by Sheeley (1969). Not surprisingly, a connection between MMFs and the magnetic field of sunspots was proposed a few years later by Harvey \\& Harvey (1973). The phenomenological properties of MMFs are relatively well know by now (Shine \\& Title 2001), but their origin and magnetic structure remain a mystery. The situation is more favorable in the case of sunspots (Solanki 2003). According to current views, sunspot penumbrae are formed by at least two magnetic components with different field inclinations. The more horizontal component drives the Evershed flow from the inner to the outer penumbra, where it partly dives down below the photosphere or continues into the canopy. The Evershed flow is not steady.  Shine et al.\\ (1994), Rimmele (1994), and Rouppe van der Voort (2003) have demonstrated the existence of velocity packages, called ''Evershed clouds'', that propagate radially outward and show an irregular repetitive behavior on a timescale of the order of 10--15~minutes. Evershed clouds evolve and remain coherent until they go through the outer penumbral border, where they seem to vanish.  Some investigators have pointed out the possible link between MMFs and the penumbral component carrying the Evershed flow (Mart\\'{\\i}nez Pillet 2002; Schlichenmaier 2002; Thomas et al.\\ 2002). In fact, recent observations seem to support this idea. Using Michelson Doppler Imager magnetograms, Sainz Dalda \\& Mart\\'{\\i}nez Pillet (2005) discovered magnetic filaments extending from the midpenumbra into the moat. Several bipolar MMFs were observed to leave the sunspot following the paths traced by these filaments. Zhang et al.\\ (2003) and Penn et al.\\ (2006) also detected the passage of MMFs through the outer penumbral boundary and their subsequent evolution across the moat. Bonet et al.\\ (2004) applied a phase diversity restoration to G-band filtergrams of a sunspot. They concluded that most G-band bright points, which are associated with MMFs (Beck et al.\\ 2006), are born close to the continuation of dark penumbral filaments. Hence, they pointed out the possible relationship between G-band bright points (or MMFs) and the more horizontal component of the penumbra. Finally, Kubo (2005) analyzed Advanced Stokes Polarimeter measurements. He found significant correspondence between the magnetic field of MMFs and the penumbral uncombed structure.  In this Letter, the connection between the Evershed flow and MMFs is established. Using high resolution spectropolarimetric measurements in two different spectral ranges, we track the evolution of Evershed clouds from the midpenumbra to well outside the outer penumbral boundary, demonstrating that they become MMFs once they leave the sunspot. ",
        "conclusions": "\\label{sec:con}    \\begin{figure*} \\centering \\epsscale{1.15} \\plotone{f4.eps} \\caption{Cospatial Stokes $I$ ({\\em top}) and $V$ ({\\em bottom}) profiles of \\ion{Fe}{1} 630.15, \\ion{Fe}{1} 630.25, \\ion{Fe}{1} 1564.85, and \\ion{Fe}{1} 1565.2 emerging from a pixel inside feature A with abnormal circular polarization signals.  The vertical solid lines represent the rest wavelengths of the spectral lines.} \\label{fig:fig4} \\end{figure*}  We have used simultaneous spectropolarimetric measurements in two spectral regions to demonstrate the relation between the penumbral Evershed flow and moving magnetic features (MMFs). Our observations combine high cadence (3.9 minutes) with one of the highest spatial resolutions achieved in solar spectropolarimetry ($\\sim 0\\farcs7$).  One after the other, two Evershed clouds are seen to move radially outward along a penumbral filament characterized by larger field inclinations to the LOS than its surroundings. They appear in the center-side penumbra, exhibiting blueshifts and positive polarity (i.e., the polarity of the spot). Eventually, the clouds cross the outer penumbral boundary and become MMFs. Once in the moat region, the flux patch closer to the spot shows negative polarity. The two MMFs are separated by about 2\\arcsec and move away from the spot with an average speed of 0.5~km~s$^{-1}$. We have examples of other Evershed clouds crossing the outer penumbral boundary and becoming MMFs, so the process just described may be quite common in sunspots.  Our observations demonstrate that at least some MMFs are a manifestation of the penumbral Evershed flow in the sunspot moat. This, in turn, gives support to earlier results about the magnetic connection between MMFs and the penumbra (e.g., Zhang et al.\\ 2003; Sainz Dalda \\& Mart\\'{\\i}nez Pillet 2005). The association between Evershed clouds and MMFs found here suggests that the physical structures responsible for the existence of the former inside the penumbra are also present in, or propagate to, the sunspot moat. This may help refine theoretical models of sunspot penumbrae.  We have also shown, for the first time, that abnormal circular polarization profiles with three or four lobes occur in MMFs. The presence of abnormal profiles hampers direct interpretations of magnetogram polarities in terms of magnetic field inclinations. As a result, we do not know if the vector magnetic field is really pointing downward in the negative-polarity MMF and upward in the the positive-polarity MMF, or if the field azimuth provides any evidence of magnetic connection between the two flux concentrations. To investigate these issues, our efforts will concentrate on the inversion of the observed Stokes spectra."
    },
    "0609/astro-ph0609436_arXiv.txt": {
        "abstract": "We analyze a volume-limited sample of massive bulge-dominated  galaxies with data from both the Sloan Digital Sky Survey and the Galaxy Evolution Explorer (GALEX) satellite. The galaxies have central velocity dispersions greater than 100 km/s and  stellar surface  mass densities that lie above the value where galaxies transition from actively star forming to passive systems. The sample is limited to  redshifts $0.03 < z < 0.07$. At these distances, the SDSS spectra sample the light from the bulge-dominated central regions of the galaxies. The GALEX  NUV data provide high sensitivity to low rates of {\\em global} star formation in these systems. Our sample of bulge-dominated  galaxies exhibits a much larger dispersion in NUV-$r$ colour than in optical $g-r$ colour. The  dispersion increases for  galaxies with smaller central velocity dispersions and nearly all of the galaxies with bluer NUV-$r$ colours are AGN. Both  GALEX images  and  SDSS colour profiles demonstrate that the excess UV light is nearly always associated with an extended disk. When comparing fibre-based estimates of stellar age to global ones, we find that galaxies with red outer regions almost never have a young bulge or a strong AGN. Galaxies with blue outer regions have bulges and black holes that span a wide range in age and accretion rate. Galaxies with young bulges and strongly accreting black holes almost {\\em always} have blue outer disks.  The black hole growth rate correlates much more strongly with the age of the stars in the bulge than in the disk. Our suggested scenario is one in which the source of gas that builds the bulge and black hole is a low mass reservoir of cold gas in the disk. The presence of this gas is a necessary, but not sufficient condition for bulge and black hole growth. Some mechanism must transport this gas inwards in a time variable way. The disk gas itself is likely to be the result of the accretion of gas from an external source. As the gas in the disk is converted into stars, galaxies will turn red, but  further inflow can  bring them back into the blue NUV-$r$ sequence. ",
        "introduction": "\\subsection{Ultraviolet Data}  Since its launch in April, 2003, GALEX has been conducting several surveys of the UV sky. Details on the GALEX mission and surveys are given in Martin et al (2005) . In this paper we make use of the GALEX Medium Deep Survey (MIS). MIS exposures are typically 1500 seconds long and reach a limiting magnitude $m_{AB} \\sim 23$. Our data are taken from the GALEX MIS IR1.1(+GR1) data release. The GALEX data include far-ultraviolet (FUV; $\\lambda_{eff}=1528$~\\AA, $\\Delta\\lambda=268$~\\AA) and near-ultraviolet (NUV; $\\lambda_{eff}=2271$~\\AA, $\\Delta\\lambda=732$~\\AA) images with a circular field of view with radius $\\sim 38$ arcminutes. The spatial resolution is $\\sim 5$ arcseconds. Details of the GALEX satellite and data characteristics can be found in Morrisey et al (2005). The data were processed through the GALEX reduction pipeline at the California Institute of Technology. The pipeline reduces the data and automatically detects, measures, and produces catalogs of FUV and NUV fluxes for sources in the GALEX images.   \\subsection{Optical data}  The GALEX catalogs were then matched to the SDSS Fourth Data Release (DR4; Adelman-McCarthy et al. 2006) spectroscopic sample. The matching procedure is described in Seibert et al (2005). Matching to SDSS provides a variety of photometric parameters in the $u,b,v,r,i,z$ photometric passbands as well as spectroscopic redshifts for a sample of 51,246 galaxies with UV detections. In this analysis, we  exclude galaxies located more than 30 arcminutes from the MIS field centers , because the photometry is not as reliable near the boundaries.  In addition to the SDSS photometric parameters, a large number of galaxy parameters derived from the SDSS spectra are available in the value-added catalogs at the SDSS website at the Max Planck Institute (http://www.mpa-garching.mpg.de/SDSS). From these catalogs we use the stellar masses, stellar velocity dispersions, the spectral indices D$_n$(4000) and H$\\delta_A$, the equivalent width of the H$\\alpha$ emission line (EQW(H$\\alpha$)), and the AGN classifiers. For detailed information about the derivation of these quantities , see Kauffmann et al (2003a,c); Brinchmann et al (2004) and Tremonti et al (2004). Throughout this paper we have assumed   $H_0=70$~km~s$^{-1}$ Mpc$^{-1}$, $\\Omega_m=0.3$, and $\\Omega_{\\Lambda}=0.7$.  The aim of the present analysis  is to study a volume-limited sample of bulge-dominated galaxies. It is thus important to make sure that our sample, which is initially selected in the UV, is not biased {\\em against}  galaxies with no ongoing star formation. In Figure 1, we show how the NUV-$r$ and $g-r$ colours of galaxies depend on the fraction of the stellar mass that was formed in the past Gyr. These curves were generated using the Bruzual \\& Charlot (2003) GALAXEV code assuming a solar metallicity stellar population and smooth, exponentially declining star formation histories with varying e-folding times. As can be seen from the plot, as the fraction of young stars in the galaxy goes to zero, the NUV-$r$ colours are predicted to saturate at a value of around 6.5. The figure also shows that the NUV-$r$ colour remains sensitive to a much lower level of recent star formation than the optical $g-r$ colour. Since our GALEX sample is limited at $m_{AB} =23$, Figure 1 suggests that an $r$-band magnitude limit of 16.5 (i.e. 1.2 magnitudes brighter than the limit of the main spectroscopic sample) is required in order to ensure that our sample is not biased against galaxies with the very oldest stellar populations.  The top left panel in Figure 2 shows the NUV-$r$ colour distribution of galaxies with $r< 16.5$ and NUV$< 23$ in our sample. We use SDSS model magnitudes, which are well matched to the AUTO/Kron  magnitudes output by the GALEX pipeline. The magnitudes have also been corrected for foreground extinction using the values of E(B-V) given by Schlegel et al (1998) and the Cardelli et al (1998) reddening curve. As can be seen, the colour distribution is clearly bimodal. The population of red galaxies is peaked at NUV-$r$ $\\sim 6$ and there are very few  galaxies with NUV-$r$$> 6.5$. This gives us confidence that our sample is indeed complete in NUV-$r$ colour space. In the next two panels, we plot the stellar masses and the stellar velocity dispersions (measured within the 3 arcsec fiber aperture) of the red galaxies with NUV-$r$ $>4.5$  as a function of their redshifts. These red objects are the most difficult to detect in the UV and they also have the highest $r$-band stellar mass-to-light ratios, so they determine the effective completeness limit of our sample. Figure 2 shows that if  we limit the sample to lie in the redshift range $0.03 < z < 0.07$, we are complete down to $\\log M_* \\sim 10.4$ and $\\log \\sigma \\sim 2.05$.  We note that at $z \\sim 0.05$, the SDSS fibre aperture subtends a physical scale of only 3.2 kpc. The physical quantities that are measured from the spectra are thus relevant to the {\\em central bulge-dominated regions of the galaxy}. Throughout this paper, we will interpret spectral quantities such as D$_n$(4000), H$\\delta_A$ and EQW(H$\\alpha$) as measures of the age of the stellar population {\\em in the bulge}. The NUV-$r$ and $g-r$ colours , on the other hand, are sensitive to the age of the stellar population of the galaxy as a whole.  In the bottom right panel of Figure 2 we plot the stellar surface mass density $\\mu_*$  as a function of stellar velocity dispersion for galaxies with $\\log M_* > 10.4$, $\\log \\sigma > 2.05$,  $r< 16.5$, NUV$<23$ and $0.03 < z < 0.07$ . This sample includes 1375 galaxies. The surface mass density $\\mu_*$ is defined as $\\mu_* = (0.5 M_*)/(\\pi R50z)^2$, where $R50z$ is the half-light radius of the galaxy in $z$-band. As discussed in Kauffmann et al (2003b; 2006), there is a sharp transition in the mean stellar age of galaxies at  a stellar surface mass density $\\mu_*$ of around $10^{8.5} M_{\\odot}$ kpc$^{-2}$. Galaxies with surface densities lower than this value have on average formed their stars at a constant rate for a Hubble time; galaxies with surface densities higher than this value have little or no ongoing star formation. Figure 2 shows that essentially all galaxies in our sample defined at $\\log \\sigma > 2.05$ have surface densities greater than the transition value. The range of densities spanned by the objects in our sample is also very similar to that of elliptical galaxies (Bernardi et al 2003)  and  it is hence meaningful to refer to these objects as ``bulge-dominated'' systems.  We also find that 85\\% of the galaxies in our sample have $r$-band  concentration indices (C=R90/R50) greater than 2.6. Strateva et al (2001) and Shimasaku et al (2001) have studied how this index relates to \"by eye\" classification into Hubble type. These authors find that C=2.6 marks a reasonably robust division between early-type galaxies (E/S0/Sa) and the late-type spiral population. Note that we choose not to apply further cuts to exclude the minority of  later-type galaxies in our sample. The aim of the present paper is to examine possible evolutionary links between different galaxy populations  at fixed bulge velocity dispersion or mass. It is thus critical to ensure that  our sample is  {\\em complete} at a given value of $\\sigma$. ",
        "conclusions": ""
    },
    "0609/astro-ph0609600_arXiv.txt": {
        "abstract": "{ A  new method is   presented for determining  the  Point Spread Function (PSF) of images  that lack bright and isolated stars.   It is based  on the same  principles as the MCS (Magain, Courbin, Sohy, 1998) image deconvolution algorithm. It uses the information contained in all stellar images to achieve the double task of  reconstructing the PSFs for single or multiple exposures of the same  field and to extract the photometry of all point sources  in  the field of view.  The use of the full information available allows to construct an accurate PSF.  The possibility to simultaneously consider several exposures makes it very  well suited  to  the measurement of the light curves   of blended point  sources from  data that would be very difficult or even impossible to analyse with traditional PSF fitting  techniques. The  potential  of  the method  for the analysis of ground-based    and space-based data is  tested on artificial images and illustrated by several examples, including HST/NICMOS images of a lensed quasar and VLT/ISAAC images of a faint blended  Mira star in  the  halo of the  giant elliptical galaxy   NGC~5128   (Cen   A). ",
        "introduction": "In recent years, the importance of point sources CCD photometry has continuously grown.  Among other applications, let us mention the determination of colour-magnitude diagrams for stellar clusters, the study of variable stars in clusters or in external galaxies and the search for microlensing events and for planetary transits in light curves, by experiments such as OGLE (e.g. Udalski et al.\\cite{Udalski}).  When the stars are sufficiently isolated, aperture photometry may provide quite reliable results.  However, this is seldom the case, and techniques based on Point Spread Function (PSF) fitting have to be used when the stellar images overlap.  The most popular of these crowded field photometry methods is undoubtedly DAOPHOT (Stetson \\cite{Stetson}).  In such methods, the PSF is generally determined from the shape of sufficiently isolated point sources.  However, it may happen that the crowding of the field is such that no star is sufficiently isolated to allow a reliable PSF determination. This may severely affect the photometric precision as any error in the PSF will translate in photometric errors for blended stars, which will be accurately separated only if the PSF used in the fitting procedure is the correct one.  In such very crowded fields, it appears therefore necessary to use a PSF determination method which is able to give accurate results even in situations where all stars are blended to some degree.  This is the aim of the present paper, which combines image deconvolution and PSF determination for reaching better performances in both areas and, as a consequence, for providing excellent photometric accuracy.  Indeed, it has also recently become clear that astronomical  image deconvolution  is a   useful tool for   reaching  higher spatial resolution  and extracting  quantitative  and precise information from the data.  It should not simply be considered as a  way to improve low quality images, or as a competitor to other techniques but rather as a complementary method,  which better allows  to take full  advantage of the existing or future observational techniques.  For example,  the  Hubble Space  Telescope  (HST),  which has a  Point Spread Function (PSF) with a very compact core, still suffers from low frequency blurring due to its extended and complex wings. This is also the case for PSFs obtained with  adaptive optics intruments and will be true with the   future Extremely Large Telescopes, all  using adaptive optics. These   extended wings (and  in  fact  also  the high frequency signal due to the core of the PSF)  can often be efficiently removed  by  adequate  deconvolution methods,   at    a cost  which  is negligible compared  to the investment  in the instrument itself.  Similarly to crowded field photometry, the quality of the deconvolution critically depends on the accuracy of the PSF determination.  The present paper describes a method which simultaneously allows to obtain accurate PSFs and to solve the deconvolution problem in fields dominated by point sources, even when the crowding is so severe that all stars are significantly blended. ",
        "conclusions": "The deconvolution-based algorithm presented here has   been  designed to compute accurate PSFs in fields very crowded with point sources.  It not only computes the PSF but also provides the  photometry and astrometry of all point sources in the field of view.   The algorithm can be applied to  single images or to  a set of  images of a   given field.  In the latter case, the images are processed  simultaneously, in the sense of the MCS deconvolution  algorithm: a PSF is computed  for each image, by considering all images simultaneously. The photometry of  all points sources  is obtained for all images in the dataset, i.e., light curves are directly produced.  The method is clearly useful when few or no isolated PSF stars are available  in the field of  view, in the  case of extreme crowding and in the case of strong PSF variation  accross the field (in which case the PSF has to be determined from stars very close to the target). It is also very efficient in extracting  photometric information from datasets of very   heterogeneous quality (e.g., with  a  broad range  of seeing values). In this case, the astrometry of the points sources in the best seeing data effectively constrains the astrometry  of the bad seeing data, as long as all the data are processed simultaneously.  Although the present method  is  not designed to   handle images that consist in a mixture of  point sources  and extended ones,  it can cope  with extended objects  which  are faint in comparison to the  point sources, by running the algorithm in an iterative way."
    },
    "0609/astro-ph0609093_arXiv.txt": {
        "abstract": "\\citet{just06} recently suggested that black-hole low-mass X-ray binaries (BHLMXBs) with short orbital periods may have evolved from black-hole intermediate-mass X-ray binaries (BHIMXBs). In their model the secondaries in BHIMXBs are assumed to  possess anomalously high magnetic fields, so that magnetic braking can lead to substantial loss of angular momentum. In this paper we propose an alternative mechanism for orbital angular momentum loss in BHIMXBs. We assume that a small fraction $\\delta$ of the transferred mass from the donor star form a circumbinary disc surrounding the binary system. The tidal torques exerted by the disc can effectively drain orbital angular momentum from the binary. We have numerically calculated the evolutionary sequences of BHIMXBs, to examine the influence of the circumbinary disc on the binary evolution. Our results indicate when $\\delta\\la 0.01-0.1$ (depending on the initial orbital periods), the circumbinary disc can cause secular orbital shrinking, leading to the formation of compact BHLMXBs, otherwise the orbits always expand during the evolution. This scenario also suggests the possible existence of luminous, persistent BHLMXBs, but it suffers the same problem as in \\citet{just06} that, the predicted effective temperatures  of the donor stars are significantly higher than  those of the observed donor stars in BHLMXBs. ",
        "introduction": "If the gravitational mass of a compact star exceeds the maximum value $\\sim 2-3M_{\\odot}$ of a neutron star \\citep[e.g.][]{rhoa74}, this object should be taken as a black hole (BH) candidate. There exist currently around twenty stellar-mass BH candidates \\citep[see][]{lee02,oros02,casa06}, all of them are located in binary systems where their dynamical masses can be availably estimated. Nine of these systems are defined as compact BH X-ray binaries (BHXBs) with short orbital periods ($\\la$ 0.5 d) and low-mass donors ($< 1M_{\\odot}$) \\citep{lee02,ritt03,pods03}. It was estimated that there may exist $\\ga1000$ compact BHXBs in the Galaxy \\citep{wije96,roma98}.  The short orbital periods of BH low-mass X-ray binaries (BHLMXBs) imply that they must have undergone secular orbital angular momentum loss. If their progenitor systems contains a low-mass secondary initially, it is not clear whether the secondary star has enough energy to eject the envelope of the black hole progenitor during the common envelope evolution phase \\citep{pods03,just06}. This difficulty can not be dealt with by assuming a intermediate-mass secondary, which have radiative envelope and are not expected to be subject to magnetic braking \\citep{kawa88}, since the binary orbits will be widen when mass is transferred from the less massive secondary to the more massive BH \\citep[see][]{huan63}. \\citet{eggl86} suggested that these systems may be evolved from triple systems via the merging of the two massive components. \\citet{pods95} proposed that massive Thorne-\\.{Z}ytkow objects could lead to the birth of BHs with a low-mass donor star. \\citet{just06} summarized the previous proposals for the formation of BHLMXBs. They instead suggested that BHXBs containing intermediate-mass Ap and Bp donor stars, which possess strong magnetic fields, may be driven to compact BHLMXBs via magnetic braking with an irradiation-driven wind from the donor star \\citep[see however,][]{yun06}.   Are there additional angular momentum loss mechanisms for BHXBs besides magnetic braking? Here we explore an alternative possibility. It is well known that some fraction of the transferred matter from the donor star may leave the system in various ways during the mass exchange process  \\citep{heu94}. In particular, for cataclysmic variables (CVs), \\citet{spru01} and \\citet{taam01} suggested that part of the outflow may be in the form of a circumbinary (CB) disc, and investigated the influence of the CB disc on the evolution of CVs. The similar idea was adopted by \\citet{mey01} for BHXBs to regulate the mass transfer rates, in which the CB disc could result from the remnants of the previous common-envelope phase. In our previous works we have investigated the evolution of neutron star (NS)  LMXBs and Algol type binaries with a CB disc \\citep{chen06a,chen06b}. These works indicate that the presence of a CB disc can accelerate the evolution process of the binary systems, enhance the mass transfer rates, and lead to secular orbit shrinking under certain conditions.    The aim of this work is to study the influence of the CB disc on the evolution of BH intermediate-mass X-ray binaries (BHIMXBs), and explore the possibility of BHIMXBs being the progenitor systems of short period BHLMXBs. The structure of this paper is as follows. In section 2 we describe the adopted orbital angular momentum loss mechanisms in the evolution model of BHIMXBs. In section 3 we present the numerically calculated results for the evolutionary sequences. We summarize and discuss the uncertainties in our model in section 4.   \\section[]{Description of the Model}  Although some BHXBs may form dynamically in globular clusters \\citep{rapp01}, we only consider those evolved from the primordial massive binary systems. We focus on the evolution of BHXBs consisting of an intermediate-mass donor star (of mass $M_{\\rm d}\\sim 3-5 M_{\\odot}$) and a black hole (of mass $M_{\\rm BH}$). Neglecting the spin angular momenta of both components, we consider three types of angular momentum loss from the binary system (in all calculations, magnetic braking mechanism is not included), which are described as follows.  \\subsection[]{Gravitational wave radiation} For a compact BHXBs  gravitational wave radiation is able to carry away orbital angular momentum effectively and lead to mass transfer. The angular momentum loss rate due to gravitational radiation is given by \\citep{lan62}: \\begin{equation} \\dot{J}_{\\rm GR}=-\\frac{32}{5}\\frac{G^{7/2}}{c^{5}}\\frac{M_{\\rm BH}^{2}M_{\\rm d}^{2}M^{1/2}}{a^{7/2}}, \\end{equation} where $G$ is the gravitational constant, $c$ the speed of light, $M=M_{\\rm BH}+M_{\\rm d}$ the total mass of the binary, $a$ the binary separation given by the Kepler's third law $a=(GM/\\Omega^{2})^{1/3}$, where $\\Omega$ is the orbital angular velocity of the binary system.  \\subsection[]{Isotropic winds} Mass and orbital angular momentum loss may occur during rapid mass transfer phase, since the mass accretion rate of a black hole is limited by Eddington mass-accretion rate\\footnote{We do not consider mass and angular momentum loss due to the wind mass loss from the donor star, which is not important for intermediate-mass stars. Irradiation effect on the stellar winds is also not included.}. For spherical accretion, this maximum accretion rate can be derived from the equation that gravity can balance the radiation pressure: \\begin{equation} \\dot{M}_{\\rm Edd}=\\frac{4\\pi GM_{\\rm BH}}{\\kappa c\\eta_{\\rm E}}, \\end{equation} where $\\kappa=0.2(1+X)\\rm cm^{2}g^{-1}$ is the electron scattering opacity for a composition with hydrogen mass fraction X \\citep{kipp90}. The energy release efficiency $\\eta_{\\rm E}$ of disc accretion onto black hole can be approximately written as: \\begin{equation} \\eta_{\\rm E}=1-\\sqrt{1-\\left(\\frac{M_{\\rm BH}}{3M^{0}_{\\rm BH}}\\right)^{2}} \\end{equation} as $M_{\\rm BH}<\\sqrt{6}M^{0}_{\\rm BH}$, where $M^{0}_{\\rm BH}$ is the initial mass of the BH  \\citep[see, e.g.][]{bard70,king99}. From the above equations, \\citet{pods03} obtained the expression of the Eddington mass-accretion rate for a disc-fed BH: \\begin{equation} \\dot{M}_{\\rm Edd}\\simeq 2.6\\times 10^{-7} M_{\\odot} {\\rm yr^{-1}} \\left(\\frac{M_{\\rm BH}}{10M_{\\odot}}\\right) \\left(\\frac{0.1}{\\eta_{\\rm E}}\\right)\\left(\\frac{1.7}{1+X}\\right). \\end{equation}  In our calculations we assume that the transferred matter in excess of the Eddington accretion rate is ejected in the vicinity of the BH in the form of isotropic winds, carrying away the specific orbital angular momentum of the BH. We introduce the parameter $f$ to describe the fraction of mass loss from the binary system, defined by the following relations: \\begin{equation} \\dot{M}=\\dot{M}_{\\rm d}f, \\end{equation} and \\begin{equation} \\dot{M}_{\\rm BH}=\\dot{M}_{\\rm d}(f-1), \\end{equation} where $\\dot{M}$ is the total mass-loss rate from the system, $\\dot{M}_{\\rm d}$ is the mass transfer rate from the secondary, and $\\dot{M}_{\\rm BH}$ is the mass accretion rate of the black hole, respectively. If the mass transfer rate $|\\dot{M}_{\\rm d}|\\leq \\dot{M}_{\\rm Edd}$, $f=0$, else $f=1+\\dot{M}_{\\rm Edd}/\\dot{M}_{\\rm d}$.  Based on the above equations and assumptions, the angular momentum loss rate via isotropic winds is \\begin{equation} \\dot{J}_{\\rm IW}=f\\frac{\\dot{M}_{\\rm d}M_{\\rm d}}{MM_{\\rm BH}}J, \\end{equation} where \\begin{equation} J=a^{2}\\mu\\Omega \\end{equation} is the total orbital angular momentum and $\\mu=M_{\\rm d}M_{\\rm BH}/{M}$ is the reduced mass of the binary system, respectively.  \\subsection[]{CB disc}  We further assume that a small fraction $\\delta$ of the material overflowed from the donor star's Roche-lobe {\\em always} feeds into the CB disc rather accretes onto the BH. Then $\\dot{M}_{\\rm d}$ in Eqs.~(5)-(7) should be replaced by $\\dot{M}_{\\rm d}(1-\\delta)$, the BH mass accretion rate is $\\dot{M}_{\\rm BH}=-(1-f)(1-\\delta)\\dot{M}_{\\rm d}$, and the total mass loss rate $\\dot{M}=\\dot{M}_{\\rm d}(f+\\delta-f\\delta)$. Tidal torques are then exerted on the CB disc extracting orbital angular momentum from the binary system \\citep[see][]{taam01}.   Similar as in \\citet{taam01}, \\citet{spru01}, and \\citet{chen06b}, the viscous torque exerted at the inner edge $r_{\\rm i}$ of the CB disc\\footnote{In this subsection we use the subscript i to denote quantities evaluated at the inner edge $r_{\\rm i}$ of the CB disc.} on the BHXB can be shown to be : \\begin{equation} T_{\\rm i}\\equiv\\dot{J}_{\\rm CB}=\\gamma\\Omega a^{2}\\delta\\dot{M}_{\\rm d}\\left(\\frac{t}{t_{\\rm vi}}\\right)^{1/3}, \\end{equation} where $\\gamma^{2}=r_{\\rm i}/a$, $t$ is the time since the onset of Roche lobe overflow (RLOF). In the standard $\\alpha$ viscosity prescription \\citep{Shak73}, the viscous timescale $t_{\\rm vi}$ at the inner edge in the CB disc is given by: \\begin{equation} t_{\\rm vi}=\\frac{4\\gamma^{3}}{3\\alpha\\Omega\\beta^{2}}, \\end{equation} where $\\beta=H_{\\rm i}/r_{\\rm i}$, $\\alpha$ and $H_{\\rm i}$ are the viscosity parameter and the scale height of disc, respectively.  Combine Eqs.~(9) and (10), the angular momentum loss rate via the CB disc can be written as: \\begin{equation} \\dot{J}_{\\rm CB}=\\eta\\dot{M_{\\rm d}}j_{\\rm CB}, \\end{equation} where $\\eta=\\delta(t/t_{\\rm vi})^{1/3}$, and $ j_{\\rm CB}=\\gamma J/\\mu$ is the specific orbital angular momentum of the disc material at $r_{\\rm i}$ \\citep[e.g.][]{sobe97}.  The orbital evolution of BHXBs is governed by the change of the orbital angular momentum $J$ of the system caused by the above mentioned three mechanisms. Differentiating Eq.~(8) we get \\begin{equation} \\frac{\\dot{J}}{J}=-\\frac{\\dot{\\Omega}}{3\\Omega}+\\frac{\\dot{M_{\\rm d}}}{M_{\\rm d}}\\left[1-(1-f-\\delta)\\frac{M_{\\rm d}}{M_{\\rm BH}}-(f+\\delta)\\frac{M_{\\rm d}}{3M}\\right], \\end{equation} where we have neglected the terms of $f\\delta$. Neglect the angular momentum loss due to gravitational wave radiation, the orbital evolution is governed by \\begin{equation} \\frac{\\dot{J}_{\\rm CB}}{J} = -\\frac{\\dot{\\Omega}}{3\\Omega}+\\frac{\\dot{M_{\\rm d}}}{M_{\\rm d}}\\left[1-\\frac{M_{\\rm d}}{M_{\\rm BH}} + \\frac{2(f+\\delta)M_{\\rm d}}{3M}\\right]. \\end{equation} Insert Eq.~(11) into Eq.~(13), orbit shrinking, i.e. $\\dot{\\Omega}>0$, will occur when \\begin{equation} \\eta  > \\frac{M_{\\rm BH}}{\\gamma M}\\left[1-\\frac{M_{\\rm d}}{M_{\\rm BH}}+\\frac{2(f+\\delta)M_{\\rm d}}{3M}\\right]. \\end{equation} In this paper we take $r_{\\rm i}/a=\\gamma^{2}=1.7$ \\citep[see][]{arty94,muno06}, then $\\eta> 0.77 $ when $M_{\\rm d}\\ll M_{\\rm BH}$, and $f,\\delta\\ll 1$ (see Fig. 1).  In Fig.~1 we plot the expected relation between $\\eta$ and $M_{\\rm d}$  when the orbital period is constant for BH masses $M_{\\rm BH}=7M_{\\odot},10M_{\\odot}$, and $20M_{\\odot}$, respectively. Generally a larger $\\eta$ is required for smaller $M_{\\rm d}$ and larger $M_{\\rm BH}$. This can be satisfied with an adequate mass input rate $\\delta$ and sufficiently long time $t$ of RLOF.   \\begin{figure} \\begin{centering} \\epsfig{figure=f1.eps, width=8cm} \\caption{\\label{EqmMdot} The critical $\\eta=\\delta(t/t_{\\rm vi})^{1/3}$ as a function of the mass $M_{\\rm d}$ of donor star for a constant binary orbit. The BH masses are taken to be 7 $M_{\\odot}$ (dashed curve), 10 $M_{\\odot}$ (solid curve), and 20 $M_{\\odot}$ (dotted curve), respectively. } \\end{centering} \\end{figure}    \\section[]{Numerical Results} We have calculated the evolution of BHIMXBs adopting an updated version of the stellar evolution code developed by \\citet{eggl71,eggl72} \\citep[see also][]{han94,pols95}. In the calculations we set initial solar chemical compositions ($X=0.7$, $Y=0.28$, and $Z=0.02$) for the donor stars, and take the ratio of the mixing length to the pressure scale height to be 2.0. We include the afore-mentioned three types of orbital angular momentum loss mechanisms during the binary evolution. For the CB disc, we take $\\alpha=0.01$ and $\\beta=0.03$ \\citep{bell04}. Setting the initial mass of the BH $M^{0}_{\\rm BH}=10M_{\\odot}$, we have performed the evolution calculations of BHXBs with an initial secondary of $M_{\\rm d}=3$, 4, and 5 $M_{\\odot}$.    \\begin{figure*} \\begin{centering} \\epsfig{figure=f2.eps, width=16cm} \\caption{\\label{1E6} Evolution of the donor masses $M_{\\rm d}$ (top), orbital periods $P_{\\rm orb}$ (middle), and mass transfer rates $\\dot{M}_{\\rm d}$ (bottom) for BHIMXBs with the initial orbital period $P_{\\rm orb}=1$ d. The left, middle, and right panels correspond to $M_{\\rm d}=3M_{\\odot}$ (solid curve $\\delta=0.0055$, dashed curve $\\delta=0.004$), $4M_{\\odot}$ (solid curve $\\delta=0.007$, dashed curve $\\delta=0.005$), and $5M_{\\odot}$ (solid curve $\\delta=0.007$, dashed curve $\\delta=0.005$), respectively. Stable and unstable mass transfer processes are plotted with thick and thin curves, respectively.} \\end{centering} \\end{figure*}  \\begin{figure*} \\begin{centering} \\epsfig{figure=f3.eps, width=16cm} \\caption{\\label{1E6} Evolution of the donor masses $M_{\\rm d}$ (left), orbital periods $P_{\\rm orb}$ (middle), and mass transfer rates $\\dot{M}_{\\rm d}$ (right) for BHIMXBs with an initial orbital period $P_{\\rm orb}=3$ d. The solid, dashed, and dotted curves correspond to the initial donor mass of $3M_{\\odot}$ ($\\delta=0.055$), $4M_{\\odot}$ ($\\delta=0.07$), and $5M_{\\odot}$ ($\\delta=0.09$), respectively.} \\end{centering} \\end{figure*}  We first estimate the possible range of the fractional mass input rate into the CB disc, $\\delta$, when the CB disc's influence on the orbital  evolution of BHIMXBs is effective. From Eq.~(10) we find the viscous timescale $t_{\\rm vi}\\sim 100$ yr at the inner edge of the CB disc for a BH binary with an orbital period of $\\sim 1$ d, and the duration of RLOF $t\\sim 10^{8}$ yr, so we get $\\delta\\ga 0.77(t_{\\rm vi}/t)^{1/3} \\sim 0.008$ from Eq.~(14). For systems with an initial orbital period of $\\sim 3$ d, $t\\sim 10^{6}$ yr, and $\\delta \\ga 0.04 $. Obviously the longer $t$, the smaller $\\delta$ required. These values of $\\delta$  seem to be in agreement with \\citet{taam01} who have shown that the CB disc is effective in draining orbital angular momentum from the system provided that $\\delta$ exceeds about 0.01. In our calculations we find that CB discs with $\\delta> \\sim 0.01-0.1$ (depending on the initial orbital periods) will cause runaway mass transfer with high mass transfer rates, while the evolutions with $\\delta$ significantly less than 0.01 is similar to those without CB discs, as already noticed in CV evolution \\citep{taam01}.  \\begin{figure} \\begin{centering} \\epsfig{figure=f4.eps, width=8cm} \\caption{\\label{EqmMdot} The evolution of the  donor masses $M_{\\rm d}$ and orbital periods $P_{\\rm orb}$ for NSLMXBs with a initial donor mass of $M_{\\rm d}=1M_{\\odot}$, and an initial orbital period of $P_{\\rm orb}=1$ d. The solid and dashed curves correspond to $\\delta=0.005$ and 0.002, respectively. } \\end{centering} \\end{figure}   Figure 2 shows the examples of the evolutionary sequence with an initial orbital period of 1 d. The solid and dotted lines in the figure correspond to two different values of $\\delta=(5.5-7)\\times 10^{-3}$ and $\\delta=(4-5)\\times 10^{-3}$ respectively. As can seen in the figure, mass transfer is initially driven by nuclear evolution of the secondary, leading to expansion of the orbits. But this tendency is held up when the angular momentum loss via the CB disc becomes sufficiently strong. The mass transfer drops into a ``plateau\" phase at a rate $\\sim 10^{-9}M_{\\odot}\\rm yr^{-1}$ for a few $10^8$ yr. These features are different from those in \\citet{pods03} for the standard evolution of BHIMXBs, in which the orbits always increase secularly. After that the mass transfer rates increase sharply as the secondary ascends the giant branch, but the final orbital evolution depends on the adopted value of $\\delta$. With the higher value of $\\delta$, a compact BHLMXB will be finally produced after a few $10^8$ yr mass transfer. If matter does not leave the CB disc, we can estimate from Fig.~2 the mass of CB disc $M_{\\rm CB}\\sim 0.015$, 0.028, $0.035M_{\\odot}$ with the initial donor mass of 3, 4, $5M_{\\odot}$, respectively.   The calculated results for slightly wider initial orbital period of 3 d are presented in Fig.~3. Generally the mass transfer rates are higher for higher donor masses and wider initial orbits - the peaks of the mass transfer rates can reach $\\ga 10^{-5} M_{\\odot}\\rm yr^{-1}$. Within less than 1 Myr the donor masses decrease to be $\\la 1 M_{\\odot}$. Short orbital periods can also be attained, but generally a larger $\\delta$ (up to $\\la 0.1$) is required compared with those in Fig.~2. In this case the the CB disc is also more massive, $M_{\\rm CB}\\sim 0.15$, 0.28, $0.45M_{\\odot}$ with the initial donor mass of 3, 4, $5M_{\\odot}$, respectively.   One may expect that the CB disc may play a similar role in the evolution of NSLMXBs. In Fig.~4, we plot the evolutionary sequence of NSLMXBs with a donor star of $M_{\\rm d}=1M_{\\odot}$ and an initial orbital period of $P_{\\rm orb}=1$ d, to show the similar influence of changing $\\delta$ on the final orbital evolution as in BH binaries. In our previous work on the formation of the binary radio pulsar PSR J1713$+$0747 \\citep{chen06a} we have found that for a given final orbital period, including the CB disc can allow a more evolved donor star at the onset of mass transfer, so that the X-ray lifetime can be considerably decreased.  \\begin{figure} \\begin{centering} \\epsfig{figure=f5.eps, width=8cm} \\caption{\\label{EqmMdot} The evolutionary tracks of BHIMXBs in the effective temperature of the secondary vs. the orbital period plane when the initial orbital period $P_{\\rm orb,i}=1$ d. The solid, dashed, dotted, dash-dotted, and dash-dot-dotted curves correspond to the initial donor star masses of $1M_{\\odot}$ ($\\delta=0.002$), $1.5M_{\\odot}$ ($\\delta=0.0042$), $2M_{\\odot}$ ($\\delta=0.0048$), $3M_{\\odot}$ ($\\delta=0.0055$), and $4M_{\\odot}$ ($\\delta=0.007$), respectively. The orbital periods of eight compact BHXBs and the earliest spectral types of the optical companions \\citep{casa06} are marked with crosses.} \\end{centering} \\end{figure}  All of the BHLMXBs with an orbital period of $\\la0.5$ d should be transient sources \\citep{casa06}. It is currently believed that the transient behavior in X-ray binaries is due to the thermal-viscous instability of accretion discs, when the mass transfer rates lie below a critical value $\\dot{M}_{\\rm cr}$ so that the surface temperature at the outer edge of accretion disc is lower than the hydrogen ionization temperature \\citep{king96}. This critical mass transfer rate depends primarily on $M_{\\rm BH}$, $M_{\\rm d}$, and $P_{\\rm orb}$, and is given by \\citep{para96,dubu99}: \\begin{equation} \\dot{M}_{\\rm cr}\\simeq 8.6\\times 10^{-9}\\left(\\frac{M_{\\rm BH}}{10M_{\\odot}}\\right)^{0.5}\\left(\\frac{M_{\\rm d}}{1M_{\\odot}}\\right)^{-0.2}\\left(\\frac{P_{\\rm orb}}{1 \\rm d}\\right)^{0.5} M_{\\odot}\\rm yr^{-1} \\end{equation} In Fig.~2 the thick and thin curves correspond to stable and unstable mass transfer during the evolution of BHXBs. However, the mass transfer would be always stable for the wide orbit system with an initial orbital period of 3 d (see Fig.~3). It is interesting to see that the model BHLMXBs are likely to be persistent X-ray sources. This is not compatible with the observations of Galactic BHLMXBs, suggesting that $\\delta$ does not need to be constant but changes during the evolution (see discussion below)\\footnote{{\\em Chandra} observations of the 12.6 hr ultraluminous X-ray source in the elliptical galaxy NGC3379 suggest that the current on phase has lasted $\\sim 10$ yr \\citep{fabb06}. This source could be a soft X-ray transient with long duration of outbursts. However, a persistent X-ray source as suggested by our calculations is an alternative possibility.}. ",
        "conclusions": "In this paper we have examined the influence of a CB disc on the evolution of BHIMXBs. Assuming that a fraction $\\delta$ of the transferred material from the secondary forms a CB disc surrounding the binary system and extracts the orbital angular momentum from the binary, we have performed evolution calculations for BHXBs containing a donor star of mass in the range from 3 to 5 $M_{\\odot}$ with an initial orbital period of 1 d and 3 d. The calculations show that the evolution of BHIMXB is sensitive to $\\delta$. Generally a larger $\\delta$ is required for a wider initial orbit since the mass transfer itself during the evolution always causes orbital expansion. The results indicate that the orbits of BHXBs would show secular shrinkage when the values of $\\delta$ lie in the range of a few $10^{-3}- 10^{-1}$. Thus our CB disc scenario suggests a new evolutionary channel for the formation of BHLMXBs from BHIMXBs.  However, our results encounter the similar difficulty as in \\citet{just06}, that the calculated effective temperatures $T_{\\rm eff}$ are not consistent with those of the observed donor stars in BHLMXBs. We compare the calculated results of both BH-LMXBs and IMXBs with the observations in the $T_{\\rm eff}$ - $P_{\\rm orb}$ diagram in Fig.~5. One can see that the observed results seem to be more consistent with the evolutionary tracks of original BHLMXBs. \\citet{just06} have already pointed out that this tempreature discrepancy seems to be a generic difficulty with any formation scenario that invokes primordially intermediate-mass donor stars. If this effective temperature problem can be solved, the CB disc mechanism may provide a plausible solution to the BHLMXB formation problem, without requiring anomalous magnetic fields in the donor stars.   The mechanism feeding the CB disc is still unclear. It has been argued that during mass exchange in binary systems, some of the lost matter which possesses high orbital angular momentum may form a disc surrounding the binary system rather than leave the binary system \\citep{heu94}. This part of matter may come from the stellar wind from the donor star, wind and/or outflow from the accretion disc, or mass lost from the outer Lagrangian point. The values of $\\delta$ adopted here seem to rule out stellar wind as the origin of the CB disc, which requires a unreasonably high wind mass loss rate in intermediate- and low-mass stars. We speculate that the disc wind/outflow may play a more important role in feeding the CB disc, as the X-ray irradiation on the accretion disc in BHXBs may evaporate the disc much more efficiently than in CVs. This may explain why the values of $\\delta$ are about $2-3$ orders of magnitude larger than those for CV evolution \\citep{spru01}.  \\citet{dubu02} investigated the structure and evolution of a geometrically thin CB disc to calculate its spectral energy distributions, and discussed the prospects for the detection of such discs in the infrared and submillimeter wavelength regions. \\citet{dubu04} searched for excess mid-infrared emission due to CB disc material in CVs. But direct detection of the CB discs in CVs by infrared continuum studies has so far been elusive, partly because of the lack of accurate disc atmosphere models. Recently, \\citet{muno06} studied the blackbody spectrum of BHLMXBs A 0620-00 and XTE J1118+480, and found that the inferred excess mid-infrared emitting areas are $\\sim 2$ times larger than the binary orbital separations. Therefore, the detection of excess mid-infrared emission from these BHLMXBs provides evidence of the existence of CB disc around some BHLMXBs. These observations may set useful constraints on the evolution of the CB discs. The masses of the model CB disc are significantly larger than those ($\\sim 10^{-9}M_{\\odot}$) estimated by \\citet{muno06}. Combined with the non-detection of discs in CVs by \\citet{dubu04}, this fact suggests that the presence of the CB disc may not always accompany the RLOF processes. For example, (X-ray) nova bursts may destroy the CB disc before it has developed to be massive enough. Or perhaps $\\delta$ only needs to be high for a short time, while normal magnetic braking takes over the CB disc during the majority of the evolution. The latter may partly account for the discrepancy between the lifetimes of BHLMXBs in Fig.~3 ($\\sim 0.5-1.5\\times 10^{6}$ yr) and those expected for observed BHLMXBs ($\\sim 10^{8}-10^9$ yr). Since the CB disc can promote mass transfer very efficiently, BHLMXBs are more likely to be observed when the value of $\\delta$ becomes extremely small."
    },
    "0609/astro-ph0609746_arXiv.txt": {
        "abstract": "{\\\\ We present a detailed analysis of the soft X-ray spectrum of the Seyfert 1 galaxy Markarian 509 taken with the \\textit{XMM-Newton} Reflection Grating Spectrometer. An underlying power-law continuum and three warm absorber phases provide a good fit to the data along with a number of broad and narrow emission lines. Our three warm absorber phases each have different ionization parameters and column densities. We identify a low ionization, log\\,$\\xi$=0.89, high outflow velocity phase producing an Fe M-shell UTA feature along with absorption from O\\,VI and N\\,VI. There is an intermediate phase, log\\,$\\xi$=2.14, showing absorption from H-like carbon and nitrogen and He-like neon and oxygen. The third, high ionization, log\\,$\\xi$=3.26, low outflow velocity phase contains absorption from O\\,VIII, Ne\\,X and highly ionized iron. All phases are blueshifted with respect to the systemic velocity with flow velocities ranging from -60\\,km\\,s$^{-1}$ to -510\\,km\\,s$^{-1}$. The observed broad emission features have an RMS velocity of 8000$\\pm$3000\\,km\\,s$^{-1}$ for the C\\,VI and O\\,VII lines.\\\\ \\\\ \\textbf{Key words.} Galaxies: active -- Galaxies: Seyfert -- Galaxies: individual: Mkn 509 -- X-rays: galaxies --  Techniques: spectroscopic -- Line: identification} ",
        "introduction": "\\label{intro}  Soft X-ray absorption by ionized gas was first identified in Active Galactic Nuclei (AGN) by~\\citet{halpern84} and recent observations indicate that over half of Seyfert 1 galaxies may have these `warm absorbers' (\\nocite{reynolds97}Reynolds 1997, \\nocite{george98}George et al. 1998). Since the launch of the \\textit{XMM-Newton} and \\textit{Chandra} high resolution grating spectrometers our ability to probe soft X-ray spectra of AGN has been greatly increased, enabling astronomers to do in-depth studies of the X-ray absorbing and emitting photoionized gas and to examine how they relate to the system as a whole (\\nocite{blustin05}e.g. Blustin et al. 2005). Gas properties are examined to improve our understanding of the kinematics, ionization and composition of these systems. The warm absorber is often characterised by its ionization parameter, $\\xi$. This determines the ionization state of the material and is defined as  \\begin{equation} \\xi=\\frac{L}{nr^2} \\end{equation} where $L$ is the source luminosity (in erg\\,s$^{-1}$), $n$ is the gas density (in cm$^{-3} $) and $r$ is the distance of the absorber from the centre of the system (in cm) (\\nocite{tarter69}Tarter et al. 1969).  Markarian 509 was classified as a Seyfert 1 galaxy by~\\citet{kopylov74} but is more like a Seyfert 1/quasi-stellar object (QSO) transition galaxy due to its compact appearance and high luminosity of $1.3-2.6\\times 10^{44}\\,$erg\\,s$^{-1}$ in the 2-10\\,keV range (\\nocite{weaver01}Weaver et al. 2001). Its close proximity, z=0.034397 (\\nocite{fisher95}Fisher et al. 1995), compared to other AGN of similar luminosity, and the low Galactic column density in the line of sight (N$_{H}=4.4\\times10^{20}\\,$cm$^{-2}$; \\nocite{murphy96}Murphy et al. 1996), make it an ideal candidate for studying the inner regions of such objects.  Markarian 509 has well-studied intrinsic UV absorption (\\nocite{kriss00}Kriss et al. 2000; ~\\nocite{kraemer03}Kraemer et al. 2003) but much is still to be learnt from its high resolution X-ray spectrum. Markarian 509 was first identified as an X-ray source by~\\citet{cooke78} and has since been observed with a number of instruments including \\textit{EXOSAT} (\\nocite{morini87}Morini et al. 1987), \\textit{GINGA} (\\nocite{singh90}Singh et al. 1990), and \\textit{BeppoSAX} (\\nocite{perola00}Perola et al. 2000).  \\citet{pounds01} studied the \\textit{XMM-Newton} Reflection Grating Spectrometer (RGS) spectrum of Markarian 509 taken in October 2000. They modelled it with a blackbody continuum as well as a weak Fe M-shell UTA region, Ne IX absorption and O\\,VII narrow emission lines. ~\\citet{yaqoob03} performed a thorough analysis of the \\textit{Chandra} High-Energy Transmission Grating data. They  modelled the spectrum with two absorption edges from O\\,VII and O\\,VIII along with Ne\\,IX, Ne\\,X and O\\,VIII line absorption. They fitted these spectral features using one phase of absorption through a photoionized gas with a column density 2.06$^{+0.39}_{-0.45}\\times 10^{21}\\,$cm$^{-2}$ and an ionization parameter of log\\,$\\xi=1.76^{+0.13}_{-0.14}$. The gas was found to be outflowing at $200\\pm$170\\,km\\,s$^{-1}$ and a curve of growth analysis determined an approximate turbulent velocity of 100\\,km\\,s$^{-1}$.  ~\\citet{kriss00} analysed the UV spectrum of Markarian 509 obtained with the \\textit{Far Ultraviolet Spectroscopic Explorer}. They modelled seven components for the UV absorber, which they grouped into two velocity phases with outflow velocities 0\\,km\\,s$^{-1}$ and 370\\,km\\,s$^{-1}$. They identified one component as having an ionization parameter and column density that could be compatible with an X-ray absorbing gas. With the exception of this component, however, they predicted that the UV and X-ray spectra are produced in separate zones within the AGN. ~\\citet{kraemer03} analysed the UV spectrum of Markarian 509 using data from the Space Telescope Imaging Spectrograph echelle gratings on the \\textit{Hubble Space Telescope}. They identified eight components in the UV absorber, with a similar range of velocities to Kriss et al. They did not find evidence for a component with an ionization and column density compatible with the X-ray absorbing phase. However, Kraemer et al. concluded from the kinematic association that the UV and X-ray absorbers may co-exist spatially, for example in the form of a two-phase clumpy medium.  In this paper we investigate the combined RGS spectral data from two observations by \\textit{XMM-Newton}.  In Section~\\ref{analysis} we describe the observations and data reduction, and the spectral fitting is reported in Section~\\ref{fitting}. We discuss the results in Section~\\ref{discussion} and present our conclusions in Section~\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  We have used the high resolution X-ray spectra obtained with the Reflection Grating Spectrometer onboard \\textit{XMM-Newton} to study the characteristics of the Seyfert 1 galaxy Markarian 509 and present the most detailed analysis to date of the X-ray absorbing and emitting photoionized gas present in this system.  The continuum has been fitted with a simple power law. The analysis suggests the presence of a warm absorber consisting of three phases: Phase 1 contains the high outflow velocity, low ionization O\\,VI and N\\,VI absorption lines along with an Fe M-shell UTA; Phase 2 contributes the medium ionization O\\,VII, N\\,VII, C\\,VI and Ne\\,IX absorption features, and Phase 3 produces the low outflow velocity, high ionization O\\,VIII, Ne\\,X and highly ionized iron absorption features. However, we cannot rule out the possibility of a continuous ionization change throughout one gas system. It is possible that Phases 1 and 2 may coexist in pressure equilibrium, but it is unlikely that Phase 3 is also part of the same structure. We have identified Phase 1 as being compatible with a UV phase identified in previous studies.  In addition to the absorption features identified here, there is also evidence of emission in the spectrum of Markarian 509. We identify two narrow emission lines due to O\\,VII at 22.1\\AA, and O\\,VIII at 18.8\\AA. We also identify two broad emission lines, due to O\\,VII at 21.7\\AA, and C\\,VI at 33.7\\AA, which have an RMS velocity of 8000$\\pm$3000\\,km\\,s$^{-1}$ consistent with originating in the broad line region."
    },
    "0609/astro-ph0609811_arXiv.txt": {
        "abstract": "{Gamma-ray bursts, black holes, neutron stars, white dwarfs}  I consider various possibilities for making gamma--ray bursts, particularly from close binaries. In addition to the much--studied neutron star + neutron star and black hole + neutron star cases usually considered good candidates for short--duration bursts, there are other possibilities. In particular neutron star + massive white dwarf has several desirable features. These systems are likely to produce long--duration GRBs, in some cases definitely without an accompanying supernova, as observed recently. This class of burst would have a strong correlation with star formation, and occur close to the host galaxy. However rare members of the class need not be near star--forming regions, and could have any type of host galaxy. Thus a long--duration burst far from any star--forming region would also be a signature of this class. Estimates based on the existence of a known progenitor suggest that this type of GRB may be quite common, in agreement with the fact that the absence of a supernova can only be established in nearby bursts. ",
        "introduction": "Accretion on to a black hole or neutron star is the most efficient way of extracting energy from normal matter. Since GRBs are briefly the brightest objects in the Universe, accretion on to one or other of these objects must be involved in making them.  Constructing GRB models then becomes a matter of arranging that a suitable mass falls on to the accretor on the right timescale.  The total emitted energy of GRBs implies that a significant fraction of a stellar mass is involved. Since we want the burst to be sudden, we presumably have to start with some equilibrium situation involving a stellar mass, and destabilize it. Observations again tell us the infall timescale is, at least initially, extremely short. This suggests that we can pick out potential stellar reservoirs by considering their dynamical timescales.  In the collapsar picture a rapidly--rotating massive star is both accretor and reservoir. During core collapse the central regions try to form a black hole, while the rest of the core accretes on to it and makes the burst. The dynamical timescale of the core suggests that this produces a long--duration burst.  The other generic situation in which a star is disrupted on a very short timescale is in a close binary. In the rest of this paper I shall consider the various possibilities in turn. ",
        "conclusions": ""
    },
    "0609/astro-ph0609485_arXiv.txt": {
        "abstract": "{In this paper we present an interference detection toolbox consisting of a high dynamic range Digital Fast-Fourier-Transform spectrometer (DFFT, based on FPGA-technology) and data analysis software for automated radio frequency interference (RFI) detection. The DFFT spectrometer allows high speed data storage of spectra on time scales of less than a second. The high dynamic range of the device assures constant calibration even during  extremely powerful RFI events. The software uses an algorithm which performs a two-dimensional baseline fit in the time-frequency domain, searching automatically for RFI signals superposed on the spectral data. We demonstrate, that the software operates successfully on computer-generated RFI data as well as on real DFFT data recorded at the Effelsberg 100-m telescope. At 21-cm wavelength RFI signals can be identified down to the $4\\sigma_\\mathrm{rms}$ level. A statistical analysis of all RFI events detected in our observational data revealed that: (1) mean signal strength is comparable to the astronomical line emission of the Milky Way, (2) interferences are polarised, (3) electronic devices in the neighbourhood of the telescope contribute significantly to the RFI radiation. We also show that the radiometer equation is no longer fulfilled in presence of RFI signals.} {} ",
        "introduction": "The sky is not still any longer ``dark'' for radio astronomers. Tele\\-com\\-muni\\-cation and radio broadcast are present at different power levels everywhere on Earth. The protected radio astronomical bands have offered the possibility to study the distribution of matter in space down to a certain sensitivity threshold. Today the sensitivity of the radio telescopes and their receiver systems has improved to detect continually fainter signals from distant sources. On the other hand, the protected bands for modern radio telescopes are becoming in parallel more polluted by signals from a variety of electronic devices which produce Radio Frequency Interference (RFI) across the entire electromagnetic frequency range.  In example, the leakage rate of a properly working microwave oven operating at 2.45\\,GHz is typically about $1\\,\\mathrm{mW\\,cm}^{-2}$ at a distance of 5\\,cm. Considering the ITU-RA 769 recommendation for 21-cm line spectroscopy sets a threshold of $-196\\,\\mathrm{dB(W\\,m^{-2})}$, all RFI-signals should be below this recommended power-level. Assuming that the first sub-harmonic contains only $1>p>10^{-3}$ of that flux density, the linear distance to detect this signal from a microwave oven is between 100\\,km and 3\\,km. Within this area several hundred houses are located in Central Europe. Moreover, we have also to consider the aging of the components of the electronic equipment, which might fulfil the strong regulations at the time of the first usage, but fails later on. Moreover, an increasing fraction of technical equipment is needed to operate a state-of-the-art radio telescope. From a TFT-display to a GPS-clock all supplementary devices are located in the immediate vicinity of a radio telescope today and are at the same time potential sources of RFI emission.  Much of the contemporary activity in radio astronomy is focused on the exploration of the high redshift universe. These studies require observations at low frequencies with the highest possible sensitivities, far away from the protected radio bands, where most of the telecommunication, radio, and television broadcast stations are active. RFI mitigation thus is of at most importance to future radio telescopes like LOFAR (http://www.lofar.org) and the SKA (http://www.skatelescope.org).    In this paper we show that RFI mitigation is one of the major challenges in future radio astronomy. As a specific example, we studied the interference situation using the Effelsberg 100-m telescope at 21-cm wavelength. Our results indicate that RFI signals with highly variable intensities degrade the astronomical spectra across the whole protected band. Somewhat surprisingly, we find indications, that at least one major source of the RFI signals lies in the direction of the telescope control building. It is likely that the computer system and/or supplemental electronic devices produce harmonics across the 21-cm wavelength band. One way to overcome such RFI degradation of the scientific data is to develop new data reduction methods.     Several methods have been developed today to detect or even mitigate RFI. In general, three different approaches can be identified. First is the ``classical'' method, a manual search for RFI signals in already recorded spectra.  Second, using sophisticated algorithms, it is possible to search for RFI signals in recorded data automatically (Bhat et al. 2005). Third, one can use real-time applications which have to be implemented into the signal chain of the telescope. Various approaches are under consideration, e.g. adaptive filters (Bradley \\& Barnbaum 1996), post-correlators, or real-time higher order statistics (HOS; see Fridman 2001).    In this work we follow the second approach using snapshot data obtained with the high dynamic range DFFT spectrometer at the Effelsberg telescope (Stanko, Klein \\& Kerp 2005). The method itself and its application to artificial spectra polluted by computer-generated RFI is presented in Sect.~\\ref{secmethod}. In order to account for the quickly changing signal signature of most RFI events it is necessary to record the spectra on short time scales, at most a few seconds of integration. It became obvious during test observations, that the RFI signal variations at Effelsberg occur on the order of less than $100\\,\\mbox{ms}$. Such short integration times yield very high data rates. In practice, one must find a compromise between reasonable time-resolution for RFI detection and the amount of data. Moreover, the read-out time of the spectrometer is a technical constraint for RFI mitigation. The spectrometer and general experimental setup is discussed in Sect.~\\ref{secsetup}.  In Sect.~\\ref{secresults} we present the results of our RFI identification campaign at the Effelsberg 100-m telescope. Section~\\ref{secdiscuss} compiles our summary and outlook for the future. To avoid confusion we introduce some naming conventions throughout the paper. A single spectrum is denoted as a \\emph{scan}. A \\emph{subset} denotes several spectra that refer to a specific position on the sky. All spectra from one observation (or all subsets) form a \\emph{dataset}. ",
        "conclusions": "\\label{secdiscuss} We developed a new algorithm for an automated RFI signal detection in existing 21-cm line data. These had to be recorded using short integration times of less than 1\\,s. Most of the existing spectrometer backends do not allow such a fast storage of spectra. The new DFFT spectrometer which was developed at the University of Bonn by (Stanko et al. 2005) fulfils this need. Furthermore it provides a large dynamic range of 14\\,bit so that it is best qualified to do RFI analysis. But in principle the proposed algorithm could easily be adapted to other spectrometers, if they were capable of delivering short-integrated spectra. The time variability of most interferences demands analysis on short time scales. This in return means a huge amount of data to handle which can hardly be done manually. This becomes especially true in case of surveys where several millions of spectra have to be analysed if RFI mitigation on such time scales is aimed for.  The application of the algorithm to simulated data revealed its capability to find very different types of interferences with a high detection rate. It could be shown that even at low signal intensities of about $4\\sigma_\\mathrm{rms}$ a large fraction of all events was found. Probably most remarkable is that without RFI mitigation whether automatically or not, the radiometer equation is no longer fulfilled. The proposed algorithm could greatly improve this issue. It is likely that we will find other morphological types of RFI (radar chirps etc.) in upcoming observations. If these turned out to be not well-detectable by our current software, further workaround procedures would be easy to implement due to its modular design.   We developed a tool for doing statistical analysis of RFI events, which turned out to be very useful for working out properties of interferences and helpful in finding possible interference sources. Analysis of the data recorded at the 100-m Effelsberg telescope yielded some interesting results. Most of the RFI signals are linearly polarised. We could show that the majority of RFI signals enters the telescope via the telescope antenna and not via the intermediate frequency chain. Also, most of the detected RFI signals are produced by the electronic devices within control building of the 100-m telescope. The most extreme RFI signals appear to enter the telescope from randomly distributed directions.  Note, that the large counting rate from the observatory in Effelsberg is not due to bad technical conditions. Radiation at such low power levels is quite common. Soon most of electronic equipment at Effelsberg will be placed in a Faraday cage. We are curious how this will improve the situation."
    },
    "0609/gr-qc0609121_arXiv.txt": {
        "abstract": "The dynamics of ``dipolar particles'', \\textit{i.e.} particles endowed with a four-vector mass dipole moment, is investigated using an action principle in general relativity. The action is a specific functional of the particle's world line, and of the dipole moment vector, considered as a dynamical variable. The first part of the action is inspired by that of a particle with spin moving on an arbitrary gravitational background. The second part is intended to describe, at some effective level, the internal non-gravitational force linking together the ``microscopic'' constituents of the dipole. We find that some solutions of the equations of motion and evolution of the dipolar particles correspond to an equilibrium state for the dipole moment in a gravitational field. Under some hypothesis we show that a fluid of dipolar particles, supposed to constitute the dark matter, reproduces the modified Newtonian dynamics (MOND) in the non relativistic limit. We recover the main characteristics of a recently proposed quasi-Newtonian model of ``gravitational polarization''. ",
        "introduction": "\\label{secI} \\subsection{Astrophysical motivation}\\label{secIA} It is recognized (see Refs.~\\cite{Milgrev, SandMcG02, Bekrev, Sandrev} for reviews) that the modified Newtonian dynamics or MOND, which has been proposed by Milgrom~\\cite{Milg1, Milg2, Milg3} as an alternative to dark matter, works extremely well at predicting the form of rotation curves of galaxies from the observed distribution of stars and gas. In addition MOND explains naturally the relation between the luminosity of galaxies and their asymptotic rotation velocity --- the so-called Tully-Fisher law~\\cite{TF}. Though there might be some examples where MOND does not fully account for the observed kinematics of galaxies~\\cite{Gentile}, overall the fit achieved by MOND of the rotation curves of most galaxies is very impressive and calls for a possible physical explanation. On the other hand we know that the mass discrepancy of clusters of galaxies is not completely accounted for by MOND~\\cite{GerbalDurret}, and that at the cluster scale there is still empirical evidence for unseen dark matter~\\cite{Clowe}.  In the usual interpretation, MOND is viewed as a modification of the fundamental law of gravity, without the need of dark matter. Several relativistic extensions of MOND, sharing this view of modifying the sector of gravity, postulate the existence of supplementary fields associated with the gravitational force, in addition to the metric tensor of general relativity~\\cite{BekM84, BekS94, Sand97, Bek04, Sand05}. So far the most developped of the relativistic MOND theories is the tensor-scalar-vector theory of Bekenstein \\& Sanders~\\cite{Sand97, Bek04, Sand05}.  At this stage we seem to face two alternatives to the issue of dark matter: \\begin{itemize} \\item[(1)] Either accept the existence of cold dark matter particles, \\textit{e.g.} predicted by super-symmetric extensions of the standard model of particle physics (see~\\cite{BHS05} for a review). However these particles are yet to be discovered, and the simplest models of cold dark matter fail to reproduce in a natural way the flat rotation curves of galaxies~\\cite{SandMcG02}. \\item[(2)] Or postulate an alteration of our fundamental theory of gravity (namely MOND and its relativistic extensions). But the motivation for altering the law of gravity is \\textit{ad hoc}, and one can argue that the relativistic MOND theory~\\cite{Sand97, Bek04, Sand05} does not in fact explain the phenomenology of MOND --- at least untill the extra fields find a well-motivated explanation coming from fundamental physics. \\end{itemize} \\noindent In the present paper, following~\\cite{B06mond}\\,\\footnote{Hereafter Ref.~\\cite{B06mond} will be referred to as paper~I.}, we propose a third alternative. Namely we: \\begin{itemize} \\item[(3)] Keep the standard law of gravity, \\textit{i.e.} general relativity and its Newtonian limit, but we add to the distribution of ordinary matter some specific form of dark matter in such a way as to naturally explain MOND. The dark matter consists of ``polarization masses'' associated with a medium of dipole moments aligned in the gravitational field of ordinary matter. (But, in this paper, we shall leave aside the problem of clusters of galaxies~\\cite{GerbalDurret, Clowe}.) \\end{itemize}  Our basic motivation is that MOND can be naturally interpreted as resulting from an effect of gravitational polarization of a dipolar medium. Paper~I argued, on the basis of a simple quasi Newtonian model, that the polarization tends to enhance the magnitude of the gravitational field of ordinary galaxies in a way consistent with MOND. This effect constitutes the gravitational analogue of the electric polarization of a dielectric material (whose atoms can be modelled by electric dipoles) by an applied electric field~\\cite{Jackson}. Thus the phenomenology of MOND results from the non-standard influence of the dark matter on ordinary matter.  Arguably the proposal (3) does not really provide an explanation for MOND. Indeed the models of paper~I and this paper are only \\textit{effective} --- the fundamental nature of the dipolar particles (\\textit{i.e.} their ``internal'' structure) will not be elucidated. However the proposal is conceptually simple and fits naturally with the MOND phenomenology of the flat rotation curves and the Tully-Fisher empirical relation. In addition the model cannot be said to be \\textit{ad hoc} because it invokes the (gravitational analogue of the) well-known physical mechanism of polarization by an external field.  The quasi-Newtonian model of paper~I is based on a microscopic description of the dipole moment using negative gravitational-type masses (or gravitational charges). As a result the motion of dipolar particles in this model violates the equivalence principle. In the present paper we elaborate a model of dipolar particles and gravitational polarization in the standard general relativity theory,\\,\\footnote{The concept of gravitational polarization at the quadrupolar order in relativistic gravity theories has been investigated in Ref.~\\cite{Bel}.} without negative (passive) gravitational masses and consistent with the equivalence principle. Consequently, we shall find that the equations of motion of the dipolar particles in the non-relativistic (NR) limit of the present model are different from those of paper~I. However we shall recover the main characteristics of the quasi-Newtonian model of paper~I, and notably the interesting connection with the phenomenology of MOND.  \\subsection{Concept of dipole moment in general relativity}\\label{secIB} In theories satisfying the equivalence principle the mass dipole moment of a mass distribution, say $\\pi^i_\\mathrm{g}=\\sum m_\\mathrm{g}\\,x^i$ (adopting a Newtonian picture), is proportional, by equivalence between the gravitational and inertial masses, $m_\\mathrm{g}=m_\\mathrm{i}$, to the position of the center of mass of the system, $C^i=\\sum m_\\mathrm{i}\\,x^i/M_\\mathrm{i}$ (where $M_\\mathrm{i}=\\sum m_\\mathrm{i}$). Similarly, the current dipole moment, $\\mu^i_\\mathrm{g}=\\sum m_\\mathrm{g}\\,\\varepsilon^{ijk}x^j v^k$, is equivalent to the spin angular momentum $S^i=\\sum m_\\mathrm{i}\\,\\varepsilon^{ijk}x^j v^k$. Defining the mass dipole moment of a particle (supposed to be composed of some sub-particles) is \\textit{a priori} delicate because one faces the problem that $\\pi^i_\\mathrm{g}=0$ in the center-of-mass frame where $C^i=0$ by definition; however the current dipole moment $\\mu^i_\\mathrm{g}$ equivalent to the spin $S^i$ is admissible. Thus, while the notion of a particle carrying a mass dipole moment seems to be possible only at the price of violating the equivalence principle (like in the model of paper~I), the notion of a particle endowed with a current dipole moment or spin is perfectly legitimate.  Particles with spins have been the subject of many fundamental works in general relativity~\\cite{papa51spin, Traut58, Dixon, BOC75, BOC79, BI80}, with applications to the problem of dynamics of spinning black holes in binary systems~\\cite{KWW93, K95, TOO01, FBB06spin, BBF06spin}.  The general formalism has been encapsulated in an action principle for a particle with spin moving on an arbitrary gravitational background. The action, due to Bailey \\& Israel~\\cite{BI80}, is derived by expressing the standard action for a non-spinning particle (\\textit{i.e.} the integral of the proper time), as defined with respect to some ``eccentric'' world-line located at the edge of some composite particle, in terms of the quantities belonging to a different world-line, which is viewed as the ``reference'' or ``central'' world-line located at the center of the particle. Crucial to the formalism is the vector separation between the two world-lines, defined as the gradient of the geodesic distance separating them.\\,\\footnote{The geodesic distance is nothing but the world function in Synge's formalism~\\cite{Synge}; it is a bi-scalar in the general theory of bi-tensors~\\cite{dWB60}.} The spin of the particle is then given by the anti-symmetric product between the latter separation vector and the particle's linear momentum. Hence the separation vector is the ``lever arm'' associated with the spin angular momentum, as defined with respect to the fiducial central world-line. The point for our purpose is that such lever arm appears essentially to be the analogue of a mass-type dipole moment. We thus see the possibility of describing dipolar particles in general relativity by starting from (some variant of) the Bailey--Israel action~\\cite{BI80}.  The previous paradox concerning the dipole moment which should vanish in the frame of the center-of-mass will be solved because (roughly speaking) there are two notions of dipole moments. The first moment is the one we mentionned earlier, $\\pi^i_\\mathrm{g}$, but let us call it now the ``passive'' dipole moment $\\pi^i_\\mathrm{p}$ to emphasize the fact that the gravitational mass it contains is really the \\textit{passive} gravitational mass $m_\\mathrm{p}=m_\\mathrm{g}$ (\\textit{i.e.} that mass which enters in the right-hand-side of the law of motion, in factor of the gravitational field). Assuming as before that the dipolar particle is composed of sub-particles, its passive dipole moment $\\pi^i_\\mathrm{p}\\sim\\sum m_\\mathrm{p}\\,x^i$ will indeed be zero in the particle's center of mass, by equivalence between the inertial and passive gravitational masses of the sub-particles, $m_\\mathrm{p}=m_\\mathrm{i}$. However the dipole moment we shall consider in this paper is different: this is the moment parametrizing the dipolar part of the stress-energy tensor $T^{\\mu\\nu}$ of the particles (notably its $00$ component or energy density). In general relativity the stress-energy tensor is the source for the gravitational field, so we can rightly refer to this moment as the \\textit{active} dipole moment. In conclusion, we shall have $\\pi^i_\\mathrm{p}=0$ by the equivalence principle, but we still have at our disposal the ``active gravitational'' version of the moment, call it $\\pi^i_\\mathrm{a}$, which enters the stress-energy tensor of the particles, say through a term of the form $\\sim -\\partial_i(\\pi^i_\\mathrm{a}\\delta)$ in the energy density at the NR approximation (where $\\delta$ is the Dirac function). In the following we shall investigate a specific model for the relativistic dynamics and evolution of particles endowed with a four-vector dipolar moment $\\pi^\\mu\\equiv\\pi^\\mu_\\mathrm{a}$ of the active type.\\,\\footnote{Hereafter we no longer mention that this moment can be regarded as the active one. Note that the names we are giving here and below for different types of masses and dipoles are useful for the discussion but are not required in the present formalism, which is entirely based on the action~\\eqref{S}.}  By analogy with the model of paper~I we shall denote by $2m$ the total inertial mass of the dipolar particle, equivalent to its total passive gravitational mass, \\textit{i.e.} $M_\\mathrm{i}=M_\\mathrm{p}=2m$. However, this equivalence does not mean that the motion of the dipolar particle is geodesic. Indeed, motivated by paper~I, we shall introduce a force $F^\\mu$, which will be considered as ``internal'' to the dipolar particle, and is aimed at ``stabilizing'' the dipole moment imbedded in an exterior gravitational field. The force $F^\\mu$ has a non gravitational origin, and will derive from a scalar potential function $V$ in the action. Because of the presence of this force reflecting its internal structure, the dipolar particle is not a test particle: its motion is not geodesic, its four-acceleration is non zero; the particle can be thought of as a ``rocket'', self-accelerated by the internal force $F^\\mu$. We shall find an approximate solution in which $F^\\mu$ accelerates the particle in such a way as to compensate for the local gravitational field, so that the dipolar particle stays essentially at rest in a gravitational field (like in the model of paper~I).  The monopolar part of the stress-energy tensor defines what can be regarded as the particle's active gravitational mass $M_\\mathrm{a}$. Intuition from paper~I would lead to expect that for a dipolar particle $M_\\mathrm{a}$ is zero, \\textit{i.e.} the particle does not generate any monopolar gravitational field. In the present model we shall find that this is not possible: $M_\\mathrm{a}$ cannot be zero, and in fact we shall have $M_\\mathrm{a}=V/c^2$, where $V$ is the potential function from which derives the internal force $F^\\mu$. However, thanks to the explicit factor $1/c^2$ it contains, $M_\\mathrm{a}$ turns out to be very small since it vanishes in the NR limit, \\textit{i.e.}  $M_\\mathrm{a}=\\mathcal{O}(c^{-2})$ when $c\\rightarrow\\infty$. Thus, the dipolar particle will indeed be ``purely dipolar'' in the NR limit.\\,\\footnote{Notice that the fact that the particle carries some different active and passive gravitational masses, $M_\\mathrm{a}$ and $M_\\mathrm{p}$, does not contradict the principle of conservation of the linear momentum (and the law of action and reaction). Indeed the particle's stress-energy tensor contains in addition to the monopolar part a dipolar contribution (parametrized by $\\pi^\\mu$) which ensures the satisfaction of the usual conservation laws.}  Applying the standard general relativistic coupling to gravity, we prove that the Einstein field equations reduce to a Poisson equation with a dipolar source term in addition to the density of ordinary matter, having the structure $\\sim -\\partial_i(\\pi^i_\\perp\\delta)$ in the limit $c\\rightarrow\\infty$ (where $\\pi^i_\\perp$ refers to an appropriate orthogonal projection). We are then close to the MOND equation; to recover MOND it suffices that the dipole moment be aligned with the gravitational field, and polarized in a certain way. In the present paper we shall find, under some hypothesis, some solutions which correspond to gravitational polarization, \\textit{i.e.} we shall show that the polarization scenario yielding MOND is consistent with our equations. There will be essentially a one-to-one correspondence between the internal potential scalar function $V$ entering the action and the Milgrom~\\cite{Milg1, Milg2, Milg3} function $\\mu$ (which is linked in the interpretation of paper~I with the gravitational susceptibility $\\chi$ of the dipolar medium by $\\mu=1+\\chi$).  The paper is organized as follows. We first deal with the relativistic model based on the action presented in Sect.~\\ref{secIIA}. The general equations of motion and evolution, and the stress-energy tensor, are derived in Sect.~\\ref{secIIB}. In Sect.~\\ref{secIIC} we restrict ourselves to a particular solution which corresponds to some ``equilibrium'' state for the dipole moment. For that solution we obtain in Sect.~\\ref{secIIIA} the dynamics of dipolar particles in the NR approximation where $c\\rightarrow\\infty$. By neglecting the tidal gravitational field we find in Sect.~\\ref{secIIIB} that when the fluid of dipole moments is at rest with respect to the ordinary matter distribution the medium is polarized and the MOND equation follows. Finally, we investigate in Sect.~\\ref{secIV} (still in the NR limit) the case of a stationary fluid of dipole moments aligned in the central gravitational field of a point mass. Section~\\ref{secV} summarizes and concludes the paper. In the Appendix~\\ref{appA} we show on general grounds how to vary a dipolar action functional. ",
        "conclusions": "\\label{secV}  Motivated by the quasi-Newtonian model of paper~I~\\cite{B06mond}, we interpret the phenomenology of MOND as resulting from an effect of gravitational polarization of a medium made of dipole moments aligned in the gravitational field of ordinary masses. We propose an action principle, based on the matter action~\\eqref{S}, to describe the dynamics of dipolar particles and the evolution of the dipole moments in standard general relativity. The action involves a kinetic-like term for the evolution of the dipole moment and a scalar function $V$ which is supposed to describe (at some effective level) the non-gravitational interaction between the constituents of the dipole.  The dynamical variables are the dipolar particle's space-time position $x^\\mu$, and the dipole moment four vector $\\pi^\\mu$, which are varied independently yielding the two basic equations~\\eqref{Pdot} and~\\eqref{Omegadot}. Variation with respect to the metric yields the stress-energy tensor~\\eqref{Tmunu}. We find a particular class of solutions, defined by the constraint~\\eqref{Lambda1}, corresponding to the intuitive idea of a dipole moment in equilibrium. For this class of solutions the physical dipole moment variable is the projection $\\pi_\\perp^\\mu$ orthogonal to the particle's four velocity.  The non-relativistic (NR) limit of the model is investigated next. The basic equations in this limit are (i) the equation of motion~\\eqref{eom} of the dipolar particle, (ii) an equilibrium condition~\\eqref{eveq} for the dipole moment $\\pi_\\perp^i$, (iii) the conservation law~\\eqref{continuity} of the number of particles $n$, and (iv) the Poisson equation~\\eqref{poisson} for the gravitational potential $U$. The equations (i) and (ii) are different from those of the quasi-Newtonian model of paper~I (indeed the model of paper~I violates the equivalence principle and cannot result from the NR limit of a general relativistic model). In addition we have the equations of motion of ordinary massive particles~\\eqref{eomord} and of photons~\\eqref{eomph}.  Assuming that the tidal gravitational field $\\partial_{ij}U$ can be neglected, we find that there is a solution for which the dipole moments are aligned with the gravitational field. For that solution the Poisson equation for the gravitational field reduces to the MOND equation like in paper~I, confirming the close relation between the phenomenology of MOND and the dipolar dark matter.  Finally the equations in the NR limit have been integrated in the idealized case where the fluid of dipolar particles is stationary and the dipole moments are exactly aligned in the gravitational field generated by a point mass $M$. The polarization of the medium is plotted in Fig.~\\ref{fig4} and the density of polarization masses (which represents in this model the density of dark matter accumulated around $M$) is shown in Fig.~\\ref{fig5}. There is a correspondence between the potential function $V$ in the action and the MOND function $\\mu$.  To conclude, we usually face two alternatives to the issue of dark matter: Either accept the existence of cold dark matter particles but which fail to reproduce in a natural way the rotation curves of galaxies, or postulate an \\textit{ad hoc} alteration of the fundamental theory of gravity (MOND and its relativistic extensions). In the present paper (following paper~I) we proposed a third alternative: Keep the standard law of gravity but add to the ordinary matter some non-standard dark matter in order to ``explain'' MOND. More precisely we invoke a mechanism of gravitational polarization, in which the ordinary masses (galaxies) are ``anti-screened'' by polarization masses associated with gravitational dipoles playing the role of dark matter. The dipolar dark matter particles are in equilibrium in the gravitational field because of their internal structure driven by some postulated non-gravitational force (whose fundamental origin is unknown)."
    },
    "0609/astro-ph0609399_arXiv.txt": {
        "abstract": "Thomson scattering of the Cosmic Microwave Background (CMB) on moving electrons in the outflows of Lyman Break Galaxies (LBGs) at redshifts 2--8 contributes to the small-scale CMB anisotropies. The net effect produced by each outflow depends on its level of deviation from spherical symmetry, caused either by an anisotropic energy injection from the nuclear starburst or quasar activity, or by an inhomogeneous intergalactic environment.  We find that for plausible outflow parameters consistent with spectroscopic observations of LBGs, the induced CMB anisotropies on arcminute scales reach up to $\\sim 1 \\mu$K, comparable to the level produced during the epoch of reionization. ",
        "introduction": "Several experiments to observe the Cosmic Microwave Background (CMB) anisotropies on arcminute scales are currently, or will soon be, underway \\citep{Kosowsky03,Ruhl04, Lo05}.  These experiments plan to measure the CMB power-spectrum for a spherical harmonic multipole index of $10^3 \\le \\ell \\le 10^4$ at several frequencies centred around the thermal Sunyaev-Zel'dovich null of 217 \\GHz. Photon diffusion damps the CMB anisotropies on these small scales during cosmological recombination at redshift $z\\sim 10^3$ \\citep{Silk68}, and so any observed signal must originate at much lower redshifts.  Indeed, the above experiments plan to constrain the epoch of reionization and the growth of structure in the low redshift universe \\citep{Zahn05}.  The primary physical mechanism which is responsible for the small scale CMB anisotropies is Thomson scattering of CMB photons off moving electrons. Any peculiar velocity induces a Doppler anisotropy of the scattered radiation along the direction of motion.  This accounts for the CMB anisotropies produced by the peculiar velocities of clusters (the so-called {\\it kinetic Sunyaev Zel'dovich effect}) \\citep{Zeldovich80}, by peculiar velocities of linear overdensities in the intergalactic medium (the so-called {\\it Ostriker-Vishniac effect}) \\citep{Ostriker86, Vishniac87} and by the peculiar velocities of the fluctuations in the ionization fraction during patchy reionization \\citep{Gruzinov98}.  In this {\\it Letter} we examine the contribution of outflows in Lyman Break Galaxies (LBGs) to the small-scale CMB anisotropies. LBGs are believed to be the ancestors of present-day luminous elliptical galaxies. They are observed to produce gas outflows with velocities of several hundred \\kms (see \\cite{Giavalisco02} for a comprehensive review). In contrast with the traditional kinetic Sunyaev-Zel'dovich effect where the bulk velocity of the virialized gas is responsible for the induced CMB anisotropy, we focus here on the Doppler effect of the outflowing gas and ignore any bulk motion of the LBG as a whole (which produced a smaller effect at the redshifts of interest). This bulk effect is included in standard calculations of the non-linear generalization of the Ostiker-Vishniac effect \\citep{Hu00}.  The contribution from a single LBG to the fractional temperature fluctuation of the CMB can be expressed as \\be \\label{los_int} \\frac{\\Delta T}{T} = - \\int dl \\sigma_T n_e\\frac{\\vec{n} \\cdot \\vec{v}}{c}, \\ee where $\\sigma_T = 6.65 \\times 10^{-25}$ \\cm is the Thomson cross section, $n_e$ is the electron number denisty, $\\vec{v}$ is the electron peculiar velocity, $c$ is the speed of light and $\\hat{n}$ is the observer's line-of-sight toward the LBG \\citep{Zeldovich80}. The integration traces the photon's path through the LBG outflow.  The radial extent of the outflow is found by solving the corresponding hydrodynamics equations. These coupled non-linear partial differential equations can be reduced to a single ordinary differential equation \\citep{Tegmark93, Furlanetto03} under the assumption that the gas swept-up by the outgoing blast wave lies in a thin shell behind the propagating shock front (the so-called {\\it thin shell approximation}). The validity of this approximation is illustrated by the self-similar Sedov-Taylor-von Neumann solution for a point explosion in which $90\\%$ of the swept-up mass resides in a shell of thickness $10\\%$ of the outflow's radius (see Ostriker \\& McKee 1988 and Ikeuchi et al. 1983 for additional discussion on this approximation).  The thin shell approximation allows us to treat the radiative transfer of CMB photons through the shock front in a plane parallel geometry. When the thickness of the shock front is small compared with the shock front's radius of curvature, the path length through the shock front can be expressed as \\be \\delta l \\approx \\frac{\\delta R}{| \\hat{n}\\cdot \\hat{v} |}, \\ee where $\\delta R$ is the thickness of the shock front. Using conservation of mass and the density compression ratio for a strong adiabatic shock, one gets $\\delta R/R = (\\gamma -1)/(3\\gamma + 3)$, where $R$ is the radius of the outflow and $\\gamma$ is the adiabatic index of the gas \\citep{Ostriker88}.  Within the thin shell approximation the line-of-sight integration in Eq. (\\ref{los_int}) is simplified to \\be \\left(\\frac{\\Delta T}{T} \\right)_i = - \\frac{\\sigma_T~\\delta R}{c}~[n_e(\\vec{r}_1) v(\\vec{r}_1) - n_e(\\vec{r}_2) v(\\vec{r}_2)], \\ee where $\\vec{r}_1$ and $\\vec{r}_2$ are the location, in a coordinate system centred on the LBG, on the shock front where the line-of-sight respectively enters and exits the shock front. The subscript $i$ labels the contribution from a particular LBG.  If the outflow is spherically symmetric then there is an exact cancellation of the anisotropies produced at the entry and exit points of the line-of-sight through the LBG shock front\\footnote{Note that even if the outflow had a perfect spherical shape, the finite light crossing time through the outflow would produce a net non-zero signal because the flow parameters are time-dependent.  This effect would produce a signal $\\sigma_T n_e R (\\Delta v)/c \\sim \\sigma_T n_e (dv/dt) (R/c)^2$ that is extremely small and is ignored here.}. However, observations of low-redshift starburst galaxies, which serve as analogs of the higher redshift LBGs, show evidence for highly non-spherical outflow geometries \\citep{Martin99}. It is therefore reasonable to expect that the early stages of LBG outflows, whether they are driven by starburst or quasar activity\\footnote{In this {\\it Letter} we will only include the feedback driven by supernovae. Energy input from quasars will lead to enhancements in the bubble size and outflow velocity [see \\cite{Furlanetto01} for a description of quasar outflows].}, would produce CMB anisotropies.  For simplicity, we will assume that the outflow is axisymmetric about some axis $\\hat{z}$ (possibly the rotation axis of the galactic disc or possibly the jet axis of a central quasar) and expand both the outflow velocity and shock front density in Legrendre polynomials, $ \\vec{v}(\\vec{r}) = \\hat{r}~\\sum_{\\ell} v_{\\ell}(r) P_{\\ell}(\\hat{z} \\cdot \\hat{r})$ and $n_e(\\vec{r}) = \\sum_{\\ell} n_{\\ell}(r) P_{\\ell}(\\hat{z} \\cdot \\hat{r})$, where we have assumed that the velocity field is radial ($\\hat{v} \\equiv \\hat{r}$).  Even if the outflow began highly collimated, it will eventually isotropize as it propagates into the intergalactic medium (IGM; an analogous tendency exists in relativistic flows, see Ayal \\& Piran 2001). On longer timescales, IGM inhomogeneities, such as filaments and voids, will once again make the outflows non-spherical. Numerical simulations are needed in order to properly model these effects. Here we parameterize the level of asphericity in the outflow by a coefficient $\\epsilon$ (see Eq. \\ref{eq:dT2} below).  The outline of this {\\it Letter} is as follows.  In \\S 2 we analyze the CMB anisotropy induced by a single LBG with an arbitrarily non-spherical outflow. In \\S 3 we calculate the resulting CMB power spectrum. In \\S 4 we present numerical results and in \\S 5 we summarize our conclusions. Throughout our discussion, we will assume the WMAP3 cosmological model \\citep{Spergel06}. ",
        "conclusions": "The contribution of outflows from Lyman-break galaxies to the CMB power-spectrum on arcminute scales is proportional to the square of their characteristic level of deviation from sphericity, $\\epsilon^2$. Future CMB experiments could therefore calibrate the intricate feedback process of galactic outflows on the IGM. Most tools used to study these feedback processes focus on the inner few \\kpc ~of the LBGs, even though the shock front is typically located at several of tens or hundreds of \\kpc. The secondary CMB anistropies calculated in this work provide a unique probe of these extended perturbed regions around starburst galaxies at high redshifts.  Even though the amplitude of the power spectrum produced by this effect is small, the signal has distinctive spectral and spatial characteristics. The power spectrum has the same scaling with Fourier multipole ($\\propto \\ell^2$) as radio or IR point sources. However, the frequency dependence of the anisotropies produced by our effect has the standard blackbody spectrum, whereas the radio point sources have a power law frequency spectrum. This distinct feature of our effect may allow future small scale experiments, many of which have excellent frequency coverage, to separate the anisotropies produced by LBGs from radio point sources.  \\bigskip  {\\bf Acknowledgements}  This work was supported in part by Harvard university grants. D.B. thanks the Harvard Institute for Theory and Computation for its hospitality when this work was completed, and acknowledges helpful conversations with Re'em Sari and Niayesh Afshordi. We also thank Matt McQuinn for making available the numerical results from his work."
    },
    "0609/astro-ph0609350_arXiv.txt": {
        "abstract": "We study convective overshooting by means of local 3D convection calculations. Using a mixing length model of the solar convection zone (CZ) as a guide, we determine the Coriolis number (Co), which is the inverse of the Rossby number, to be of the order of ten or larger at the base of the solar CZ. Therefore we perform convection calculations in the range Co = 0...10 and interpret the value of Co realised in the calculation to represent a depth in the solar CZ. In order to study the dependence on rotation, we compute the mixing length parameters $\\alpha_{\\rm T}$ and $\\alpha_{\\rm u}$ relating the temperature and velocity fluctuations, respectively, to the mean thermal stratification. We find that the mixing length parameters for the rapid rotation case, corresponding to the base of the solar CZ, are 3-5 times smaller than in the nonrotating case. Introducing such depth-dependent $\\alpha$ into a solar structure model employing a non-local mixing length formalism results in overshooting which is approximately proportional to $\\alpha$ at the base of the CZ. Although overshooting is reduced due to the reduced $\\alpha$, a discrepancy with helioseismology remains due to the steep transition to the radiative temperature gradient.  In comparison to the mixing length models the transition at the base of the CZ is much gentler in the 3D models. It was suggested recently (Rempel \\cite{Rempel2004}) that this discrepancy is due to the significantly larger (up to seven orders of magnitude) input energy flux in the 3D models in comparison to the Sun and solar models, and that the 3D calculations should be able to approach the mixing length regime if the input energy flux is decreased by a moderate amount. We present results from local convection calculations which support this conjecture. ",
        "introduction": "Convection poses a difficult problem for stellar structure modelling. One-dimensional stellar structure models require a parameterization of convection in order to be able to yield the thermal stratification within the CZ. Since convection in the stellar envelopes is in general highly efficient, the stratification is close to adiabatic in much of the CZ and thus a detailed description of convection is not needed there. The most often used way to parameterize convection is to use the mixing length concept in a form introduced by \\cite{Vitense1953}; see also \\cite{BoehmVitense1958}, which considers convective elements to lose their identity after rising (or descending) the so-called mixing length which is proportional to the local pressure scale height, i.e. \\begin{equation} l = \\alpha H_{\\rm p}\\;.\\label{equ:mlt} \\end{equation} Using this basic assumption, it is possible to derive equations that relate the velocity and temperature fluctuations to the mean stratification, and thus, to compute the convective energy flux (see, e.g. Chapter 6 of Stix \\cite{Stix2002}).  Among other conceptual problems, the mixing length formalism neglects the effects of rotation although convection can be significantly influenced by it in regions of the CZ where the turnover time is longer than the rotation period. In a recent study, K\\\"apyl\\\"a et al. (\\cite{Kaepylaeea2005}) estimated the Coriolis number, which is the inverse of the Rossby number, in the solar CZ from \\begin{equation} {\\rm Co} = {\\rm Ro}^{-1} = 2\\,\\Omega_\\odot \\tau = 2\\,\\Omega_\\odot \\alpha H_{\\rm p}/u\\;, \\label{equ:Co} \\end{equation} where $\\alpha = 1.66$, $\\Omega_\\odot = 2.6 \\cdot 10^{-6}$\\,s$^{-1}$, and $u$ the convective velocity obtained using a local mixing length model. This relation gives values of the order of $10^{-3}$ near the solar surface and of the order of ten or larger near the base of the CZ (see figure \\ref{fig:comlt}). We compute the mixing length parameters relating the velocity and temperature fluctuations to the mean stratification from local 3D convection calculations in the range ${\\rm Co} = 0 \\ldots 10$, which coincides with the range expected in the Sun. Thus it is possible to probe the influence of rotation on the mixing length relations and take the effect implicitly into account in a solar model employing a non-local formulation of the mixing length concept in order to study overshooting below the CZ.  \\begin{figure} \\begin{centering} \\includegraphics[height=3.5in,width=3.5in]{comlt.eps} \\caption{Coriolis number in the solar convection zone according to (\\ref{equ:Co}).}\\label{fig:comlt} \\end{centering} \\end{figure}  There is a striking difference between the almost adiabatic overshooting with a very sharp transition to the radiative gradient seen in non-local mixing length models in comparison to the much more subadiabatic and smoother overshooting seen in 3D convection models. In a recent paper Rempel (\\cite{Rempel2004}) suggested that this discrepancy arises due to the fact that the two models are simply working in different parameter regimes in the sense that the input energy flux in the 3D models is usually up to $10^7$ times larger in comparison to the non-local mixing length models and the Sun. We have performed 3D numerical calculations in which we decrease the input flux by two orders of magnitude in order to study this effect in more detail.  The remainder of the paper is organised as follows: in~\\S\\,\\ref{sec:nummodel} a brief description of the model is given and in~\\S\\,\\ref{subsec:mltrot} and \\S\\,\\ref{subsec:inpute} the results concerning the effects of rotation and input energy flux on convective overshooting are presented. Finally, the main results of the study and remaining problems are summarised in~\\S\\,\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl} Three dimensional local convection calculations were used to probe the effects of rotation on mixing length coefficients relating the temperature and velocity fluctuations to the mean thermal stratification. It was found that when the rotational influence on the flow is comparable to that expected in the deep layers of the solar CZ, the mixing length parameters are reduced by a factor of three to five. If a depth-dependent mixing length $\\alpha$ is introduced into a solar model, the overshooting at the base of the CZ is reduced approximately in proportion to the reduction of $\\alpha$. Although the depth of the solar CZ can be correctly reproduced in this way, the steep transition to the radiative gradient should still be visible in the helioseismic inversions.  The overshooting in 3D convection calculations is much more subadiabatic with a smooth transition, which is due to the much higher input energy flux used to meet the time step constraints. In the present study we show that decreasing the input flux in 3D calculations leads to more adiabatic overshooting and sharper transition at the base of the CZ. Although this result seems to suggest that the 3D calculations will approach the mixing length regime when the flux is reduced enough, one must bear in mind that in the present models the stratification is rather weak in comparison to the Sun. Thus the effects of compressibility are likely to be underestimated, and lead to too large filling factor for the downflows. Furthermore, the spatial size of the downflow plumes is restricted by the grid size, and it is probable that the filling factor of downflows further decreases when the resolution is increased.  To summarize, we stress the point that rotation should be taken into account in models of convective overshooting since it can exert considerable influence on convection already in slowly rotating stars such as the Sun. Furthermore, high resolution numerical studies of deep convection are needed in order to study whether the convective overshooting is due to very few strong downflows, producing nearly adiabatic overshoot region with a steep transition, or whether downflows of different strengths penetrate into the stable region in a larger area producing smooth overshooting required by helioseismology."
    },
    "0609/astro-ph0609621_arXiv.txt": {
        "abstract": "Using ray tracing for simple analytic profiles, we demonstrate that the lensing cross section for producing giant arcs has distinct contributions due to arcs formed through image distortion only, and arcs form from the merging of two or three images.  We investigate the dependence of each of these contributions on halo ellipticity and on the slope of the density profile, and demonstrate that at fixed Einstein radius, the lensing cross section increases as the halo profile becomes steeper.  We then compare simulations with and without baryonic cooling of the same cluster for a sample of six clusters, and demonstrate that cooling can increase the overall abundance of giant arcs by factors of a few.  The net boost to the lensing probability for individual clusters is mass dependent, and can lower the effective low mass limit of lensing clusters.  This last effect can potentially increase the number of lensing clusters by an extra $50\\%$.  While the magnitude of these effects may be overestimated due to the well known overcooling problem in simulations, it is evident that baryonic cooling has a non-negligible impact on the expected abundance of giant arcs, and hence cosmological constraints from giant arc abundances may be subject to large systematic errors. ",
        "introduction": "Soon after the discovery of giant arcs \\citep[][]{lyndspetrosian86,soucailetal87} and their subsequent interpretation as gravitational lensing events \\citep[][]{paczynski87}, \\citet[][]{grossmannarayan88} realized that the statistics of these relatively rare events can be used as a cosmological probe.  This idea was greatly expanded by \\citet[][]{miraldaescude93} and \\citet[][]{wuhammer93} who used arc statistics to constrain the mass distribution of clusters and the high redshift galaxy population \\citep[][]{miraldaescude93b,bezecourtetal98}.  Attempts to derive cosmological parameters from arc abundances soon followed, and concluded that giant arc abundances had a relatively mild cosmological dependence \\citep[][]{wumao96}.  These early studies used simple analytical models of cluster mass distributions, which turn out to severely under-predict lensing probabilities relative to clusters in numerical simulations \\citep[][]{bartelmannweiss94,bartelmannetal95}.  Nevertheless, it seemed reasonable to expect that the {\\em scaling} of arc abundances with cosmology should be broadly consistent with predictions from analytic prescriptions, so it was surprising when the first numerical study to investigate arc abundances in various cosmologies \\citep[][]{bartelmannetal98} found not only that there were order of magnitude differences in the predictions from different cosmologies, but also that the abundance of giant arcs expected in the now-standard $\\lcdm$ cosmology is an order of magnitude lower than the observed abundance \\citep[][]{lefevreetal94}.  The claimed sensitivity of giant arc abundances to cosmology and the strong inconsistency between theory and observations found by \\citet[][]{bartelmannetal98} generated an explosion of both theoretical and observational studies.  On the observational front, new arc samples with more clearly defined selection functions became available \\citep[][]{luppinoetal99,zaritskygonzalez03,gladdersetal03}. At the same time, the cosmological community mounted an extensive theoretical effort to search for possible systematics and/or theoretical uncertainties in order to determine whether the disagreement was real.  The theoretical program has revealed that the problem of giant arc statistics is extremely rich, with predictions being sensitive to details of both background galaxies and cluster properties. Concerning sources, many models for giant arc abundances have assumed circular sources, but source ellipticity can significantly boost the number of giant arcs in the sky \\cite[][]{bartelmannetal95,keeton01b}. Assumptions about the distribution of source radii appear to be important \\citep[][]{miraldaescude93b,oguri02,howhite05}.  The redshift distribution of sources plays a significant role \\citep[][]{wambsganssbodeostriker04,dalalholderhennawi04} and in fact contributes much of the theoretical uncertainty in modern predictions \\citep[][]{hennawietal05}.  There are also several systematics associated with the lens population.  For instance, while many studies have assumed spherically symmetric mass distributions \\citep[e.g.,][]{hattorietal97,hamanafutamase97,molikawaetal99, cooray99,cooray99b,williamsetal99}, many others have pointed out that halo ellipticity significantly increases giant arc abundances \\citep[e.g.,][]{grossmannarayan88,bartelmannetal95,oguri02,meneghettietal03}. Moreover, halo triaxiality can give rise to order of magnitude variations in the lensing probability of a given cluster with viewing angle \\citep[][]{ogurietal03,dalalholderhennawi04}, and therefore must be treated properly in arc abundance predictions.  Lensing probabilities for individual clusters appear to be significantly enhanced during mergers \\citep[e.g.,][]{torrietal04,meneghettietal05b,fedelietal06}, though such episodes appear to be rare and transient enough that the overall lens populations is not significantly biased towards a cluster merging population \\citep[][]{hennawietal05}.  One final related trend, noted in both analytical and numerical works, is that lensing probabilities are very sensitive to the radial density profiles of lensing clusters \\citep[e.g.,][]{miraldaescude93,hattorietal97,ogurietal01, meneghettietal05,hennawietal05}.  Consequently, clear predictions for the distribution of cluster density profiles must be in place for arc abundance predictions to become robust.  Systematics that do {\\em not} seem to have significant effects on lensing probabilities have also been identified.  For instance, studies in which a central galaxy is ``painted'' on the center of simulated clusters suggest that this additional mass component has minimal impact on the expected giant arc abundance \\citep[][]{molikawahattori01,meneghettietal03b,dalalholderhennawi04}. Likewise, the presence of substructures and/or galaxies in the cluster also seems to have a negligible effect on lensing probabilities \\citep[][]{grossmannarayan88,floresetal00,meneghettietal00,hennawietal05}. Line of sight projection effects \\citep[][]{wambsganssetal05} also appear to be unimportant \\citep[][]{hennawietal05}.  At present, there seems to be general agreement between theory and observations \\citep[][]{dalalholderhennawi04,hennawietal05,horeshetal05}, although the situation is not completely clear \\citep[][]{lietal05}.  The new agreement arises from the coherent contributions of various effects, with no single systematic explaining the order of magnitude effects observed by \\citet[][]{bartelmannetal98}.  Recent analytical \\citep[][]{kaufmannstraumann00,bartelmannetal03} and numerical \\citep[][]{meneghettietal05} studies suggest that variations in the arc abundance between different cosmologies are factors of a few rather than an order of magnitude, further relaxing any real tension between theory and observations.  What seems clear at this time is that arc abundances are closer to old cosmology than to precision cosmology: agreement is to be understood at the factor of two level, even if current observations did not suffer from small number statistics.  Given the complexity of the problem, it is important to consider whether there are any additional systematic effects that alter theoretical predictions.  The goal of this paper is to study one such possibility: the impact of baryonic cooling on cluster mass profiles. Numerical simulations that include cooling have shown that once baryons have cooled and sunk to the center of the halo to form the cluster's central galaxy, the dark matter halo itself contracts in response to the increased central density \\citep[][]{blumenthaletal96,kochanekwhite01, gnedinetal04,sellwoodmcgaugh05}.  The effects of this contraction are significant, and include a steepening of the inner regions of the halo profile, a boost to the mass enclosed within a fixed radius $r$, and an increase of the degree of spherical symmetry of the mass distribution \\citep[][]{kazantzidisetal04}.  Since arc abundances are known to be sensitive to cluster density profiles we expect that baryonic cooling may significantly affect the predictions.  Indeed, recent work by \\citet[][]{puchweinetal05} has shown that baryonic cooling enhances the lensing probability of very massive clusters ($M\\gtrsim 10^{15}\\ \\msun$) by a factor of $\\sim$2--4. \\citet[][]{puchweinetal05} did not consider, however, the effect of baryonic cooling on smaller clusters, which is particularly significant given that the median mass of the cluster lens sample \\citep[see][]{hennawietal05} is well below that considered in \\citet[][]{puchweinetal05}.  In this work, we extend the analysis of the impact of baryonic cooling on arc statistics to clusters of considerably smaller masses.  In particular, we demonstrate that baryonic cooling affects lensing probabilities for low mass clusters more than for high mass clusters, both because low mass clusters have a larger cooled mass fraction, and because lensing in low mass clusters is more confined to the central regions that are strongly affected by cooling.  We also demonstrate that baryonic cooling has a strong isotropization effect for both the lensing probability of individual cluster as a function of viewing angle, and the spatial distribution of arcs around the cluster center. In order to understand {\\em why} baryonic cooling can have such significant impact on the lensing cross section, we develop a theoretical framework that allows us to interpret the lensing probability associated with each cluster as a function of the minimum length to width ratio of the arcs considered.  We have found that our theoretical framework can be used to better interpret and contextualize results from previous works, providing an intuitive base upon which the sensitivity of arc statistics to various cluster parameters can be deduced in a qualitative fashion. ",
        "conclusions": "\\label{sec:summary}  Using simple analytic halo density profiles, we have investigated how lensing cross sections depend on various halo parameters.  In particular, we found that the lensing cross section has three distinct contributions corresponding to arcs formed through image distortion only, merging of two images, and merging of three images.  Moreover, $\\sigma(\\lw)$ exhibits a knee at the scale which the cross section goes from being distortion dominated to image merging dominated, and the latter component falls exponentially fast with increasing length to width ratio.  We then proceeded to investigate how these various contributions to the lensing cross section depend on halo properties.  In particular, we found that while the image merging contribution to the cross section increases with increasing halo ellipticity, the distortion contribution remains relatively constant, except for the fact that it gets truncated at a smaller transition scale $\\lw_t$.  In other words, $\\lw_t$ decreases with increasing ellipticity.  We also investigated the impact of the slope of the profile on the lensing cross section, and found that at fixed Einstein radius, steeper profiles result in more effective lenses, and more circular distribution of arcs about the cluster center.  We also noted that the ratio of cusp arcs to fold arcs appears to be very sensitive to the slope of the density profile.  Based on these observations, we argued that the probability for finding arcs of a given length to width ratio should quickly rise with mass below some mass threshold $M_{min}$, and that above this mass cut the scaling should become much flatter.  We also discussed how the distinction between arcs formed through image merging and arcs formed through distortion leads to a non-trivial dependence of the lensing cross section on the assumed source radius.  We also investigated the impact of baryonic cooling on the lensing cross section, and found that the steepening of the halo density profile in response to the central baryon condensation enhances the lensing cross section at low $\\lw$ by about a factor of three, and lowers the mass threshold $M_{min}$ for halos to become effective lenses.  The latter effect helps further increase arc abundances by about $50\\%$.  Moreover, cooling clusters had a significantly larger probability of forming fold arcs than adiabatic clusters, and their lensing cross section was much less dependent upon the line of sight projection axis than adiabatic clusters.  Both of these effects are also explained by the steepening of the halo profile brought about by the contraction of the halo in response to the condensation of cooling baryons at the center.  Finally, we also argued that substructures did not appear to significantly enhance lensing cross sections, and that the multiplicity ratio, i.e.  the number of lenses that exhibit multiple arcs, is only moderately affected by the details of the mass distribution of the lenses.  Given the sensitivity of arc abundances to baryonic processes, it is difficult to see this observable becoming a tool for precision cosmology in the near future.  Conversely, as the range of cosmological parameters get narrower from other observations, this same sensitivity can be used to probe cluster properties and/or the properties of high redshift galaxies.  Concerning this last possibility, it seems clear that a systematic analysis similar to the one carried here aimed at understanding how various source properties affect lensing cross section would be a fruitful endeavor which we intend to carry out.  \\it Acknowledgments: \\rm E.R. would like to thank Scott Dodelson for a careful reading and helpful suggestions throughout the development of this work.  This work was supported in part by the Kavli Institute for Cosmological Physics through the grant NSF PHY-0114422.  DN is supported by the Sherman Fairchild Postdoctoral Fellowship at Caltech. AVK is supported by the National Science Foundation (NSF) under grants No. AST-0206216 and AST-0239759 and by NASA through grant NAG5-13274. The authors would also like to thank KICP Thunch, where many conversations related to this topic took place."
    },
    "0609/astro-ph0609767_arXiv.txt": {
        "abstract": "We present a simplified solution to orbit determination of a binary system from astrometric observations. An exact solution was found by Asada, Akasaka and Kasai by assuming no observational errors. We extend the solution considering observational data. The generalized solution is expressed in terms of elementary functions, and therefore requires neither iterative nor numerical methods. ",
        "introduction": "In the seventeenth century, Kepler discovered the laws of the motion of celestial objects. Since then, a fundamental problem has been how to determine the orbital elements (and the mass) of a binary from observational data of its positions projected onto a celestial sphere. In fact, astrometric observations play an important role in astronomy through determining a mass of various unseen celestial objects currently such as a massive black hole (Miyoshi et al. 1995), an extra-solar planet (Benedict et al. 2002) and two new satellites of Pluto (Weaver et al. 2006).  The orbit determination of {\\it resolved double stars} was solved first by Savary in 1827, secondly by Encke 1832, thirdly by Herschel 1833 and by many authors including Kowalsky, Thiele and Innes (Aitken 1964 for a review on earlier works; for the state-of-the-art techniques, e.g, Eichhorn and Xu 1990, Catovic and Olevic 1992, Olevic and Cvetkovic 2004). Here, resolved double stars are a system of two stars both of which can be seen. The relative vector from the primary star to the secondary is in an elliptic motion with a focus at the primary. This relative vector is observable because the two stars are seen. On the other hand, an {\\it astrometric binary} is a system of two objects where one object can be seen but the other cannot like a black hole or a very dim star. In this case, it is impossible to directly measure the relative vector connecting the two objects, because one end of the separation of the binary, namely the secondary, cannot be seen. The measures are made in the position of the primary with respect to unrelated reference objects (e.g. quasars )whose proper motion is either negligible or known.  There are two major problems in determination of a binary orbit: (1) observational errors and (2) the conditional equations which connect the observable quantities with the orbital parameters. The conditional equations are not only non-linearly coupled but also transcendental because of Kepler's equation (e.g., Goldstein 1980, Danby 1988, Roy 1988, Murray and Dermott 1999, Beutler 2004).  As a method to determine the orbital elements of a binary, an analytic solution in an explicit form has been found by Asada, Akasaka and Kasai (2004, henceforth AAK) by assuming no observational errors. This solution is given in a closed form by requiring neither iterative nor numerical methods. The purpose of this brief article is to extend the AAK solution considering observational data. We will summarize the AAK solution in $\\S$ 2, next extend the solution considering observational data in $\\S$ 3, and make a computational test of the extended solution in $\\S$ 4. ",
        "conclusions": "In this paper, as a simplified method to determine a binary orbit, we extend the solution considering observational data. We expect that the present formula will be used for orbit determinations by some planned astrometric space missions such as SIM\\footnote{http://sim.jpl.nasa.gov/} (Shao 2004), GAIA\\footnote{http://astro.estec.esa.nl/GAIA/} (Mignard 2004, Perryman 2004) and JASMINE\\footnote{http://www.jasmine-galaxy.org/} (Gouda et al. 2004), which will observe with the accuracy of possibly a few micro arcseconds a number of binaries whose components are a visible star and a companion ranging from a black hole to an extra-solar planet. Our numerical computations correspond to a case that by these missions we observe a binary system at a distance of 1 kpc from us, where the semimajor axis of the primary's orbit is around 1AU."
    },
    "0609/astro-ph0609284_arXiv.txt": {
        "abstract": "The premise that Ultraluminous X-ray sources (ULXs) exist beyond the optical extent of nearby galaxies is investigated. A published catalog containing 41 ULX candidates located between 1 and $\\sim$3 times the standard \\DTF\\ isophotal radius of their putative host galaxies is examined. Twenty-one of these sources have spectroscopically-confirmed distances. All 21 are background objects giving a 95\\% probability that at least 37 of the 41 candidates are background sources. Thirty-nine of the 41 sources have X-ray-to-optical flux ratios, $-1.6<\\log(F_{\\rm X}/F_{\\rm O})<+1.3$, consistent with those of background active galactic nuclei. (The remaining two are not detected in optical to a weak limit of $m_B \\sim 21.5$~mag corresponding to $\\log(F_{\\rm X}/F_{\\rm O})\\gtrsim 1.6$.) The uniform spatial distribution of the sample is also consistent with a background population. This evidence suggests that ULXs rarely, if at all, exist beyond the distribution of luminous matter in nearby galaxies and, as a consequence, there is no correlation between the population of ULXs and halo objects such as old globular clusters or Pop~III remnants. ",
        "introduction": "%  Ultraluminous X-ray sources (ULXs) are point-like non-nuclear sources in nearby galaxies with apparent X-ray luminosities $L_{\\rm X} \\gtrsim 10^{39}$~\\ergl. Their nature is still a mystery: if they are accreting sources at the distance of their putative host galaxies, then their high luminosities require either beamed emission geometries, or super-Eddington emission rates, or accretion onto compact objects more massive than predicted by stellar evolution models. See Fabbiano (2006) for a review of ULX theory, observation, and their importance in studies of extragalactic X-ray source populations.  Distances to candidate ULXs can, in principal, be determined spectroscopically if they have bright optical counterparts. Because spectroscopy requires high contrast between the optical counterpart and the underlying galaxy light, most spectroscopic measurements made to date have been of bright counterparts located in the low surface brightness regions at relatively large distances from the centers of their putative host galaxies. The fact that all these spectroscopically-examined objects have been confirmed\\footnote{Confirmation formally applies only to the optical source but, under these circumstances, the probability of a chance positional coincidence is small and it is likely the X-ray and optical sources are physically related.} to be background sources instead of true ULXs suggests ULXs may not exist beyond the luminous regions of galaxies and hence that there is no association between ULXs and halo objects. ",
        "conclusions": "For the sample analysed here, the overwhelming majority, and perhaps all, of the ULX candidates located beyond the optical extent of their host galaxies are unrelated background objects. This, in itself, is not remarkable since it must be true when \\fdtf$\\gg$1. What is of interest is the strong distinction between sources inside and outside of \\fdtf$\\sim$1. The physical processes that favor production of ULXs within the optical extent of a galaxy must % operate rarely, if ever, in galactic halos. Since the most relevant X-ray emission mechanism here is accretion onto a compact object, the lack of halo ULXs means either compact objects are rare or suitable accretion reservoirs are absent in galactic halos. What are the implications of this for ULX models?  There are two classes of object that have been recognized as potential sources of halo compact objects and in particular as candidate ULXs. These are the remnants of zero-metallicity Population~III stars and the (possibly massive) black holes nurtured in the cores of globular clusters.  Remnants of the first generation of stars may be massive black holes and many could exist in galactic halos (Madau \\& Rees 2001; Islam \\etal\\ 2004a). In order to appear as ULXs, these remnants must be accreting at high rates~--~through a thin disk instead of by Bondi-Hoyle accretion (King \\etal\\ 2001)~--~from the interstellar medium or from their own bound relic minihalos. Their large masses, $\\gtrsim$300~\\msun, imply they should be strong UV/optical sources with $\\log(F_{\\rm X}/F_{\\rm O})\\ll1$ (Islam \\etal\\ 2004b; Volonteri \\& Perna 2005). There are no ULX candidates in the present sample that meet these conditions.  Globular clusters are known to host X-ray binaries; in fact, dynamical interactions in the dense cluster environment enhances close binary formation. ULXs in globular clusters would, in general, have $\\log(F_{\\rm X}/F_{\\rm O})\\gtrsim2$ even for the fainter X-ray sources (e.g., Kundu \\etal\\ 2002; Sarazin \\etal\\ 2003). But this is close enough to the optically-faint end of the observed range for \\fdtf$>$1 objects that a globular cluster host cannot be completely excluded from the present sample.  The importance of these two ULX scenarios is that Pop~III remnants and the dense stellar environments of globular clusters are two potential sites for the elusive intermediate-mass black holes sought to fill the gap between stellar-mass and supermassive black holes (van~der~Marel 2004). In Pop~III stars, these form from direct collapse of massive stars. In globular clusters, they may form through the gradual accrual of mass over the long lifetime of the cluster, primarily in the form of accreted black holes  (Miller \\& Hamilton 2002) or gas accretion (Kawakatu \\& Umemura 2005), onto a seed black hole of initial mass large enough to prevent ejection from the cluster through recoil. Alternatively, they may form at the birth of the cluster through runaway stellar or proto-stellar mergers (Portegies~Zwart \\etal\\ 2004). (It is unclear if these latter clusters survive to become old globular clusters in galactic halos or if their compact objects would appear as ULXs at this late evolutionary stage). This work has shown, whether such objects exist or not, they rarely, if ever, appear as ULXs."
    },
    "0609/astro-ph0609666.txt": {
        "abstract": "The far-infrared (FIR) emissivity of dust is an important parameter characterizing the physical properties of the grains. With the availability of stellar databases and far-infrared data from Infrared Space Observatory (ISO) it is possible to compare the optical and infrared properties of dust, and derive the far-infrared emissivity with respect to the optical extinction. % In this paper we present the results of a systematic analysis of the FIR emissivity of interstellar clouds observed  with ISOPHOT (the photometer onboard ISO) at least at two infrared wavelengths, one close to $\\sim$100$\\mu$m and one at 200$\\mu$m. We constructed FIR emission maps, determined dust temperatures, created extinction maps using 2MASS survey data, and calculated far-infrared emissivity for each of these clouds. We present the largest homogeneously reduced database constructed so far for this purpose. During the data analysis special care was taken on possible systematic errors. We find that far-infrared emissivity has a clear dependence on temperature. % The emissivity is enhanced by a factor of usually less than 2 in the low dust temperature regime of 12K$\\le$T$_{\\rm d}$$\\le$14K. This result suggests larger grain sizes in those regions. However, the emissivity increase of typically below 2 restricts the possible grain growth processes to ice-mantle formation and coagulation of silicate grains, and excludes the coagulation of carbonaceous particles on the scales of the regions we investigated. % In the temperature range \\lele{14\\,K}{T$_{\\rm d}$}{16\\,K} a  systematic decrease of emissivity is observed with respect to the  values of the diffuse interstellar matter. Possible scenarios for this behaviour are discussed in the paper. % ",
        "introduction": "%Dust emissivity  The far-infrared (FIR) emissivity of interstellar grains predominantly emitting at wavelengths in excess of 100\\,$\\mu$m can now be determined in many molecular and moderate density regions thanks to the increasing availability of far-infrared, sub-mm, and massive stellar data sets. % Several studies have reported on the detection of an enhancement of the FIR emissivity in the 100-200\\,$\\mu$m wavelength range in regions of higher column density compared to the dust emissivity in the diffuse interstellar medium associated with HI (Bernard et al., 1999; Cambr\\'esy et al., 2001; Juvela et al., 2002; Stepnik et al., 2003; Lehtinen et al., 2004, Rawlings et al., 2005). % The largest sample of different regions is presented in del Burgo et al. (2003) presenting a study of eight translucent clouds.  The observations suggest a trend where the emissivity increases with decreasing temperature. % The variation is attributed to a change in grain properties which is expected to take place in denser environments. % In particular, the increase in emissivity is interpreted as a signature of an enhancement in grain size.  However, the trend shows a large scatter which might be due to the observational limitations which increase the uncertainties. % % note: there may be a scatter due to the nature of the % sources and regions investigated. % On one hand, star counts statistically probe only a limited extinction range and, on the other hand, the determination of the grain temperatures in the infrared can have large uncertainties. % % The variation was inferred from the comparison % of the ratio between far-infrared or submm opacity % and visual extinction in different regions. % % Although simple two-compenent models predict a continous % increase of emissivity with the decreasing temperature, % observations show only a modest increase of $\\epsilon$ % for low ($\\sim$12K) temperatures. %  In this study we present a large sample based on Infrared Space Observatory (ISO) data of cloud regions with reliable values of the FIR dust emissivity and optical extinction data.  We examine in detail the possible observational and data processing errors to ensure the reliability of our results. A general good agreement was found between our results and previous studies of individual regions. Due to the large number of data points we are able to put stronger constraints on the emissivity changes with temperature.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": ""
    },
    "0609/astro-ph0609551_arXiv.txt": {
        "abstract": "We give two examples of spiral galaxies that show non-circular gas motions in the inner kiloparsecs, from SAURON integral field spectroscopy. We use harmonic decomposition of the velocity field of the ionized gas to study the underlying mass distribution, employing linear theory. The higher order harmonic terms and the main kinematic features of the observed data are consistent with an analytically constructed simple bar model. We also present maps of a number of strong absorption lines in M 100, derive simple stellar populations and correlate them with features in the gas kinematics. ",
        "introduction": "\\label{Intro}  Although circular rotation is the dominant kinematic feature of galactic disks, it is clear from observations and theory that in many galaxies non-circular motions are also present (e.g. Freeman et al. 1965, Shlosman et al. 1989). Non-axisymmetric motion can, in fact, be a crucial ingredient in determining the evolution of a galaxy, by triggering star formation, or by building a central mass concentration.  We know that many galaxies show non-axisymmetric mass-distributions: near-infrared photometric studies show that at least 70\\% of the galaxies in the nearby Universe show strong or weak bars (Eskridge et al. 2000, Knapen et al. 2000a). Such asymmetric features in the mass-distribution will produce non-circular motions, although also other effects, such as interactions with other galaxies, can cause them. Gas, being more responsive to perturbations and instabilities than stars, will show these non-circular motions more easily. Athanassoula (1992) shows how gas streaming manifests itself in galaxies with strong bars. Her calculations show that one can use the motion of the gas, as measured in two dimensions, to derive information about the mass distribution in galaxies, especially if these are relatively isolated. This has not been done very often in the literature up to now, since good, high resolution data is scarce, and up to recently a good theoretical framework was lacking.  Recently, Wong, Blitz \\& Bosma (2004) have quantified  the strength of bars and spiral arms using HI and CO velocity fields of a number of nearby galaxies. They used a harmonic decomposition formalism developed by Schoenmakers et al. (1997), and compared the results with analytical  models. Here we do the same with much higher resolution optical emission line data, obtained with the Integral Field Spectrograph SAURON, for NGC~5448, an early-type spiral galaxy. More details can be found in Fathi et al. (2005; {\\bf F05}).  From the SAURON data one can not just measure emission lines, but also absorption lines. In the second part of this paper we present absorption line maps of the central regions of M100, together with some simple stellar population models, recently published by Allard et al. (2006; {\\bf A06}). Absorption lines, contrary to integrated colours, are relatively insensitive to dust absorption, so that more accurate ages and other stellar population parameters can be obtained.  In the future we plan to perform a harmonic decomposition (also called kinemetry, see Krajnovic et al. 2006) on the full sample SAURON early-type spirals, presented in Falc\\'on-Barroso et al. (2006), and analyse statistically the non-axisymmetric motions in the centers of nearby spirals. ",
        "conclusions": "\\label{Discussion}  In the previous sections we have shown two applications of integral field spectroscopy aimed at understanding the nature of the central regions of spiral galaxies. We show that, in contrast to long slit spectroscopy, we can detect streaming motions of the ionized gas, and use them to measure the potential. Although something similar can also be done from the photometry, there are more assumptions involved here, since one has to look through the dust, and make assumptions for {\\it M/L} ratios as a fuction of position. In a galaxy like M100, where the stellar populations are strongly varying, estimating M/L purely from colours and absorption line strengths will not be easy at all. Up to now we have only analyzed the Sa galaxy NGC 5448 in detail. Once we have analyzed the full sample of Falc\\'on-Barroso et al. (2006) we will be able to derive statistics on the number of bars, and their strengths, based for the first time on kinematic data. These statistics will be important, among other topics, in assessing whether AGN activity is related to the presence of bars (e.g., Shlosman et al. 1989).   In the second application we show line strength maps of a galaxy with a circumnuclear ring. Since most Sb galaxies are very dusty in their inner regions one cannot learn too much from broad band colours, unless the galaxy has a special orientation (e.g., Peletier et al. 1999). From three line strength maps we find that the galaxy has zones with different ages. In the ring itself we show that apart from the young stars there is also an underlying old stellar population.  M 100 is a beautiful case where many phenomena can be seen that play a role in the inner regions of spiral galaxies. The outer and inner parts of its bar have been detected both in near-IR (ground-based and Spitzer) photometry and in the gas kinematics, and a ring separating the two. Abundant star formation occurs within and near the ring, with 2 hot spots at the end of the inner bar. Young stars are formed where the velocity dispersion of the gas is low (Allard et al. 2005). At present we are studying the stellar populations of the sample of Sa galaxies of Falc\\'on-Barroso et al. (2006) in the same way. Many galaxies of that sample show young stellar populations in circumnuclear rings or disks, showing that for most galaxies the central regions are not dominated by a bright, old stellar bulge. More details can be found in Peletier et al. (in preparation)."
    },
    "0609/astro-ph0609237_arXiv.txt": {
        "abstract": "We use Subaru Suprime-Cam and VLT FORS1 photometry of the dwarf galaxy DDO210 to study the global stellar content and structural properties of a transition-type galaxy (with properties intermediate between dwarf irregular and dwarf spheroidal systems). This galaxy is sufficiently isolated that tidal interactions are not likely to have affected its evolution in any way. The colour-magnitude diagrams of DDO210 show a red giant branch (RGB) population (with an RGB bump), a bright asymptotic giant branch population, a red clump, young main sequence stars and blue-loop stars. The youngest stars formed within the last 60\\,Myrs and have a distinct radial distribution compared to the main population. Whereas the overall stellar spatial distribution and HI spatial distribution are concentric, the young stars are offset from the center of DDO210 and are coincident with a `dent' in the HI distribution. The implied recent star formation rate required to form the young population is significantly higher than the derived current star formation rate, by a factor of $> 10$.  Most of the stars in DDO210 are found in a red clump, and its mean $I$-band magnitude suggests that the majority of stars in DDO210 have an average age of $4^{+2}_{-1}$\\,Gyr. Given this age, the colour of the RGB implies a mean metallicity of [Fe/H] $\\simeq -1.3$. By comparing the shape of the red clump with models for a variety of star formation histories, we estimate that an old ($\\ge 10$\\,Gyr) stellar population can contribute $\\sim 20 - 30$\\% of the stars in DDO210 at most. The unusual star formation history of DDO210, its low mass estimate and its isolated nature, provide insight into how star formation proceeds in the lowest mass, unperturbed, dwarf galaxy haloes. ",
        "introduction": "The stellar, chemical and structural evolution of dwarf galaxies is believed to depend strongly on environment. In the Local Group, dwarf spheroidal (dSph) galaxies are preferentially found as satellites to M31 and the Milky Way (MW), whereas dwarf irregular (dIrr) galaxies are preferentially found in isolated locations. This was first highlighted by \\cite{einasto1974} (see also \\citealt{vandenbergh1999}), and implies that the presence of a nearby large galaxy causes morphological changes in its satellites, primarily via the loss of gas by a combination of tidal stirring and ram-pressure stripping (\\citealt{mayer2001a,mayer2001b,grebel2003}). Tidal effects from the large galaxies may also be responsible for star formation episodes in their satellites, by triggering compression of their gas. However, the orbital properties of the MW satellites are difficult to determine, and so correlations between star formation histories (SFHs) and orbital phase have yet to be robustly demonstrated.  The stellar evolution of dwarf galaxies is also interesting from a cosmological viewpoint, in particular regarding the missing satellites problem. A long standing problem faced by theories of galaxy formation in the cold dark matter paradigm is that purely gravitational simulations predict that there should be around an order of magnitude more dwarf galaxies in the Local Group, particularly around the MW and M31 (\\citealt{kauffmann1993,klypin1999,moore1999}). As recent results are demonstrating, it appears as if there are numerous low surface brightness dwarf galaxies which have yet to be discovered (\\citealt{zucker2004a,zucker2006a,zucker2006b,willman2005}; see also \\citealt{willman2006,belokurov2006}) although it is unclear if these discoveries will increase the observed number by the required order of magnitude. One potential solution to this difficulty is that not all dwarf galaxy-scale dark matter haloes can accrete and retain gas to form stars, and instead reionisation can prevent certain haloes from forming a luminous component (\\citealt{bullock2000,kravtsov2004,ricotti2005,gnedin2006}). \\cite{susa2004} show that star formation is suppressed significantly by the effects of reionisation, but that sufficiently massive dwarf galaxies are able to form stars eventually because the baryons are self-shielded.  Even in these cosmological scenarios, the influence of tidal effects from the MW and M31 can play important roles in determining the evolution of a dwarf galaxy. Indeed, tidal effects present one of the greatest uncertainties in our interpretation of dwarf galaxy evolution, and it is therefore of considerable interest to study dwarf galaxies unaffected by the presence of an external tidal field. DDO210 ($20^{\\rm h}\\,46^{\\rm m}\\,51.8^{\\rm s}$, $-12^\\circ\\,50'\\,53''$) is one of the most isolated galaxies in the Local Group. It was discovered at the limits of detectability on the Palomar Sky Survey by \\cite{vandenbergh1959}, in the constellation of Aquarius. \\cite{fisher1975} detected it in HI 21cm emission, and later assigned it an arbitrary distance modulus of $\\left(m - M\\right) = 25$\\,mags, based upon some general similarities to the Local Group galaxy WLM. For a long time DDO210 remained very poorly studied, and \\cite{marconi1990} were the first to resolve its individual stars, to a limiting magnitude of $V = 23.5$\\,mags.  Despite the lack of any distance measurement, DDO210 was considered to be a member of the Local Group from soon after its discovery (eg. \\citealt{yahil1977,vandenbergh1979,devaucouleurs1983}). \\cite{greggio1993} were the first to estimate an independent distance to this galaxy, $\\left(m - M\\right) \\simeq 28$\\,mags ($d \\simeq 4.3$\\,Mpc), placing it well beyond the zero velocity surface of the Local Group. However, more recent distance measurements by \\cite{lee1999} (L99), \\cite{karachentsev2002} (K02) and \\cite{mcconnachie2005a} (M05), all of which are based upon the tip of the red giant branch (TRGB; \\citealt{lee1993a}), show that DDO210 is, in fact, located at $\\sim 1$\\,Mpc from the Milky Way (MW), on the periphery of the Local Group. This places DDO210 at slightly over 1\\,Mpc from M31. The free-fall timescale of DDO210 into M31 or the MW is  \\begin{equation} t_{ff} \\simeq 14.4$\\,Gyr\\,$\\left(\\frac{r}{1\\,{\\rm Mpc}}\\right)^{\\frac{3}{2}}\\,\\left(\\frac{M}{10^{12}\\,{\\rm M_\\odot}}\\right)^{-\\frac{1}{2}}~, \\label{tff} \\end{equation}  \\noindent ie. of order a Hubble time. It is clearly not, and never has been, a satellite of either of the dominant members of the Local Group.  \\cite{mateo1998a} classifies DDO210 as a transition-type galaxy, with properties intermediate to those which are usually ascribed to dSph and dIrr galaxies. Its magnitude is $M_V \\simeq -10.6$\\,mags (L99), making it the faintest dwarf galaxy in the Local Group with a significant gaseous component. Despite the presence of this gas, an $H\\alpha$ study of this galaxy (\\citealt{vanzee1997}) shows that there is no current star formation. This is rather unusual, given that L99 show that this galaxy contains a significant number of young stars, and that the central portion of this galaxy has undergone a recent enhancement in its star formation rate (SFR) in the last few hundred Myrs.  \\cite{begum2004} conducted deep, high velocity resolution HI 21cm observations of DDO210 with the Giant Meterwave Radio Telescope. They find that the velocity field of this galaxy is quite regular, and has a corrected peak rotation velocity of $\\sim 16$\\,km\\,s$^{-1}$. When incorporated into mass models of DDO210, this implies a total dynamical mass of order $10^{7 - 8}$\\,M$_\\odot$, comparable to many of the dwarf satellites of M31 and the MW.  Since DDO210 is so isolated, it offers an excellent opportunity to study the stellar evolution of a dwarf galaxy in the absence of any significant external tidal fields.  In this paper, we analyse deep, global, multi-colour photometry of DDO210, obtained with Subaru Suprime-Cam and VLT FORS1. We analyse the stellar content, structural properties and SFH of this galaxy, and show that it is dominated by an intermediate age stellar population with an average age of $\\sim 4$\\,Gyr. In Section~2, we discuss our Subaru Suprime-Cam and VLT FORS1 photometric datasets and the data-reduction procedure. In Section 3, we present deep, global, CMDs of the stellar populations of DDO210, and use these to derive the properties of its multiple stellar populations. We compare the CMDs with synthetic models of various SFHs and use this to estimate the possible contribution from an ancient ($\\gtrsim 10$\\,Gyr) stellar component. In Section 4, we analyse the global structure of DDO210, in terms of its integrated properties and the radial profiles of the various stellar populations, and we compare these to the global distribution of HI from previous Very Large Array (VLA) observations by \\cite{young2003}. We discuss the nature of DDO210 in light of these new results in Section 5, and we summarise our findings in Section 6. We assume an extinction value towards DDO210 of $E\\left(B - V\\right) = 0.052$ (\\citealt{schlegel1998}). ",
        "conclusions": "DDO210 is dominated by an intermediate age stellar population. In this respect, it resembles Leo~A, another relatively isolated Local Group galaxy which has been shown to consist of a predominantly young ($< 2$\\,Gyr) stellar population (\\citealt{tolstoy1998}). Leo~A is located $\\sim 690\\,$kpc from the MW (\\citealt{mateo1998a}), and it is conceivable that it is on a highly radial orbit around the MW. Interactions may have affected its SFH. In particular, the presence of RR~Lyrae stars in Leo~A shows that it has an old component, and may well have had ongoing, low-level, star formation until a few Gyr ago, when the star formation rate increased. This may have been brought on by a close encounter with the MW which triggered enhanced star formation.  DDO210, on the other hand, is even more isolated than Leo A, and external tidal effects from the MW and M31 cannot have played a significant role in its evolution (Equation~\\ref{tff}). Interactions are not likely to have affected its observable properties, and so its unusual SFH must be explained without recourse to interactions. The total mass of DDO210 is important to consider as well, since its circular velocity is $\\sim 16$\\,km\\,s$^{-1}$, implying a mass of $\\sim 10^{7 - 8}$\\,M$_\\odot$ (\\citealt{begum2004}). DDO210 is an intrinsically low mass galaxy which has not interacted with large masses, and yet it did not form most of its stars until $\\sim 4$\\,Gyr ago.  What seems clear is that DDO210 must have accreted its gas prior to reionisation, since it is far too low mass to successfully accrete gas post-reionisation. Since the virial temperature of DDO210's halo will not have been vastly higher than the temperature of the accreted gas, the gas will probably have remained in a relatively diffuse configuration in the dark matter halo. What star formation occurred at these early times (if any) will have been small and will have been confined to localised gaseous overdensities, rather than acting on a global scale. Deeper observations are required to determine the exact fraction of old stars in this galaxy. Since the gas never really collapsed all the way to the center of this system, then a bona-fide disk will not have had a chance to form. This would explain why we, along with others, conclude that it is unlikely that the dominant structure in DDO210 is a disk.  Why are most stars intermediate age? It could be that DDO210 continued to accrete mass (but not gas) post-reionisation, such that eventually the gas collapsed sufficiently that star formation proceeded at a reasonable rate and formed the dominant intermediate age component. Alternatively, it may be that the gas eventually collapsed sufficiently without the overall mass of DDO210 increasing, either through random perturbations or via normal cooling mechanisms.  The above is, of course, speculation, but a scenario such as this seems plausible given our results and the current cosmological paradigm. As well as the dominant intermediate age population, however, the recent SFH of DDO210 is also unusual. The current SFR in DDO210 is basically zero ($< 0.000003\\,M_\\odot\\,{\\rm yr}^{-1}$; \\citealt{vanzee1997,young2003}), even although this galaxy was forming stars as recently as a few tens of Myrs ago. We quantify this apparent inconsistency by assuming DDO210 has formed all of its young stars within the past 500\\,Myrs. We compare the number of young stars observed in the colour -- magnitude range $\\left(V - I\\right) < 0.5$, $I < 24.3$\\,mags with the corresponding number in an IAC-STAR simulation, convolved with our error distribution, which has a constant SFR over the last 500\\,Myrs. In order to match the number of young stars observed, the average SFR over the last 500\\,Myrs must be $\\sim 0.0002\\,M_\\odot\\,{\\rm yr}^{-1}$ ie. nearly two orders of magnitude larger than the derived upper limit on the current SFR. It could be that the recent SFH of DDO210 is best characterised by a very low, relatively constant SFR, interrupted by periods of inactivity. This `gasping' star formation scenario has already been implied for several other Local Group galaxies (eg. \\citealt{tosi1991}).  The young stars in DDO210 are not centered in the middle of the galaxy, but are offset a few arcmins to the east. This location is coincident with a slight dent in the HI, and perhaps suggests that some recent event, such as an interaction, could have triggered off-center star formation in this galaxy (although what could have interacted with DDO210 is a mystery!). Alternatively, the off-center star formation may be due to recently captured gas, which has not yet become fully virialised.  The fact that the young stars are still clumped on a scale of $\\sim 1$\\,arcmin ($\\sim 0.3$\\,kpc) suggests that they were probably born with a relatively small velocity dispersion ($\\sigma \\sim 3 - 6$\\,km\\,s$^{-1}$). We note that \\cite{dohmpalmer2002} showed that the recent star formation in Sextans~A was not confined to the center of the galaxy, and seemed to migrate with time. They argued that star formation in one part of a cloud could cause star formation in another part of the cloud, by triggering instabilities in the gas. In this case, no interactions would be required to explain the spatial location of the young stars, which would reconcile our finding with the splendid isolation of DDO210.  Finally, we note that the presence of a spatially extended intermediate age population and a centrally concentrated young population is compelling evidence for the presence of multiple structural components in DDO210. The systematic change in ellipticity of DDO210 as a function of radius supports this claim. These findings provide yet more evidence that dwarf galaxies are not simple single component stellar systems (eg. Sculptor, \\citealt{tolstoy2004}; Leo~A, \\citealt{vansevicius2004}; NGC6822, \\citealt{demers2005}; Andromeda~II, \\citealt{mcconnachie2006b})."
    },
    "0609/astro-ph0609001_arXiv.txt": {
        "abstract": "We present a spectroscopic study and dynamical analysis of $\\sim$2600 M7 dwarfs.  We confirm our previous finding that the fraction of magnetically active stars decreases with vertical distance from the Galactic plane.  We also show that the mean luminosity of the H$\\alpha$ emission has a small but statistically significant decrease with distance.  Using space motions for $\\sim$1300 stars and a simple one-dimensional dynamical simulation, we demonstrate that the drop in the activity fraction of M7 dwarfs can be explained by thin disk dynamical heating and a rapid decrease of magnetic activity at a mean stellar age of $\\sim$6-7 Gyr. ",
        "introduction": "The most ubiquitous stars in the Galaxy are low-mass dwarfs.  Because they are plentiful, they provide an excellent probe of the structure and dynamics of the local thin disk stellar population.  In addition, many low-mass dwarfs are magnetically active (as identified by H$\\alpha$ emission), especially the mid to late M-dwarfs.  Various studies have suggested that there is a relation between the magnetic activity of an M-dwarf and its age.  These studies have used the Galactic motions of M-dwarfs (Wilson \\& Woolley 1970; Wielen 1974; Hawley et al. 1996), M-dwarf activity in stellar clusters (Hawley et al. 1999; Hawley et al. 2000; Gizis, Reid \\& Hawley 2002) and active M-dwarfs in binary systems (Silvestri, Hawley \\& Oswalt 2005; Silvestri et al. 2006). By utilizing the activity of M-dwarfs, and the largest spectroscopic sample of M-dwarfs ever assembled, we can augment dynamical and positional data with the additional parameter of time.  Previous studies have thoroughly investigated the vertical dynamics and structure of the Galaxy. Wielen (1977) found that the velocity dispersion of stars increases with time to the 0.5 power (t$^{0.5}$). Spitzer \\& Schwarzchild (1951, 1953) suggested that massive gas clouds would cause an increase in the velocity dispersion of stars through cloud-star interactions.  Subsequent observations of molecular clouds confirm that enough interstellar gas exists to dynamically heat disk stars.  However, current evidence indicates that molecular clouds alone are not responsible for the total amount of heating.  Black holes and spiral density waves have been suggested as other possible contributors to the heating of the Galactic disk (Jenkins 1992; H\\\"anninen \\& Flynn 2002).  Recently, many groups have investigated the structure and age of the Galactic Disk by examining the observed and simulated kinematics of nearby stars (Kuijken \\& Gilmore 1989; Dehnen \\& Binney 1998; Rocha-Pinto \\& Maciel 1998; Binney, Dehnen \\& Bertelli 2000; Fuchs et al. 2001; H\\\"anninen \\& Flynn 2002; Holmberg et al. 2003; Siebert, Bienayme \\& Soubiran 2003).  The values reported for the exponent ($\\alpha$) of the time dependence (t$^{\\alpha}$) of the velocity dispersion have not converged and range from 0.26-0.59 (in most cases the uncertainties in $\\alpha$ are smaller than this spread).  Despite the discrepancies, these studies have aided in constraining the age of the solar neighborhood, the surface mass density of the Galactic plane and the time evolution of dynamical heating.  Most of this work has not utilized low-mass stars due to the small numbers of M-dwarfs included in magnitude-limited samples.  Aside from the initial work by Wielen (1974, 1977) and the more recent PMSU and 100 Parsec surveys (Reid, Gizis \\& Hawley 2002; Bochanski et al. 2005), few Galactic studies have examined the low-mass stellar constituents.  The advent of large survey science has vastly increased the number of low-mass dwarfs available for analysis.  Using several hundred M7 dwarfs from the Sloan Digital Sky Survey (SDSS), West et al. (2004; hereafter W04) suggested that the fraction of activity in M7 stars decreases as a function of height above (or below) the Plane.  They suggested that this phenomenon would naturally result from an age-activity relation; the older stars are on average farther away and are no longer active, thus decreasing the active fraction.  However, despite using hundreds of stars in their analysis, the W04 results had large uncertainties.  In this paper, we focus on the activity and dynamics of a much larger sample of M7 dwarfs.  We describe our selection criteria and data in \\S2.  In \\S3 we confirm the findings of W04, namely that the fraction of activity decreases as a function of vertical distance from the Plane. We also investigate how the magnitude of activity (as quantified using L$_{\\rm{H\\alpha}}$/L$_{bol}$) changes with distance. We utilize proper motions from the USNO-B catalog (Monet et al. 2003) and a new set of low-mass star templates for measuring radial velocities (Bochanski et al. 2006) to derive space motions. In \\S4 we use a simple dynamical simulation to explain the decline in activity fraction as due to dynamical heating effects, and derive a timescale for the existence of magnetic activity in M7 dwarfs.  Our results are discussed and plans for future projects proposed in \\S5. ",
        "conclusions": "Using over 2600 M7 dwarf spectra, we confirm the result of W04, that the activity fraction does decrease as a function of vertical distance from the Galactic Plane.  A possible explanation for this observation is that activity ceases in older stars and that due to dynamical heating, these older stars are on average further away.  Our model simulation using simple dynamical arguments and the assumption that activity ceases at a given stellar age adequately reproduces the observations and validates the feasibility of our hypothesis.  A more detailed dynamical simulation should be carried out to aid in the physical interpretation of the observations. The importance of our simulation is simply to demonstrate that our hypothesis is reasonable.  The extrapolation of the simple log-linear fit of the age at which activity ceases vs. spectral type (color, mass) derived from a few open clusters in Hawley et al (2000) would predict a very large age for the M7 dwarfs to cease their activity.  Our sample has allowed us to actually probe these late-type dwarfs, which are not accessible in old clusters, and provides the first measured estimate of the time dependence of their magnetic activity.  The investigation of the entire M dwarf sequence by our method will allow us to further refine and extend the relation from the open clusters (Hawley et al 2006 in prep.).  The L$_{\\rm{H\\alpha}}$/L$_{bol}$ data show a large scatter that may be due to: 1) an intrinsic variation in the activity of a sample of stars; or 2) significant variation in the H$\\alpha$ line emission during the lifetime of an individual M7 dwarf.  A large sample of low-mass stars with large time baselines is needed to distinguish between these possibilities.  If individual stars cannot produce the variation we see in large single epoch data, then the scatter is likely the intrinsic scatter of the population.  The L$_{\\rm{H\\alpha}}$/L$_{bol}$ data suggest that there is a small but robustly detected decrease in the amplitude of magnetic activity as a star ages but the large spread at all distances and small decrease of the mean furthers the notion that these stars have some range of magnetic activity during their ``active'' lifetimes. The mechanism behind this age-activity relation is currently unknown. Future work will statistically explore the spread in L$_{\\rm{H\\alpha}}$/L$_{bol}$ to determine if the data support the rapid decrease of activity hypothesis.   Our results are important for placing constraints on models of magnetic dynamos in low-mass dwarfs.  Many recent studies have investigated the mechanisms controlling magnetic field generation in low-mass stars (e.g. Dobler, Sitx \\& Brandenburg 2006; Chabrier \\& K\\\"uker 2006; Bercik et al. 2006; Donati et al. 2006).  Future dynamo models should include an age dependence in order to match our observations.  Work is underway to expand our M7 analysis to include \\emph{all} M dwarf spectral types.  Hawley et al. (in prep) are using the largest spectroscopic sample of low-mass stars ever assembled (from SDSS Data Release 5) to examine how the activity fractions, space motions and L$_{\\rm{H\\alpha}}$/L$_{bol}$ vary as a function of Galactic height through the M spectral sequence.  This will allow us to examine how the dependence of age on activity varies with mass, putting even stronger constraints on dynamo models.  One element that we do not address in this paper is how metallicity varies as a function of Galactic height.  The sample size of known M7 subdwarfs is quite small and does not allow for a robust statistical study.  However, the Hawley et al. sample will also allow for an exploration of how M-dwarf abundances vary with vertical height from the Plane.  A metallicity gradient, together with strong age constraints, will provide important information on the chemical evolution of the Galactic disk.  Almost 30 years have passed since Wielen (1977) used low-mass dwarfs to probe the kinematics and structure of the Milky Way.  The advent of large surveys exemplified by SDSS and 2MASS, coupled with the large abundance of low-mass dwarfs, allows us to continue probing the structure of the Galaxy with its smallest and most numerous constituents."
    },
    "0609/astro-ph0609371_arXiv.txt": {
        "abstract": "We present results from a study of the globular cluster luminosity function (GCLF) in a sample of 89 early-type galaxies observed as part of the ACS Virgo Cluster Survey. Using a Gaussian parametrization of the GCLF, we find a highly significant correlation between the GCLF dispersion, $\\sigma$, and the galaxy luminosity, $M_{B,{\\rm gal}}$, in the sense that the GC systems in fainter galaxies have narrower luminosity functions. The GCLF dispersions in the Milky Way and M31 are fully consistent with this trend, implying that the correlation between sigma and galaxy luminosity is more fundamental than older suggestions that GCLF shape is a function of galaxy Hubble type. We show that the $\\sigma - M_{B,{\\rm gal}}$ relation results from a bonafide narrowing of the distribution of (logarithmic) cluster masses in fainter galaxies. We further show that this behavior is mirrored by a steepening of the GC mass function for relatively high masses, ${\\cal M} \\ga 3\\times10^5\\,{\\cal M}_\\odot$, a mass regime in which the shape of the GCLF is not strongly affected by dynamical evolution over a Hubble time. We argue that this trend arises from variations in initial conditions and requires explanation by theories of cluster formation. Finally, we confirm that in bright galaxies, the GCLF ``turns over\" at the canonical mass scale of ${\\cal M}_{\\rm TO}\\simeq 2\\times 10^5\\,{\\cal M}_\\odot$. However, we find that ${\\cal M}_{\\rm TO}$ scatters to lower values ($\\approx$ 1-2$\\times 10^5\\,{\\cal M}_\\odot$) in galaxies fainter than $M_{B,{\\rm gal}}\\ga -18.5$, an important consideration if the GCLF is to be used as a distance indicator for dwarf ellipticals. ",
        "introduction": "\\label{sec:intro}  The luminosity function of globular clusters (GCs) represents one of the most remarkable features of these stellar systems. The distribution of GC magnitudes, commonly referred to as the GC luminosity function (GCLF), shows a turnover, or peak, at $M_V \\simeq -7.5$ mag, corresponding to a mass of ${\\cal M} \\simeq 2 \\times 10^5 {\\cal M}_\\odot$. Observations have shown that this turnover is nearly invariant across and within galaxies, prompting its widespread use as a distance indicator (see, e.g., Harris 2001). Accounting for this nearly universal mass scale remains an open problem for theories of GC formation and evolution. It follows that establishing whether or not the GCLF as a whole is universal --- i.e., whether its overall form depends on host galaxy properties --- can help guide and constrain theories for the formation and evolution of galaxies and GC systems.  In this {\\it Letter}, we present results from a study of the GCLFs of 89 early-type galaxies observed by HST as part of the ACS Virgo Cluster Survey (ACSVCS; C\\^ot\\'e et al.~2004). We find the clearest evidence to date for a correlation between the width (i.e., Gaussian dispersion) of the GCLF and the luminosity of the host galaxy; we also show that there is some downward scatter in the {\\it mass} scale of the GCLF turnover in galaxies fainter than $M_{B,{\\rm gal}}\\ga -18.5$. Focusing on the observed steepening of the GCLF at the bright (high-mass) end in the faint galaxies, we argue that this behavior was probably imprinted at the time of GC formation. A more detailed discussion of the whole GCLF, including the faint (low-mass) end and the role that long-term dynamical evolution plays in that regime, is deferred to a subsequent paper (Jord\\'an et~al.~2006, hereafter J06). That paper presents our data in full and gives details of our analysis techniques, including modeling of the GCLFs with a new, non-Gaussian, physically motivated fitting function. ",
        "conclusions": "\\label{sec:discussion}  An obvious question prompted by Figure \\ref{fig:sigma_B} is whether the correlation between $\\sigma_z$ and $M_{B,{\\rm gal}}$ was established at the time of cluster formation or built up afterwards as GCLFs were modified by the dynamical destruction of GCs over a Hubble time. We favor the first interpretation.  Star clusters can be destroyed over Gyr timescales as a result of mass loss driven by stellar evolution, dynamical friction, gravitational shocks, and internal two-body relaxation (evaporation) --- processes that have been studied in detail by several groups. Recent discussions, centered specifically on how these affect the GCLF, can be found in Fall \\& Zhang (2001) and Vesperini (2000, 2001). Fall \\& Zhang in particular show that, while stellar evolution and gravitational shocks certainly deplete the total number of GCs in a galaxy, they do not significantly alter the overall shape of the GCLF. Evaporation, on the other hand, {\\it can} change the shape of the GCLF, but significantly so only for cluster masses ${\\cal M}\\la 2$-$3\\times10^5\\,{\\cal M}_\\odot$, i.e., below the typical GCLF turnover mass.  In the theoretical treatments of Fall \\& Zhang, Vesperini, and many others, the evaporation rate is independent of cluster mass, which ultimately drives the {\\it low-mass} side of the GCLF to a universal shape: a simple exponential $dN/dz \\propto 10^{-0.4\\,z}$ for the number of GCs per unit magnitude fainter than the turnover (equivalent to a flat distribution for the number of GCs per unit linear luminosity or mass). But fitting a Gaussian model to the GCLF, as we have done here, tacitly assumes that the distribution is symmetric. The results in Figure \\ref{fig:sigma_B} might therefore seem to imply that both the bright side {\\it and} the faint side of the GCLF become progressively steeper in fainter galaxies. However, various observational uncertainties make it difficult to determine precisely the form of the faintest tail of the GCLF. Thus, in J06 we show that good fits to our GCLFs can also be obtained using an alternate model with a universal exponential shape at magnitudes fainter than the turnover --- and that the downward scatter in ${\\cal M}_{\\rm TO}$ for faint galaxies persists in such a model and so is not an artifact of any assumed Gaussian symmetry. Here we concern ourselves only with the brighter half of the GCLF, which is observationally better defined.  We have performed maximum-likelihood fits of exponential models $dN/dz \\propto 10^{0.4 (\\beta_z-1)\\,z}$ (corresponding to power-law {\\it mass} distributions, $dN/d{\\cal M}\\propto {\\cal M}^{-\\beta_z}$) to the GCLFs at absolute magnitudes $-8.7 \\ga z \\ga -10.8$ (cluster masses $\\simeq\\! 3\\times10^5$--$2\\times10^6\\ {\\cal M}_\\odot$) in 66 of our galaxies. Such distributions accurately describe the bright sides of giant-galaxy GCLFs (Harris \\& Pudritz 1994; Larsen et al. 2001), and with $\\beta_z\\simeq2$ they also give good matches to the mass functions of young star clusters in nearby mergers and starbursts (Zhang \\& Fall 1999).  Figure~\\ref{fig:power} shows the results of this exercise. There is a clear steepening in the power-law exponent, from $\\beta_z\\simeq 1.8$ in bright galaxies to $\\beta_z\\simeq3$ in the faintest systems. However the faint side of the GCLF behaves in detail, the bright side alone suggests that smaller galaxies were unable to form very massive clusters in the same {\\it relative} proportions as giant galaxies.  \\begin{figure} \\epsscale{1.0} \\plotone{f3.eps} \\caption{Slope of the power law that best fits our $z$-band GCLF data, $\\beta_z$, for masses $3\\times 10^5 \\lae ({\\cal M}/{\\cal M}_\\odot) \\lae 2\\times 10^6$, plotted against host galaxy absolute magnitude, $M_{B,{\\rm gal}}$. The star and triangle show $\\beta$ values for the Milky Way and M31 respectively, measured in the same mass regime using the data from Harris (1996) and Reed et~al. (1994) assuming a $V$-band mass-to-light ratio $M/L_V = 2$. The bright side of the GCLF is steeper in fainter galaxies. \\label{fig:power}} \\end{figure}    A potential complication here is dynamical friction. A cluster of mass ${\\cal M}$ on an orbit of radius $r$ in a galaxy with circular speed $V_c$ will spiral in to the center of the galaxy on a timescale $\\tau_{\\rm df} \\propto {\\cal M}^{-1} r^2 V_c$ (Binney \\& Tremaine 1987). In the Milky Way and larger galaxies, $\\tau_{\\rm df} > 13$~Gyr for all but the very most massive clusters at small radii, and thus dynamical friction does not significantly affect their GCLFs (e.g., Fall \\& Zhang 2001). In dwarfs with low $V_c$, however, $\\tau_{\\rm df}$ can be interestingly short for smaller GCs at larger $r$ --- suggesting, perhaps, that the process might significantly deplete the bright side of the GCLF in small galaxies and contribute to the type of trend seen in Figure~\\ref{fig:power}. However, Vesperini (2000, 2001) has modeled the GCLF evolution over a Hubble time in galaxies with a wide range of mass, and his results strongly suggest that dynamical friction does {\\it not} suffice to explain our observations. In particular, the widths of the Gaussian GCLFs in his models do not decrease, even in dwarf galaxies, to anywhere near the extent seen in the data. Thus, any significant galaxy-to-galaxy variations in the shape of the GCLF above the turnover mass probably reflect initial conditions (see J06 for further discussion).  In summary, the gradual narrowing of the GCLF as a function of galaxy luminosity --- or the steepening of the mass distribution above the classic turnover point --- presents a new constraint for theories of GC formation and evolution. In our view, it is the cluster formation process in particular that is likely to be most relevant to the observed behavior at the high-mass end of the GCLF. Exactly what factors might lead to more massive galaxies forming massive clusters in greater relative numbers, is an open question of some interest."
    },
    "0609/astro-ph0609692_arXiv.txt": {
        "abstract": "We examine whether massive-star accretion disks are likely to fragment due to self-gravity.  Rapid accretion and high angular momentum push these disks toward fragmentation, whereas viscous heating and the high protostellar luminosity stabilize them.  We find that for a broad range of protostar masses and for reasonable accretion times, massive disks larger than $\\sim 150$ AU are prone to fragmentation.  We develop an analytical estimate for the angular momentum of accreted material, extending the analysis of Matzner and Levin to account for strongly turbulent initial conditions.  In a core-collapse model, we predict that disks are marginally prone to fragmentation around stars of about four to 15 $\\Msun$ -- even if we adopt conservative estimates of the disks' radii and tendency to fragment.  More massive stars are progressively more likely to fragment, and there is a sharp drop in the stability of disk accretion at the very high accretion rates expected above 110 solar masses.  Fragmentation may starve accretion in massive stars, especially above this limit, and is likely to create swarms of small, coplanar companions. ",
        "introduction": "Advances in submillimeter telescopy have enabled the discovery of flattened structures, in some cases clearly Keplerian disks, surrounding massive ($\\gtsim 10\\Msun$) protostars \\citep{Ces2005,Ch2004,Pat2005,Beuth}.  The logical inference -- that high mass star formation (HMSF) proceeds through disk accretion -- raises a question: can such disks process rapid mass accretion, or do they fragment to produce secondary stars?  We seek to answer this question by estimating the criterion for disk fragmentation in the vicinity of a massive star.  If massive-star disks typically do not fragment, then disk accretion poses little barrier to massive star formation.  Conversely if disk fragmentation occurs, then accretion onto the central star may be (partially) choked off, as suggested by \\cite{2005MNRAS.362..983T} in the context of low-luminosity active galactic nuclei.   Moreover each massive star that forms from a fragmenting disk may be surrounded by smaller stars that began as disk fragments.  It is therefore important to evaluate models for HMSF in light of the disk fragmentation criterion.  The current work uses and extends the results of \\citeauthor{ML2005} (\\citeyear{ML2005}, hereafter ML05), who showed that protostellar disks around low mass stars are strongly stabilized against fragmentation by a combination of viscous heating and irradiation by the central protostar.  Massive star formation is fundamentally different, however, in three important ways.  The rapid rise of luminosity with mass implies that stellar irradiation is far more intense; this tends to stabilize disks against fragmentation.  However, massive stars must accrete quickly to form at all \\citep[e.g.,][]{WC1987}, and rapid accretion favors fragmentation.   These effects compete to set the critical radius outside of which disks fragment.  Whether fragmentation actually occurs depends on the initial disk radius, which itself depends on the physical state of the gas prior to accretion. We discuss disk fragmentation in \\S \\ref{Dfrag}, considering the stabilizing effect of viscous heating (\\S \\ref{visc}) before incorporating irradiation by the central star (\\S \\ref{irrad}).  This combination allows us to identify (\\S \\ref{nomodel}) the disk radius at which fragmentation sets in. The \\cite{MT2003} core collapse model is examined in more detail in \\S \\ref{coremod}: we compute angular momentum scales in \\S \\ref{angmom} using formulae derived in the Appendix. In \\S \\ref{coreval} we calculate expected fragmentation radii for a range of masses in the core collapse model.  Turning to the consequences of fragmentation, we examine in \\S \\ref{fragments} the likely properties of stars born within fragments and the possibility that fragmentation limits accretion.  In \\S \\ref{observations} we compare our results with observed regions of HMSF. ",
        "conclusions": "\\label{discussion} Our primary conclusion (\\S \\ref{nomodel}) is that massive-protostar disks that accrete more slowly than $\\sim 1.7\\times 10^{-3} \\Msun$~yr$^{-1}$ are subject to fragmentation at disk radii beyond about 150 AU.  This critical radius is set primarily by the viscous heating of the disk midplane as it accretes, with reprocessed starlight playing an equal or secondary role for most stellar masses. As all the disks we consider are optically thick, the critical radius depends on the Rosseland opacity law $\\kappa_R(T)$ within dusty disk gas.  Comparing to our conservative estimate of the disk radius in the \\cite{MT2003} model for massive star formation by the collapse of turbulent cores (\\S \\ref{coremod}), we find that fragmentation is marginal for stars accreting four to 15 solar masses; higher-mass stars are increasingly afflicted by disk fragmentation.  Although the mass at which fragmentation sets in is sensitive to our somewhat uncertain fragmentation criterion (\\S \\ref{criteria}) and angular momentum calculation (\\S \\ref{angmom} and the Appendix), we have been conservative in five ways: (1) by adopting a fragmentation temperature lower than that implied by the \\cite{2005MNRAS.364L..56R} simulations; (2) by adopting a low estimate of the specific angular momentum that determines the disk radius; (3) by adopting a relatively opaque model for the disk's Rosseland opacity; (4) by ignoring the shielding effect of a moderately opaque infall envelope, and (5) by adopting a low estimate for the cooling time, and a hard equation of state.  All of these approximations should, if anything, lead us to underestimate the prevalence of disk fragmentation.  Along with the existence of turbulent fluctuations in $j$ (\\S \\ref{coreval}), these points ensure that {\\em some} massive stars above $\\sim 10\\Msun$ experience disk fragmentation.   As noted in \\S \\ref{finitemass}, we cannot rule out the possibility that global instabilities flush disk material fast enough to suppress fragmentation, but we consider it unlikely that this prevents {\\em all} fragmentation.  We therefore expect multiple, coplanar, low-mass (M5 to G5, \\S \\ref{fragmass}) companions to form around many O (and some B) stars. Given initial separations of order 100-200 AU, their photospheric emission is not observable with present techniques.  Disk migration, followed by mass transfer or common-envelope evolution, may however make them evident as the primary evolves (\\S \\ref{observations}).  Even when disks fragment, we expect some accretion onto the central star -- if only because of material that falls within the fragmentation radius (\\S \\ref{diskeff}).  Although we cannot yet quantify the disk efficiency parameter $\\varepsilon_d$, we expect it to be significantly less than unity for those early O stars ($M_{\\star f}\\gtsim 50 \\Msun$) whose disks are most prone to fragment.   \\subsection{Imprint on the initial mass function} \\label{IMFlimit}  The stabilizing effect of viscous heating is absent in disks that accrete more rapidly than $1.7\\times 10^{-3}\\Msun$\\,yr$^{-1}$, thanks to a sharp drop in dust opacity at $\\sim 1050$ K (see \\S \\ref{coreval}).  As this affects stars of mass greater than 87$(t_{\\rm acc}/10^5$yr$)\\Msun$, or about 110 $\\Msun$ in the fiducial core model, it may be related to the cutoff of the initial mass function (IMF) at about 120 $\\Msun$.  Several other explanations for this cutoff have been proposed, all involving the increasing bolometric luminosity, ionizing luminosity, or outflow force emitted by the central star.  Starvation by disk fragmentation has the distinctive feature that it becomes much more severe at a specific accretion rate.  For this reason we expect it to produce a sharper IMF cutoff.  (The transition to super-Eddington luminosities could also produce a sharp cutoff, but \\citealt{2005ApJ...618L..33K} argue that this can be overcome by asymmetric radiation transfer.)  Note also that disk accretion is {\\em destabilized} by rapid accretion, whereas rapid accretion quenches the effects of direct photon force and of the ionizing radiation \\citep{WC1987}.  Disk fragmentation may close an avenue by which very massive stars would otherwise form.   \\subsection{Proto-binary disks} \\cite{PS2006} state that close massive binaries $< 10 $ year periods) are more likely to have nearly equal masses. More generally, the binarity fraction overall among massive stars is higher than their low mass counterparts ($1.5 \\mbox{ versus } 0.5$, \\cite{2005IAUS..227...12B}). Due to the high fraction of roughly equal mass binaires, their effect on disk dynamics must be addressed in future work. Moreover, the stellar densities in regions of HMSF are high, suggesting that other cluster stars might be close enough to interfere with disks stretching out to $\\sim$ 100 AU \\citep{2003MNRAS.339..577B}.  It is not currently clear whether disk accretion can preferentially grow a low-mass companion until its mass rivals that of the primary star, as \\cite{1996ApJ...467L..77A} suggest.  If so, then disk fragmentation may be relevant in the production of equal-mass binaries; if not, they must form by another mechanism.  In any case, the multiplicity of the center of gravity should be accounted for in future work.  It seems unlikely, though, that a binary with $\\sim 10$ year period will stabilize the fragmenting regions whose periods are $\\gtsim 300$ years.    \\noindent {\\bf Acknowledgments:} We are pleased to thank the referee, Jonathan Tan, for very useful suggestions, as well as Ray Jayawardhana, Scott Kenyon, Chris McKee, Stefan Mochnacki, Roman Rafikov, Debra Shepherd, and Yanqin Wu for helpful discussions, and Alyssa Goodman for a leading question. We would also like to thank Susana Lizano for insightful suggestions regarding disk observations. C.D.M.'s research is funded by NSERC and the Canada Research Chairs program. K.M.K.\\ is supported by a U. Toronto fellowship."
    },
    "0609/astro-ph0609316.txt": {
        "abstract": "We present  first results from a large-area ($\\sim$80$^\\circ$$\\times$20$^\\circ$), sparsely sampled two-filter (B,R) imaging survey towards the Canis Major stellar over-density, claimed to be a disrupting Milky Way satellite galaxy. Utilizing stellar colour-magnitude diagrams reaching to B$\\sim$22\\,mag, we provide a first delineation of its surface density distribution using main sequence stars. It is located below the  Galactic mid-plane,  and can  be discerned to at least b=-15$^\\circ$. Its projected shape is highly elongated, nearly parallel to the Galactic plane, with an axis ratio of $\\ga$ 5:1, substantially more so than what Martin et  al. originally found. We  also provide a first map of a prominent over-density of blue, presumably younger main sequence stars, which  appears to have a maximum near [l,b $\\sim$ 240$^\\circ$, -7$^\\circ$ ] and extends in latitude to b$\\sim$ -10$^\\circ$. The young population  is markedly more localized. We estimate an {\\it upper} limit  on the  line-of-sight depth, $\\sigma_{\\rm los}$, of the old population based on the main sequence width, by  comparing with a simple stellar population (Pal 12), obtaining $\\sigma_{\\rm los}$  $<$ 1.8 $\\pm$ 0.3\\,kpc, at an adopted  D$_\\odot$ = 7.5 $\\pm$ 1\\,kpc. For the young stellar population, we find  $\\sigma_{\\rm los}$ $<$ 1.5\\,kpc. The overall picture presented is one of a young stellar population that is less extended, both in terms its line-of-sight depth and angular size, than the older population. In the literature  three  explanations for the CMa over-density have been invoked, namely (a) a partially disrupting dwarf galaxy on a low-latitude orbit,  (b) a projection of a warped outer Galactic disk, and (c) a projection of an out-of-plane spiral arm. While the data provide no firm arguments against the less well-defined third scenario,  they have clear implications for each of the  others: (a) We infer from the strong elongation of the over-density in longitude, and simulations in the literature, that the CMa over-density  is unlikely to be a gravitationally bound system at the present  epoch, but may well be just a recently disrupted satellite remnant. A complexity for the satellite origin may arise from the `flattening' of the young MS population, which is possibly more pronounced than the older one.  (b) Based on modeling, we find that the line-of-sight depth of the MS over-density in old stars is clearly inconsistent with published locally axi-symmetric descriptions of the warped Galactic disk, such as those considered in Momany et al. (2006). Without detailed modeling, the data-set itself does not provide sufficient leverage to distinguish between an interpretation as  sub-structure in the warped outer Galactic  disk or a disrupted satellite. ",
        "introduction": "The current paradigm in which  large disk galaxies like our Milky Way form (e.g. White \\& Rees 1978; White \\&  Frenk 1991), is based on the successive coalescence and accretion of smaller systems of dark matter,  stars,  and diffuse interstellar  gas into larger assemblies. Some of these smaller systems may be satellite dwarf galaxies that can be disrupted by tidal shocks and evaporation in a Galactic potential, and can spawn tidal tails. Numerical simulations exist that model distinct tidal stellar streams in and around large galaxies, indicating that at r$\\ga 10$kpc such streams should remain detectable as coherent stellar over-densities for billions of years (Johnston et al. 1999; Ibata \\& Lewis 1998;  Mart\\'inez-Delgado et al. 2004; Pe\\~narrubia et al. 2006).   Recently, much work has focused on a possible low-latitude stellar stream, the so-called Monoceros stream, encircling the Milky Way (MW) at a galacto-centric distance of around 20\\,kpc, which was found through colour-magnitude diagrams (CMDs) from SDSS\\footnote{\\it Sloan Digital Sky Survey} (Newberg et al. 2002; Yanny et al. 2003); also see  Ibata et al. (2003),   Conn et al. (2005) \\& Martin et al. (2006). For stellar over-densities at low galactic latitudes it is not obvious  a priori whether such a stream is of external origin, i.e. the tidal debris of  a now (partially) disrupted satellite or, alternatively, a distorted part of the pre-existing outer stellar disk.   Arguments in favour of the dwarf satellite hypothesis have come from dynamical modeling (Pe\\~narrubia et al. 2005), showing that it is possible to find a plausible model  of the Monoceros stream that explains all detected parts of this low-latitude stream  as the wrapped tidal debris tails of a disrupting (model) dwarf galaxy. In this dynamical model the position of the CMa over-density is  not an initial constraint, and yet the main body of the disrupting model satellite happens to located in the direction of CMa at the present epoch. All kinematics, including the subsequently determined proper motion (Dinescu et al. 2005) of main sequence (MS) stars in a small area (0.25\\,deg$^2$) towards the CMa stellar over-density are consistent with the  dynamical model. Additional empirical support appears to come from the apparent bifurcation of suspected Monoceros stream stars at right ascension $\\sim$ 120$^\\circ$ in Fig.1 of Belokurov et al. (2006), as could be expected in general based on the above-mentioned Pe\\~narrubia et al. model.  Other evidence  comes from the discovery of an over-density of possible  Monoceros stream RR Lyrae stars (Vivas et al. 2006), whose distances closely match the model expectation in Pe\\~narrubia et al. (2005).   The position and survival of the  progenitor of this stellar stream is still under debate. The best candidate is a seemingly well-defined stellar over-density of stars discovered in the direction of Canis Major (CMa) using 2MASS red giants (Martin et al. 2004a). This discovery has spawned a lively debate in the literature on whether this apparent over-density of red giants could be part of the distorted outer stellar disk (Momany et al. 2004; 2006; see also Rocha-Pinto et al. 2006) or an accreted and possibly disrupting satellite  [Martin et al. 2004b; Martinez-Delgado et al 2005 (MD05); Bellazzini et al. 2006 (B06)]. However, a firm large-area kinematical, spatial and chemical link between the three known stellar  components of the overall debate (i.e. the young and old stellar over-densities, and the Monoceros stream)  is lacking.   From deep, visible-band photometry there are signs of at least two different star formation episodes in the direction of CMa [Bellazzini et al. (2005), B05; MD05; Carraro et al. (2005)]. For the sake of clarity, we illustrate this in an annotated colour-magnitude diagram, Fig.~\\ref{Fig_Select_CMD_single}, which corresponds to a field near the presumed centre of the CMa over-density, and where the data are taken from this survey (see Sec.~\\ref{obs_datared}).      Whether the most recent burst of star formation activity occurred 1-2\\,Gyr ago, as touted in B05, has been debated, in favour of a much younger stellar population of $\\la$ 100\\,Myr-old stars (Carraro et al. 2005). The CMa over-density  itself, which is taken to lie near D$_\\odot$=7.5\\,kpc (B05,MD05), comprises a predominantly older population of stars (4-10\\,Gyr; B05) and Moitinho et al. (2006) suggest that it may be consequence of viewing a  local arm-like MW stellar sub-structure in projection. For a different and independent study of the stellar populations, we will make use of detailed CMD-fitting in a forth-coming paper (de Jong et al., in prep.).  To determine the full angular extent of the young and old main sequence stars towards CMa and to test whether they stem from a satellite galaxy, we conducted a large-area ($\\sim$80$^\\circ$$\\times$20$^\\circ$), sparsely sampled visible-band imaging survey of the CMa region, drawing on the wide field imager (WFI) at the MPG/ESO 2.2\\,m telescope on La Silla. The survey sub-area considered in this paper is 230$^\\circ$ $\\la$ l  $\\la$ 260$^\\circ$, -20$^\\circ$ $\\la$ b $\\la$ 15$^\\circ$.  \\subsection{Aim and design of the WFI survey} The principal goal of our imaging survey is to provide a database that can address ultimately whether the CMa stellar over-density results from a dynamical distortion of the outer Galactic disk or whether it is the remains of a formerly large dwarf galaxy.  As a step towards this goal, we aim to map its  full angular extent through its star-count profile.  Recent reports (B06) based on 2MASS red clump  stars suggest that the discernable part of the CMa over-density may extend across several 100deg$^2$ on the sky. Our survey goal is to cover this over-density through sparse sampling, but with relatively deep visible-band imaging. Specifically, we want to obtain CMDs that reach  well below the turn-off magnitude of an ancient stellar population; our target limiting magnitude is B, R = 23\\,mag (S/N = 10). We chose a wide colour baseline (B-R) to have sensitivity to age- and/ or [Fe/H]-dependent CMD features (young MS, old MS turn-off and the red clump). The key difference (and merit) of the present CMa survey over previous large-area surveys, which have been based exclusively on 2MASS red-giant stars and MW models, is that we base our analysis on colour-magnitude diagrams that exhibit a well-defined stellar main sequence at a defined distance range. This provides an independent and  high-contrast way of tracing the full  angular extent of the  CMa over-density. The chief disadvantage of visible-band data is that Galactic extinction in many parts of this area is high, A$_{\\rm V}$$\\ga$0.5\\,mag.     \\subsection{Aim of this paper} In this paper we present the data from the first observing phase of our imaging survey, covering the region around [l,b = 240$^\\circ$, -8$^\\circ$] that had originally been identified as the over-density in Martin et al. (2004a). As it turns-out, these data provide only partial spatial coverage of the over-density. Therefore, we  refrain from a wide-spread quantitative comparison with Galaxy star-count models, except  (i) to use a Besancon model CMD as a visual guide for the reader to help interpret the general morphology in our control CMDs, and (ii) to aid the discussion in Sec.~\\ref{interp}.  We report our observations, and present the data and their reduction in Sec.~\\ref{obs_datared}. The issue of photometry completeness and uncertainties is assessed in Sec.~\\ref{complt_tst}. The crucial issue of reddening by dust extinction and its correction is presented in Sec.~\\ref{extinc}. In Sec.~\\ref{controlfields}, we explain the choice of control fields, which  is followed by a description of the stellar content and morphology of the CMDs in Sec.~\\ref{stel_content}. In   Sec.~\\ref{ms_counts} we provide a star-count analysis of young and old CMa main sequence stars. We estimate an upper limit on the line-of-sight size of the CMa stellar overdensity in Sec.~\\ref{los}. We briefly discuss our results in  Sec.~\\ref{discussion} and  summarize the key results in Sec.~\\ref{conclusions}. ",
        "conclusions": "We have presented initial results from our imaging survey of the Canis Major stellar over-density, providing a large, representative sub-set of colour-magnitude diagrams from our survey area, a depth estimation for the young and old MS populations, and a first analysis of their surface density distribution.  In particular, our key findings are these:  \\begin{itemize}  \\item Using `old' MS stars, we can delineate  an over-density of MS stars elongated along galactic longitude at a distance of D$_\\odot$$\\sim$7.5\\,kpc. It coincides with the over-density of 2MASS red-giants discovered in Martin et al. (2004a), but our surface density mapping reveals it to be markedly more elongated than initially thought. Its projected aspect ratio is probably $\\ga$ 5:1,  which is consistent with the more recent  2MASS analyses  (B06; Rocha-Pinto et al. 2006). We also map the angular distribution of the much bluer MS stars (the `young' MS).   \\item The distributions of young and old  stars are approximately co-spatial in projection, but the young MS density profile is markedly more localized, with a shallow maximum near  (l, b) $\\sim$ (240$^\\circ$, -7$^\\circ$).   \\item We report the clear detection of young and old MS stars up to 1.3\\,kpc and 2\\,kpc respectively below the Galactic midplane at l$\\sim$240$^\\circ$.   \\item We derive an upper limit on the line-of-sight depth of the (old) stellar over-density, by assuming that its stellar population is simple (as, e.g. globular cluster Pal 12), and that its MS width soley reflects a distance spread. We obtain  $\\sigma_{\\rm los}$ $<$ 1.8 $\\pm$  0.3\\,kpc (or $\\sim$ 4\\,kpc, FWHM$_{\\rm los}$) at the adopted D$_\\odot$ = 7.5 $\\pm$ 1\\,kpc. The young MS stars are consistent with $\\sigma_{\\rm los}$  $\\la$ 1.5\\,kpc.    \\end{itemize}  We discussed these results in the context of three broad explanations put forth in the literature for the CMa over-density. (a) We infer from the strong elongation of the over-density in longitude and simulations in the literature that, if it is a satellite galaxy on a near-circular, low-latitude orbit, it is unlikely to be still gravitationally bound at the present  epoch; it would have to be a recently disrupted satellite. (b) The distance and line-of-sight depth of the over-density is in disagreement with all published `smooth' (or locally axisymmetric)  models of the Galactic warp. Those produce a density profile that is markedly more extended along  the line-of-sight and reaches a maximum significantly  closer.  Lastly, of an out-of-plane spiral arm hypothesis for the young MS stars, we note that the presence of young  out-of-plane stars is not uncommon in the MW, and such stars in the literature have been attributed to possible accretion events, based on stellar kinematics.  Without detailed modeling the data themselves  are not yet sufficient to discriminate between an interpretation as sub-structure in the pre-existing warped outer Galactic  disk  or a disrupted satellite."
    },
    "0609/astro-ph0609302.txt": {
        "abstract": "We find 9 nuclear cluster candidates in a sample of 14 edge-on, late-type galaxies observed with HST/ACS.  These clusters have magnitudes ($M_I \\sim -11$) and sizes ($r_{eff} \\sim 3$~pc) similar to those found in previous studies of face-on, late-type spirals and dE galaxies.  However, three of the nuclear clusters are significantly flattened and show evidence for multiple, coincident structural components.  The elongations of these three clusters are aligned to within $\\sim$10$^\\circ$ of the galaxies' major axes.  Structurally, the flattened clusters are well fit by a combination of a spheroid and a disk or ring, with the disk preferred in two of three cases.  The nuclear cluster disks/rings have F606W-F814W ($\\sim$V-I) colors 0.3-0.6 magnitudes bluer than the spheroid components, suggesting that the stars in these components have ages $<$1~Gyr.  In NGC~4244, the nearest of the nuclear clusters, we further constrain the stellar populations via spectroscopy and multi-band photometry.  This nuclear cluster is equally well fit by single-stellar populations with ages of either $\\sim$100~Myr or $\\sim$1~Gyr, and with masses of 2-5$\\times$10$^6$ M$_\\odot$.  However, significantly better fits to the spectroscopy and photometry are obtained by combining two or more stellar populations.  Exploiting emission lines that appear to originate $\\sim$1\\arcsec\\ from the NGC~4244 nucleus, we determine a lower limit on the dynamical mass of 2.5$^{+1.7}_{-1.2} \\times 10^{6}$M$_\\odot$ within 19~pc, typical of values found for other nuclear clusters.  We also present tentative evidence that another of the nuclear clusters (in NGC~4206) may also host a supermassive black hole.  Based on our observational results we propose an {\\it in situ} formation mechanism for nuclear clusters in which stars form episodically in compact nuclear disks, and then lose angular momentum or heat vertically to form an older spheroidal structure.  We estimate the period between star formation episodes to be $\\sim$0.5 Gyr and discuss possible mechanisms for tranforming the disk-like components into spheroids.  We also note the connection between our objects and massive globular clusters (e.g. $\\omega$~Cen), ultra-compact dwarfs, and supermassive black holes. ",
        "introduction": "Due to their deep gravitational potential wells and unique dynamics, the centers of galaxies frequently host exotic objects such as supermassive black holes, active galactic nuclei, and massive stellar clusters.  Of these, only nuclear star clusters provide a visible record of the accretion of stars and gas into the nucleus.  Therefore, studies of nuclear star clusters hold the promise to %may shed light on the formation of central massive objects of all types.  Bright nuclear star clusters have been observed in $\\sim$75\\% of the 77 face-on galaxies of type Scd and later studied by \\citet{boker02}. Similar clusters appear to be frequent in earlier type spirals as well, although their study is complicated by the presence of bulges and other nuclear phenomena \\citep{carollo02}.  Dwarf elliptical galaxies, which may originate from low-mass spiral galaxies \\citep[e.g.][]{mastropietro05,beasley06}, also frequently possess nuclear star clusters \\citep{grant05,cote06}. %added mastropietro reference  In late-type spiral galaxies with little or no bulge, nuclear star clusters are very prominent and can be studied in detail.  From their sample of face-on, late-type spiral galaxies, \\citet{boker04a} found that nuclear clusters have dimensions similar to Milky Way globular clusters, with typical half-light radii of 3.5 pc.  However, these clusters are significantly brighter than globular clusters, with $I$ band magnitudes ranging between $-8$ and $-16$ \\citep{boker02}.  Their high luminosities are at least in part due to their large masses; \\citet{walcher05} have used high-resolution spectra to derive typical masses $\\sim$3$\\times10^6$M$_\\odot$ for nine galaxies in the \\citet{boker02} sample, similar in mass to the most massive Milky Way cluster, $\\omega$~Cen \\citep{meylan86}.  However, the brightness of nuclear clusters %may also results in part from the presence of young stellar populations. Photometry and spectra of a number of clusters suggest that nuclear clusters in both spiral and dwarf ellipticals have populations much younger than globular clusters \\citep[e.g.][]{ho95,lotz04}.  In late-type galaxies, most clusters appear to have stars with ages $<$100~Myr \\citep{walcher06}. %added in Walcher reference here Furthermore, recent spectral studies have shown that nuclear clusters are made up of composite stellar populations, with most having substantial old ($\\gtrsim$1~Gyr) stellar component \\citep{long02,sarzi05,walcher06,rossa06}. %Furthermore, spectra of nuclear %clusters in M33 \\citep{long02} and other galaxies \\citep{sarzi05} show %signatures of composite stellar populations.  Not only are nuclear clusters present in a large fraction of present day spirals and dwarf ellipticals, they also have been invoked as the progenitors of massive globular clusters and ultra-compact dwarfs. In a hierarchical merging scenario, many galaxies will be tidally stripped during mergers leaving behind only their dense nuclei \\citep[e.g.][]{bekki04}.  Nuclear cluster properties match the half-light radii and exceptional luminosity of massive globulars and UCDs \\citep{mackey05,phillipps01,jones06}.  Furthermore, the multiple stellar populations seen in $\\omega$~Cen and G1 \\citep[e.g.][]{bedin04,meylan01}, and the evidence for possible self-enrichment among the brightest globular clusters in ellipticals \\citep[][]{harris06} strengthen the association between these bright compact objects.   Finally nuclear clusters appear to be intimately connected to the formation of supermassive black holes (SMBHs) at the centers of galaxies.  For both spiral and elliptical galaxies, a clear correlation is seen between nuclear cluster mass and bulge mass \\citep{cote06,wehner06,ferrarese06,rossa06}.  This correlation appears to fall along the well-known $M_\\bullet-\\sigma$ relation between black hole mass and bulge or galaxy mass \\citep[e.g.][]{ferrarese00b,ferrarese02}.  Both SMBHs and nuclear clusters appear to contain the same fraction of the total galaxy mass, supporting the idea that they are two related types of %form a single class of central massive object \\citep{wehner06}.  Whether a central massive object is a nuclear cluster or a SMBH appears to depend on mass; early-type galaxies more massive than $\\sim$10$^{10}$~M$_\\odot$ appear to have SMBHs, while lower mass systems commonly host nuclear clusters \\citep{ferrarese06}.  These findings strongly suggest a similar formation mechanism for nuclear clusters and SMBHs.  Increasing evidence also points to massive star clusters as sites of intermediate-mass black hole (IMBH) formation \\citep[e.g.][]{miller04,gebhardt05,patruno06}.  These IMBHs may also follow the $M_\\bullet-\\sigma$ relation \\citep{gebhardt02}.  In this context, a detailed understanding of how nuclear clusters formed may reveal the general mechanism driving the formation of central massive objects of all types.  Two basic scenarios have been suggested \\citep[e.g.][]{boker04a}: (1) nuclear clusters form from multiple globular clusters accreted via dynamical friction \\citep[e.g.][]{tremaine75,lotz01}, and (2) nuclear clusters form {\\it in situ} from gas channeled into the center of galaxies \\citep{milosavljevic04}.  In this paper we present observational evidence that favors the {\\it in situ} formation scenario.  In particular, we find nuclear clusters that are elongated along the major axis of the galaxies and have two coincident components. %multiple components aligned %with the plane of the galaxy. These clusters appear to consist of a younger flattened disk or ring and an older spheroid. %added ring Based on the evidence we present, we suggest a detailed formation mechanism in which nuclear clusters form in repeated star formation events from very compact nuclear gas disks.  In \\S\\ref{obssec} we describe the imaging and spectroscopic observations used in the paper.  \\S\\ref{morphsec} focuses on an analysis of the morphology and luminosities of the clusters as revealed by HST observations.  The stellar populations of the multiple-component clusters are examined in \\S\\ref{popsec}.  In \\S\\ref{masssec} we derive a dynamical mass for the NGC~4244 nuclear cluster, and then in \\S\\ref{sbhsec} present tentative %added tentative evidence that the NGC~4206 nucleus may harbor both a star cluster and a supermassive black hole.  We present a detailed formation mechanism and discuss it in the broader context in \\S\\ref{dissec}, and then summarize our findings in \\S\\ref{sumsec}. ",
        "conclusions": "\\label{dissec}  We now discuss how our data constrains the formation mechanisms of nuclear clusters.  We suggest a formation mechanism in which stars in nuclear clusters form episodically in disks and then over time lose angular momentum or are heated vertically, ending up in a more spheroidal distribution. %added heated After discussing this mechanism in some detail, we examine the relation of nuclear clusters to other massive star clusters and supermassive black holes.  \\subsection{Formation Mechanisms}  Our observations of multi-component clusters place significant constraints on nuclear cluster formation mechanisms.  The presence of young NC disks ($<$1~Gyr) aligned with the disks of the host galaxies strongly suggests that they are formed {\\it in situ} from gas accreted into the nuclear regions.  We have shown in \\S\\ref{morphsec}.2 that these disks may be a common feature of nuclear clusters.  Such disks cannot be formed via the accretion of globular clusters by dynamical friction.  While globular cluster accretion very likely occurs, this mechanism has been shown to be relatively inefficient in late-type galaxies \\citep{milosavljevic04}.  We propose here a scenario in which the {\\it dominant} formation mechanism for nuclear clusters is via {\\it in situ} formation in disks. %added cautionary sentence Although the mechanism will need to be tested with detailed simulations, it is consistent not just with our data, but with both the evidence for young and composite stellar populations in many nuclear clusters \\citep[e.g][]{walcher06}, and the direct detection of a molecular gas disk coincident with the nuclear cluster in IC~342 \\citep{schinnerer03}. %added walcher reference  %made new paragraph and added most sentences On a large scale ($\\sim$1~kpc), both dissipative merging of molecular clouds and magnetorotational instabilities have been shown to be capable of supplying the nucleus of a late-type galaxy with sufficient gas to form nuclear clusters \\citep{milosavljevic04,bekki06}.  Our proposed {\\it in situ} formation mechanism is concerned with what occurs once the gas reaches the center. %added sentence Our model is also fully consistent with the suggestion by \\citet{walcher06} that an initial seed star cluster is required for subsequent nuclear star cluster formation. %added final two sentences  The details of an {\\it in situ} formation mechanism can be constrained by our observations.  To explain presence of redder/older spheroids in the multi-component clusters and the absence of disk components in a majority of our nuclear clusters, we propose a model in which nuclear cluster disks form episodically and are transformed into a rounder distribution over time due to loss of angular momentum or dynamical heating.  There are two important timescales in this model: (1) the period of time between star formation episodes, and (2) the time it takes stars in a disk morphology to end up in a spheroidal distribution.  We examine these two timescales in detail below.   \\subsubsection{Period Between Star Formation Episodes}  A number of authors have previously suggested that star formation in the nuclear regions of late-type spirals may be episodic. The physical motivation for episodic star formation is that a massive star formation event will feedback on the local environment preventing star formation for the next $\\sim$10$^8$ years \\citep{milosavljevic04}.  Although this feedback may destroy the gas disk and shut off star formation, it does not disrupt the disk of recently formed stars, whose lifetime we discuss below in \\ref{dissec}.1.2.  Assuming that the present epoch is not a special time in the life of the nuclear clusters, one can estimate the period on which star formation recurs.  Observationally, \\citet{davidge02} argue for a period of $\\sim$0.3~Gyr based on the detection frequency of Pa$\\alpha$ emission in galaxy nuclei. \\citet{schinnerer03} also determine a recurrence period of a few hundred Myr to 1 Gyr based on the gas accretion rates observed in IC~342's nucleus.  More recently, from analysis of the youngest stellar components in the spectrum, \\citet{walcher06} has found the typical time between star formation bursts to be $\\sim$70~Myr with a typical mass of roughly $1.6 \\times 10^5$~M$_\\odot$.  Our observations of the NGC~4244 nuclear cluster are consistent with these previous estimates.  Based on the $\\sim$10$^5$~M$_\\odot$ of stars in the young component (\\S\\ref{specpopsec}), approximately 25 episodes of disk formation must have occurred to build up the 2.5$\\times$10$^6$~M$_\\odot$ of the cluster (\\S\\ref{masssec}).  If this occurs over a Hubble time, then the period between bursts is $\\sim$0.5~Gyr.  \\subsubsection{Timescale of Disk Transformation}  Given that nuclear cluster disks are observed in $<$50\\% of nuclear clusters, and that these disks all have stellar populations younger than a Gyr, a mechanism must exist that destroys or disrupts this disk on a timescale shorter than $\\sim$1~Gyr.  We suggest that angular momentum loss or dynamical heating causes the disks to be transformed over time, thus creating the more spheroidal, older components seen in our nuclear clusters.  Although mechanisms of angular momentum loss have not been studied previously in nuclear clusters, these objects are dynamically similar to globular clusters, for which the theory of rotation and rotational evolution has been well-developed \\citep{agekian58,king61,shapiro76,fall85,davoust90,longaretti97}.  As they age, star clusters flattened due to rotation will become increasingly round due to the preferential evaporation of high angular momentum stars \\citep{agekian58}.  We can compare the timescale of this angular momentum loss to the expected lifetimes of our disks. The timescale over which the ellipticity of a rotating stellar cluster decreases is $\\gtrsim$40$\\times$ the median relaxation time, $\\tau_{rh}$ \\citep{shapiro76,fall85}.  For globular clusters in the Milky Way, $\\tau_{rh}$ is typically $\\sim$10$^8$ years.  However, this timescale is longer for more massive clusters like $\\omega$~Cen, with $\\tau_{rh} = 3.5$~Gyr \\citep{shapiro76}.  For the NGC~4244 NC we have calculated a median relaxation time for the whole cluster of $\\sim$2 Gyr \\citep[Eq. 8-72][]{binney87}.  For just the disk component, assuming it has a mass of $\\sim$10$^5$~M$_\\odot$ and a half mass radius of $\\sim$3 pc, we get $\\tau_{rh}=4\\times10^8$ years. These relaxation times suggest that the timescale for evaporation of angular momentum in NGC~4244 is in the range of 16-80~Gyr.  This timescale is significantly longer than the 100s of Myrs over which we expect the disk to be transformed into a spheroidal component.  It therefore appears that angular momentum loss due to processes internal to the clusters are unlikely to transform the disk components into spheroids on sufficiently short timescales.  Instead, we must look to external processes that can affect the the dynamical evolution of the nuclear cluster.  We %are currently investigating the response %of nuclear cluster disks to the overall potential of the galaxy, and have identified a few alternative mechanisms for transforming the stellar disk component into a spheroid.  These include: (1) small angular offsets of the nuclear disk from the plane of the galaxy, and (2) slight positional offsets of the cluster from the center of the galaxy potential, (3) torques on the stellar disk from successive gas accretion. %Initial simulations suggest that these effects can %cause strong evolution in the shape of the cluster on relatively short %timescales. A forthcoming paper will discuss this problem in more detail.  \\subsubsection{Rings}   We have been assuming thus far that stars are formed initially in disks, as suggested by our observations of NGC~4244 and NGC~4206. However, in \\S\\ref{morphsec}.1.3, we showed that one of the three multi-component NCs, in IC~5052, was best fit by a spheroid $+$ ring model with inner and outer ring radii of 6.8 and 11.8~pc.  We discuss here how such a ring could form and examine rings in context of our proposed nuclear cluster formation model.  Based on an elongated near-infrared component in IC~342, \\citet{boker97} suggested that the observed molecular ring (with radius\\footnote{The radius of their ring was recalculated assuming the Cepheid distance to IC~342 of 3.3~Mpc derived by \\citet{saha02}.} of $\\sim$65~pc)  results from an inner Lindblad resonances (ILR) with a stellar bar.  If a similar mechanism is at work in forming the ring in IC~5052, it would require the star cluster itself to have a barred disk morphology.  From HST imaging of nuclear clusters in face-on spirals, there is no evidence for bars on these scales \\citep{boker04a}.  A ring could also result from tidal stripping of an infalling molecular cloud or star cluster.  Assuming a pre-existing cluster with mass of $3 \\times 10^6$~M$_\\odot$ and a $10^5$~M$_\\odot$ molecular cloud, the molecular cloud would be tidally stripped between radii of 8 and 20~pc for molecular cloud number densities ranging between $5 \\times 10^4$ and $5 \\times 10^3$~cm$^{-3}$.  Thus, tidal stripping of a molecular cloud is a feasible mechanism for creating a ring such as the one in IC~5052.  Episodic star formation occuring at least occasionally in rings is consistent with the proposed formation model.  The ring in IC~5052 is internal to the cluster itself, so after loss of angular momentum or dynamical heating, its stars would still contribute to the stellar distribution of the spheroid.  The radius of the ring depends on the current mass of the star cluster in both mechanisms for ring formation described above.  This correlation might provide a natural explanation for the clear correlation of nuclear cluster size and luminosity observed in dwarf elliptical galaxies by \\citet{cote06}.  A similar explanation could work for formation in disks if the size of the disks depends on the mass of the nuclear cluster.  Detailed modelling of disk/ring formation and dissipation would be required to test the feasibility of this idea.  \\subsubsection{Other Formation Scenarios}  Our proposed scenario naturally explains the formation of multi-component clusters.  However other scenarios may also be consistent with our observations.  While our observations seem to require some {\\it in situ} formation to explain the presence of young stellar disks, they do not rule out the possibility that the multi-component nuclear clusters were formed in part from the merging of globular clusters \\citep[e.g.][]{tremaine75,lotz01}.  In this ``hybrid'' scenario, the young disk components may be a late-time veneer on an already massive cluster formed from merged globular clusters.  However, we note that \\citet{milosavljevic04} has calculated that the timescale for dynamical friction in late-type spirals is $\\gtrsim 3 \\times 10^{9}$ years even for relatively massive clusters (10$^6$~M$_\\odot$) located within the central kpc of the galaxy.  This quantity increases for less massive and more distant clusters.  He concludes that this timescale is thus too long to be the dominant mechanism for nuclear cluster formation, and shows that magnetorotational instabilities can transport sufficient quantities of gas into the galaxy center to account for the observed NCs in late-type spirals.    \\subsection{Relationship to Other Massive Star Clusters}  Our observations have implications beyond just the formation of nuclear star clusters in late-type galaxies.  We briefly review here the possible connections between late-type spiral nuclei and dwarf elliptical nuclei, and the evidence that some massive globular clusters and ultra-compact dwarfs are stripped nuclear clusters.   \\subsubsection{The relation between spiral and dE nuclei}  The recent study by \\citet{cote06} shows that the nuclei of dE and late-type spiral galaxies have comparable sizes, luminosities, and frequencies of detection (see also Fig.~\\ref{magfig}).  They argue that this similarity suggests a nuclear cluster formation mechanism that does not depend strongly on galaxy properties. However, the formation scenario suggested by our observations for nuclear clusters in late-type galaxies cannot take place in the gas-free dE galaxies.  We therefore interpret the similarity of the nuclei in these galaxies as support for dE galaxies evolving from low-mass spiral and irregular galaxies \\citep[e.g.][]{davies88}.  In this scenario, a dE,N galaxy would result from a nucleated late-type spiral galaxy after depletion of gas through star formation or galaxy harrassment \\citep{moore96}.  If so, then the ages of the nuclear clusters in dE galaxies can constrain the epoch when these systems were stripped.   \\subsubsection{Massive star clusters as stripped galaxy nuclei}  There is direct evidence that some massive globular clusters are associated with stripped galaxies.  The galactic globular cluster M54 appears to be the nucleus of the Sag dSph \\citep{layden00}, and six of the most massive globular clusters in NGC~5128 appear to be surrounded by extratidal light \\citep{harris02,martini04}.  Stripped nuclear clusters have also been invoked to explain the multiple stellar populations observed in the most massive nearby globular clusters \\citep[e.g][]{gnedin02}.  These observations include the large metallicity spread \\citep{smith04} and double main-sequence \\citep{bedin04} in $\\omega$~Cen, and the spread in red giant branch colors for G1 in Andromeda \\citep{meylan01}.  More recently, \\citet{harris06} have observed a trend towards redder colors with increasing luminosity for the brightest blue globular clusters in massive elliptical galaxies.  This trend is interpreted as evidence for self-enrichment in these clusters pointing to an extended star formation history more typical of nuclear clusters than globular clusters.  A similar redward trend is seen for dE nuclei by \\citet{cote06}.  Self-enrichment is naturally explained in an {\\it in situ} formation mechanism, but is not easily explained if these clusters are formed by merging less massive clusters via dynamical friction.  The ultracompact dwarf galaxies recently found in the Fornax \\citep{phillipps01,drinkwater03} and Virgo \\citep{jones06} clusters may also be related to nuclear star clusters.  One possible mechanism for their formation is ``galaxy threshing'', in which nucleated galaxies lose their envelopes leaving behind only their nuclei as UCDs \\citep{bekki01a}.  In the $M_I$ vs. log($r_{eff}$) diagram in Figure~\\ref{magfig}, NGC~4206's NC has sufficient size and luminosity, even after modest age fading, to be a UCD progenitor.  The galaxy threshing scenario has also been shown to be feasible by \\citet{cote06} for dwarf elliptical nuclei, although \\citet{depropris05} suggests that dE,N are in general more compact than UCDs.  One of the galaxy nuclei they study does have similar properties to the UCDs, as well as possessing spiral arms which indicate the galaxy may be an anemic spiral.   \\subsection{Connection to Black Holes}  As residents of the same neighborhoods, it is natural to suspect that there may be some connection between nuclear clusters and supermassive black holes (SMBHs).  The presence of SMBHs in the nuclei of $>$30 galaxies have been confirmed through dynamical measurements of their masses \\citep[reviewed by][]{ferrarese05}.  The masses of SMBHs are known to correlate with the mass and/or velocity dispersion of the bulge component \\citep[the $M_\\bullet-\\sigma$ relation,][]{ferrarese00b,gebhardt00}.  Recent work on nuclear clusters in both spiral and elliptical galaxies have shown a remarkably similar trend \\citep{rossa06,cote06,ferrarese06,wehner06}. For a large sample of nucleated elliptical galaxies, \\citet{cote06} have shown that the nucleus accounts for 0.30$\\pm$0.04\\% of the total galaxy mass.  This fraction is very similar to the fraction of galaxy mass found in the SMBHs \\citep{ferrarese06}.  For this reason, \\citet{cote06}, \\citet{ferrarese06} and \\citet{wehner06} suggest that nuclear clusters and SMBHs be thought of as a single-class of {\\it Central Massive Objects}.  They suggest a transition at $\\sim$10$^7$ M$_\\odot$ between nuclear clusters and SMBHs.  However, the detection of numerous lower mass SMBHs at the centers of early and late-type dwarf galaxies \\citep{barth05,greene05}, and the co-existence of nuclear star clusters and SMBHs in NGC~4395 \\citep{filippenko03} and perhaps in NGC~4206 (\\S\\ref{sbhsec}), suggest that this transition may not be abrupt.  The formation of black holes in the centers of nuclear clusters may be related to the intermediate mass black holes (IMBHs) found in the centers of massive non-nuclear star clusters.  A 1.7$\\times$10$^4$~M$_\\odot$ black hole has been reported by \\citet{gebhardt05} in the globular cluster G1 in Andromeda (which is also a possible stripped nuclei), based on stellar dynamics.  A 10$^2$-10$^4$~M$_\\odot$ black hole has also been identified in one of M82's young massive clusters, based on X-ray observations \\citep{patruno06}.  These black holes are thought to form either by early merging of massive young stars within the cluster, or by dynamical interaction in the core of the clusters over long time periods \\citep[][and references therein]{miller04}.  If IMBH formation is common in dense massive clusters, then nuclear clusters should frequently host black holes as well.  Subsequent mergers of many IMBHs could then form supermassive black holes (SMBHs) \\citep{miller04}.  As the mass of the SMBH increases, these mergers will likely destroy the nuclear clusters \\citep[see \\S5.2.2 of ][]{cote06}, distributing their stars into the bulge."
    },
    "0609/hep-ph0609158_arXiv.txt": {
        "abstract": "Neutrino masses are likely to be a manifestation of the right-handed, or sterile neutrinos.  The number of sterile neutrinos and the scales of their Majorana masses are unknown.  We explore theoretical arguments in favor of the high and low scale seesaw mechanisms, review the existing experimental results, and discuss the astrophysical hints regarding sterile neutrinos. ",
        "introduction": "The term {\\em sterile neutrino} was coined by Bruno~Pontecorvo, who hypothesized the existence of the right-handed neutrinos in a seminal paper~\\cite{pontecorvo}, in which he also considered vacuum neutrino oscillations in the laboratory and in astrophysics, the lepton number violation, the neutrinoless double beta decay, some rare processes, such as $\\mu \\rightarrow e \\gamma$, and several other questions that have dominated the neutrino physics for the next four decades.   Most models of the neutrino masses introduce sterile (or right-handed) neutrinos to generate the masses of the ordinary neutrinos via the seesaw mechanism~\\cite{seesaw}.  The seesaw lagrangian \\begin{equation} {\\cal L} = {\\cal L_{\\rm SM}}+\\bar N_{a} \\left(i \\gamma^\\mu \\partial_\\mu \\right ) N_{a} - y_{\\alpha a} H \\,  \\bar L_\\alpha N_{a} - \\frac{M_{a}}{2} \\; \\bar {N}_{a}^c N_{a} + h.c. \\,, \\label{lagrangianM} \\end{equation} where $ {\\cal L_{\\rm SM}}$ is the lagrangian of the Standard Model, includes some number $n$  of singlet neutrinos $N_a$ with Yukawa couplings $ y_{\\alpha a}$.  Here $H$ is the Higgs doublet and $L_\\alpha$ ($\\alpha=e,\\mu,\\tau$) are the lepton doublets. Theoretical considerations do not constrain the number $n$ of sterile neutrinos. In particular, there is no constraint based on the anomaly cancellation because the sterile fermions do not couple to the gauge fields. The experimental limits exist only for the larger mixing angles~\\cite{sterile_constraints}. To explain the neutrino masses inferred from the atmospheric and solar neutrino experiments, $n=2$ singlets are sufficient~\\cite{2right-handed}, but a greater number is required if the lagrangian (\\ref{lagrangianM}) is to explain the LSND~\\cite{deGouvea:2005er}, the r-process nucleosynthesis~\\cite{r}, the pulsar kicks~\\cite{ks97,Kusenko:review}, dark matter~\\cite{dw,Fuller,shi_fuller,nuMSM}, and the formation of supermassive black holes~\\cite{biermann_munyaneza}.  The scale of the right-handed Majorana masses $M_{a}$ is unknown; it can be much greater than the electroweak scale~\\cite{seesaw}, or it may be as low as a few eV~\\cite{deGouvea:2005er,deGouvea:2006gz}.   The seesaw mechanism~\\cite{seesaw} can explain the smallness of the neutrino masses in the presence of the Yukawa couplings of order one if the Majorana masses $M_a$ are much larger than the electroweak scale. Indeed, in this case the masses of the lightest neutrinos are suppressed by the ratios $ \\langle H \\rangle/M_a$.  However, the origin of the Yukawa couplings remains unknown, and there is no experimental evidence to suggest that these couplings must be of order 1. In fact, the Yukawa couplings of the charged leptons are much smaller than 1. For example, the Yukawa coupling of the electron is as small as $10^{-6}$.  One can ask whether some theoretical models are more likely to produce the numbers of order one or much smaller than one.  The two possibilities are, in fact, realized in two types of theoretical models.  If the Yukawa couplings arise as some topological intersection numbers in string theory, they are generally expected to be of order one~\\cite{Candelas:1987rx}, although very small couplings are also possible~\\cite{Eyton-Williams:2005bg}. If the Yukawa couplings arise from the overlap of the wavefunctions of fermions located on different branes in extra dimensions, they can be exponentially suppressed and are expected to be very small~\\cite{Mirabelli:1999ks}.  In the absence of the fundamental theory, one may hope to gain some insight about the size of the Yukawa couplings using 't~Hooft's naturalness criterion~\\cite{tHooft}, which states essentially that a number can be naturally small if setting it to zero increases the symmetry of the lagrangian. A small breaking of the symmetry is then associated with the small non-zero value of the parameter.  This naturalness criterion has been applied to a variety of theories; it is, for example, one of the main arguments in favor of supersymmetry. (Setting the Higgs mass to zero does not increase the symmetry of the Standard Model.  Supersymmetry relates the Higgs mass to the Higgsino mass, which is protected by the chiral symmetry.  Therefore, the light Higgs boson, which is not natural in the Standard Model, becomes natural in theories with softly broken supersymmetry.) In view of 't~Hooft's criterion, the \\textit{small} Majorana mass is natural because setting $M_a$ to zero increases the symmetry of the lagrangian (\\ref{lagrangianM})~\\cite{Fujikawa:2004jy,deGouvea:2005er}.  One can ask whether cosmology can provide any clues as to whether the mass scale of sterile neutrinos should be above or below the electroweak scale.  It is desirable to have a theory that could generate the matter--antimatter asymmetry of the universe. In both limits of large and small $M_a$ one can have a successful leptogenesis: in the case of the high-scale seesaw, the baryon asymmetry can be generated from the out-of-equilibrium decays of heavy neutrinos~\\cite{Fukugita:1986hr}, while in the case of the low-energy seesaw, the matter-antimatter asymmetry can be produced by the neutrino oscillations~\\cite{baryogenesis}.  The Big-Bang nucleosynthesis (BBN) can provide a constraint on the number of light relativistic species in equilibrium~\\cite{bbn}, but the sterile neutrinos with the small mixing angles may never be in equilibrium in the early universe, even at the highest temperatures~\\cite{dw}.  Indeed, the effective mixing angle of neutrinos at high temperature is suppressed due to the interactions with plasma~\\cite{high-T}, and, therefore, the sterile neutrinos may never thermalize.  High-precision measurements of the primordial abundances may probe the existence of sterile neutrinos and the lepton asymmetry of the universe in the future~\\cite{Smith:2006uw}.  While many seesaw models assume that the sterile neutrinos have very large masses, which makes them unobservable, it is worthwhile to consider light sterile neutrinos in view of the above arguments, and also because they can explain several experimental results.  In particular, sterile neutrinos can account for cosmological dark matter~\\cite{dw}, they can explain the observed velocities of pulsars~\\cite{ks97,Kusenko:review}, the x-ray photons from their decays can affect the star formation~\\cite{reion}.  Finally, sterile neutrinos can explain the LSND result~\\cite{deGouvea:2005er,Sorel:2003hf,LSND_sterile_decay}, which is currently being tested by the MiniBooNE experiment. ",
        "conclusions": "The underlying physics responsible for the neutrino masses is likely to involve right-handed, or sterile neutrinos.  The Majorana masses of these states can range from a few eV to values well above the electroweak scale.  Theoretical arguments have been made in favor of both the high-scale and low-scale seesaw mechanisms: the high-scale seesaw may be favored by the connection with the Grand Unified Theories, while the low-scale seesaw is favored by 't~Hooft's naturalness criterion.  Cosmological considerations are consistent with a vast range of mass scales.  The laboratory bounds do not provide significant constraints on the sterile neutrinos, unless they have a large mixing with the active neutrinos.  The atmospheric and solar neutrino oscillation results cannot be reconciled with the LSND result, unless sterile neutrinos (or other new physics) exist.  There are several indirect astrophysical hints in favor of sterile neutrinos at the keV scale.  Such neutrinos can explain the observed velocities of pulsars, they can be the dark matter, and they can play a role in star formation and reionization of the universe.  The preponderance of indirect astrophysical hints may be a precursor of a major discovery, although it may also be a coincidence.  One can hope to discover the sterile neutrinos in the x-ray observations. The mass around 3~keV and the mixing angle  $ \\sin^2 \\theta \\sim 3\\times 10^{-9}$ appear to be particularly interesting because the sterile neutrino with such parameters could simultaneously explain the pulsar kicks and dark matter (assuming the sterile neutrinos are produced at the electroweak scale).  However, it is worthwhile to search for the signal from sterile dark matter in other parts of the allowed parameter space shown in Fig.~\\ref{fig:range}.  The existence of a much lighter sterile neutrino, with a much greater mixing angle can be established experimentally if MiniBooNE confirms the LSND result.   \\ack This work was supported in part by the DOE grant DE-FG03-91ER40662 and by the NASA ATP grants NAG~5-10842 and NAG~5-13399.  The author thanks Theory Unit at CERN and ITP, EPFL, for hospitality."
    },
    "0609/astro-ph0609153_arXiv.txt": {
        "abstract": "We present the first results from three-dimensional spherical shell simulations of magnetic dynamo action realized by turbulent convection penetrating downward into a tachocline of rotational shear.  This permits us to assess several dynamical elements believed to be crucial to the operation of the solar global dynamo, variously involving differential rotation resulting from convection, magnetic pumping, and amplification of fields by stretching within the tachocline.  The simulations reveal that strong axisymmetric toroidal magnetic fields (about 3000 G in strength) are realized within the lower stable layer, unlike in the convection zone where fluctuating fields are predominant.  The toroidal fields in the stable layer possess a striking persistent antisymmetric parity, with fields in the northern hemisphere largely of opposite polarity to those in the southern hemisphere.  The associated mean poloidal magnetic fields there have a clear dipolar geometry, but we have not yet observed any distinctive reversals or latitudinal propagation.  The presence of these deep magnetic fields appears to stabilize the sense of mean fields produced by vigorous dynamo action in the bulk of the convection zone. ",
        "introduction": "The solar global dynamo responsible for the observed 22-year cycles of magnetic activity is likely to require several dynamical processes operating at differing sites within the convection zone and below its base.  The primary elements  were first discussed within the conceptual ``interface dynamo'' proposed by Parker (1993) and then explored by Charbonneau \\& MacGregor (1997).  The key dynamical processes involve (a) the generation of magnetic fields by the intense turbulence influenced by rotation within the deep solar convection zone, (b)~the transport or pumping of these fields downward into the tachocline of shear at the base of this zone, (c) the stretching of fields there by differential rotation to form strong toroidal fields, and (d) the magnetic buoyancy instability of such structures leading to field loops ascending toward the surface.  Fully self-consistent magnetohydrodynamic (MHD) simulations of the complete solar global dynamo are not yet feasible, given the vast range of dynamical scales involved in these turbulent processes.  However, many of the elements have now begun to be studied singly with reasonable fidelity using three-dimensional numerical modelling, often turning to localized planar domains to obtain sufficient spatial resolution within the simulations (see reviews by Ossendrijver 2003, Fan 2004, Charbonneau 2005, Miesch 2005).  Yet some aspects require using full spherical shell geometries in the modelling. Early simulations of dynamo processes in spherical shells provided insights into the coupling of convection, rotation and magnetism, but were limited by spatial resolution (e.g., Gilman 1983; Glatzmaier 1985a, b). Recent studies of the interaction of turbulent convection with rotation (e.g., Miesch et al. 2000; Brun \\& Toomre 2002; Miesch, Brun \\& Toomre 2006) now yield global differential rotation profiles in close accord with helioseismic deductions (e.g., Thompson et al. 2003) in the bulk of the convection zone.  Similarly, related MHD modelling of dynamo processes in such deep shells of turbulent convection has revealed that strong magnetic fields are produced without diminishing the differential rotation significantly (Brun, Miesch \\& Toomre 2004, hereafter BMT04).  Fluctuating magnetic fields with strengths of order 5000 G are realized that possess complex structures, with radial fields concentrated in downflow lanes and toroidal fields appearing as twisted ribbons extended in longitude.  However, the associated mean fields are relatively weak and do not exhibit the systematic latitudinal propagation or periodic polarity reversals seen in the Sun.  Whereas BMT04 only considered dynamical element (a), we report here on also incorporating elements (b) and (c) into global simulations of dynamo action achieved by turbulent convection able to penetrate downward into a tachocline of rotational shear. Such modeling is the next step in refining our intuition about the operation of the interface dynamo. ",
        "conclusions": ""
    },
    "0609/astro-ph0609686_arXiv.txt": {
        "abstract": "A review of the study of superclusters based on the 2dFGRS and SDSS is given. Real superclusters are compared with models using simulated galaxies of the Millennium Run.  We show that the fraction of very luminous superclusters in real samples is about five times greater than in simulated samples. Superclusters are generated by large-scale density perturbations which evolve very slowly. The absence of very luminous superclusters in simulations can be explained either by non-proper treatment of large-scale perturbations, or by some yet unknown processes in the very early Universe. ",
        "introduction": "Superclusters are the most extensive density enhancements in the Universe of common origin.  Investigation of large systems of galaxies was pioneered by the study of the {\\em Local Supercluster} by de Vaucouleurs \\cite{deV53}, and by Abell \\cite{abell61} using rich clusters of galaxies by Abell \\cite{abell} and Abell et al. \\cite{aco}. Superclusters consist of galaxy systems of different richness: single galaxies, galaxy groups and clusters, aligned in chains (J\\~oeveer, Einasto \\& Tago \\cite{jet78}, Gregory \\& Thompson \\cite{gt78}, Zeldovich, Einasto \\& Shandarin \\cite{zes82}).  New deep galaxy surveys, such as the Las Campanas Galaxy Redshift Survey, the 2 degree Field Galaxy Redshift Survey (2dFGRS, Colless et al.  \\cite{col01}, \\cite{col03}) and the Sloan Digital Sky Survey (SDSS, Adelman-McCarthy et al. \\cite{dr4}) cover large areas in the sky and are almost complete up to fairly faint apparent magnitudes. Thus these surveys are convenient to detect superclusters using both galaxy and cluster data.  This possibility has been used by Basilakos \\cite {bas03}, Basilakos et al. \\cite{bpr01}, Erdogdu et al.  \\cite{erd04}, Porter and Raychaudhury \\cite{pr05}, and Einasto et al.  \\cite{e03a}, \\cite{e03b}, \\cite{ein05}, \\cite{e06a}.  The goal of the present review is to analyse properties of superclusters based on the 2dF Galaxy Redshift Survey and the Sloan Digital Sky Survey Data Release 4 by Einasto et al. \\cite{e06a}, \\cite{e06b} and \\cite{e06c}, and to compare properties of real superclusters with theoretical models. For comparison we use the superclusters found for the Millennium Run mock galaxy catalogue by Croton et al. \\cite{croton06}, that itself based on the Millennium Simulation of the evolution of the Universe by Springel et al. \\cite{springel05}.  \\begin{figure*}[ht] \\centering \\resizebox{0.49\\textwidth}{!}{\\includegraphics*{lum_2dF_N7_0.8s.eps}} \\resizebox{0.49\\textwidth}{!}{\\includegraphics*{lum_2dF_S7_0.8s.eps}}\\\\ \\caption{The high-resolution 2-dimensional density fields of the Northern and Southern parts of the 2dF redshift survey are shown in the left and right panels, respectively.  The samples are conical, i.e. their thickness increases with distance, thus at large distances from the observer we see many more systems of galaxies. The richest supercluster in the Northern region is SCL126 in the list by Einasto et al. \\cite{e1997}, also called the Sloan Great Wall by Vogeley et al. \\cite{vogeley04}; the richest Southern supercluster is SCL9 by Einasto et al. \\cite{e1997}, or the Sculptor Supercluster.} \\label{fig:1} \\end{figure*} ",
        "conclusions": "\\begin{enumerate}  \\item Geometric properties of superclusters are well explained by current models.  \\item There are much more very rich superclusters than models predict.  \\item Large perturbations evolve very slowly and represent the fluctuation field at the epoch of inflation.  \\item The difference between observations and models can be explained in two ways: \\\\ large-scale perturbations are not incorporated in the models, i.e. models need   improvement; \\\\ there ocurred presently unknown processes during  inflation.  \\end{enumerate}   The present review is based on talks held in Budapest on April 20, 2006 in Detre Centenarium, in Uppsala University on April 27, 2006 and in Aspen Workshop on Cosmic Voids on June 6, 2006.  I thank my collaborators Maret Einasto, Enn Saar, Erik Tago and Volker M\\\"uller for permission to use results of our common work in this review.  We are pleased to thank the 2dFGRS and SDSS Teams for the publicly available data releases.  The present study was supported by Estonian Science Foundation grants No.  4695, 5347 and 6104, and Estonian Ministry for Education and Science support by grant TO 0060058S98.  I thank Astrophysikalisches Institut Potsdam (using DFG-grant 436 EST 17/2/05) and Uppsala University for hospitality where part of this study was performed.  2dFGRS supercluster catalogues are available at {\\tt http://www.aai.ee/~maret/2dfscl.html}, Sloan DR4 supercluster catalogues at {\\tt http://www.aai.ee/~maret/SDSSDR4scl.html}."
    },
    "0609/astro-ph0609015_arXiv.txt": {
        "abstract": "{Recent improvements in the modeling of solar convection and line formation led to downward revisions of the solar photospheric abundances of the lighter elements, which in turn led to changes in the radiative opacity of the solar interior and hence to conflicts with the solar convection zone depth as inferred from helioseismic oscillation frequencies. An increase of the solar Ne/O abundance to values as observed for nearby stars has been proposed as a solution.  Because of the absence of strong neon lines in the optical, neon abundances are difficult to measure and the correct solar and stellar Ne/O abundances are currently hotly debated. Based on X-ray spectra obtained with \\emph{XMM-Newton}, we determine a reference value of Ne/O for the inactive, solar-like star $\\alpha$~Cen (primarily $\\alpha$~Cen~B, which is the dominant component in X-rays), with three independent, line-based methods, using differential emission measure reconstruction and an emission measure-independent method. Our results indicate a value of $\\approx 0.28$ for $A_{Ne}$/$A_O$ in $\\alpha$~Cen, approximately twice the value measured for the Sun, but still below the average value obtained for other stars. The low Ne/O abundance of the Sun is peculiar when compared to $\\alpha$ Cen and other stars; our results emphasize the necessity to obtain more and accurate Ne/O abundance measurements of low activity stars.} ",
        "introduction": "The canonical values of the solar photospheric elemental abundances have recently become far less canonical.  Starting from the work of \\citet{Anders_Grevesse}, a standard reference for many years, the absolute abundance values of the more abundant trace elements like carbon, nitrogen, oxygen, neon, or iron have been significantly reduced over the past decade, initially by \\citet{Grevesse_Sauval}, and again more recently \\citep[and references therein]{Asplund}. These lighter elements considerably influence the physics of the solar interior since they provide a substantial contribution to its radiative opacity.  A change in the elemental abundances usually changes the depth of the solar convection zone, which can be inferred from the measured helioseismological oscillation frequencies.  With the \"old\" abundances by \\citeauthor{Grevesse_Sauval}, good agreement could be found between the standard solar model of the appropriate age and the observed oscillation spectrum, while the \"new\" but rather controversial abundances proposed by \\citeauthor{Asplund} and collaborators turned out to be inconsistent with helioseismology \\citep{Bahcall_helioseismology}.  In order to rescue the agreement between the standard solar model and helioseismology the opacity reduction by the downward revision of the CNO abundances must be sufficiently compensated by increased abundances of other elements. The only suitable element is neon, since its (photospheric) abundance is not well determined due to the absence of strong photospheric lines; rather, the solar neon abundance is obtained either from solar energetic particles or from coronal measurements in the X-ray or EUV bands. The solar neon abundance is usually quoted relative to oxygen, and the solar Ne/O abundance ratio has remained more or less constant during the course of the revisions of the solar photospheric abundances, with values ranging from 0.14 from the compilation of \\citeauthor{Anders_Grevesse}, 0.18 from \\citeauthor{Grevesse_Sauval}, and 0.15 for the new set from \\citeauthor{Asplund}. An increase of the solar Ne/O abundance by a factor of 2.5--3.5 \\citep{Antia_Basu,Bahcall} would provide sufficient opacity to reconcile the low oxygen abundance with helioseismology. Evidence for an increased neon abundance has been proposed by \\citet{Drake_Testa} in their survey of the coronal Ne/O abundance in a sample of nearby stars, finding an average value of $A_{Ne}/A_O = 0.41$.  On the other hand, a re-analysis of \\emph{SOHO} CDS spectra and re-investigation of archival solar coronal X-ray spectra confirm the long-established, \"canonical\"  lower $A_{Ne}/A_O$ values \\citep{Young_neon_to_oxygen,Schmelz_neon_to_oxygen}. Also, a closer look at the sample of stars used by \\citeauthor{Drake_Testa} reveals that most of these stars are RS~CVn systems or  well-known young and active stars, known to show the inverse FIP effect \\citep{Brinkman_HR_1099}, i.\\,e., an enhancement of elements with high first ionization potential. Since neon is (apart from helium) the element with the highest FIP and the occurrence of the inverse FIP effect is related to activity \\citep[and references therein]{Guedel_review}, the sample may be biased to higher neon abundances and not be representative of the \"true\" cosmic neon abundance.  In order to settle the issue of a possible bias, a comparison to exclusively low-activity solar-like stars is needed. Due to their low X-ray luminosity, high-resolution X-ray spectra with reasonable signal-to-noise ratio of such inactive stars can be obtained only for very few objects, like $\\epsilon$~Eri, Procyon, or $\\alpha$~Cen. Additionally, their low coronal temperatures complicate the measurement of the otherwise prominent \\ion{Ne}{ix} and \\ion{Ne}{x} He-like and H-like lines that have peak formation temperatures of $\\log T = 6.6$ and $\\log T = 6.75$ respectively; for example, the \\emph{Chandra} LETGS spectra of $\\alpha$ Cen A and B presented by \\citet{Raassen}) do not show the \\ion{Ne}{x} Ly~$\\alpha$ line.  Our new \\emph{XMM} RGS spectra of $\\alpha$~Cen provide good sensitivity and signal-to-noise to detect and measure the relevant H- and He-like lines of O and Ne to allow an accurate determination of its neon-to-oxygen abundance. ",
        "conclusions": "The results of our abundance modeling (cf., Table~\\ref{abundances}) are very robust and yield values of $A_{Ne}$/$A_O \\approx 0.28$, independent of the applied method.  This value is twice as large as the \"canonical\" $A_{Ne}$/$A_O$ for the solar corona.  $\\alpha$~Cen is probably the most suitable star for a comparison with the Sun avoiding a possible FIP/I-FIP bias. Our $DEM$ of $\\alpha$~Cen resembles that of the quiet Sun \\citep[e.\\,g.][]{Brosius,Landi_Landini_DEM}, which typically peaks around $\\log T \\approx 6.0$--6.2. However our values refer to  $\\alpha$~Cen~B, which is by far the more active component known to show flaring activity \\citep{NEXXUS,Robrade_alpha_Cen}.  Separate X-ray spectra of $\\alpha$~Cen~A and B are available from a 79~ks \\emph{Chandra} LETGS exposure. However, both spectra have extremely low signal in the wavelength range covering the \\ion{Ne}{ix} and \\ion{Ne}{x} lines. Values of $A_{Ne}$/$A_O = 0.18 \\pm 0.07$ and $0.24 \\pm 0.09$ for $\\alpha$~Cen~A and B respectively were derived from global fitting by \\citet{Raassen} and are thus based primarily on \\ion{Ne}{vii} and \\ion{Ne}{viii} lines located at longer wavelengths. Many of these lines suffer from significant blending as well as low signal and their atomic physics parameters should be considered as more uncertain than those of the H-like and He-like lines.  Formally, the Ne/O abundances of $\\alpha$~Cen~A and B are consistent with each other, and the value for the B~component is consistent with our XMM-Newton result.  Measurements of the solar Ne/O abundance ratio tend to show a broad scatter (cf. the compilation provided by \\citet{Drake_Testa} in the supplementary information, with values of $A_{Ne}$/$A_O$ ranging from 0.08 to 0.47), but the majority of them, based on miscellaneous data like solar energetic particles, X-ray or EUV spectra, are in good agreement with the \"low\" Ne/O abundances. Additionally, the most recent analyses of \\citet{Young_neon_to_oxygen} and \\citet{Schmelz_neon_to_oxygen}, based on the most recent atomic data, clearly support values as low as 0.15.  All stars in the survey of \\citeauthor{Drake_Testa} show higher values of $A_{Ne}$/$A_O$, incompatible with the \"canonical\" low solar value.  The inference of the solar Ne/O abundance from other stars is, however, problematic. Apart from the fundamental question of why the Sun should implicitly show the same abundance pattern as other stars do (and many stellar photospheric measurements show that it does not), the most severe problem is to find truly solar-like stars, i.\\,e. relatively old and inactive single stars of similar spectral type. These conditions and the basic requirement that the stars are observable in X-rays (or in the EUV) with a sufficient signal to obtain abundance measurements of neon and oxygen, are almost mutually exclusive since stars with an X-ray luminosity as low as that of the Sun can only be observed in the very solar vicinity with today's X-ray telescopes. Instead, the typical well-studied stellar coronal X-ray source is much brighter, usually consisting of an active young late-type star or even an RS~CVn system. Such objects are not appropriate for a direct comparison with the Sun, especially if one accepts the reality of the inverse FIP effect, i.\\,e. an enhancement of elements with high first ionization potential; while the physics of abundance anomalies like the inverse FIP effect and its counterpart, the FIP effect as observed on the Sun, are not fully understood, a framework to explain both effects has been provided by \\citet{Laming}.  A correlation seems to exist in the sense that the FIP effect turns into the inverse FIP effect, with the inverse FIP effect becoming stronger with increasing activity. Clearly, the sample used by \\citeauthor{Drake_Testa} is then strongly biased. The only star in their sample of truly solar-like activity is Procyon, where a value of 0.42 from \\citet{coronal_photospheric} was used, which was obtained from three combined \\emph{Chandra} LETGS spectra but still low signal at the wavelengths of the \\ion{Ne}{ix} and \\ion{Ne}{x} lines. For part of these data \\citet{Raassen_Procyon} obtained $A_{Ne}$/$A_O = 0.22$ with a global fitting approach, a value also found by the same authors from 91~ks of \\emph{XMM} RGS and MOS data; thus the correct Ne/O abundance ratio of Procyon remains an open issue.  Moderately active K dwarfs seem to have values of $A_{Ne}$/$A_O \\sim 0.37$, i.e,  slightly lower than the median obtained by \\citeauthor{Drake_Testa}, but still at approximately twice the level of the Sun as shown by  \\citet{Wood_Linsky}, who investigated Ne/O abundance ratios of $\\epsilon$~Eri,  36~Oph and 70~Oph.  As pointed out above, our Ne/O value also refers to a K star, $\\alpha$~Cen~B, which is less active than any of the stars studied by \\citeauthor{Wood_Linsky}. Therefore, none of the stars studied so far has a Ne/O abundance as low as observed for the Sun.  Therefore, the question remains why the Sun is so special. Or do we first have to find a real solar-like star?"
    },
    "0609/astro-ph0609824_arXiv.txt": {
        "abstract": "Using a set of hydrodynamical simulations of 9 galaxy clusters with masses in the range $1.5 \\times 10^{14} M_{\\odot} < M_{\\rm vir} < 3.4 \\times 10^{15} M_{\\odot}$, we have studied the density, temperature and X-ray surface brightness profiles of the intra-cluster medium in the regions around the virial radius. We have analyzed the profiles in the radial range well above the cluster core, the physics of which are still unclear and matter of tension between simulated and observed properties, and up to the virial radius and beyond, where present observations are unable to provide any constraints. We have modeled the radial profiles between $0.3 R_{200}$ and $3 R_{200}$ with power laws with one index, two indexes and a rolling index. The simulated temperature and [0.5--2] keV surface brightness profiles well reproduce the observed behaviours outside the core. The shape of all these profiles in the radial range considered depends mainly on the activity of the gravitational collapse, with no significant difference among models including extra-physics. The profiles steepen in the outskirts, with the slope of the power-law fit that changes from -2.5 to -3.4 in the gas density, from -0.5 to -1.8 in the gas temperature, and from -3.5 to -5.0 in the X-ray soft surface brightness. We predict that the gas density, temperature and [0.5--2] keV surface brightness values at $R_{200}$ are, on average, $0.05$, $0.60$, $0.008$ times the measured values at $0.3 R_{200}$. At $2 R_{200}$, these values decrease by an order of magnitude in the gas density and surface brightness, by a factor of 2 in the temperature, putting stringent limits on the detectable properties of the intracluster-medium (ICM) in the virial regions. ",
        "introduction": "\\label{sect:intro}  Galaxy clusters form in correspondence of the highest peaks of the fluctuations in the cosmic dark matter density field and collapse under the action of gravitational attraction to define structures with a mean density enhanced with respect to the cosmic value by a factor of few hundreds. They reach the virial equilibrium over a volume with a typical virial radius that indicates, as a first approximation, the regions where the pristine gas accretes on the dark matter halo through gravitational collapse and is heated up to millions degrees through adiabatic compression and shocks. This intracluster-medium (ICM), that represents about 80 (15) per cent of the baryonic (total) cluster mass \\citep[e.g.][]{ettori2004}, becomes, then, X-ray emitter mainly through bremsstrahlung processes, thus allowing one to trace and characterize the distribution of the baryonic and dark matter components.  From an observational point of view, the measurement of the properties of the ICM is possible only where the X-ray emission can be well resolved against the background (both instrumental and cosmic). At the present, data permit to firmly characterize observables like surface brightness \\citep[see][]{mohr1999, ettori1999} and gas temperature \\citep[see][]{markevitch1998, degrandi2002, pointecouteau2005, vikhlinin2006, zhang2006} out to a fraction ($\\sim 0.5-0.6$) of the virial radius. Only few examples of nearby X-ray bright clusters with surface brightness estimated out to the virial radius are available, thanks to the good spatial resolution, the large field-of-view and the low instrumental background of the {\\it ROSAT/PSPC} instrument \\citep{vikhlinin1999, neumann2005}. However, since the regions not-yet resolved occupy almost 80 per cent of the total volume, they retain most of the information on the processes that characterize the accretion and evolution within the cluster of the main baryonic component \\citep{molendi2004}.  \\begin{table*} \\begin{center} \\caption{ List of main physical parameters of the 9 clusters taken from the \\ovisc\\ runs. The objects are divided into 2 subsamples according to their virial mass. Note that $T_{\\rm vir}$ is the average mass-weighted temperature of the gas up to $R_{\\rm vir}$ while $T_{200}$ is the temperature of the gas at $R_{200}$.} \\begin{tabular}{lccccccccccc} & \\multicolumn{4}{c}{Sample $A$: $M_{\\rm vir}>10^{15} M_\\odot$} &  &  & \\multicolumn{5}{c}{Sample $B$: $M_{\\rm vir}<10^{15} M_\\odot$} \\\\ Cluster name & {\\it g1} & {\\it g8} & {\\it g51} & {\\it g72} &  &  & {\\it g676} & {\\it g914} & {\\it g1542} & {\\it g3344} & {\\it g6212} \\\\ \\hline $M_{\\rm vir}$ ($10^{15} M_\\odot$) & 2.12 & 3.39 & 1.90 & 1.96 &  & & 0.15 & 0.17 & 0.15 & 0.16 & 0.16  \\\\  $R_{\\rm vir}$ (Mpc)               & 3.37 & 3.94 & 3.25 & 3.28 &  & & 1.40 & 1.46 & 1.40 & 1.43 & 1.43  \\\\  $T_{\\rm vir}$ (keV)               & 6.49 & 9.89 & 5.78 & 5.44 &  & & 1.25 & 1.28 & 1.18 & 1.26 & 1.22  \\\\  \\hline $R_{200}$ (Mpc)                   & 2.48 & 2.91 & 2.39 & 2.40 &  & & 1.04 & 1.08 & 1.04 & 1.05 & 1.05  \\\\  $T_{200}$ (keV)                & 4.18 & 7.52 & 4.18 & 4.09 &  & & 0.82 & 1.03 & 0.83 & 0.88 & 0.84  \\\\ \\hline \\label{tab:9clusters} \\end{tabular} \\end{center} \\end{table*}  \\begin{figure*} \\includegraphics[width=0.78\\textwidth]{mass_100_095.ps} \\includegraphics[width=0.78\\textwidth]{temp_100_095.ps} \\includegraphics[width=0.78\\textwidth]{xray_100_095.ps} \\caption{ From top to bottom: maps of the projected gas density, mass-weighted temperature and soft (0.5-2 keV) X-ray emission of the {\\it g1} cluster (\\csf\\ model) for the 100 per cent (left) and 95 per cent (right) volume-selection scheme. The circles indicate the virial radius. The size of the side of each map is 12 Mpc, so they cover roughly up to $2.5 R_{200}$. } \\label{fig:maps} \\end{figure*}  In the present work, we study the distribution of the ICM density, temperature and surface brightness in a set of simulated galaxy clusters by looking at their radial profiles in the range between $0.3 - 3 \\times R_{200}$\\footnote{ In this paper we define the typical physical quantities of clusters as follows: $R_{200}$ ($R_{500}$) is the radius of a sphere centered in a local minimum of the potential and enclosing an average density of $\\rho = 200 (500) \\rho_{\\rm cr}$, with $\\rho_{\\rm cr}$ being the critical density of the Universe. The virial quantities are defined using the density threshold obtained by the spherical collapse model with the cosmological parameters of the simulations. So the virial radius $R_{\\rm vir}$ encloses an average density of $\\rho_{\\rm vir} \\simeq 104 \\rho_{\\rm cr}$ \\citep[see][eq.~4]{eke1998}. Consequently the virial mass $M_{\\rm vir}$ is the total mass (DM plus baryons) of the cluster up  to $R_{\\rm vir}$.}. We avoid considering the regions where the clusters evolve a core with active cooling and suspected feedback regulation, being this balance still matter of tension between the present simulated constraints and the observed counterparts \\citep[e.g.][]{borgani2004}. Therefore, we concentrate our analysis on cluster volumes where the gas density and temperature suggest that radiative processes are not dominant and the main physical driver is just the gravitational collapse, that is expected to be well described in the present numerical N-body simulations. We consider objects extracted from a dark matter only cosmological simulation that were resimulated at much higher resolution by including the gas subjected to gravitational heating and other physical treatments like cooling, star formation, feedback, thermal conduction and an alternative implementation of artificial viscosity.  We organize this paper in the following way: in the next Section, we describe our dataset of simulated clusters and their different physical properties; in Section \\ref{sect:results} we discuss the results on the slopes of the gas temperature, density and surface brightness profiles in the outskirts (i.e. $r > 0.3 R_{200}$) of the simulated objects and compare our results with the present constraints obtained from X-ray observations of bright clusters. We summarize and discuss our results in Section \\ref{sect:concl}. ",
        "conclusions": "\\label{sect:concl}  Using a set of hydrodynamical simulations, performed with the Tree+SPH code {\\tt GADGET-2}, composed by 9 galaxy clusters covering the mass range $1.5 \\times 10^{14} M_{\\odot} < M_{\\rm vir} < 3.4 \\times 10^{15} M_{\\odot}$ and adopting 4 different physical prescriptions with fixed metallicity, we have studied the behaviour in the outer regions of the gas density, temperature, soft (0.5--2 keV) and hard (2--10 keV) X-ray surface brightness profiles paying particular attention to the logarithmic slope. Outside the central regions, where a not-well defined heating source can partially or totally balance the radiative processes, the physics of the X-ray emitting intracluster plasma is expected to be mainly driven by adiabatic compression and shocks that take place during the collapse of the cosmic baryons into accreting dark matter halos. These processes are properly considered in cosmological numerical simulations here analyzed, allowing us to investigate on solid basis the expected behaviour of the cluster outskirts.  Our main findings can be summarized as follows\\footnote{ Note that, as in the rest of the paper, the values of the slopes here quoted are in fact negative slopes.}:  \\begin{enumerate}  \\item the behaviour of the profiles in the external regions of clusters ($r \\gtrsim R_{200}$) does not depend significantly on the presence or absence of cooling and SN feedback. Only the inclusion of a model of thermal conduction can introduce non-negligible changes in the values of the slopes of the profiles. This shows that the mean behaviour of the X-ray emitting plasma in the cluster outskirts is mainly due to the gravitational force.  \\item The gas density profile steepens in the outskirts changing the slope of the power-law from $\\sim 2.5$ to $\\sim 3.4$ at about $1.2 R_{200}$. The temperature profile varies the slope of the power-law like functional form from $ 0.4$ in non-radiative models and $ 0.6$ in models with cooling and feedback, to about $ (1.7-1.9)$ at $\\sim1.3-1.6 R_{200}$.  \\item The surface brightness in the outskirts profiles appear to be dominated by the presence of subclumps of cold and dense material that we exclude from our analyses. After removing these objects, the surface brightness in the [0.5--2] keV presents a profile with a slope that steepens towards the external regions. Fitting the profiles with a power-law function in different radial ranges between 0.3 and 2.7 $R_{200}$ we obtain slopes between 4 and 5.5 in the 4 most massive systems and between $ 3.8$ and $ 6.5$ in the 5 objects with masses $\\approx 10^{14} M_{\\odot}$.  \\item We introduce a new rolling-index power-law functional (eq. \\ref{eq:roll_pow_law}) which fits well the X-ray surface profiles in the interval $0.3 \\leq r/R_{200} \\leq 2.7$. From the results of these fits, we obtain that our sample of groups of galaxies (sample $B$) has soft X-ray profiles shallower in the centre ($b_{\\rm min, B} \\sim 2.6$, $b_{\\rm min, A} \\sim  3.3$) and steeper in the outskirts ($b_{\\rm max, B} \\sim 7.4$, $b_{\\rm max, A} \\sim 5.8$) than our sample of clusters (sample $A$). This is mainly an effect of the drop out from the pass band of the soft X-ray emission, due to the low ICM temperature of the objects in sample $B$ at the virial radius.  \\item We compare our results on the shape of the temperature and soft X-ray surface brightness profiles with the present observational constraints. By using the functional form of \\cite{vikhlinin2006} (eq. 9) extrapolated beyond $R_{500}$, we verify that there is a good agreement with the simulated profiles over the radial range $0.3 - 2 R_{200}$. On the surface brightness profile in the [0.5--2] keV band, we measure a slope in the outskirts that is consistent within $1 \\sigma$ with the values estimated from {\\it ROSAT/PSPC} data in \\cite{neumann2005}.  \\item From the overall shape of the simulated profiles, we predict that (see Table~\\ref{tab:prof_values}) (1) the ICM temperature profile decreases with radius reaching at $R_{200}$ average values of $0.60$ times the value measured at $0.3 \\times R_{200}$, (2) the estimated gas density at $R_{200}$ is $4.8 \\times 10^{-3}$ times the gas density measured at $0.3 R_{200}$, (3) the surface brightness in the soft band is about $10^{-12}$ erg s$^{-1}$ cm$^{-2}$ deg$^{-2}$ at $R_{200}$ at $z$=0, that is a factor of few below the extragalactic unresolved X-ray background \\citep[see][]{hickox2006}.  \\end{enumerate}  This last item shows that the perspective to resolve the X-ray emission around the cluster virial regions is extremely challenging from an observational point of view, even at low redshift, unless tight constraints on all the components of the measured background can be reached \\citep[see also][]{molendi2004}. Considering the similarity of the surface brightness profiles of the objects of the same mass, the observational technique of stacking the images of different clusters \\citep[e.g.][]{neumann2005} can help in enhancing the statistics of the profiles in the outer regions.  Overall, our results show how the study of simulated objects can give indications on the requested capabilities of future experiments in order to resolve the cluster virial regions. On the other hand, since the behaviour of the slopes have shown to be mainly independent of the physical models adopted in our simulations, any discrepancy between our results and observations of the clusters virial regions could provide useful hints on the presence of additional physics (e.g. turbulence, magnetic fields, cosmic rays) influencing the thermodynamics of the ICM.  Finally, the work presented here further demonstrates the complementarity between numerical and observational cosmology in describing the properties of highly non-linear and over dense structures."
    },
    "0609/astro-ph0609223_arXiv.txt": {
        "abstract": "We investigate the nature of the soft excess below 1~keV observed in AGN. We use the {\\it XMM-Newton} data of the low redshift, optically bright quasar, PG~1211+143, and we compare it with the Narrow Line Seyfert 1 galaxy, 1H~0707-495, which has one of the strongest soft excesses seen. We test various ideas for the origin of the soft X-ray excess, including a separate spectral component (for example, low temperature Comptonized emission), a reflection-dominated model, or a complex absorption model. All three can give good fits to the data, and $\\chi^2$ fitting criteria are not sufficient to discriminate among them. Instead, we favour the complex absorption model on the grounds that it requires less extreme parameters. In particular the geometry appears to be more physically plausible as the reflected component in the smeared absorption model is no longer dominant, and relativistic distortions, while still clearly present, are not tremendously larger than expected for a disc around a Schwarzchild black hole. ",
        "introduction": "The X-ray spectra of Active Galactic Nuclei (AGN) and Quasars have long been known to often contain a 'soft excess' component at energies below $\\sim$~1~keV (e.g. Turner \\& Pounds 1988). This rises rather smoothly above the extrapolation of the power law continuum seen in the 2--10~keV band, and seems to connect onto the peak of the UV accretion disc spectrum (Zheng et al. 2001; Czerny et al. 2003). This shape is suggestive of low temperature, high optical depth Comptonization of the inner accretion disc, in addition to the high temperature, low optical depth Comptonization required to produce the high energy power law (e.g. Page et al. 2004a). However, the temperature of this additional component seems remarkably constant in high mass accretion rate objects ($L/L_{Edd}>0.1$), despite a large range in black hole mass and hence inferred accretion disc temperature (Czerny et al. 2003; Gierli\\'nski \\& Done 2004, hereafter GD04; Crummy et al. 2006). The nature of this soft excess becomes even more puzzling when compared to the spectra of Galactic black hole (GBH) binary systems. These can indeed show a low temperature, high optical depth Comptonized component in addition to the disc and high energy power law spectrum (in the very high state).  However, while it is initially tempting to identify this as the GBH counterpart of the soft excess seen in AGN (Page et al. 2004b; Murashima et al. 2005), the parameters required are rather different, most obviously in the fact that the temperature of this component is {\\em variable} in the GBH (Kubota, Makishima \\& Ebisawa 2001; Kubota \\& Done 2004).  AGN with the largest soft excesses, generally Narrow Line Seyfert 1 galaxies (NLS1; Boroson 2002), also often show a strong deficit at $\\sim 7$~keV which again has no obvious identification (Boller et al. 2002; Fabian et al. 2002, 2004, 2005; Tanaka, Ueda \\& Boller 2005). These objects are often highly luminous, close to the Eddington limit, yet their underlying high energy 2--10~keV spectrum can appear to be rather hard (e.g. GD04; Crummy et al. 2006). All Galactic black holes at $L/L_{Edd}\\sim 0.5$--1 have intrinsically soft power law Comptonized emission (e.g. Remillard \\& McClintock 2006), though occasional super Eddington source spectra can be harder (e.g. V404 Cyg: \\.Zycki, Done \\& Smith 1999).  The constancy of temperature for the soft excess emission could plausibly point to an origin in atomic processes. In particular, there is an abrupt increase in opacity in partially ionized material between $\\sim 0.7$--3~keV due to OVII/OVIII and Fe L transitions which can produce a large increase in reflected or transmitted flux below 0.7~keV, making the apparent soft X-ray excess. The corresponding rise in reflected or transmitted flux above 2--3~keV could also artificially harden the high energy spectrum, while the same material might also have strong iron features, giving a simultaneous explanation of the 7~keV deficit.  By contrast, models where the soft excess is a separate continuum component require some other explanation for the origin of the iron features and hard spectra, such as partial covering of the source by cold absorbing material (Boller et al. 2002; Tanaka et al. 2004). However, the neutral iron edge at 7.1~keV in these models is not big enough to fit the strongest observed features unless iron is very overabundant, at least 5--7 times the Solar value (Tanaka et al. 2004; Gallo et al. 2004; Gallo 2006). More generally, such models seem contrived as the iron features and hard spectra appear to be correlated with the soft excess, pointing to a common origin for the hard and soft spectral complexity (Boller et al. 2003; Reeves, Porquet \\& Turner 2004; Tanaka et al. 2005). This connection is strongly supported by spectral variability. The {\\it rms} spectra from these objects typically show more variability in the 0.7--3~keV range than at higher or lower energies (e.g. Gallo et al. 2004; Ponti et al. 2004, 2006). The spectrum below 0.7~keV, which is dominated by the soft excess, typically has the same variability as the spectrum above 3~keV, showing that they are almost certainly linked to the same component. The amplified variability in the 0.7--3~keV region is compelling evidence for an atomic origin for the spectral complexity, as it is exactly over this energy range that the atomic features dominate (Gierli\\'nski \\& Done 2006), but both reflection and absorption can match the variability spectra seen (Ponti et al. 2006; Gierli\\'nski \\& Done 2006).  One obvious problem with the atomic models is that they predict strong, sharp atomic features from the partially ionized material. Such characteristic narrow absorption lines and sharp edges are seen in AGN (warm absorbers) but these can be easily isolated using high spectral resolution data (see e.g. Blustin et al. 2005 and references therein), and a strong {\\em smooth} soft excess remains. Large velocity smearing is the only way to keep an atomic origin of the soft excess. If at least some of the material is moving at moderately relativistic speeds, the characteristic line/edge features become strongly broadened into a quasi--continuum. Such speeds naturally arise only close to the black hole, so the atomic models predict that the soft excess is formed in regions of strong gravity, irrespective of whether it arises from partially ionized reflection or absorption. However, it is important to distinguish between a reflection or absorption origin for the soft excess, as they give very different geometries and have different implications for our understanding of the X-ray source. In reflection, the partially ionized material is optically thick and out of our line of sight. The accretion disc is the obvious identification but in order to produce the strongest soft excesses a reflection model requires that the intrinsic continuum is strongly suppressed, perhaps by the disc fragmenting or by extreme lightbending (Fabian et al. 2002, 2004, 2005; Miniutti \\& Fabian 2004). The velocity structure should be Keplarian, and the large inferred smearing implies a very small inner disc radius, and hence extreme Kerr space-time and/or extremely centrally concentrated emissivity (e.g. Miniutti \\& Fabian 2004).  By contrast, in an absorption origin, the partially ionized material is optically thin, and seen in our line of sight. This implies that it is some sort of wind from the disc, so its velocity structure is not well defined, and cannot easily give constraints on the space-time (GD04; Chevallier et al. 2006; Schurch \\& Done 2006).   \\begin{figure*} \\centering \\includegraphics[width=8.8cm,bb=32 144 545 444,clip]{reflion.abundance.ps} \\includegraphics[width=8.8cm,bb=32 144 545 444,clip]{reflion.amplitude.ps} \\caption{Strength of the soft excess in a reflection model (with a photon index of underlying continuum of $\\Gamma = 2.0$ and $2.4$) as a function of ionisation parameter and (a) iron abundance for $\\Omega/2\\pi = 1$, (b) amplitude of reflection for Solar abundance. The curves correspond to Fe abundance (or $\\Omega/2\\pi$) of 0.1 (dotted), 0.5 (short-dashed), 1 (solid), 3 (long-dashed), and 7 (dot-dashed).} \\label{fig:sereflion} \\end{figure*}   The structure of the paper is as follows. In Sec.~2 we try to distinguish between the reflection and absorption models for the soft excess. We calculate the strength of the soft excess in both models, showing that the largest observed soft excesses can be produced in the absorption model by simply changing the column of material in the line of sight.  In Sec.~3 we explicitly compare reflection and absorption models by fitting data from PG~1211+143 and 1H~0707-495, representative of the largest soft excesses seen in Quasars and NLS1's, respectively. Both models give comparably good fits to the data in terms of $\\chi^2$ criteria, as does the more phenomenological model of a separate soft component. Statistics cannot discriminate between them, and depending on $\\chi^2$ alone may be somewhat misleading as there are systematic uncertainties in the models (such as range of ionisation states present) which probably dominate the residuals.  How then can we distinguish between them? The {\\it XMM--Newton} data is fairly high signal--to--noise, moderate spectral resolution and fairly broad bandpass (0.3--10~keV), yet neither energy nor variability spectra can clearly discriminate between these models. In Sec.~4 we suggest that simultaneous data at higher energies may be able to break these spectral model degeneracies. We also review the physical plausibility of the models, which leads us to favour the smeared absorption as the origin of the soft excess. ",
        "conclusions": "We tried to determine the origin of the soft excess seen in AGN by fitting 0.3--10~keV {\\it XMM-Newton} spectra of two representative objects with all currently proposed models, namely a separate continuum, reflection and absorption. All these give comparably good fits to the overall continuum given the model uncertainties (such as range of ionisation states). Here we review the physical plausibility of each model to help distinguish between them.  A separate soft excess component involves an emission from an unknown physical process, with unknown 'fine-tuning' mechanism fixing its typical energy. Moreover, this model cannot simultaneously explain the 7~keV feature, nor is it easy to imagine how this can produce the spectral dependence of the variability, where the {\\it rms} spectra typically peak in 0.7--3~keV band (though it is possible to arrange: Page et al. 2004b). By contrast, both reflection and absorption scenarios are much more plausible, as they give a physical reason for the fixed energy of the soft excess, can reproduce the structure around iron K and give apparently hard spectra in the 2--10~keV band, while the intrinsic spectrum remains soft.  The reflection models can successfully fit the spectra, but the strongest soft excesses require that the intrinsic continuum is suppressed, while the large velocity shear to smear out the atomic features implies extreme spin and/or emissivity (e.g. Crummy et al. 2006).  The alternative absorption model can reproduce the range of observed soft excess strengths by simply changing the column density ($\\sim 10^{22-23}$~cm$^{-2}$), without any requirement for the intrinsic source to be suppressed. Thus, only a quantitative change in column is required, rather than a qualitative change in geometry, as in the case of reflection.  Similarly to reflection, a large velocity shear is required in absorption model, but since this is now identified with the wind rather than the disc, it does not directly constrain the space-time or energy extraction process. In pure absorption form the model does not well fit the most extreme sharp iron features seen at ${\\sim}7$~keV.  However, emission and scattering as well as absorption in a wind can produce a P Cygni profile in He-- and H--like iron K$\\alpha$ resonance lines. Including this gives an excellent fit to the feature at 7~keV.  Both reflection and absorption require that the ionisation state is 'fine-tuned' such that Oxygen is partially ionized. There is no readily apparent reason for this in the reflection models. In fact, the ionisation instability may actually give a physical argument against such partially ionized species dominating reflected spectra, making it difficult for a disc in hydrostatic equilibrium to produce any soft excess in reflection (Done \\& Nayakshin, in preparation). By contrast, this same ionisation instability may rather naturally produce the required range of ionisation parameter for the absorption model (Chevallier et al. 2006). In addition, on physical grounds we expect strong winds from high $L/L_{\\rm Edd}$ accretion discs, especially in AGN featuring a disc flux that peaks in the UV region, where there can be strong line driving (e.g. Proga, Stone \\& Kallman 2000). The resulting, messy environment appears rather more physically plausible than a clean reflecting disc at high $L/L_{\\rm Edd}$,  and there is growing evidence for (less extreme) relativistic outflows in the UV spectra of NLS1's (Leighly 2004, Green 2006). Similar material may also be present in GBH binaries at high $L/L_{Edd}$. This might contribute to spectral complexity at energies around iron line, but it is unlikely to produce a soft excess at similar temperatures as in AGN as Oxygen is probably always completely ionized due to the much higher accretion disc temperature.  Thus, an absorption origin seems favoured in terms of physical plausibility, though we note that reflection from the disc should also contribute to the spectrum at some level (Chevallier et al. 2006). Higher energy observations soon to be available from {\\em Suzaku} may provide a more sensitive test for objects with the soft excess/iron feature not so large as to require an optically thick wind."
    },
    "0609/astro-ph0609709_arXiv.txt": {
        "abstract": "Recent observations with the Spitzer Space Telescope show clear evidence that star formation takes place in the surrounding of young massive O-type stars, which are shaping their environment due to their powerful radiation and stellar winds. In this work we investigate the effect of ionising radiation of massive stars on the ambient interstellar medium (ISM): In particular we want to examine whether the UV-radiation of O-type stars can lead to the observed pillar-like structures and can trigger star formation. We developed a new implementation, based on a parallel Smooth Particle Hydrodynamics code (called IVINE), that allows an efficient treatment of the effect of ionising radiation from massive stars on their turbulent gaseous environment. Here we present first results at very high resolution. We show that ionising radiation can trigger the collapse of an otherwise stable molecular cloud. The arising structures resemble observed structures (e.g. the pillars of creation in the Eagle Nebula (M16) or the Horsehead Nebula B33). Including the effect of gravitation we find small regions that can be identified as formation places of individual stars. We conclude that ionising radiation from massive stars alone can trigger substantial star formation in molecular clouds. ",
        "introduction": "In the surroundings of hot OB-Associations filamentary substructures on different scales are observed (see e.g. \\cite{2002ApJ...565L..25S} and references therein). As observational resolution increases, more and more sub-millimeter sources, which could trace the birth of future stars are detected (e.g. \\cite{2006MNRAS.369.1201W}). It has long been suggested that radiation driven implosion of molecular clouds can explain the morphology and the star formation in these regions (e.g. \\cite{1995ApJ...451..675E}).  Recent simulations (see e.g. \\cite{2006ApJ...647..397M}, \\cite{2005MNRAS.358..291D}, \\cite{2003MNRAS.338..545K}) demonstrate the importance of massive stars for the subsequent evolution of their parental molecular clouds. The ionising radiation is a vital ingredient to understand the disruption of molecular clouds and their star formation efficiency.  Our goal is to investigate the morphology of molecular clouds and the formation of protostars in much greater detail. To do so we use very high resolution simulations of a small region of a molecular cloud ionised by a massive nearby star. ",
        "conclusions": "We developed a fully parallel treatment for the ionising radiation of young massive stars. Ionising Radiation alone is sufficient to explain the morphology observed in the surroundings of hot OB-clusters. In our simulations of turbulent ISM exposed to UV-radiation characteristic trunks similar to the ones observed in M16 evolve. Further studies including gravity will show whether the UV-radiation from young massive stars is sufficient to trigger gravitational collapse within these pillars. \\begin{figure} \\centering \\includegraphics[width=12cm]{./init_box_bw.eps} \\caption{Initial density distribution after the turbulence has decayed to Mach 5. The box dimension are $2pc$ in each direction.}\\label{fig_init} \\end{figure}  \\vspace{.5cm}  \\begin{figure} \\centering \\includegraphics[width=12cm]{./end_box_bw.eps} \\caption{Final stage of evolution after $t \\approx 300kyrs$: the morphology has clearly evolved. The size of the most prominent pillar is roughly $1pc$ as it is observed in e.g. M16 }\\label{fig_end} \\end{figure}"
    },
    "0609/astro-ph0609545_arXiv.txt": {
        "abstract": "Using results from high-resolution galaxy formation simulations in a standard $\\Lambda$CDM cosmology and a fully conservative multi-resolution radiative transfer code around point sources, we compute the energy-dependent escape fraction $\\fesc$ of ionizing photons from a large number of star forming regions in two galaxies at five different redshifts from $z=3.8$ to 2.39. All escape fractions show a monotonic decline with time, from (at the Lyman-limit) $\\sim6-10\\%$ at $z=3.6$ to $\\sim1-2\\%$ at $z=2.39$, due to higher gas clumping at lower redshifts. It appears that increased feedback can lead to higher $\\fesc$ at $z\\gsim3.4$ via evacuation of gas from the vicinity of star forming regions and to lower $\\fesc$ at $z\\lsim2.39$ through accumulation of swept-up shells in denser environments. Our results agree well with the observational findings of \\citet{inoue..06} on redshift evolution of $\\fesc$ in the redshift interval $z=2-3.6$. ",
        "introduction": "In the redshift interval $2<z<6$ most ionizing photons in the Universe are thought to originate in PopII stars in normal ($V_c\\ga30-50\\kms$) galaxies \\citep{nagamine......06}. For reionization calculations at $z>6$ photons from lower mass halos should be taken into account \\citep{iliev...06}, whereas at $z\\sim 2$ active galactic nuclei start to play a dominant role \\citep{madau..99}. The fact that the metagalactic UV field peaks at $2<z<4$ \\citep[and ref. therein]{inoue..06} points to a peak in the galactic star formation (SF) rate at such redshifts \\citep[see also][]{panter...06}. A major uncertainty in modeling the effect of this SF on the thermal state of the intergalactic medium (IGM) is the value of the escape fraction $\\fesc$ of ionizing UV photons, which is a function of redshift, and galaxy type and mass.  Best observational estimates of $\\fesc$ come from the detailed multiwaveband studies of the local Universe. \\citet{leitherer...95} find that less than $3\\%$ of ionizing photons from low-redshift starburst galaxies escape into the IGM. More recently, for several local starburst galaxies, \\citet{heckman.....01} estimate $\\fesc\\lsim 6\\%$, noting that inclusion of dust will further reduce $\\fesc$, and \\citet{bergvall..06} find $\\fesc\\sim 4-10\\%$ for another local starburst.  The situation is markedly different at $z\\gsim3$, where the Lyman continuum (LyC) detection points at much higher $\\fesc$ \\citep{steidel..01}. By comparing the direct observations of the LyC from galaxies to the mean cosmic UVB intensity, \\citet{inoue..06} conclude that the average $\\fesc$ increases from $1-2\\%$ at $z=2$ to $\\sim10\\%$ at $z\\gsim3.6$.  Most theoretical estimates of $\\fesc$ to date have been based on models with a smooth distribution of gas \\citep{ricotti.00} or semi-analytical models with expanding shells and superbubbles in a disk galaxy \\citep{dove..00,fujita...03}. In this letter we present ab initio calculations of the energy-dependent $\\fesc$ of ionizing radiation from SF regions in proto-galaxies found in numerical simulations at five different redshifts from $z=2.39$ to 3.8. We define $\\fesc(\\nu,r)$ simply as a fraction of photons of energy $h\\nu$ which reach a shell of radius $r$ around a given source; by definition $\\fesc(\\nu,0)=1$. We postprocess high-resolution simulation datasets with radiative transfer (RT) on nested grids, with the maximum grid resolution of $15\\pc$. Each galaxy contains from several hundred to several thousand distinct stellar sources representing individual, young SF regions. Our interstellar gas distribution is extremely clumpy due to an array of physical processes including feedback from SF. We also estimate the effect of the dynamical expansion of HII regions on the escape of ionizing photons. ",
        "conclusions": "Fig.~\\ref{escapeFraction} shows spectral dependence of $\\fesc$ at $r=100\\kpc$ in each model, averaged over all sources. All $\\fesc$ have been corrected for the boundary of the volume. For individual sources the escape of photoionizing radiation is similar to a phase transition: it is very sensitive to conditions at the source and therefore can vary by a sizable factor for the same galaxy over the course of its evolution. We find that while there are variations among the models, the escape fractions tend to decrease with lower redshift, reflecting the fact that with time more gas cools into the higher density clouds, as well the declining SFR at $z<2.95$. With normal feedback in both galaxies $\\fesc$ at the Lyman edge drops from $8-10\\%$ at $z=3.8$ to $\\sim1-2\\%$ at $z=2.39$.   \\begin{figure} \\epsscale{1.}\\plotone{f3.eps} \\caption{Spectral dependence of $\\fesc$ at $r=100\\kpc$ for all five input galaxy models at (where transfer has been computed) $z=3.8$ (red), $z=3.6$ (magenta), $z=3.4$ (blue), $z=2.95$ (cyan), and $z=2.39$ (green). For each model the redshift evolution is from left to right, and each curve goes from $13.6\\ev$ to $135\\ev$. All solid lines are models without the HII region expansion, with the SFR distributed uniformly among all stars younger than $34\\Myrs$. For K15-06-8 the dashed lines show models with the HII region expansion. For K15-1-8 we also computed the models with the {\\it same} SFR distributed among stars younger $3.4\\Myrs$ (dotted) and $10\\Myrs$ (dash-dotted lines).} \\label{escapeFraction} \\end{figure}  Not surprisingly, in all our models $\\fesc$ of photons capable of doubly ionizing helium is very low, due to the softness of stellar radiation. The escape fraction of photons capable of single He ionization is comparable at all redshifts to that of H-ionizing photons.   Increasing SN feedback in K15 evacuates more gas from the vicinity of the SF regions leading to a threefold rise in $\\fesc$ at $z=3.6$. Lower local gas densities translate into a $40\\%$ reduction in the SFR (Table 1). At $z=2.95$, a denser environment near the sources produces a much smaller difference in the SFR and $\\fesc$, when we increase the strength of feedback, and, at $z=2.39$, $\\fesc$ actually drops by $25\\%$ as a larger fraction of gas swept-up by the shocks stays in the vicinity of the SF regions.   Next we compare our normal and high resolution results, at $z=3.8-3.4$, and find a good agreement between the two. We see it as an additional test of our algorithm, since other properties of the two sets of simulations (SFR, etc.) also match.  Since our choice of distributing SF over the $t_{\\rm up}=34\\Myrs$ youngest stars is somewhat arbitrary, in galaxy K15-1-8 we also experimented with the same amount of SF distributed over all stars younger than $t_{\\rm up}=10\\Myrs$ and $3.4\\Myrs$. Ideally, these young stars would produce most of the ionizing UV photons. Since we do not calculate transfer and hydro simultaneously, we feel that applying continuous SF with the constant rate is a better computational approach than assuming instantaneous SF which will have no lasting effect on the thermal state of the IGM after few Myrs. With continuous SF we expect our results to converge as we increase the number of star particles. The results at $t_{\\rm up}=10\\Myrs$ and $34\\Myrs$ are in fact very close to each other at all three redshifts (Fig.~\\ref{escapeFraction}), whereas the $3.4\\Myrs$ results are somewhat off, likely related to the small number of sources used to represent the SF in this latter case.  The dashed lines in K15-06-8 in Fig.~\\ref{escapeFraction} show the energy-dependent $\\fesc$ at $r=100\\kpc$ corrected for the HII region expansion. At most redshifts expansion raises $\\fesc$ by $10-30\\%$, although at $z=2.95$ $\\fesc$ jumps by more than a factor of two. Further analysis shows that most of this change can be attributed to removal of the gas from the highest refinement level cells which are hosting the sources, not the surrounding cells.   In Fig.~\\ref{radial} we show how $\\fesc$ changes with distance from sources in galaxy K15-1-8 at $z=3.6$. A large fraction of photons reaches the radius $r=100\\pc$. As we do not have any reliable information on gas clumping on sub-resolution scales at these redshifts, the effective radius of absorption is not firmly established at this point. In our current models most absorption occurs within few hundreds pc of each star particle, well inside the virial radius of $\\sim$45 kpc. At $z=3.6$, the proto-galaxy consists of several star forming regions as well as a significant amount of the intergalactic HI, all of which account for some fractional absorption beyond $1\\kpc$.   \\begin{figure} \\epsscale{1.1}\\plotone{f4.eps} \\caption{Lyman-limit escape fraction distribution for all sources in galaxy K15-1-8 at $z=3.6$ (without HII region expansion), at four different radii, from 100 pc to 100 kpc.} \\label{radial} \\end{figure}  In this study we neglected the effect of dust extinction. \\citet{benson...06} point out that in the case of a smooth ISM attenuation by dust can reduce $\\fesc$ by a factor of few. For a clumpy ISM at high redshifts the magnitude of dust attenuation should be much lower, particularly that most ionizing photons are produced in star bursts, as opposed to a quiescent disk.  In conclusion, our models with exact ionizing continuum RT around a large number of SF regions in young protogalaxies show a monotonic decline in $\\fesc$ of ionizing photons from higher to lower redshifts. With the normal feedback strength $\\fesc$ at the Lyman edge drops from $\\sim6-10\\%$ at $z=3.6$ to $\\sim1-2\\%$ at $z=2.39$. At higher redshifts SF on average occurs at slightly lower densities resulting in easier escape of UV photons into the IGM. This result agrees well with the observational findings of \\citet{inoue..06} on redshift evolution of $\\fesc$. Note that although the two galaxies in our study have a mass ratio of $\\sim2.5$, their $\\fesc$ are very similar. One could suggest that perhaps, once star formation begins, it is the physical conditions in the clumpy gas clouds that determine the escape of UV photons, rather than the overall properties of their host galaxies. However, one would need a much larger statistical sample to test this speculation. In our next paper we will further examine the role of the fine structure of the clumpy ISM in the escape of ionizing photons, as well as the importance of hydrodynamical effects."
    },
    "0609/gr-qc0609028_arXiv.txt": {
        "abstract": "LISA is a planned space-based gravitational-wave (GW) detector that would be sensitive to waves from low-frequency sources, in the band of roughly $(0.03 - 0.1)\\,{\\rm mHz} \\lesssim f \\lesssim 0.1\\,{\\rm Hz}$. This is expected to be an extremely rich chunk of the GW spectrum --- observing these waves will provide a unique view of dynamical processes in astrophysics.  Here we give a quick survey of some key LISA sources and what GWs can uniquely teach us about these sources. Particularly noteworthy science which is highlighted here is the potential for LISA to track the moderate to high redshift evolution of black hole masses and spins through the measurement of GWs generated from massive black hole binaries (which in turn form by the merger of galaxies and protogalaxies).  Measurement of these binary black hole waves has the potential to determine the masses and spins of the constituent black holes with percent-level accuracy or better, providing a unique high-precision probe of an aspect of early structure growth.  This article is based on the ``Astrophysics Tutorial'' talk given by the author at the Sixth International LISA Symposium. ",
        "introduction": "Stochastic GWs are ``random'' waves, arising from the superposition of many discrete, uncorrelated sources.  They typically combine to produce a broadband spectrum that is nearly flat, or weakly peaked near some fiducial frequency.  Some of the most anticipated stochastic sources are cosmological in origin.  By far the most eagerly sought waves are GWs arising from primordial fluctuations in the universe's spacetime metric which have been parametrically amplified by inflation.  In energy density units (commonly used to parameterize these waves; see {\\cite{allen96}} for a very readable discussion), these waves are flat across an extremely wide band, ranging from frequencies $f \\sim 10^{-16}\\,{\\rm Hz}$ to roughly $10^{10}\\,{\\rm Hz}$.  The lower end of this scale corresponds to the inverse Hubble scale when the universe became matter dominated; below this frequency, the spectrum grows as $f^{-2}$.  At the high end, the spectrum is cut off by the (short but finite) time scale over which inflation ends and the universe enters a hot, radiation-dominated phase; see Fig.\\ 12 of {\\cite{allen96}} and associated discussion.  What's particularly exciting about these waves is that the level of this spectrum is set by --- and thus encodes --- the potential which drives inflation.  Thus, measuring these waves would {\\it directly probe inflationary physics}.  Unfortunately, these waves have such small amplitudes that direct detection with LISA (or any other near term GW experiment) is completely out of reach.  LISA should be able to detect a background at a level $\\sim 10^{-11}$ of the universe's closure density; current estimates suggest that this is at least 4 or 5 orders of magnitude from the sensitivity needed to measure these waves.  Direct detection at the relatively high frequencies of man-made facilities requires finding a relatively quiet band (without too many ``foreground'' sources) and a much more sensitive instrument than LISA.  At the lowest frequencies, detection is plausible in the next decade or so through cosmic microwave background (CMB) studies.  The tensor nature of GWs has a unique impact on the polarization of CMB photons {\\cite{zs97,kks97}}.  The signal is weak compared to other signals in the CMB, and potentially contaminated by foreground effects. Searching for it is nonetheless one of the most important directions in CMB physics today.  Although inflationary waves are out of LISA's reach, other cosmological backgrounds may not be.  Backgrounds are associated with phase transitions that occur as the universe expanded and cooled.  For example, when the mean temperature was roughly 1 TeV, the electroweak interaction separated into electromagnetic and weak interactions. This separation was not spatially homogeneous --- some regions underwent this transition before others.  Different regions were at different densities; boundaries between regions collided and coalesced, producing a GW background {\\cite{allen96,kmk02}}.  More recently, there has been a flurry of work in superstring theory suggesting that cosmic strings may have been produced in the early universe and then expanded to cosmic sizes (see, for example {\\cite{p05}} for an overview).  A network of such objects could produce a background accessible to LISA; see Hogan's contribution to these proceedings {\\cite{craig}} for detailed discussion.  Though a relatively speculative source, cosmological stochastic backgrounds are easily searched for, and there is a tremendous amount of discovery space.  The probability of payoff may be low, but the potential reward for discovery is extraordinarily high. ",
        "conclusions": "LISA is an astrophysics mission that will make possible unprecented {\\it precision} measurements of dynamic, strong-gravity phenomena. Perhaps most exciting will be the ability to track the mergers of massive black holes, precisely weighing their masses and spins.  Such measurements will directly probe the early growth of black holes, opening a new window onto structure growth.  \\begin{theacknowledgments} Support for this work is provided by NASA Grants NAG5-12906 and NNG05G1056; work on gravitational-wave science generally is supported by NSF Grants PHY-0244424 and PHY-0449884.  We also gratefully acknowledge support from MIT's Class of 1956 Career Development Fund. \\end{theacknowledgments}"
    },
    "0609/astro-ph0609259_arXiv.txt": {
        "abstract": "{Results of photometric surveys have brought to light the existence of a population of giant planets orbiting their host stars even closer than the hot Jupiters (HJ), with orbital periods below 3 days. The reason why radial velocity surveys were not able to detect these very-hot Jupiters (VHJ) is under discussion. A possible explanation is that these close-in planets are short-lived, being evaporated on short time-scales due to UV flux of their host stars. In this case, stars hosting transiting VHJ planets would be systematically younger than those in the radial velocity sample. We have used the UVES spectrograph (VLT-UT2 telescope) to obtain high resolution spectra of 5 faint stars hosting transiting planets, namely, OGLE-TR-10, 56, 111, 113 and TrES-1. Previously obtained CORALIE spectra of HD189733, and published data on the other transiting planet-hosts were also used. The immediate objective is to estimate ages via Li abundances, using the Ca~II activity-age relation, and from the analysis of the stellar rotational velocity. For the stars for which we have spectra, Li abundances were computed as in Israelian et al. (2004) using the stellar parameters derived in Santos et al. (2006). The chromospheric activity index $S_{US}$ was built as the ratio of the flux within the core of the Ca II H \\& K lines and the flux in two nearby continuum regions. The index $S_{US}$ was calibrated to Mount Wilson index $S_{MW}$ allowing the computation of the Ca II H \\& K corrected for the photospheric contribution. These values were then used to derive the ages by means of the Henry et al. (1996) activity-age relation. Bearing in mind the limitations of the ages derived by Li abundances, chromospheric activity, and stellar rotational velocities, none of the stars studied in this paper seem to be younger than 0.5 Gyr. ",
        "introduction": "Since the discovery of 51Peg b by \\cite{Mayor-1995} the number of new extra- solar planets  has been rapidly growing. More than about 160 extra-solar planets\\footnote{For a  continuously  updated list see table at http://obswww.unige.ch/Exoplanets} have been found and  yet  none of them look like our paradigm for planetary formation, i.e., the Solar  System.  On the contrary, these new discoveries provide evidence that a large variety of orbital  characteristics probably related to their formation process are actually  possible.  Due to the considerable number of giant planets found orbiting their host stars  at short  periods (typically below 10 days) it was soon realized that migration has to  play a key  role in their formation, since a solid core massive enough to accrete gas is not  able to grow at such a short distance from the parent star \\citep{Pollack-1996}.% With the increasing number of discoveries, statistical trends are arising,  revealing  interesting imprints thought to be left by the migration mechanism which brought  these  planets close-in \\citep{Udry-2003}. An annoying point when dealing with  migration is  how to stop it. The pile-up of planets with $P_{orb}\\sim3-3.5$ days in the the orbital period distribution was understood  as an observational clue indicating the migration parking point. Although appealing, this interpretation was severely challenged by the discovery  of the  transiting planets. Some of these new transiting planets have more than a  Jupiter mass  but nevertheless they revolve around their host stars in an even shorter orbital  period  than the so-called Hot Jupiters (HJ), typically $P_{orb}\\la1-2$ days. For this  reason,  they have been called Very Hot Jupiters (VHJ).  Why these planets have not been detected more often by the radial velocity  surveys is a  matter of debate. \\cite{Gaudi-2005} advocate that the non-detection of VHJ by  the radial  velocity surveys is a statistical effect due to the much lower frequency of VHJ  as compared  to that of HJ.  Given their proximity to their host stars, the VHJ may experience some degree of  evaporation  due to the heating by stellar UV photons. Vidal-Madjar et al. (2004) showed that HD209458b is indeed evaporating, but due to saturation of the spectral lines, only an upper limit on the mass loss of 10$^{10}\\rm\\,g\\,s^{-1}$ can be derived. According to theoretical calculations, the extreme case of an efficient heating of the planet atmosphere may lead to an energy-limited evaporation $~10^{12}\\rm\\,g\\,s^{-1}$ for HD209458b (Lammer et al. 2003), and hence to a very efficient evaporation, possibly to the total dissociation of some planets \\citep{Etangs-2004,Baraffe-2004}. Conversely, cooling in the atmosphere may well lead to a much more limited mass loss close to the lower limit derived by Vidal-Madjar et al. \\citep{Yelle-2004,Hubbard-2006}.  Is a catastrophic fate a possible explanation for the  lack of planets with periods shorter that 3 days? The existence of the VHJ is an important constraint to this type of scenario. Evaporation scenarios must be  able to explain the observed evaporation rate of HD 209458b, while at the same  time accounting for the long-term survival of the VHJ on even closer orbits. One possibility is that VHJ are not viable in the long term and are fated to  catastrophic evaporation. In this case the three observed VHJ would simply be  too young to have evaporated yet.  The aim of the present article is to examine the likelihood of the hypothesis that the three VHJ were found around parent stars that are very young compared to typical planet host stars and to the time scales of evaporation. For this purpose, three indicators of young stellar ages are combined: i) chromospheric activity  \\citep[e.g.][]{Henry-1996} ii) Lithium abundances and iii) rotational velocity.  The manuscript is organized as follows. Our observations and data reduction are described in Sec. 2. In Sec. 3 our index $S_{US}$ is built and calibrated with respect to the index $S$ of Mount Wilson $(S_{MW})$. In Sec. 4, we derive the Li abundances for the observed stars. Projected rotational velocities available in the literature for stars hosting transiting planets are compiled in Sec. 5. Using the different observables we derive ages in Sec. 6 and discuss the implications of the derived ages for the planetary evolution under evaporation in Sec. 7. ",
        "conclusions": "We have shown that the analysis of high resolution spectra of stars with planets could be used to determine lower limits on their age using Li abundances, the Ca~II activity-age relation and from the analysis of the stellar rotational velocity. Applied to six stars known to host transiting giant planets, OGLE-TR-10, 56, 111, 113, TrES-1, and HD189733, we were able to show that none appears to be younger than 1\\,Gyr at the 2-sigma level.  {\\bf In principle, this could be used to constrain the amount of mass loss experienced by the planets, since high evaporation rates should yield smaller average ages for the stars with very close-in giant planets. Because of the various uncertainties and small size of the sample studied here, the results remain inconclusive concerning evaporation. However, it highlights the importance of a careful determination of the ages of stars with planets, in particular for the stars harboring very close-in planets (``Pegasids''), but also, as a means of comparison, for stars with planets on orbits with longer-periods. }  We hence recommend to pursue the fine determination of the ages of stars with planets, in priority for the transiting planets as they are progressively being detected \\citep[for a direct application see][]{Guillot-2006}, for non-transiting very close-in giant planets (at less than 0.1 AU from their star), and then for a sample of stars with planets that are orbiting at larger distances."
    },
    "0609/astro-ph0609773_arXiv.txt": {
        "abstract": "None of N-body gravitating systems have been considered to emit periodic gravitational waves because of their chaotic orbits when $N=3$ (or more). We employ a figure-eight orbit as a specific model for a 3-body system in order to illustrate that some of triple stars are capable of generating periodic waves. This illustration would imply that a certain class of N-body gravitating systems may be relevant to the gravitational waves generation. We show also that the total angular momentum of this 3-body system is not carried away by gravitational waves. A waveform generated by this system is volcano-shaped and thus different from that of a binary system. Finally, by evaluating the radiation reaction time scale, we give an order-of-magnitude estimate of merging event rates. The estimate suggests that figure-eight sources, which require carefully prepared initial states, may be too rare to detect. ",
        "introduction": "Gravitational waves represent one of the great challenges of fundamental physics and will open a new window in the astronomy (e.g. Misner et al. 1973, Centrella 2003, Mio 2006a, 2006b). No one has ever detected directly the ripples in a curved space-time generated by accelerated masses. A (quasi-)periodic source is the most promising candidate for the first detection. Hence, (quasi-)periodic sources have been studied for the last few decades in numerous articles, where they assume a single star in rotation/oscillation (e.g. Kokkotas and Schmidt 1999, Stergioulas 2003) or double stars in binary motion (e.g. Asada and Futamase 1997, Blanchet 2006). In particular, event rates of an inspiraling and finally merging compact binary (Curran et al. 1995, van den Heuvel et al. 1996, Arzoumanian et al. 1999, Kalogera et al. 2001, Kim et al. 2003, Burgay et al. 2003) are discussed by taking account of large-scale detectors such as LIGO, VIRGO, GEO600, TAMA300 and LISA (Centrella 2003, Mio 2006a, 2006b). Furthermore, gravitational waves generated by a binary system have been intensively studied by perturbations such as the post-Newtonian approximation (e.g. Asada and Futamase 1997, Blanchet 2006, Faye et al. 2004, Berti et al. 2006), or by full numerical simulations for neutron stars (Shibata 1999) and for black holes (Pretorius 2005, Campanelli et al. 2006b, Baker et al. 2006, Diener et al. 2006). However, much less attention has been paid to N-body gravitating systems because of their chaotic orbits when $N\\geq 3$. There are existing works which are limited in the sense that they study a binary system affected by the existence of a third body, namely perturbations induced by the third object (Ioka et al. 1998, Wardell 2002, Campanelli et al. 2006a). Here, one may ask whether or not N-body systems can generate (quasi-)periodic gravitational waves. The purpose of this paper is to answer this question.  The main results of this paper are two; (1) For the first time as far as we know, we show that some of triple stars are capable of generating periodic waves if a certain condition is satisfied. This illustration would imply that a class of N-body gravitating systems may be relevant to the gravitational waves generation. (2) We show also that a solution employed as a specific model does not radiate the angular momentum. A well-known example in which no angular momentum is radiated away is two non-spinning black holes in a head-on collision (Smarr 1979, Price and Pullin 1994), where this system is axisymmetric and thus has a rotational Killing vector. Therefore, we can easily understand no angular momentum loss. On the other hand, the specific system employed in this paper has no Killing vector as seen below.  Our paper is organized as follows. In this paper, we restrict ourselves within 3-body systems for convenience. In section 2, we assume the Newtonian gravity for motion of a 3-body system. In section 3, the waveform and the loss rate of the energy and the angular momentum are estimated by using the so-called quadrupole formula for simplicity. Finally, we discuss the the time scale for radiation reaction in order to estimate merging event rates. ",
        "conclusions": "A figure-eight orbit for a 3-body system illustrated that some of triple stars are capable of generating periodic waves. We showed also that the total angular momentum of this 3-body system is not carried away by gravitational waves. A waveform generated by this system is volcano-shaped and thus different from that of a binary system. Our result will call our further attention to, for instance, % N=4 (or more) cases.  Before closing this paper, we mention the possibility of gravitational waves from 3-body systems in a `periodic' orbit such as a figure-eight. Clearly the possibility seems extremely low though it is quite uncertain. A quick order-of-magnitude estimate is made as follows. As a new outcome of binary-binary scattering, the figure-eight orbit was discussed for presenting a way of detecting such an orbit in numerical computations (Heggie 2000). According to the numerical result, the probability of the formation of figure-eight orbits is a tiny fraction of one percent. Here, we assume rather optimistically one per galaxy, which implies $10^6$ figure-eight systems within 100Mpc. If $\\ell$ distributes randomly with its mean of O(1AU), we use Eq. $(\\ref{tGW})$ to obtain $t_{GW} \\sim 10^8$ Gyr for $\\ell \\sim 1$AU. As a result, merging event rates are of the order of $10^{-2}$ event/Gyr within 100 Mpc, which is extremely low. Therefore, the above estimate suggests that figure-eight sources may be less relevant to LIGO and other detectors."
    },
    "0609/astro-ph0609403_arXiv.txt": {
        "abstract": "The High Resolution Fly's Eye (HiRes) experiment has observed the GZK cutoff.  HiRes' measurement of the flux of cosmic rays shows a sharp suppression at an energy of $6 \\times 10^{19}$ eV, exactly the expected cutoff energy.  We observe the ``ankle'' of the cosmic ray spectrum as well, at an energy of $4 \\times 10^{18}$ eV.  However we cannot claim to observe the third spectral feature of the ultrahigh energy cosmic ray regime, the ``second knee''.  We describe the experiment, data collection, analysis, and estimate the systematic uncertainties.  The results are presented and the calculation of a five standard deviation observation of the GZK cutoff is described. ",
        "introduction": "K. Greisen \\cite{g} and G. Zatsepin and V. Kuzmin \\cite{zk} predicted in 1966 that there would be an end to the cosmic ray spectrum.  Caused by the production of pions in collisions between cosmic ray protons and photons of the cosmic microwave background radiation (CMBR), this strong suppression is called the GZK cutoff and should occur at about $6 \\times 10^{19}$ eV.  This energy is set by the mass and width of the $\\Delta$ resonance and the temperature of the CMBR.  Several experiments \\cite{early,pravdin} whose exposures were too small to measure the flux of cosmic rays at the cutoff energy nevertheless observed one event each above $10^{20}$ eV, raising questions about the validity of Greissen's and Zatsepin and Kuzmin's prediction.  The Akeno Giant Air Shower Array (AGASA) was the first experiment large enough to measure the flux of cosmic rays at $10^{20}$ eV. AGASA consisted of 111 scintillation counters placed in an array of spacing about 1 km, and completed its data collection in 2004.  They found that the spectrum continued unabated beyond the GZK cutoff energy: their main evidence consisted of 11 events above $10^{20}$ eV \\cite{shinozaki}.  These results formed a serious challenge to the expectation that the cosmic ray flux should be cut off.  Since the prediction of the existence of the cutoff is based on accepted physics principles this was a truly startling result.  The High Resolution Fly's Eye (HiRes) experiment collected data for nine years (1997-2006), and now has a much higher exposure than AGASA near the cutoff energy.  HiRes is a fluorescence experiment; i.e., it observes cosmic ray showers by collecting the fluorescence light emitted by nitrogen molecules excited by shower particles.  By using this different technique HiRes has much better resolution and better control of systematic uncertainties than AGASA.  A fluorescence detector collects much more information about cosmic ray showers (for example it sees the shower development), measures the energy more accurately, and because it depends only on what goes on in the center of the shower one can simulate a fluorescence detector's performance very accurately.  This is an important point that has not been lost on new experiments being constructed: the Auger and Telescope Array experiments have both fluorescence detectors and ground arrays.  HiRes sees the GZK cutoff.  In this article we describe the measurement of the flux of cosmic rays in which this observation occurs, show the cosmic ray spectrum that we observe, and discuss of our estimates of systematic uncertainties. ",
        "conclusions": "We have measured the flux of ultrahigh energy cosmic rays by the fluorescence technique, over the energy range $10^{17.2}$ to $10^{20.5}$ eV.  We observe the GZK cutoff and the ankle of the cosmic ray spectrum.  We cannot claim to observe the second knee.  The statistical power of our observation of the GZK cutoff is just over $5\\sigma$.  The systematic uncertainty in our energy scale is about 17\\%, and the uncertainty in the flux is about 30\\%."
    },
    "0609/astro-ph0609039_arXiv.txt": {
        "abstract": "We present maps, source catalogue and number counts of the largest, most complete and unbiased extragalactic submillimetre survey:  the $850\\,\\mu\\mathrm{m}$ SCUBA HAlf Degree Extragalactic Survey (SHADES).  Using the Submillimetre Common-User Bolometer Array (SCUBA) on the James Clerk Maxwell Telescope (JCMT), SHADES mapped two separate regions of sky:  the Subaru/\\textit{XMM-Newton} Deep Field (SXDF) and the Lockman Hole East (LH).  Encompassing 93 per cent of the overall acquired data (i.e.~data taken up to 2004 February 1), these SCUBA maps cover $720\\,\\mathrm{arcmin}^{2}$ with an RMS noise level of about $2\\,\\mathrm{mJy}$ and have uncovered $>100$ submillimetre galaxies.  In order to ensure the utmost robustness of the resulting source catalogue, data reduction was independently carried out by four sub-groups within the SHADES team, providing an unprecedented degree of reliability with respect to other SCUBA catalogues available from the literature.  Individual source lists from the four groups were combined to produce a robust 120-object SHADES catalogue; an invaluable resource for follow-up campaigns aiming to study the properties of a complete and consistent sample of submillimetre galaxies.  For the first time, we present deboosted flux densities for each submillimetre galaxy found in a large survey.  Extensive simulations and tests were performed separately by each group in order to confirm the robustness of the source candidates and to evaluate the effects of false detections, completeness, and flux density boosting.  Corrections for these effects were then applied to the data to derive the submillimetre galaxy source counts.  SHADES has a high enough number of detected sources that meaningful differential counts can be estimated, unlike most submillimetre surveys which have to consider integral counts.  We present differential and integral source number counts and find that the differential counts are better fit with a broken power-law or a Schechter function than with a single power-law; the SHADES data alone significantly show that a break is required at several mJy, although the precise position of the break is not well constrained.  We also find that an $850\\,\\mathrm{\\mu m}$ survey complete down to $2\\,\\mathrm{mJy}$ would resolve 20--30 per cent of the Far-IR background into point sources. ",
        "introduction": "Deep blank-field surveys conducted with the Submillimetre Common-User Bolometer Array (SCUBA) on the James Clerk Maxwell Telescope (JCMT) have resolved as much as 50 per cent (depending on the survey depth) of the Far-Infrared Background (FIB) into discrete, high redshift-sources with flux density levels of $S_{850}\\gsim2\\,\\mathrm{mJy}$ (\\citealt{Smail97}, \\citealt{Hughes98}, \\citealt{Barger98}, \\citealt{Blain99}, \\citealt{Barger99}, \\citealt{Eales}, \\citealt{Cowie}, \\citealt{Scott}, \\citealt{Webb}, \\citealt{Borys2003}). Intensive campaigns to study these submillimetre galaxies (SMGs) and millimetre galaxies at other wavelengths have shown that they are primarily powered by star-formation, although many do harbour an Active Galactic Nucleus (AGN; \\citealt{alexander05}).  The implied star-formation rates are extremely high (100--$1000\\,\\mathrm{M}_{\\odot}\\,\\mathrm{yr}^{-1}$), and since the well-known local radio/Far-Infrared correlation \\citep{condon1992} apparently holds for these higher redshift sources (e.g~\\citealt{kovacs2006}), the high submillimetre luminosities result in correspondingly high radio luminosities, which are detected in deep 1.4\\,GHz images.  The resolution afforded by radio-interferometers results in precise optical identifications of $\\gsim50\\,\\mathrm{per cent}$ of all known SMGs, and in this manner \\cite{Chapman2005} derive a median redshift of $\\left\\langle z \\right\\rangle\\sim2.2$ for the population using deep Keck spectroscopy.  While much of what is known about SMGs is based on the radio detected subset \\citep{Ivison}, other studies (e.g.~\\citealt{Pope}) find no significant differences in the radio-undetected population.  Many investigations (e.g.~\\citealt{Smail97}; \\citealt{Barger}, \\citealt{Ivison}; \\citealt{Smail2002}; \\citealt{Webb}; \\citealt{Borys2004}; \\citealt{Wang}; \\citealt{Greve}) have suggested that SMGs are likely to be associated with an early-phase in the formation of massive galaxies. The intensity of their starbursts, the resulting high metallicity, along with their large dynamical masses, high gas fractions and inferred strong clustering (\\citealt{gravy}; \\citealt{2004ApJ...617...64S}; \\citealt{BCSI04}) are all suggestive of a close link to the formation phase of the most massive spheroids (e.g.~\\citealt{2004ApJ...616...71S}).  Despite this progress, the samples being used in any given study are typically quite small, since SCUBA could only detect roughly one SMG per night in good weather conditions.  And while the total number of SMGs detected over the lifetime of the instrument\\footnote{SCUBA was operational between 1997 and 2005} is now $\\sim400$, these are mainly drawn from small ($\\lsim100$ square arcminute) fields spread all over the sky, each observed and reduced by different groups using different techniques and source identification criteria.  The desire to obtain a well characterised sample of hundreds of SMGs in a large, contiguous area was the motivation for the SCUBA HAlf-Degree Extragalactic Survey (SHADES; \\citealt{Mortier}, \\citealt{vanKampen}).  SHADES is an ambitious wide extragalactic submillimetre survey, split evenly between 2 separate regions of sky:  the Subaru/\\textit{XMM-Newton} Deep Field (SXDF) and the Lockman Hole East (LH).  The aim was to map $0.5\\,\\mathrm{deg}^{2}$ ($\\simeq7$ times the area of its predecessor, the SCUBA 8-mJy Survey; \\citealt{Scott}) to a comparable depth of $\\sigma \\simeq 2\\,\\mathrm{mJy}$ at $850\\,\\mu\\mathrm{m}$, which is roughly three times the $1\\,\\sigma$ confusion limit imposed by the underlying sea of fainter unresolved sources in the coarse SCUBA beam (\\citealt{Hogg01}; \\citealt{Hughes98}; \\citealt{Cowie}).  The survey began in late 2002, and finished when SCUBA was decomissioned in late 2005.  SHADES is the largest survey in terms of observing time carried out with SCUBA.  Many scientifically powerful results from SHADES will come from comparisons of the properties of sources found in the SHADES catalogue with observations at other wavelengths.  Some of these include radio identifications (Ivison et al. in preparation), FIR-radio photometric redshift estimates (Aretxaga et al. in preparation), and submillimetre-Spitzer-based SEDs and photometric redshifts (Clements et al., Eales et al. and Serjeant et al. in preparation).  In this paper we present new constraints on the numbers of submillimetre sources as a function of $850\\,\\mathrm{\\mu m}$ flux density based on the SHADES catalogues.  A complementary analysis effort to constrain the numbers of faint submillimetre sources directly from our maps without the intermediate step of making a catalogue, a so-called $P(D)$ approach (e.g.~\\citealt{Condon}, \\citealt{Scheuer}), will be reported separately.  The main results presented in this paper are the SHADES catalogue and 850\\mum\\ number counts.  The observations are summarised in Section~\\ref{obsdr}.  Section~\\ref{dr} provides details about the data reduction.  For the first time in a submillimetre survey the data are processed by four independent data reduction pipelines in order to increase the robustness of the results.  In Section~\\ref{compareall}, the amalgamation of the four source lists into a master $850\\,\\mathrm{\\mu m}$ SHADES catalogue is described.  In Section~\\ref{complete} two different approaches to derive the number counts are described, compared and contrasted.  The SHADES differential count measurements are provided here.  In Section~\\ref{numcts}, models are fit to the differential counts, and the cumulative source counts are computed in order to compare them with previous data.  Concluding remarks regarding what was learned about data analysis by comparing the different reductions, as well as advice for future surveys, are given in Section~\\ref{conc}. ",
        "conclusions": "\\label{conc}  This is the first instance that a SCUBA data set has been subjected to a detailed comparison of differing independent reductions.  As a result, we believe we have produced the most robust and reliable submillimetre-selected source catalogue to date.  We have learned that using monthly average FCFs is just as good as using nightly calibration observations, but with the advantage of less time lost to calibration observations.  Also, using larger pixels (3\\,arcsec as opposed to 1\\,arcsec pixels) to make maps of these data does not seem to increase the positional uncertainty in the centroiding of detected sources.  The SHADES survey is the final legacy of SCUBA, being the largest extragalactic survey undertaken to investigate the nature and redshift distribution of a complete and consistent sample of $>100$ SMGs.  Here, we have presented $850\\,\\mathrm{\\mu m}$ source catalogues and number counts (differential and integral) using information from four independent reductions of the data, complete to about $720\\,\\mathrm{arcmin}^{2}$, down to an RMS noise of $\\simeq2.2\\,\\mathrm{mJy}$.  We also provide $3\\,\\sigma$ upper limits to the $450\\,\\mathrm{\\mu m}$ flux densities of each SHADES source.  With approximately double the area of all blank-field surveys combined, this work has produced the most robust submillimetre source catalogue for multi-wavelength studies, as well as the most accurate measurement of SMG number counts.  We find that the differential counts are better fit with a broken power-law than a single power-law, the location of the break in the source counts being at several mJy.  Using the ratio of the counts between the two SHADES fields, we are unable to provide a useful constraint on SMG clustering.  Overall, the increased accuracy of the source counts due to the new measurement by SHADES will allow us to study the evolutionary nature of SMGs in much more detail than has previously been possible.  SMG number counts have been able to put severe constraints on galaxy evolution models, where new ingredients or assumptions had to be adopted in order to reproduce them (e.g.~\\citealt{Kaviani}, \\citealt{Granato}, \\citealt{Baugh}).  The maps and catalogues that have resulted from this work, in combination with multi-wavelength measurements from deep radio, far-, mid-, and near-infrared, optical and X-ray imaging, will yield the photometric redshifts necessary to investigate the SMG contribution to the cosmic star-formation history of the Universe and to discover the true nature of the SMG population.  Further work on the SHADES data will include a $P(D)$ fluctuation analysis (which should be relatively straightforward because of the fairly uniform noise), and clustering estimates (again, significantly easier for SHADES than for previous surveys), as well as follow-up studies, focussed on estimating redshifts and determining luminosities and other physical properties for the sample.  A detailed analysis and modelling of the number counts will be presented in future paper (Rowan-Robinson et al. in preparation).  Together these will provide much tighter constraints on galaxy evolution models.  We note that wider surveys, such as those planned for SCUBA-2, the LMT, \\textit{Planck} or \\textit{Herschel}, will be able to place better constraints on the bright end of the number counts, since these SHADES maps are insufficiently wide to detect rare bright SMGs.  On the other hand, deeper surveys, such as those also planned for SCUBA-2, are needed to push into the JCMT's confusion regime, and eventually ALMA will allow us to study these SMGs with the desired angular resolution."
    },
    "0609/astro-ph0609513_arXiv.txt": {
        "abstract": "HD149026b is a short-period, Saturn-mass planet that transits a metal- rich star.  The planet radius, determined by photometry, is remarkably small compared to other known transiting planets, with a heavy element core that apparently comprises $\\sim 70$\\% of the total planet mass.  Time series spectra were obtained at Keck before and during transit in order to model the Rossiter-McLaughlin effect. Here, we make use of these observations to carry out a differential comparison of spectra obtained in and out-of-transit to search for signatures of neutral atomic lithium and potassium from the planet atmosphere.  No signal was detected at the 2\\% level; we therefore place upper limits on the column density of these atoms. ",
        "introduction": "Extrasolar planets that transit their host stars offer a unique opportunity to search for constituents in the planet atmosphere. As stellar flux passes through the planet atmosphere, additional absorption features will be imprinted on the (otherwise pure) stellar spectrum (Seager \\& Sasselov 2000, Brown 2001, Hubbard \\etal 2001). This offers a possibility for detection of atmospheric transmission features when working differentially with in-transit and out-of-transit spectra. We have used this fact to search for additional absorption in alkali lines from the atmosphere of HD 149026b during transit.   Transiting planets that reside close to their host stars are strongly irradiated and have atmospheres with temperatures (but not gravities) similar to those of late type L dwarfs.  In the cool atmosphere of sub-stellar objects like brown dwarfs and giant exoplanets, the more refractory metals condense and rain out-of-the atmosphere leaving the less refractory alkali metals in atomic form (Burrows \\& Sharp 1999; Lodders 1999). Thus, the strongest absorption features are expected for alkali metals such as sodium and potassium (Seager \\& Sasselov 2000, Sudarsky, Burrows \\& Pinto 2000).  Since the discovery of HD 209458b (Charbonneau et al. 2000; Henry et al. 2000), this star has been the subject of many ground-based and space-based studies.  HD 209458 is a bright star with $V=7.65$, permitting spectroscopy with high resolution and high S/N. The transiting planet has an anomalously large radius of 1.3 \\rjup, making it an excellent candidate to search for atmospheric signatures. In contrast, searching for atmospheric features in other transiting exoplanets may be more difficult, because the planet to star radius is generally smaller and the host stars have a lower intrinsic brightness. However, if other exoplanets have tenuous, but extended atmospheres, then it may be possible to identify atmospheric features from the planet.  The most successful searches for exoplanet atmospheres have been carried out from space where telluric contamination is avoided and higher S/N has enabled more sensitive searches.  Based on observations with the \\emph{Hubble Space Telescope} (HST), Charbonneau \\etal (2002) detected absorption from the sodium resonance doublet at 5893 \\AA\\ in the dense lower atmosphere of HD 209458b, and Vidal-Madjar \\etal (2003) detected an escaping, trailing extended exosphere of hydrogen.  In contrast, it is more challenging to search for exoplanet atmospheres with ground-based telescopes. Bundy \\& Marcy (2000) carried out a differential analysis of spectra for HD 209458 obtained in and out-of-transit at the Lick and Keck Observatories.  They also studied the spectra of 51 Peg, which is not a transiting planet.  Because of contamination from ${\\rm H_2O}$ and ${\\rm O_2}$ features, they were only able to place weak upper limits on Na I and K I line depth variations. Deming et al. (2005) searched for carbon monoxide absorption features in the transmission spectrum of HD 209458b at $\\sim2.3$ $\\micron$ using NIRSPEC on Keck II telescope.  Their analysis required 1077 high-resolution spectra to achieve comparable sensitivity to the HST sodium observations. No changes in the CO features in and out-of-transit were measured, implying the presence of high clouds in the planet atmosphere. ",
        "conclusions": "The detection of absorption features in the lower atmosphere of a transiting giant planet requires outstanding precision.  Charbonneau \\etal (2002), who detected neutral atomic sodium, and Deming \\etal (2005), who placed upper limits on CO absorption, obtained a precision 100 times better for HD209458 than we have in our study. They were looking for 0.02\\% effects.  These studies made use of the \\emph{Hubble Space Telescope} and over 1000 ground-based spectra, respectively.  However, The detection of species in an extended exosphere need not necessarily require this same precision, because the atmosphere's cross-sectional area is significantly larger (Vidal- Madjar \\etal 2003).  If one thinks of our 2\\% precision as the smallest cross-sectional atmospheric area that we could observe, this corresponds to 2.1 $R_ {\\rm J}$ on top of the opaque-at-all-wavelengths 0.725 $R_{\\rm J}$ planet.  This large distance makes it immediately clear that we would only be sensitive to species in a large extended exosphere, such as that observed by Vidal-Madjar \\etal (2003).  Therefore, only if K and Li had a large enough column abundance to remain optically thick past 2.1 $R_{\\rm J}$ would these atoms have been detected.  Furthermore, these species must remain neutral and atomic, such that they would absorb stellar flux at our specified wavelengths.  While these species are likely neutral in the lower atmosphere around $P \\gtrsim$1 mbar, where Charbonneau \\etal (2002) detected Na in HD 209458b, all exospheric models for highly irradiated giant planets predict temperatures on the order of 10$^4$ K (Lammer \\etal 2003, Yelle 2004, Tian \\etal 2005).  Based on calculations by Yelle (2004) for an HD 2094358b-like planet at 0.04 AU, 2.1 $R_{\\rm J}$ likely corresponds to $P \\sim 10^{-10}$ bar and $T \\sim 2 \\times 10^4$ K. This high temperature makes it likely that Li and K will be ionized. Examining K specifically, Lodders \\& Fegley (2006) review the chemistry of alkalis in substellar atmospheres and plot a curve where K I and K II have an equal abundance.  This curve has only weak pressure dependence.  At the lowest pressure considered, 1 mbar, K I and K II are equal in abundance at a temperature of 1900 K.  Our exospheric temperatures are likely 5-10 times hotter, meaning K II is strongly preferred.  Photodissociation is also likely important and will further reduce the abundances of Li I and K I.  Fortney \\etal (2003) examined ionization of Na in the atmosphere of HD 209458b and found that photoionization of this species on the planet's limb could proceed to pressures as high as $\\sim$ 1 mbar.  The atoms Li and K have even lower ionization potentials.  Due to the high temperatures and strong photoionizing flux, it appears unlikely that neutral atomic Li and K will be found in significant abundance in the extended atmosphere of HD 149026b.  This is consistent with the observations of Charbonneau \\etal (2002), who observed Na I in the lower atmosphere, but not the exosphere, of HD 209458b.  R.~S.~Freedman has kindly provided us with absorption cross-sections as a function of wavelength for K and Li across our observation bands.  These are $\\sim 4 \\times 10^{-13}$ cm$^2$ for Li and $\\sim 1 \\times 10^{-13}$ cm$^2$ for K.  Since we find that the exosphere has an optical depth $\\lessim 1$ at 2.1 $R_{\\rm J}$, this corresponds to column densities for these atoms in their neutral atomic ground state, of $\\lesssim 2 \\times 10^{12}$ cm$^{-2}$for Li, and $\\lessim 9 \\times 10^{12}$ cm$^{-2}$ for K, at this radius.   A more sensitive search from space-based telescopes could result in a detection of the planetary atmosphere. Ionized alkali species would have noble gas electron configurations and would not have prominent absorption features even for space based observations.  We advocate a search for H Lyman $\\alpha$ from space to give a first estimate of the extent of an exosphere in HD 149026b.  At the same time, since there is a clear link between heavy element mass and atmospheric metallicity in the giant planets in our own solar system, the atmosphere and exosphere of HD 149026b could be richer in metals than HD 209458b. The next elements to search for would probably be C, N, and O species.   We gratefully acknowledge the dedication and support of the Keck Observatory staff, in particular Grant Hill for support with HIRES.  We thank Paul Butler and Steve Vogt for contributing the observed spectra.  We thank Richard Freedman for providing us with opacity data.  Thanks to Jessica Lovering for reduction of the CCD images to 1-D spectra.  We thank the NOAO and NASA Telescope assignment committees for generous allocations of telescope time. Data presented herein were obtained at the W. M. Keck Observatory from telescope time allocated to the National Aeronautics and Space Administration through the agency's scientific partnership with the California Institute of Technology and the University of California.  The Observatory was made possible by the generous financial support of the W. M. Keck Foundation. We thank the Michaelson Science Center for travel support and support through the KDPA program.  DAF is a Cottrell Science Scholar of Research Corporation. We acknowledge support from NASA grant (to DAF); a NASA Postdoctoral Program Fellowship (to JJF); NASA grant NCC5-511; NASA grant NAG5-75005 (to GWM).    \\clearpage"
    },
    "0609/astro-ph0609808_arXiv.txt": {
        "abstract": "The spectral energy distributions (SEDs) of some T Tauri stars display a deficit of near-IR flux that could be a consequence of an embedded Jupiter-mass planet partially clearing an inner hole in the circumstellar disc. Here, we use two-dimensional numerical simulations of the planet-disc interaction, in concert with simple models for the dust dynamics, to quantify how a planet influences the dust at different radii within the disc. We show that pressure gradients at the outer edge of the gap cleared by the planet act as a filter - letting particles smaller than a critical size through to the inner disc while holding back larger particles in the outer disc. The critical particle size depends upon the disc properties, but is typically of the order of 10 microns. This filtration process will lead to discontinuous grain populations across the planet's orbital radius, with small grains in the inner disc and an outer population of larger grains. We show that this type of dust population is qualitatively consistent with SED modelling of systems that have optically thin inner holes in their circumstellar discs. This process can also produce a very large gas-to-dust ratio in the inner disc, potentially explaining those systems with optically thin inner cavities that still have relatively high accretion rates. ",
        "introduction": "A fraction of T Tauri stars, including TW~Hya \\citep{calvet02,uchida04}, GM~Aur \\citep{marsh92,koerner93,rice03}, CoKu Tau/4 \\citep{forrest04,quillen04, dalessio05} and DM~Tau \\citep{calvet05}, present spectral energy distributions (SEDs) that combine little excess flux (above the photospheric level) at near-IR wavelengths with robust excesses at $\\lambda \\gtrsim 10 \\mu{\\rm m}$. A similar effect is seen in at least one Herbig Ae/Be star, HD100546 \\citep{bouwman03,grady05}. Since circumstellar discs are generally most optically thick in the inner, hottest regions, this observation is noteworthy and indicative of some process that can create an optically thin inner cavity within the disc. One exciting possibility is that the cavity arises from the tidal truncation of the inner disc by a massive planet \\citep{goldreich80,lin86}, in which case the SED observations provide unique constraints on the time scale and frequency of giant planet formation. However, non-planetary explanations are also possible. Dust can grow rapidly to form solid bodies that contribute negligible opacity in any observed waveband, while larger bodies may themselves suffer destructive collisions or rapid orbital decay due to aerodynamic drag against the gaseous disc \\citep{weidenschilling77}. Photoevaporation of the outer disc by ionizing stellar radiation can also starve the inner disc of gas, yielding a short inner hole phase even in the absence of a planet \\citep{alexander06}. Protoplanetary disc models that include some or all of these effects \\citep{takeuchi01,kenyon04,dullemond05,takeuchi05} admit the formation of dusty rings and inner holes, even in the absence of planets and gas. As improved SEDs and spectra from {\\em Spitzer} become available, more detailed theoretical study of both classes of model (pure disc and disc plus embedded planet) is warranted in order to determine whether these models can be distinguished without direct imaging data.  A sufficiently massive planet (or brown dwarf) can probably create an inner cavity in which both the gas and dust surface densities are extremely low. However, of the four aforementioned T~Tauri stars -- around which inner disc `holes' are inferred from the lack of near-IR emission -- only CoKu Tau/4 appears to be entirely devoid of circumstellar material within 1~AU of the star. In the case of TW~Hya there is both ongoing low-level gas accretion \\citep{muzerolle00}, and evidence (from the existence of a silicate emission feature at 10$\\mu$m) for a sub-lunar mass of micron sized dust at small radii \\citep{calvet02}. Both GM~Aur and DM~Tau are Classical T Tauri stars, with estimated accretion rates on to the stellar photosphere variously estimated at $\\dot{M} \\approx 10^{-9} \\ M_\\odot {\\rm yr}^{-1}$ to $10^{-8} \\ M_\\odot {\\rm yr}^{-1}$ \\citep{gullbring98,hartmann98,white01}. Unfortunately \\citet{najita03} failed to detect CO emission from an inner disc around GM~Aur, so there are no observational clues as to the {\\em gas} surface density that must accompany this accretion. However, for these accretion rates protoplanetary disc models predict gas surface densities at $r \\approx 0.1 \\ {\\rm AU}$ of 200-2000~gcm$^{-2}$ \\citep{bell97}. Although disc models are subject to large uncertainties arising from the unknown efficiency of angular momentum transport within protoplanetary discs, the robust accretion signatures seen for these stars imply that an inner gas disc must be present.  Within a planetary interpretation of the observed SEDs, the presence of inner gas discs and ongoing stellar accretion is not in itself surprising. Numerical simulations \\citep{artymowicz96,lubow99} show that the tidal barrier erected by a massive planet is quite leaky. For Jupiter mass planets gas can readily overflow the gap, and while much of this gas is accreted by the planet in excess of 10\\% of the accretion rate at large radius can replenish the inner disc and ultimately accrete on to the star \\citep{lubow06}. This is consistent with GM~Aur being a Classical T~Tauri star, but the lack of dust opacity then requires that the planet needs to be relatively more effective at holding back dust -- i.e. it needs to filter the gas as it passes the tidal barrier. In this paper, we show that the required filtration may occur as a consequence of the sharp pressure gradient that exists near the outer edge of the gap. The pressure gradient leads to super-Keplerian rotation of the gas, which in turn allows standard aerodynamic coupling to solid bodies \\citep{weidenschilling77,paardekooper06} to inject angular momentum and hold back the inward flow of all but the smallest solid particles. We estimate the efficiency of this process, and argue that it allows relatively low mass planets -- which are observed \\citep{marcy05} to be much more common than those planets with $M_p > M_J$ that would start to cut off accretion entirely -- to nevertheless create dust-poor inner cavities. Within this scenario, the variation in the observed properties of T~Tauri stars with weak near-IR emission could reflect different mass embedded planets -- relatively low in the case of GM~Aur and DM~Tau, and substantially higher in the case of CoKu Tau/4. ",
        "conclusions": "Although we cannot definitively conclude that near-IR deficits in the SEDs of some TTauri stars are a consequence of an embedded Jupiter-like planet, that the filtering effect of such a planet - producing a small grain population in the inner disc that differs from the population in the outer disc - is consistent with SED modelling of these systems is suggestive. The SED modelling also requires that the inner disc be optically thin, while ongoing gas accretion requires that the inner disc still have a substantial amount of gas.  If the grain size distribution is somewhat flatter than the standard ISM size distribution \\citep{mathis77}, filtration at the outer gap edge can significantly reduce the amount of solid material reaching the inner disc, while still allowing sufficient gas through the gap. Although we don't actually know the grain size distribution in the outer discs of those systems that show inner holes, they do show evidence for grain growth. In some cases flatter size distributions do provide better fits to TTauri SEDs \\citep{dalessio01}, although this has yet to be tested on the systems we consider here. That filtration of solid particles at the outer egde of a gap opened by an embedded Jupiter-like planet can produce not only a population of small grains in the inner disc, but can potentially also make this region optically thin, makes this an attractive possibility.  The filtration process should also lead to an enhancement of larger particles at the peak of the gap edge. Particles exterior to the gap will continue to drift inwards until they reach the gap edge, at which point the change in the gas pressure gradient will prevent the larger particles from drifting any further. If the surface density of such particles becomes sufficiently high, further planet formation could occur at this location. Since `Hot' Jupiters are thought to form at modest radii and then migrate inwards via Type II migration to their present locations, one may expect in some cases to find additional lower mass planets in 2:1 resonances with these `Hot' Jupiters.  The presence of even terrestrial mass planets in resonant orbits with `Hot' Jupiters could be detected using transit timing \\citep{agol05}. Although such a detection would be consistent with the analysis presented here, it has been suggested that this could also occur if low-mass, fast migrating planets ($> 0.1$ M$_\\oplus$) were resonantly trapped by gap opening planets \\citep{thommes05}.  In the latter case, however, one may expect the trapped planets to subsequently grow into gas giants \\citep{thommes05}, while in the former case the planetesimal trapping could occur later when there may only be sufficient time for the newly forming planet to reach a terrestrial-like mass."
    },
    "0609/astro-ph0609725_arXiv.txt": {
        "abstract": "We perform test-particle simulations in a 2D, differentially rotating stellar disk, subjected to a two-armed steady state spiral density wave perturbation in order to estimate the influence of spiral structure on the local velocity field. By using Levenberg-Marquardt least-squares fit we decompose the local velocity field (as a result of our simulations) into Fourier components to fourth order. Thus we obtain simulated measurements of the Oort constants, $A$, $B$, $C$, and $K$. We get relations between the Fourier coefficients and some galactic parameters, such as the phase angle of the Solar neighborhood and the spiral pattern speed. We show that systematic errors due to the presence of spiral structure are likely to affect the measurements of the Oort constants. Moderate strength spiral structure causes errors of order 5 km/s/kpc in $A$ and $B$. We find variations of the Fourier coefficients with velocity dispersion, pattern speed, and sample depth. For a sample at an average heliocentric distance of 0.8 kpc we can summarize our findings as follows: (i) if our location in the Galaxy is near corotation then we expect a vanishing value for $C$ for all phase angles; (ii) for a hot disk, spiral structure induced errors for all Oort constants vanish at, and just inward of the corotation radius; (iii) as one approaches the 4:1 Lindblad resonances $|C|$ increases and so does its variation with galactic azimuth; (iv) for all simulations $|C|$, on average, is larger for lower stellar velocity dispersions, contrary to recent measurements. ",
        "introduction": "Jan Oort proposed \\citep{oort27a} and measured \\citep{oort27b} the now called Oort constants $A$ and $B$, over seven decades ago by assuming that stars in the Galaxy moved on circular orbits. Later on his analysis was generalized to the non-axisymmetric case (e.g., \\citealt{ogorodnikov32}) resulting in the definition of two more constants: $C$ and $K$. The importance of the Oort constants (OC) is in their simple relation (at least in the case of vanishing random motions) to the galactic potential.  Many investigations have been conducted in the attempt to measure the OC. However, not all the results are in agreement (for a review see \\citealt{kerr86}), which is mainly due to the absence of a complete proper motions catalog. Recent studies have failed to improve measurements of the OC.  Another reason for the diversity among the values derived for the OC has been discussed by \\citet{olling03}. The authors present an effect which arises from the longitudinal variations of the mean stellar parallax caused by intrinsic density inhomogeneities. Together with the reflex of the solar motion these variations create contributions to the longitudinal proper motions which are indistinguishable from the OC at $\\le$ 20\\% of their amplitude. \\citet{olling03} (hereafter O\\&D) measured the OC using proper motions from ACT/Tycho-2 catalog and corrected for the ``mode mixing\" effect described above, using the latitudinal proper motions. They found that the OC vary approximately linearly with asymmetric drift (or age) from the young-stars values to those appropriate for old populations. This trend may result from perturbations induced by the Galactic bar or spiral structure. In this paper we investigate the latter possibility.  Deviations from the axisymmetric values of the OC have not received a satisfactory interpretation although they are often assumed to arise from a non-axisymmetric Galactic potential perturbation such as spiral structure or/and the Galactic bar.  Given their preeminent importance, we feel that a better understanding of the effect of spiral perturbation on the measurements of the OC is necessary. Spiral arms perturb the local velocity field corrupting local measurements of the OC. Thus we look for the effect of spiral structure on the OC $A,B,C,$ and $K$ and their deviation from the axisymmetric limit. We hope that this would give us some constraints on the spiral structure. Our approach applies to the kinematics of the Local Standard of Rest (LSR), while at the same time it constitutes a general analysis of a spiral density wave perturbing a stellar disk.  Our simulation model is described in \\S \\ref{sec:model}. We present a derivation of the OC in \\S \\ref{sec:derive} and describe our numerical measurement technique in \\S \\ref{sec:measure}. Results are presented in \\S \\S \\ref{sec:axi_hot}$-$\\ref{sec:sample_depth}. Discussion of observational  measurements of the OC is given in \\S \\ref{sec:observ} and summary follows in \\S \\ref{sec:sum}. ",
        "conclusions": "\\label{sec:sum}  The importance of the Oort constants is in their simple relation to the galactic potential, in the case of vanishing random motions. The circular orbits have velocity $r\\partial\\Phi/\\partial r=v^2_\\phi=(A_0-B_0)^2$. If we could find the causes for the deviations, $\\Delta A$ and $\\Delta B$, from the ``true\" values, $A$ and $B$, then we would provide a direct constraint on the Galactic potential.  In this paper we have investigated the effect of spiral structure on the measurements of the OC $A$, $B$, $C$, and $K$. We have performed test-particle simulations with an initially cold or hot stellar disk and an imposed two-armed spiral density wave perturbation. Variations of simulated measurements of the OC with pattern speed, galactic azimuth, and sample depth have been explored.  We find that systematic errors due to the presence of spiral structure {\\it are} likely to affect the measurements of the Oort constants. Moderate strength spiral structure causes errors of order 5 km/s in $A$ and $B$. Axisymmetric, high velocity dispersion disks also yield errors in the simulated measurements of Oort's $A$ (as a result of the asymmetric drift), but they are smaller, of order 2 km/s (see \\S \\ref{sec:axi_hot}).  If the Sun is located near the corotation resonance then, regardless of the phase angle, we would measure $C\\approx0$ and thus get no information about the spiral structure. If, in addition, we constrain $\\phi_0\\approx45^\\circ$, or $135^\\circ$, all OC would have values within the current measurement error. As the Lindblad resonances are neared, $C$ varies strongly with phase angle and has an increased value.  In \\S \\ref{sec:sample_depth} we investigated the effect of sample depth on the measurements of the OC. We found that for an initially cold disk, spiral structure can give rise to marked variations of all OC with the average heliocentric distance of the sample. For $A$ and $B$ this results from an oscillating average $v_\\phi(r)$, whereas $C$ is sensitive to variations in the radial velocity and its gradient. We find that for phase angles placing the Sun near the convex arm negative values for the axisymmetric drift are possible for low velocity dispersion stars. Unlike the $\\sigma_u$ induced asymmetric drift in a hot axisymmetric disk, in a cold disk it is the spiral structure that causes it. It is possible the reason for never observing $v_a<0$ could be our proximity to the concave spiral arm.  Our result for the variation of $C$ with velocity dispersion is distinctly different than the observational measurement of $C$ done by O\\&D. Whereas in our simulations the absolute value of $C$ decreases with increasing velocity dispersion, the opposite trend is observed in fig. 6 by O\\&D. In the same paper the authors suggest that it is possible that the Galactic bar is responsible for this behavior. Since orbits change orientation at the 2:1 OLR of the bar, at this location the bar is expected to most strongly distort the local kinematics. As discussed by O\\&D, stars from inside the OLR should produce $C<0$. If the Sun's location is just outside the OLR then we would sample more stars giving rise to negative $C$ with increasing sample depth. To test this possibility, O\\&D analyzed only stars brighter than $V_T=10.5$ which yielded a decrease in the ``phase mixing\" corrected value of $C$ from $\\sim-10$ to $\\sim-7.5$ km/s. Another way to try to explain it is to consider the combined effect of spiral structure and bar perturbations. Other possibilities are the effects from minor mergers, or those due to a triaxial halo \\citep{kuijken94}. Future work should aim at understanding this unusual behavior for $C$."
    },
    "0609/astro-ph0609455_arXiv.txt": {
        "abstract": "We report high angular-resolution ($\\sim$1$''$) CO J=3--2 interferometric mapping, using the Submillimeter Array (SMA), of IRAS\\,22036+5306 (I\\,22036), a bipolar pre-planetary nebula (PPN) with knotty jets discovered in our HST SNAPshot survey of young PPNs. In addition, we have obtained supporting lower-resolution ($\\sim$10$''$) CO and $^{13}$CO J=1--0 observations with the Owens Valley Radio Observatory (OVRO) interferometer, as well as optical long-slit echelle spectra at the Palomar Observatory. The CO J=3--2 observations show the presence of a very fast ($\\sim$220\\,\\kms), highly collimated, massive (0.03\\ms) bipolar outflow with a very large scalar momentum (about $10^{39}$ g cm s$^{-1}$), and the characteristic spatio-kinematic structure of bow-shocks at the tips of this outflow. The H$\\alpha$ line shows an absorption feature blue-shifted from the systemic velocity by $\\sim$100\\,\\kms, which most likely arises in neutral interface material between the fast outflow and the dense walls of the bipolar lobes at low latitudes. Since the expansion age of the outflow as determined from our data is only about 25 years, much smaller than the time required by radiation pressure to accelerate the observed bipolar outflow to its current speed ($\\gtrsim10^5$ yr), the fast outflow in I\\,22036, as in most PPNs, cannot be driven by radiation pressure. The total molecular mass (0.065\\,\\ms) is much less than that derived from a previous detailed model of the near to far-infrared SED of this object (showing the presence of a large, cool dust shell of total mass 4.7\\ms), most likely due to the interferometric observations resolving out the CO flux from the dust shell due to its size. We find an unresolved source of submillimeter (and millimeter-wave) continuum emission in I\\,22036, implying a very substantial mass (0.02-0.04\\,\\ms) of large (radius$\\gtrsim$1\\,mm), cold ($\\lesssim$50\\,K) dust grains associated with I\\,22036's toroidal waist. We also find that the $^{13}$C/$^{12}$C ratio in I\\,22036 is very high (0.16), close to the maximum value achieved in equilibrium CNO-nucleosynthesis (0.33). The combination of the high circumstellar mass (i.e., in the extended dust shell and the torus) and the high $^{13}$C/$^{12}$C ratio in I\\,22036 provides strong support for this object having evolved from a massive ($\\gtrsim$4\\,\\ms) progenitor in which hot-bottom-burning has occurred. ",
        "introduction": "Pre-Planetary nebulae (PPNs) - objects in transition between the AGB and planetary nebula (PN) evolutionary phases - hold the key to one of the most vexing and long-standing problems in our understanding of these very late stages of evolution for low and intermediate mass stars. Current observational evidence shows that during the PPN phase, the dense, isotropic mass-loss which generally marks the late AGB phase changes to high-velocity, bipolar or multipolar mass-loss. The emergence of fast collimated outflows or jets has been hypothesized as the primary mechanism for this dramatic change in the geometry and dynamics of the mass-loss (Sahai \\& Trauger 1998). However, the physical mechanism for producing the fast outflows remains unknown. (Sub)millimeter-wave interferometric observations offer one of the best probes of the dynamics and energetics of the shock acceleration process, which transfers a substantial amount of directed momentum to large parts of the dense AGB wind.  IRAS\\,22036+5306 (I\\,22036) is a bipolar PPN with knotty jets discovered in our HST SNAPshot survey of young PPNs, with an estimated distance and luminosity of 2\\,kpc and 2300\\ls, respectively (Sahai et al. 2003 [SZSCM03]). Our HST imaging suggested that I\\,22036 is a key object for understanding how jets can sculpt bipolar lobes in a progenitor AGB star wind, since its jets and their working surfaces appear to be simultaneously observed. In this paper (the second one by Sahai and collaborators on this important PPN) we report CO J=3--2 observations of the PPN with the Submillimeter Array (SMA), which resolve a massive high-velocity bipolar outflow aligned with the optical nebulosity and clearly pinpoint the location and nature of the bow-shocks at the tips of this outflow. We also report supporting lower-resolution ($\\sim$10$''$) CO and $^{13}$CO J=1--0 observations with the Owens Valley Radio Observatory (OVRO) interferometer, as well as optical long-slit echelle spectra at the Palomar Observatory.  The rest of the paper is organised as follows: in \\S\\,2 we describe our observational techniques and data-reduction, in \\S\\,3 we describe our observational results, in \\S\\,4 we discuss the kinematics, age and mass of the bipolar outflow, in \\S\\,5 we derive the total molecular mass and compare it to the mass of dust in the circumstellar environment, in \\S\\,6 we derive and discuss the implications of the $^{13}$C/$^{12}$C isotope ratio in I\\,22036, in \\S\\,7 we discuss the formation of the nebula, and in \\S\\,8 we summarize our main conclusions. ",
        "conclusions": "Using the SMA, we have made high angular resolution ($\\sim 1{''}$) maps of the CO J=3--2 emission in IRAS\\,22036+5306, a bipolar pre-planetary nebula with knotty jets, discovered previously in one of our HST imaging surveys of candidate PPNs. These data have been supplemented by lower-resolution CO and $^{13}$CO J=1-0 observations with OVRO, as well as optical long-slit echelle spectroscopy. We find  \\noindent (i) The CO J=3--2 spatially integrated line profile covers a total velocity extent (FWZI) of $\\sim$500\\,\\kms. The CO and $^{13}$CO J=1--0 also show large velocity widths, although with a smaller spread due to the lower sensitivity of these data.\\\\ (ii) The CO J=3--2 emission comes from a massive high-velocity bipolar outflow aligned with the optical nebulosity. The characteristic position-velocity signature of a bow-shock resulting from a high-velocity bipolar collimated outflow/jet interacting with dense ambient material can be seen in the CO J=3--2 data in both the red and blue-shifted components of the outflow. The P-V structure has a high degree of symmetry around the center, indicating that the bipolar outflow and surrounding medium are correspondingly symmetric. The bow-shocks do not lie at the tips of the optical lobes.\\\\ (iii) The H$\\alpha$ line shows a profile with a central peak, very wide wings, and an absorption dip in the blue wing of the line, blue-shifted from the systemic velocity by about 100\\,\\kms. The absorption feature indicates the presence of a fast, neutral, outflow inside the lobes.\\\\ (iv) A central, unresolved, continuum source with a flux of 0.29$\\pm$0.04\\,Jy (8.4$\\pm$1.7\\,mJy) at $\\lambda$=0.88\\,mm (2.6\\,mm) is seen in the SMA (OVRO) data.\\\\ (v) The mass in the high-velocity outflow, i.e., at velocities offset more than $\\pm$\\,15\\,\\kms~from the systemic velocity, is $0.03$\\,\\ms, and the scalar momentum is 1.1$\\times10^{39}$ g cm s$^{-1}$. The large value of the scalar momentum, together with the small age of the high-velocity outflow, implies that it cannot be driven by radiation pressure.\\\\ (vi) The total molecular mass derived from integrating the flux of the CO J=3--2 line over its full velocity range, is $\\sim$0.065\\ms, much lower than the value ($\\sim$4.7\\ms) estimated previously from a dust-shell model of I\\,22036's near to far-infrared fluxes, assuming a gas-to-dust ratio, $\\delta$=200. This is probably due to the interferometric J=3--2 observations resolving out the flux from the large shell (inner and outer radii of $1.4\\times 10^{17}$ and $5.3\\times 10^{17}$ cm) in which the cool dust resides. Furthermore, the gas in this shell may not be sufficiently excited.\\\\ (vii) The small size of the continuum source indicates that it is associated with the dusty toroidal waist of I\\,22036. The submillimeter and millimeter-wave continuum fluxes imply the presence of a very substantial mass ($\\sim$0.02-0.04\\ms) of large (of radius $\\gtrsim$ 1\\,mm) cold grains in this waist.\\\\ (viii) The $^{13}$C/$^{12}$C ratio in I\\,22036 is 0.16, close to the maximum value value achieved in equilibrium CNO-nucleosynthesis (0.33). The enhancement is most likely due to the operation of hot-bottom-burning, which operates in stars with initial masses larger than about 4\\ms.  We thank Nick Sterling for pointing out the importance of hot-bottom-burning in enhancing the $^{13}$C/$^{12}$C ratio, and Noam Soker for reading an earlier version of this paper. RS is thankful for partial financial support for this work from a NASA/ ADP and LTSA grant. CSC's work was supported by the National Science Foundation through Grant No. 9981546 to Owens Valley Radio Observatory and the Spanish MCyT under project DGES/AYA2003-2785 and the 2005 ``Ram\\'on y Cajal\" program."
    },
    "0609/astro-ph0609349_arXiv.txt": {
        "abstract": "If the universe had a large curvature before inflation there is a deviation from the scale invariant perturbations of the inflaton at the beginning of inflation. This may have some effect on the CMB anisotropy at large angular scales. We calculate the density perturbations for both open and closed universe cases using the Bunch-Davies vacuum condition on the initial state. We use our power spectrum to calculate the temperature anisotropy spectrum and compare the results with the WMAP three year data. We find that our power spectrum gives a lower quadrupole anisotropy when $\\Omega-1 >0$, but matches the temperature anisotropy calculated from the standard Ratra-Peebles power spectrum at large $l$. The determination of spatial curvature from temperature anisotropy data is not much affected by the different power spectra  which arise from the choice of different boundary conditions for the inflaton perturbation. ",
        "introduction": "The curvature of the universe goes down exponentially after the start of inflation \\cite{guth}. If there is a residual curvature still present by the time the scales which are entering our horizon at present were leaving the inflationary horizon there will be deviation from the scale invariant perturbations  due to non-zero curvature. The corrections to the power spectrum at horizon scales are multiplicative powers of $(1\\pm K/\\beta^2)$, where the curvature $K=(\\Omega_0-1)(a_0 H_0)^2$ and $\\beta$ is the comoving canonical wavenumber. We calculate the primordial power spectrum for the case closed and open universe at the time of inflation. We choose the Bunch-Davies boundary condition to normalize the wave-functions. For the case of closed universe we obtain the following expression for the power spectrum \\begin{equation} P_{\\cal R}(\\beta)=\\frac{H_\\lambda^4}{2 \\pi^2 \\dot \\phi^2}\\frac{1}{\\left(1+\\frac{K}{\\beta^2}\\right)^2},\\quad \\quad \\frac{\\beta}{\\sqrt{K}}=3,4,5 \\cdots \\,\\,( {\\rm for}\\ K>0) \\label{closed1} \\end{equation} and for the case of inflation in an open universe \\begin{equation} P_{\\cal R}(\\beta)=\\frac{H_\\lambda^4}{2 \\pi^2 \\dot \\phi^2}\\frac{1}{\\left(1-\\frac{|K|}{\\beta^2}\\right)^2 \\left(1+\\frac{|K|}{\\beta^2}\\right)},\\quad \\quad \\frac{\\beta}{\\sqrt|K|}>1, \\,\\,( {\\rm for} \\ K<0) \\label{open1} \\end{equation} where $\\phi$ is the inflaton field. In the case of closed universe $\\beta$ takes discrete values in units of $\\sqrt{K}=R_c^{-1}$ ($R_c$ being the curvature radius), the modes corresponding to $\\beta/\\sqrt{K}=1,2$ can be eliminated by gauge transformations \\cite{abbott} so there is a large-wavelength cutoff at $\\beta_c^{-1}= R_c/3$. This large wavelength cut-off in a closed universe has been used to explain the observed low CMB anisotropy at low multi-poles \\cite{efsth} and \\cite{lasenby}. In the case of open universe only modes with $\\beta > \\sqrt{|K|}$ cross the inflationary horizon.  Our result for the power spectrum in the closed and open universe cases differs from the phenomenological power spectrum \\cite{Hu1997}, \\begin{equation} P_{\\cal R}(\\beta) = \\frac{H_\\lambda^4}{2 \\pi^2 \\dot \\phi^2} \\,\\,\\,\\frac{1} {1+\\frac{K}{\\beta^2}} \\label{HSWZ} \\end{equation} used in the calculation of CMB anisotropies in both the closed and open cases. Our results agree qualitatively with (\\ref{HSWZ}) in that the power at small $\\beta$ is suppressed in the closed universe inflation (\\ref{closed1}) and enhanced in the open universe (\\ref{open1}).  According to inflation \\cite{guth}, the curvature of the present universe $\\Omega_0-1$ is related to the curvature at any time during  inflation $\\Omega_{i}-1$ as \\begin{equation} \\frac{\\Omega_0-1}{\\Omega_i-1}=\\left(\\frac{a_i H_i}{a_0 H_0}\\right)^2. \\end{equation} If $a_i$ is the scale factor at the time during inflation when scales of the size of our present horizon were exiting the inflationary horizon then $a_0 H_0= a_i H_i$ and $\\Omega_0=\\Omega_i$. If at the beginning of inflation $(\\Omega_{start}-1)=O(1)$ then in order to have a deviation of say one-percent from unity in the present curvature, the number of e-foldings prior to the $a_i$ must be small. Putting an upper bound on the present curvature $(\\Omega_0-1)$ from observations also puts a lower bound on the number of extra e-foldings necessary in inflation in addition to the minimum number needed to solve the horizon problem \\cite{ellis}.  The geometry of the universe can be determined from the CMB anisotropy from the angular size of the acoustic peak. However the constraints on the density of the universe $\\Omega$ depend upon priors like the value $H_0$ and $\\Omega_\\lambda$. For example the combination of WMAP \\cite{wmap3} , LSS \\cite{Tonry:2003zg} and HST \\cite{Freedman:2000cf} supernovae observations  gives a constraint on the density of the universe as $(\\Omega-1)=0.06^{+0.02}_{-0.02}$ \\cite{Dunkley} which means that the curvature at one-$\\sigma$ could be as large as $K/(a_0 H_0)^2=0.08$. In the case of the closed universe the power spectrum (\\ref{closed1}) $P_{\\cal R} \\propto (1+K/\\beta^2)^{-2}$ at the scale of the horizon $\\beta=a_0 H_0$ would be suppressed by about $16\\%$  compared to the power for the flat universe.  We use our power spectrum to calculate the temperature anisotropy spectrum and compare the results with the WMAP three year data. We find that our power spectrum gives a lower quadrupole anisotropy when $\\Omega-1 >0$, but matches the temperature anisotropy calculated from the standard Ratra-Peebles power spectrum at large $l$. We also find that using the closed universe power spectrum (\\ref{closed2})  for larger  values of $ \\Omega_0$ the quadrupole anisotropy is suppressed more and fits the WMAP data better .  This supports the idea proposed in \\cite{efsth} that a positive spatial curvature should suppress the power at low $l$. ",
        "conclusions": "At the beginning of inflation the curvature $\\Omega-1 $ is expected to be of order one. By the time perturbations of our horizon size exit the inflation horizon, the curvature drops to $\\Omega_0-1$ which is the present value. A non-zero observation of the curvature will tell us whether the universe prior to inflation was open or closed (even though it is almost flat now) and put constraints on the number of extra e-foldings that must have occurred beyond the minimum number needed to solve the horizon problem. Spatial curvature is a threshold effect which can give us information on the pre-inflation universe from observations of the CMB anisotropy at large angles, similar to the effect of a possible pre-inflation thermal era \\cite{temp,gravity}. >From the power spectrum of the closed (\\ref{closed1}) and open inflation (\\ref{open1}) cases we see that if $K >0$, power is suppressed at large angles and if $K<0$ power is enhanced at large angles. The WMAP observation of a suppression of the quadrupole temperature anisotropy supports inflation in a closed universe \\cite{efsth}. The spatial curvature is determined in the CMB anisotropy mainly from the angular size of the acoustic horizon. The determination of the spatial curvature from the WMAP data is not substantially affected by the choice of the boundary condition used for the determination of the primordial power spectrum in a curved inflationary universe."
    },
    "0609/astro-ph0609663_arXiv.txt": {
        "abstract": "We present observations with the INTEGRAL/IBIS telescope of the wind nebula powered by the young pulsar B1509-58 and we discuss the spatial and spectral properties of the unpulsed emission in the $20-200$ keV energy band. The source extension and orientation along the northwest-southeast axis corresponds to the jet emission seen at keV and TeV energies. The hard X-ray spectrum is consistent with the earlier Beppo-SAX measurements. It follows a power law with a photon index $\\alpha = -2.12 \\pm 0.05$ up to $160$ keV. A possible break at this energy is found at the $2.9\\sigma$ confidence level. The 0.1-100 keV data are consistent with synchrotron aging of pairs in the jet and yield a magnetic field strength of 22-33 $\\mu$G for a bulk velocity of 0.3-0.5c. The synchrotron cut-off energy thus corresponds to a maximum electron energy of 400-730 TeV. ",
        "introduction": "The young radio pulsar B1509-58 is associated with the supernova remnant MSH 15-5{\\it 2} (G320.4-1.2) \\citep{Seward1982,Manchester1982}. Its spin parameters (a period of $150$ ms and period derivative of $1.5 \\times 10^{-12}$ s s$^{-1}$) make it one of the youngest and most energetic pulsars known, with a characteristic age $\\tau \\sim 1700$ yr and a spin-down power of $1.8 \\times 10^{37}$ ergs s$^{-1}$. It presents one of the largest magnetic field strength ($1.5 \\times 10^{13}$ G) \\citep{Kaspi1994} recorded for an isolated pulsar. Recent X-ray and $\\gamma$-ray images have revealed a very complex pulsar environment where one can study several manifestations of pulsar wind nebulae: the shocked wind in an equatorial flow, a powerful jet, and the interaction and confinement of these wind features within the surrounding remnant.  ROSAT first detected the wind nebula in X rays, showing its 17'-long tail extending to the southeast along a faint ridge of radio emission \\citep{Greiveldinger1993,Trussoni1996}, later confirmed by BeppoSAX. The tail length decreases from 17' to 8.5' (at 10 $\\%$ level of the central nebula intensity) between 1.6 keV and 10.0 keV \\citep{Mineo2001}. High-resolution Chandra images revealed further details of the non-thermal structure \\citep{Gaensler2002}; first, the axis of the elongated nebula (with a position angle of $150^{\\circ} \\pm 5^{\\circ}$ north through east) which may correspond to the pulsar spin axis; then two northern toroidal arcs of emission, 17\" and 30\" away from the pulsar, that may correspond to wisps in an ion-loaded equatorial flow as in the Crab nebula; and a 4'-long bright collimated feature along the axis which is interpreted as a jet with a velocity $v_{jet} > 0.2c$ \\citep{Gaensler2002,Tamura1996,Brazier1997}. The apparent absence of a counterjet sets a lower limit of 5 to the brightness ratio between the two sides (if two jets are powered). Doppler boosting can account for this contrast provided that the jet bulk velocity and its inclination to the line of sight verify $\\beta_{jet} cos{\\theta_{jet}} \\sim 0.28$. The wind carries away at least $0.05\\%$ of the pulsar spin down power.  The emission found by H.E.S.S between $280$ GeV and $40$ TeV presents a comparable morphology \\citep{Hess2005}. The elongated nebula extends on both sides of the pulsar along the X-ray axis, over 6' (1 $\\sigma$ level) to the northwest and southeast direction. The $\\gamma$-ray spectrum is well fitted by an $E^{-2.27 \\pm 0.23 \\pm 0.20}$ power law and can be explained by TeV electrons up-scattering the cosmological microwave background and the ambient interstellar radiation field. Fitting both the TeV emission and the BeppoSAX one gives a mean magnetic field of $17$ $\\mu$G which is in agreement with lower limits found for the jet in the equipartition assumption \\citep{Delaney2006}. At a distance of $5.2 \\pm 1.4$ kpc \\citep{Gaensler1999}, the bright part of the X-ray jet, the longer southeast X-ray tail, and the coincident TeV one extend from 5 to 30 pc on the sky.  We present observations of the PSR B1509-58 region obtained with the INTEGRAL/IBIS telescope operating in the hard X-ray range, describing the morphology and spectral properties of the emission and comparing them with the X-ray and TeV data. ",
        "conclusions": ""
    },
    "0609/astro-ph0609380_arXiv.txt": {
        "abstract": "}[2]{{\\footnotesize\\begin{center}ABSTRACT\\end{center} \\vspace{1mm}\\par#1\\par \\noindent {~}{\\it #2}}} \\newcommand{\\TabCap}[2]{\\begin{center}\\parbox[t]{#1}{\\begin{center} \\small {\\spaceskip 2pt plus 1pt minus 1pt T a b l e} \\refstepcounter{table}\\thetable \\\\[2mm] \\footnotesize #2 \\end{center}}\\end{center}} \\newcommand{\\TableSep}[2]{\\begin{table}[p]\\vspace{#1} \\TabCap{#2}\\end{table}} \\newcommand{\\FigCap}[1]{\\footnotesize\\par\\noindent Fig.\\  % \\refstepcounter{figure}\\thefigure. #1\\par} \\newcommand{\\TableFont}{\\footnotesize} \\newcommand{\\TableFontIt}{\\ttit} \\newcommand{\\SetTableFont}[1]{\\renewcommand{\\TableFont}{#1}} \\newcommand{\\MakeTable}[4]{\\begin{table}[htb]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}} \\newcommand{\\MakeTableSep}[4]{\\begin{table}[p]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}} \\newenvironment{references}% { \\footnotesize \\frenchspacing \\renewcommand{\\thesection}{} \\renewcommand{\\in}{{\\rm in }} \\renewcommand{\\AA}{Astron.\\ Astrophys.} \\newcommand{\\AAS}{Astron.~Astrophys.~Suppl.~Ser.} \\newcommand{\\ApJ}{Astrophys.\\ J.} \\newcommand{\\ApJS}{Astrophys.\\ J.~Suppl.~Ser.} \\newcommand{\\ApJL}{Astrophys.\\ J.~Letters} \\newcommand{\\AJ}{Astron.\\ J.} \\newcommand{\\IBVS}{IBVS} \\newcommand{\\PASP}{P.A.S.P.} \\newcommand{\\Acta}{Acta Astron.} \\newcommand{\\MNRAS}{MNRAS} \\renewcommand{\\and}{{\\rm and }}  {We present echelle spectroscopic data for five eclipsing binary stars and two giant stars in the field of the open cluster NGC~2243. The average cluster velocity is determined to be $+60.4\\pm 0.6~km~s^{-1}$. Four of the eclipsing binaries are very likely members of the cluster based on their observed radial velocities. The absolute parameters of cluster member NV~CMa are determined by analysing photometric and radial velocity data. We obtain $1.089\\pm 0.010 M_{\\odot}$ and $1.221\\pm 0.031 R_{\\odot}$ for the primary, and $1.069\\pm 0.010 M_{\\odot}$ and $1.178\\pm 0.037 R_{\\odot}$ for the secondary. Both components of the binary are located on the Main Sequence, about 1~mag. below the turn-off point on the cluster color-magnitude diagram. Using model age-luminosity and age-radius relations we obtain $4.35\\pm0.25$~Gyr for the age of NV~CMa. The derived age is, however, very sensitive to the adopted metallicity of the cluster. We demonstrate that a meaningful determination of the ages of objects like NV~CMa based on evolutionary models is possible only if their metallicity is know with a relative accuracy of a few percent. The distance moduli calculated for the components of NV~CMa agree closely with each other, and imply an apparent distance modulus of the cluster of $(m-M)_{\\rm V}=13.25\\pm 0.08$. } {Stars: binaries: eclipsing, binaries -- stars: individual: NV CMa -- open clusters and associations: individual: NGC~2243} ",
        "introduction": "The field of the intermediate-age open cluster NGC~2243 contains 5 known detached eclipsing binaries. Four of these are likely cluster members based on their photometric properties (Kaluzny et~al. 2006; hereafter KKTS). The components of these systems should have the same age, metallicity and heliocentric distance and the determination of their absolute parameters can provide an interesting test of evolutionary models of low mass stars. Moreover, through the use of the surface brightness method, one may obtain a direct measure of the cluster distance.  This paper is focused on the determination of the absolute properties of NV~CMa, an eclipsing binary located in the central area of NGC~2243. In addition we use spectroscopic observations to check on the membership status of 4 other binaries in the cluster field. ",
        "conclusions": "The analysis of photometric and spectroscopic observations of the eclipsing binary NV~CMA has allowed us to determine the orbital parameters and physical properties of the component stars. Our determinations have formal uncertainties of 1\\% in the masses and 3\\% in radii. Uncomfortably large uncertainties in the radii result from the degeneracy of the light curve solution for this partially eclipsing system. Comparison with model tracks for $\\rm{[Fe/H]}=-0.525$ give an age of 4.1 to 4.6 Gyr for the binary and consequently for the cluster. This determination suffers from a substantial systematic error related to the uncertain metallicity of the cluster. For the relevant range of stellar masses and ages an uncertainty in the metallicity of 0.1~dex leads to an uncertainty in the estimated age of about 0.8~Gyr.  There is still room for obtaining a better determination of the parameters of NV~CMa, including its age. First of all, it is straightforward to derive masses of the components with an accuracy of 0.5\\% or even better. We used 12 spectra only and additional radial velocity data obtained near quadrature will lead to an improvement in the estimates of the masses of the components. It is also possible to obtain better light curves than these used in the present analysis; an improvement is possible in both the quality and the phase coverage. In particular, our photometry for the $B$ band had poor coverage inside the primary eclipse. Improved light curves would better constrain the parameters of the components obtained in the light curve solution. And last, but not least, the age determination of NV~CMa and of other cluster binaries would benefit enormously from accurate metallicity determinations, preferably based on high resolution spectroscopy of an extended sample of cluster member stars.  The analysis of single spectra obtained for four other eclipsing binaries in the cluster field indicates that three of them are radial velocity members of NGC~2243. The fourth star is definitely a non-member. Further observations of member binaries would allow a better constraint on the cluster age as well as a test of evolutionary models of low-mass stars with $\\rm{[Fe/H]}\\approx -0.5$. \\begin{figure}[htb] \\centerline{\\includegraphics[height=80mm,width=120mm]{fig1.eps}} \\FigCap{Spectroscopic observations and adopted orbit for NV~CMa.} \\end{figure} \\begin{figure}[htb] \\centerline{\\includegraphics[height=90mm,width=120mm]{fig2.eps}} \\FigCap{Broadening Function extracted from a spectrum of NV~CMa obtained at orbital phase 0.61 (solid line). The dashed line shows the fit of a model BF to the observed one.} \\end{figure} \\begin{figure}[htb] \\centerline{\\includegraphics[height=80mm,width=120mm]{fig3.eps}} \\FigCap{Phased $BV$ light curves of NV~CMa.} \\end{figure} \\begin{figure}[htb] \\centerline{\\includegraphics[height=80mm,width=120mm]{fig4.eps}} \\FigCap{Dependence of the luminosity ratio $(L_{1B}/L_{2B})$ and relative radius $<r_{2}>$ on the assumed radius $<r_{1}>$.} \\end{figure} \\begin{figure}[htb] \\centerline{\\includegraphics[height=50mm,width=120mm]{fig5.eps}} \\FigCap{The residuals for the fit corresponding to the light curve solution listed in Table~5.} \\end{figure} \\begin{figure}[htb] \\centerline{\\includegraphics[height=100mm,width=120mm]{fig6.eps}} \\FigCap{Color-magnitude diagram for the turnoff region of NGC~2243. The location of NV~CMa is marked with an open triangle. The squares denote the locations of individual components of the binary.} \\end{figure} \\begin{figure}[htb] \\centerline{\\includegraphics[height=90mm,width=120mm]{fig7.eps}} \\FigCap{Theoretical age-luminosity relations for masses $M_{1}=1.089\\pm 0.010 M_{\\odot}$ (lower) and $M_{2}=1.069\\pm 0.010 M_{\\odot}$ (upper). Horizontal dashed lines mark $1\\sigma$ ranges for the observed luminosities of the components of NV~CMa. The solid lines represent the mass $\\pm~1~\\sigma$.} \\end{figure} \\begin{figure}[htb] \\centerline{\\includegraphics[height=90mm,width=120mm]{fig8.eps}} \\FigCap{Theoretical age-radius relations for masses $M_{1}=1.089\\pm 0.010 M_{\\odot}$ (lower) and $M_{2}=1.069\\pm 0.010 M_{\\odot}$ (upper). Horizontal dashed lines mark $1\\sigma$ ranges for the observed radii of the components of NV~CMa. The solid lines represent the mass $\\pm~1~\\sigma$} \\end{figure} \\begin{figure}[htb] \\centerline{\\includegraphics[height=80mm,width=120mm]{fig9.eps}} \\FigCap{Theoretical age-luminosity and age-radius relations for a star with mass $M=1.089 M_{\\odot}$ and for metallicity ${\\rm [Fe/H]=-0.606}$ (upper curve), ${\\rm [Fe/H]=-0.525}$ (middle curve) and ${\\rm [Fe/H]=-0.397}$ (lower curve).} \\end{figure} \\Acknow{ JK and WP were supported by the grant 1~P03D~001~28 from the Ministry of Science and Information Society Technologies, Poland. Research of SMR is supported by a grant from the Natural Sciences and Engineering Council of Canada while IBT is supported by NSF grant AST-0507325. We thank Alexis Brandeker who obtained one of the spectra of NV~CMa used here. It is also a pleasure to thank Tomek Plewa for enlightening hints on MINGA usage. }"
    },
    "0609/astro-ph0609105_arXiv.txt": {
        "abstract": "A new numerical code, called SFUMATO, for solving self-gravitational magnetohydrodynamics (MHD) problems using adaptive mesh refinement (AMR) is presented. A block-structured grid is adopted as the grid of the AMR hierarchy. The total variation diminishing (TVD) cell-centered scheme is adopted as the MHD solver, with hyperbolic cleaning of the divergence error of the magnetic field also implemented. \\rev{The MHD solver exhibits a second-order accuracy in convergence tests of linearized MHD waves.} The self-gravity is solved using a multigrid method composed of (1) a full multigrid (FMG)-cycle on the AMR hierarchical grids, (2) a V-cycle on these grids, and (3) an FMG-cycle on the base grid.  The multigrid method exhibits spatial second-order accuracy, fast convergence, and scalability. The numerical fluxes are conserved by using a refluxing procedure in both the MHD solver and the multigrid method.   Several tests are performed and the results indicate that the solutions are consistent with previously published results. ",
        "introduction": "Protostellar collapse is one of the most important processes in star formation; it exhibits a high dynamic range in density and spatial dimensions.  Adaptive mesh refinement (AMR) is a powerful technique for performing numerical simulations requiring a high dynamic range in the mesh schemes.  Local high-resolution is realized by employing grids of differing resolutions.  The finer grids are inserted and their location changed according to given refinement criteria. Following the introduction of the fundamental principles of AMR by \\citet{Berger84} and \\citet{Berger89}, this technique is becoming widely used in astrophysical simulations.  In star formation, along with self-gravity, the magnetic field also plays an important role.  Both self-gravity and magnetohydrodynamics (MHD) are therefore implemented in recent AMR codes (see the review by \\cite{Klein06}).  Existing self-gravitational MHD AMR codes are generally based on Cartesian grids; although block-structured grids tend to be most often frequently, alternative approaches also exist, e.g. RAMSES \\citep{Fromang06}. In block-structured AMR, numerical cells are refined in the unit of the block, and a block is itself an ordinal uniform grid.  Block-structured grids are divided into two categories: patch-oriented grids and self-similar blocks.  In the former approach, a block has variable number of cells, and the shape of the block is also changed in the course of refinement.  A large block containing many cells and a small block containing just a few cells can co-exist.  This approach originated in the study of \\citet{Berger89}, and has been adopted in many codes: e.g., the AMR code \\rev{Orion} of Berkeley \\citep{Truelove98}, Enzo \\citep{Norman99}, and RIEMANN \\citep{Balsara01}.  In the latter approach, all blocks contain the same number of cells, but the physical sizes of the blocks are different.  This approach has also been adopted by many codes: e.g., FLASH \\citep{Fryxell00,Banerjee06}, and NIRVANA \\citep{Zieglar05}. This approach offers advantages including relatively simple algorithms for refinement, parallelization, and vectorization.  This paper also follows this latter approach.  In particular, this approach enables the AMR code to implement a full multigrid (FMG)-cycle in the multigrid method for self-gravity.  This scheme is an extension of the multigrid method for the nested grid approach \\citep{Matsumoto03} to the AMR grid.  There are several ways to treat the MHD, particularly in the treatment of the $\\nabla \\cdot \\bmath{B}$ term.  These approaches may be categorized as (1) the constrained transport method with a staggered grid, (2) the projection scheme with a Poisson solver, and (3) the eight-wave formulation (see \\cite{Toth00,Balsara04} for a comparison of these schemes).  Recently, \\citet{Dedner02} proposed an alternative scheme for cleaning the $\\nabla \\cdot \\bmath{B}$ term.  This scheme has a formulation that is similar to, but distinct from, the eight-wave formulation; two additional waves transfer the $\\nabla \\cdot \\bmath{B}$ error isotropically at a given speed, and the waves decay at a given rate. As a result, the error is propagated independently of the gas motion, and so is diluted. By contrast, using the eight-wave formulation, a single additional wave transfers the error at the gas velocity.  The error is always propagated downstream of the gas motion, and can become stagnated at the shock wave and the center of the collapsing cloud. The scheme of \\citet{Dedner02} is therefore adopted here.  The author has previously developed a self-gravitational MHD code with a nested grid \\citep{Matsumoto03,Matsumoto03b,Matsumoto04,Machida04}.  In the nested grid approach, the grids are refined adaptively, while the sizes and positions of the sub-grids are fixed. The problems to which this method may be applied have therefore been restricted to those problems whose objects that require fine grid resolution are located in the center of the computational domain, e.g., \\citet{Machida06}.  A self-gravitational MHD code with AMR has now, however, been developed, and this method is presented in this paper. In \\S~\\ref{sec:equations}--\\ref{sec:refine}, the governing equations to be solved, the rules of discretization, and the refinement algorithm are presented. In \\S~\\ref{sec:HD}, the present methods for treating the hydrodynamics and MHD of the problem are described, while in \\S~\\ref{sec:MG}, the methods of the multigrid iteration for solving self-gravity are described. In \\S~\\ref{sec:test}, the results of several numerical tests are presented. The paper is concluded in \\S~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  SFUMATO, a self-gravitational MHD code applying the AMR technique is presented. The grid is configured in a block structure.  The MHD scheme is implemented so that it is of second-order spatial accuracy by means of the TVD approach.  The upwind numerical flux is obtained by the linearized Riemann solver.  The scheme is fully cell-centered, and the divergence error of the magnetic fields is cleaned. \\rev{The convergence tests of the linearized MHD waves exhibit a second-order accuracy, and the decay rates shows a third-order accuracy.  Stiffness of the scheme is confirmed by the MHD flux tube problems.}  The self-gravity is solved by a multigrid method composed of: (1) a FMG-cycle on the AMR hierarchical grids, (2) a V-cycle on these grids, and (3) a FMG-cycle on the base grid.  The FMG-cycle on the AMR hierarchical grids enables a scalable dependence on the number of cells. The multigrid method ensures that the solution converges rapidly; the residual is reduced by a factor of $10^{-3}-10^{-2}$ every iteration.  The multigrid method exhibits a second-order spatial accuracy.  Moreover, no spurious features appear at the interfaces between fine and coarse grid-levels, due to the flux conservation at the interface.   The MHD scheme and multigrid method are combined so as to have second-order temporal accuracy.  The time-marching has two modes: synchronous and adaptive time-step modes. The former mode is adopted for problems including self-gravity, while the latter mode is used for all other problems.  The AMR code was tested by considering test problems given by double Mach reflection, the MHD rotor, fragmentation of an isothermal cloud, and outflow formation in a collapsing cloud.  The present results are in good agreement with those of previous studies conducted by the present author as well as other authors.   \\bigskip  The author would like to thank K. Tomisaka and T. Hanawa for valuable discussions. The author thanks the anonymous referee for his/her valuable comments, which have been very helpful in improving the paper. Numerical computations were carried out on VPP5000 at the Center for Computational Astrophysics, CfCA, of the National Astronomical Observatory of Japan, and SX6 at Japan Aerospace Exploration Agency, JAXA. This research was supported in part by the Hosei Society of Humanity and Environment, and by Grants-in-Aid for Young Scientists (B) 16740115, 18740113, for Scientific Research (C) 17540212, and for Scientific Research (B) 17340059 awarded by the Ministry of Education, Culture, Sports, Science and Technology, Japan."
    },
    "0609/astro-ph0609333_arXiv.txt": {
        "abstract": "The massive-star IMF is found to be invariable. However, integrated IMFs probably depend on galactic mass. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609619_arXiv.txt": {
        "abstract": "We report on the identification of the lens responsible for microlensing event MACHO-LMC-20. As part of a \\textit{Spitzer}/IRAC program conducting mid-infrared follow-up of the MACHO Large Magellanic Cloud microlensing fields, we discovered a significant flux excess at the position of the source star for this event. These data, in combination with high resolution near-infrared \\textit{Magellan}/PANIC data has allowed us to classify the lens as an early M dwarf in the thick disk of the Milky Way, at a distance of $\\sim 2$ kpc. This is only the second microlens to have been identified, the first also being a M dwarf star in the disk. Together, these two events are still consistent with the expected frequency of nearby stars in the Milky Way thin and thick disks acting as lenses. ",
        "introduction": "We have been conducting near and mid-infrared follow-up of the MACHO (Massive Compact Halo Object) Large Magellanic Cloud (LMC) microlensing event fields to try to identify the lensing population. Identification of the lenses can resolve questions regarding their locations, and can support a microlensing interpretation of the events. Deep photometry using \\textit{Spitzer's} Infrared Array Camera (IRAC) of all the lensed LMC stars reported by MACHO can be used to test the hypothesis that the Milky Way has a thin or thick disk population of cool stars that are responsible for the lensing.   Such a population would be readily detectable by \\textit{Spitzer}, but because IRAC has a spatial resolution of $\\sim1.7\\arcsec$ (FWHM), not enough time has past since the original microlensing event to allow the source and lens stars to separate via proper motions. However, the presence of a cool lens can be inferred by way of an infrared (IR) excess (Von Hippel \\etal 2003). While this technique requires a number of assumptions to be made about the lens and the source star, it has had some success to date. In the case of MACHO-LMC-5 (hereafter Event 5), the lens properties inferred from the flux excess technique using IRAC (Nguyen \\etal 2004) were strikingly consistent with those obtained by direct high-resolution \\textit{HST} imaging (Alcock \\etal 2001b; Gould 2004; Drake \\etal 2004). The Event 5 lens was found to be a M5 dwarf in the disk of the Milky Way, and was the only microlens to have been identified to date. On its own it did not shed any light on the microlensing population in general.    In this letter we report the identification of one more lens from the MACHO sample, that for MACHO-LMC-20 (hereafter ``Event 20''; Alcock \\etal 2000). Event 20 was a low-magnification event ($A_{\\rm max}=2.95$) that occurred in November, 1997. It only passed the looser `criteria B' MACHO cuts which allowed events of lower signal-to-noise than the stricter `criteria A' ones (Alcock \\etal 2000). The event duration was 72.7 days, which for a typical halo model, places the lens mass at approximately $0.5M_\\sun$ (Alcock \\etal 1993). The baseline (unmagnified) $V$ magnitude of the source$+$lens system was 21.35, with $V-R=0.57$. IRAC detected a clear IR excess at the source location in both the 3.6 and 4.5$\\micron$ bands relative to the expected flux in these bands for a star of the above-mentioned magnitude and color ($\\sim0.6$ mag of excess flux), suggesting blending with a foreground cool star. This data, in combination with subsequent ground-based follow-up using the PANIC infrared imager on Magellan, has enabled us to infer that the lens is an early M dwarf approximately $2$ kpc away in the Milky Way thick disk. Together with Event 5, the two events are consistent with the number of events expected from the Milky Way disk, and thus do not signify any anomalous structure in the disk. \\S~2 briefly describes the IRAC data, \\S~3 describes the PANIC observations and analysis, and \\S~4 presents the results and a discussion of them. ",
        "conclusions": "The suggestion of microlensing by the Galactic dark halo (Paczynski 1986) was followed-up by many teams. While the MACHO collaboration reported 16 microlensing events towards the LMC (Alcock \\etal 1993, 1997, 2000), the EROS collaboration has found only 3 (Aubourg \\etal 1993; Lasserre \\etal 2000). The MACHO efficiency analysis indicates a dark halo with a MACHO fraction of 20\\% which corresponds to an optical depth of $\\tau=1.2^{+0.4}_{-0.3} \\times 10^{-7}$ due to lenses of $\\sim0.5M_{\\odot}$. This rate is \\textit{significantly} higher than what was expected from known Galactic and LMC stellar populations; the latter have $\\tau$ between $0.24 \\times 10^{-7}$ and $0.36 \\times 10^{-7}$ (Alcock \\etal 2000; Griest \\& Thomas 2005). The upper limit on the EROS microlensing optical depth is only barely consistent with that of MACHO. Recently, the EROS-2 project reported their results, and combined with the results of EROS-1, they report an optical depth toward the LMC of $\\tau<0.36\\times10^{-7}$ (95\\% confidence) due to lenses of $\\sim0.4M_{\\odot}$, corresponding to a halo mass fraction of less than 7\\% (Tisserand \\etal 2006).     \\begin{figure} \\begin{center} \\includegraphics[angle=0,scale=0.42]{jhkcolorplot.ps} \\caption{A plot in $J-H$ vs.~$H-K_s$--space showing the intrinsic relation for dwarfs (solid line) and giants (dashed line). The filled circle with error bars shows our PANIC photometry for the Event 20 system. } \\end{center} \\end{figure}   Several groups have investigated the excess of microlensing towards the LMC seen by MACHO, attributing it to as yet undetected structure in the Milky Way (Rahvar 2005; Nguyen \\etal 2004; Gates \\& Gyuk 2001; Evans \\etal 1998; Zhao 1998; Zaritsky \\& Lin 1997) or in the LMC itself (Wu 1994, Sahu 1994; Alves \\& Nelson 2000; Evans \\& Kerins 2000; Di Stefano 2000; Gyuk \\etal 2000; Zhao \\etal 2003; Mancini \\etal 2004; Alves 2004; Nikolaev \\etal 2004). However, over and above the question of the population responsible for the lensing, the discrepancy between the MACHO and EROS results have also cast doubt on whether the MACHO events are microlensing at all, or just contamination from variable stars and background supernovae. Belokurov \\etal (2003, 2004) applied a neural network analysis to the MACHO data in an attempt to develop independent selection criteria for microlensing. They selected only 10 of the candidates that MACHO picked out as bona fide microlensing. Their verdict on the microlensing events as well as the original MACHO verdict are presented in Table~4 of Bennett \\etal (2005).   \\begin{figure*} \\begin{center} \\includegraphics[angle=0,scale=0.65]{galaxyplot.ps} \\caption{(\\textit{Left}) A plot showing $J-[3.6]$ vs.~$[3.6]-[4.5]$ for stars and galaxies in the First Look Survey. Late-type M dwarfs from the Patten \\etal (2006) sample are overlaid using triangles, L dwarfs using diamonds and T dwarfs using squares. The filled circle with error bars shows are PANIC$+$IRAC photometry for the Event 20 system. (\\textit{Right}) The same objects are shown in $K-[4.5]$ vs.~$[3.6]-[4.5]$-space} \\end{center} \\end{figure*}  \\begin{figure} \\begin{center} \\includegraphics[angle=0,scale=0.45]{absmrel.ps} \\caption{A plot showing $M_J$ vs.~$J-[3.6]$ for M dwarfs (and 2 K7 dwarfs) from the Patten \\etal (2006) sample (filled circles), and some earlier type M dwarfs from M. Schuster (private communication) (diamonds).  The square with error bars shows the lens candidate, which is consistent with where early M dwarfs lie on this relation. } \\end{center} \\end{figure}   In the case of the event at hand, Belokurov \\etal (2003) used their neural network to classify it as a supernova (SN) rather than microlensing. However, Bennett \\etal (2005) pointed out that Event 20 is a good disk lens candidate because Figure~7 of Alcock \\etal (2000) shows that, like Event 5, Event 20 is quite red for its brightness, i.e. it appears to be blended with a star that is much redder than other LMC stars of similar magnitude.  Our data allows us to weigh into this debate. Figure~2 is a $J-H$ vs.~$H-K_s$ plot, showing the intrinsic relation for dwarfs (solid line) and giants (dashed line; dwarf data from Leggett 1992 and Kenyon \\& Hartmann 1995; giant data from Bessell \\& Brett 1988) . The filled circle with error bars shows our PANIC result for the lens, which appears to be consistent with an early-M dwarf star. The mean extinction at the location of Event 20 in the LMC is $A_V \\sim 0.5$ mag (Zaritsky \\etal 2004). This implies that, even if our M dwarf is at the distance of the LMC (in which case we certainly would not see it with either PANIC or IRAC), the mean extinction in $K$-band would be $A_K \\sim 0.05$. This is within our photometric errors, and thus we do not account for reddening. Note, however, that in this and all following analysis plots, the lens object is still blended with the LMC source star. At these infrared wavelengths we expect that the contamination from the source flux is minimal ($\\sim 10\\%$). The data thus seem to corroborate Bennett's prediction of this event has lensing by a low-mass star.   However, as suggested by Belokurov \\etal (2003), we also consider the possibility that the magnification in the MACHO light curve was caused by a background SN. If this were the case, at $\\sim9$ years since the event we would now be seeing the host galaxy. In Figure~3 we address this by showing the near-IR color-color space of objects in the \\textit{Spitzer}/IRAC First Look Survey (Lacy \\etal 2005) for which 2MASS $J$ and $K$ photometry is available. The panel on the left shows $J-[3.6]$ vs.~$[3.6]-[4.5]$. The points are stars and galaxies, with the blue end of the plot comprising mostly main sequence stars. The galaxies tend to be redder and extend into the diffuse plume at the red end of the $y$-axis. Overlaid are late-M dwarfs (spectral type M6 and later) (triangles), L dwarfs (diamonds) and T-dwarfs (squares) from Patten \\etal (2006). The filled circle with error bars shows the IRAC$+$ PANIC photometry for the lens. The lens lies in a part of the color-color space that is more consistent with stars than with galaxies, and within the large errors propagated through from the IRAC photometry, it is consistent with an early M dwarf. The panel on the right shows $K-[4.5]$ vs.~$[3.6]-[4.5]$. Again, the points at the blue end of the plot are mostly stars with galaxies extending into a red tail. In this color-space as well the data suggest that the lens is a M dwarf and not a background galaxy. The placement of the lens could be consistent with a blue galaxy, however, at the $JHK_s$ magnitude of the lens, we would expect that galaxies are discernible from stars in the high resolution PANIC images.   Thus our follow-up data confirm not only the microlensing nature of this event but also that it was caused by a low-mass star. Finally, Mancini \\etal (2004) have made predictions, based on detailed models of the LMC, about which of the MACHO events were likely due to lenses in the LMC itself. They picked events 6, 8, 13 \\& 14 as most likely due to self-lensing because of their locations in the LMC. Event 20 was not picked as a high probability self-lensing candidate. The Mancini \\etal (2004) predictions are tabulated and compared against the MACHO and Belokurov \\etal (2004) predictions in Bennett \\etal (2005).   Given our measurement of the spectral type of the lens, we estimate its distance from a sample of early M dwarfs using 2MASS $JHK_s$ and [3.6] photometry provided by M. Schuster (private communication). If the M dwarf is in the range M0--3.5, then its absolute magnitude would be in the range $M_J = 7\\pm0.5, M_H=7\\pm0.5, M_{K_s}=6\\pm0.4,$ and $M_{[3.6]}=6\\pm0.3$. In addition to this line of argument, in Figure~4 we show $M_J$ versus $J$-[3.6] for these M dwarfs (plus two K7 dwarfs and additional late-type M dwarfs from M. Schuster) (triangles) and for some late-type M dwarfs from Patten \\etal (2006) (filled circles). The square with error bars is the lens and it is clear that its $J$-[3.6] color is consistent with the absolute magnitude expected for an early M dwarf ($M_J\\sim7$). The range of absolute magnitudes quoted above imply a distance to the lens of $2\\pm0.7$ kpc where the error includes contributions from the range in spectral types and the photometric error in the $J$-[3.6] color. The lens is thus in the thick disk of the Milky Way\\footnote{Events 20 \\& 5 are evidence that nearby dim stars can be studied through such monitoring programs (``mesolensing'', Di Stefano 2005).}.  This is the second Milky Way disk lens in the MACHO sample, the first being that of Event 5, which was closer at $\\sim600$ pc. A total of 0.75 events were expected from the Milky Way thin and thick disks (Bennett \\etal 2005), and thus these two disk lenses are still consistent with the expected number. Event 20 did not pass the MACHO criteria set `A' and so does not affect the conclusion to the halo fraction made from it.  In a future paper we will present the full consensus of the near-IR follow-up of the MACHO LMC microlensing fields. The main result of this project thus far is that two of the microlensing events were caused by foreground disk lenses. This is consistent with what was expected from the Milky Way thin and thick disks."
    },
    "0609/astro-ph0609043_arXiv.txt": {
        "abstract": "So far, solar-like oscillations have been studied using radial velocity and/or light curve variations, which reveal frequencies of the oscillation modes. Line-profile variations, however, are also a valuable diagnostic to characterise radial and non-radial oscillations, including frequencies, amplitudes, the spherical mode wavenumbers $(\\ell,m)$ and the stellar inclination angle.  Here we present a line profile analysis of $\\epsilon$ Ophiuchi, which is a pulsating red giant.  The main differences compared to previous line profile analyses done for heat-driven oscillations are the small amplitudes and the predicted short damping and re-excitation times in red giants.  Two line diagnostics have been tested to see whether these are sensitive to the small line profile variations present in red giants. In addition, line profiles have been simulated with short damping and re-excitation times and are compared with the observations.  This comparison reveals that non-radial modes are detected in the observed line profile variations of $\\epsilon$ Ophiuchi. This is rather surprising, as theoretical predictions favours the occurrence of radial modes. ",
        "introduction": "The discovery of oscillations in the sun, about five decades ago, was the onset of seismology of stars. Oscillations probe the interior of the stars and are therefore a direct means to determine their internal structure. Solar-like oscillations are excited by turbulent convection near the surface of cool stars of spectral type F, G, K or M. They show radial velocity variations with amplitudes of typically a few cm\\,s$^{-1}$ to a few m\\,s$^{-1}$, and with periods ranging from a few minutes for main sequence stars to about half an hour for subgiants and a couple of hours for giants.  During the last decade, the refinement of the techniques to perform very accurate radial velocity observations, made it also possible to observe solar-like oscillations in distant stars.  Recently, these solar-like oscillations are indeed observed in several red giants, for example in $\\xi$ Hydrae \\cite{frandsen2002}, $\\epsilon$ Ophiuchi \\cite{deridder2006b} and $\\eta$ Serpentis \\cite{carrier2006}.  This proceedingspaper describes observations as well as a frequency analysis of different parameters and a line profile analysis of $\\epsilon$ Ophiuchi. It is organised as follows. In Section 2, we start with a discussion on the target selection, followed in Section 3 by the description of the observations. In Section 4 the data analysis is described, i.e. first, the frequency determination of different oscillation diagnostics and subsequently the mode identification using these frequencies.  In Section 5 we show simulations, including the short damping and re-excitation times present in solar-like oscillations of red giants \\cite{stello2004}.  Finally, in Section 6 we draw some conclusions. ",
        "conclusions": "We have found a different amplitude distribution across the cross correlation profile for the dominant frequencies of the pulsating red giant star $\\epsilon$ Ophiuchi.  This clearly indicates that modes with different wavenumbers ($\\ell,m$) must be present in the data. Moreover, the amplitude distribution of the dominant mode is clearly different from the one due to $m=0$ modes.  Thus, we have detected non-radial modes in the observed cross correlation profiles of this star.  We compared the observed amplitude distributions of $\\epsilon$ Ophiuchi with the one of simulated cross correlation profiles in an attempt to identify the wavenumbers of the individual modes.  Due to the fact that single realisations of different modes are in some cases comparable, e.g.\\ for $(\\ell,m)=(1,1)$ and $(2,1)$ (see Figures\\,\\ref{amplphaseveleta2i55} and \\ref{amplphaseveleta2i75}), different sets of wavenumbers are equally likely for observed modes.  It is therefore clearly not yet possible to identify the individual modes unambiguously with the method used in the present work, but clear evidence for modes with $m \\neq 0$ values is found.  \\cite{hekker2006} also analysed the line profile variations of three additional pulsating giants and came to similar conclusion than for $\\epsilon$ Ophiuchi for two of these three stars. Moreover, non-radial mixed oscillation modes were recently also found in a red giant observed with the MOST space mission (\\cite{matthews2006}, these proceedings). It therefore seems that red giant stars pulsate in both radial and non-radial modes. These observational results constitute a fruitful starting point to fine-tune theoretical predictions of stochastically excited modes in evolved stars."
    },
    "0609/astro-ph0609569_arXiv.txt": {
        "abstract": "{The luminosity function (LF) is a basic tool in the study of galaxy evolution since it constrains galaxy formation models. The earliest LF estimates in the IR and far-IR spectral ranges seem to suggest strong evolution. Deeper samples are needed to confirm these predictions. We have a useful  IR  dataset, which provides a direct link between IRAS and ISO surveys, and the forthcoming deeper Spitzer Space Telescope and  Akari cosmological surveys, to address this issue. } { This allows us to derive the 60\\,$\\mu$m local LF to sensitivity levels 10 times deeper than before, to investigate evolutionary effects up to a redshift of 0.37,  and, using the  60/15$\\mu$m bi-variate method,  the poorly known 15\\,$\\mu$m local LF of galaxies. } { We exploited our ISOCAM observations of the IRAS Deep Survey (IDS) fields (Hacking \\& Houck 1987), to correct the $60\\,\\mu$m fluxes for confusion effects and observational biases.  We find indications of a significant incompleteness of the IDS sample, still one of the deepest far-IR selected galaxy samples, below $\\simeq 80\\,$mJy (Mazzei et al. 2001). We have reliable identifications and spectroscopic redshifts for 100\\% of  a complete subsample comprising 56 sources with $S(60\\mu\\hbox{m})> 80\\,$mJy. }{ With our spectroscopic coverage we construct the 60\\,$\\mu$m LF for a sample complete down to 80 mJy. This LF extends over three orders of magnitude in luminosity, from 9 up to more than 12 in $\\log(L_{60}/L_{\\odot})$. Despite the fact that the redshift range of our sample exceeds $z=0.3$, the $V/V_{max}$ test  gives $<V/V_{max}>=0.51\\pm 0.06$, consistent with a uniform distribution of sources. A more direct test, whereby the LF was measured in each of four different redshift intervals, does not point out any  signature of evolution. On the other hand, the rest--frame 15$\\mu$m local LF we derive, extends up to $\\log(L_{15}/L_{\\odot})=12$ and predicts 10 times more sources at $\\log(L_{15}/L_{\\odot})=11$ than are seen by  Pozzi et al. (2004). } {} ",
        "introduction": "Understanding how  galaxies form and evolve is a key goal of  modern physical cosmology.  A fundamental observable of galaxies is their luminosity function (LF) which  has long been used to constrain galaxy formation models and to quantify star formation  and evolution both in luminosity and in density. The IR and far-IR  spectral ranges are the best to deepen our knowledge  of this subject since  they trace the star formation that is responsible for galaxy formation. In particular several satellite missions, in the past (IRAS and ISO), and in the present (Spitzer, Akari), provided and will provide  data  which will be complementary for a  detailed study of the LFs in such  spectral domains.  The earliest IR estimate of the  LFs, derived from the IRAS data \\citep{RRW87, Sau00}, indicated   strong evolution, so that LF increases with redshift $\\propto (1+z)^{3 \\pm 1}$. Moreover deep surveys at 15$\\mu$m carried out with ISO \\citep[i.e.,][]{Elbazetal99, Flores99, Larietal01, Metcalfe2003} seem to require, indeed,  strong evolution of 15$\\mu$m sources (Lagache et al. (2005) and references therein). Several evolutionary models were developed to explain these results, \\citep[e.g.,][]{Fetal01,RR01} and fit the IR/submillimiter source counts with different degrees of success. Nevertheless none of them is based on a local LF obtained from 15$\\mu$m data, since the only available data until  recently came from IRAS 12$\\mu$m photometry \\citep{Rush93, xu98, Fang98}. The first attempt to build up the 15$\\mu$m LF of a NEPR subsample was made by Xu (2000). However  he said that it must be considered as a preliminary work because: i) the sample of galaxies used is a incomplete sample, ii) there is a possible misidentification between the sources in the 60$\\mu$m redshift survey of Ashby et al. (1996) and the 15$\\mu$m sources in his work  (see della Valle et al. (2006), for more details), iii) the model used to interpret the data treated all IR galaxies as a single population.  Our IDS/ISOCAM sample overcomes all these issues. It  comprises a complete, 60$\\mu$m  selected sample of 56 galaxies in the North Ecliptic Polar Region (NEPR), a subsample of the original  98 IRAS Deep Survey (IDS) fields \\cite{HH87}.  The IRAS Deep Survey (IDS) sample was defined by co-adding IRAS scans of the North Ecliptic Polar Region (NEPR), representing more than 20 hours of integration time \\cite{HH87}.  It comprises 98 sources with $S(60\\mu\\hbox{m})> 50\\,$mJy over an area of 6.25 square degrees.  Mazzei et al. (2001) exploited  ISOCAM observations (range 12-18$\\mu$m) of 94 IRAS Deep Survey (IDS) fields \\citep{Aus99}, centered on the nominal positions of IDS sources, to correct the $60\\,\\mu$m fluxes for confusion effects and observational biases,   finding indications of a significant incompleteness of the IDS sample below $\\simeq 80\\,$mJy. In della Valle et al (2006) we presented spectroscopic and optical observations of candidate identifications  of our ISOCAM sources. Combining such observations with those by Ashby et al. (1996), we have reliable identifications and spectroscopic redshifts for 100\\% of the complete subsample comprising 56 sources with $S(60\\mu\\hbox{m})> 80\\,$mJy. It is the deepest complete IRAS selected sample available and still one of the deepest complete far-IR selected samples. For comparison, the IRAS {\\it Point Source Catalog} \\citep[hereafter PSC;][]{Beichman1988}, comprises about $250,000$ IR sources with a completeness limit of 0.5 Jy at 60$\\mu$m  \\citep{Soi87}. The deep ISOPHOT surveys, FIRBACK at 170$\\mu$m, (Puget et al. 1999; Dole et al. 2001), and ELAIS at 90$\\mu$m \\citep{Oliver00} are all complete down to about 100\\,mJy. Moreover, the $70\\,\\mu$m Spitzer catalog of the 8.75 sq. deg. Bootes field is flux limited to 80 mJy (Dole et al. 2004) and the Spitzer extragalactic ``main\" First Look Survey, covering about 4 sq. deg., is complete to about 20 mJy at $70\\,\\mu$m (cf. Fig. 2 of Frayer et al. 2006), but redshift measurements are available for a substantial fraction of sources (yet only 72\\%) merely for $S_{70\\mu{\\rm m}} > 50\\,$mJy.  Thanks to our ISOCAM and optical/near-IR observations  our sample, which provides a direct link between IRAS and ISO surveys, and the forthcoming deeper Spitzer Space Telescope and  Akari cosmological surveys \\footnote{New deep observations of the NEPR are planned with the Akari space mission, also known as the InfraRed Imaging Surveyor (IRIS). It will map the entire sky in four far-IR bands, from 50 to 200$\\mu$m, and two mid-IR bands, at 9 and 20$\\mu$m, with far-IR angular resolutions of 25--45 arcsec, reaching a detection limit of 44 mJy ($5 \\sigma$ sensitivity) in the 50--75$\\mu$m band \\citep{Pearson2004}. With the Akari orbit, the integration time on the NEPR will be particularly high, and, correspondingly, the detection limit significantly deeper than average.}, is one of the far-IR selected complete samples with the larger spectral coverage. In addition to the ISOCAM and to the 60$\\mu$m fluxes, most of them $\\simeq$70\\%   have 100$\\mu$m fluxes, the remaining 30\\% with upper limits, and several ($\\simeq$40\\%) have 25$\\mu$m fluxes from IRAS \\cite{Maetal01}. Optical imaging has been already performed for 62.5\\% out of such a sample  in at least one band, B or R, and for 34\\% in both the bands,  moreover Two Micron All Sky Survey (2MASS) data are available for 68\\% of the sample and VLA observations for a large fraction of these sources are also available \\citep{H89}.  In this paper, which is the  second step of  our multi-wavelength approach devoted to study  the evolution of a far-IR selected sample of galaxies on which numerous studies of the far-IR evolution of galaxies still rely, we derive  the 60$\\mu$m luminosity function (LF) of such a sample. Our sample, ten times deeper in flux density than the PSC catalog, thus less liable to the effect of local density inhomogeneity, allows us to investigate evolutionary effects up to a redshift of 0.37,  five times deeper than the PSCz catalog  \\cite[z$\\le$ 0.07,][]{Sau00}. Moreover, we use the bi-variate method to translate the 60$\\mu$m  LF to the poorly known  15$\\mu$m LF. We will compare our results with the  recent determination of the 15$\\mu$m local LF obtained by  Pozzi et al. (2004)  using the available data on the southern  fields, S1 and S2, of the ELAIS survey \\cite{Oliver00}.  The plan of the paper is the following: Section 2 focuses on far-IR properties of our complete sample, Section 3 shows our derived 60\\,$\\mu$m LF, Section 4 presents the 15\\,$\\mu$m LF based on the bi-variate 60/15\\,$\\mu$m method. In Section 5 there are our conclusions.\\\\ Here and in the following we adopt: $\\Lambda$=0.7, $\\Omega_{b}$=0.3, $H_0$=70 km/s/Mpc. ",
        "conclusions": "Combining our observations \\cite{DV05} with those  by Ashby et al. (1996) we have reliable identifications and spectroscopic redshifts for 100\\% of the complete  far--IR selected subsample comprising 56 IDS sources with $S_{60}>80$mJy \\cite{Maetal01}. The  redshift distribution shows a tail extending up to $z\\simeq 0.37$, in particular  $\\simeq$ 26\\% of the sources have redshifts  $z>0.1$.   To fully exploit the potential of this sample, ten times deeper than the IRAS PSC, thus less liable to the effect of local density inhomogeneity, for investigating galaxy evolution, we calculate the 60$\\mu$m LF using the $1/V_{max}$ method. Current estimates are based on rather shallow samples. Even though our  sample is five times deeper in redshift than the PSCz  \\cite{Sau00} used by Takeuchi et al. (2003), our LF agree with their determination \\cite{tsu, tsu04} and with that by Frayer et al. (2006) based on a complete sample of galaxies with $S_{70} > 50$~mJy drawn from the Spitzer {\\it extragalactic First Look Survey}. Despite the fact that our redshift range exceeds $z=0.3$, and our 60$\\mu$m LF   extends up to $\\log L_{60}=12$, whereas  that by Frayer et al.(2006) up to $\\log L_{60}\\simeq 11$,  it does not show any evidence of  evolution. The $V/V_{max}$ test gives a value consistent with an uniform distribution ($V/V_{max} = 0.51 \\pm 0.06$). Moreover, a more direct test for evolution has been performed by splitting the LF in different redshift bins assuming no evolution.  The results   are fully consistent with this assumption.  We present the bi-variate 15$\\mu$m LF,  one of the few determinations based on ISO data, by convolving the 60$\\mu$m LF with the luminosity ratio distribution, $L_{15}/L_{60}$ of our sample. This extends to luminosity 100 times higher than before. Our result agrees with the recent determination by Pozzi et al. (2004) in the common range of luminosity, i.e., from 7.8  up to  10.6 in $\\log(L_{15}/L_{\\odot})$. However, above $\\log(L_{15}/L_{\\odot})=10.6$, the two parametric LFs diverge so that at  $\\log(L_{15}/L_{\\odot})=11$  we expect  10 times more sources than Pozzi et al. (2004).  In order to investigate  the role of galaxy populations on such a result, we disentangle starbursts and  spirals on the basis both of their far-IR dust temperature, and of their $L_{15}/L_{R}$ ratios, as assumed by Pozzi et al. (2004). Such criteria select galaxies with different dust properties, as we point out in \\S$\\,$\\ref{sbssect}. Nevertheless, we find that both galaxy populations contribute to the faint end of the rest-frame 15$\\mu$m LF, even though we adopt the same criterion as Pozzi et al. (2004). Moreover, in this case, above $\\log(L_{15}/L_{\\odot})=9$ our findings are that starbursts and spirals give almost the same contribution to the mid--IR  LF whereas Pozzi et al. (2004) predict that  spirals contribute 5-10 times more than  starbursts."
    },
    "0609/astro-ph0609275_arXiv.txt": {
        "abstract": "The Magnificent Seven have all been discovered by their exceptional soft X-ray spectra and high ratios of X-ray to optical flux. They all are considered to be nearby sources. Searching for similar objects with larger distances, one expects larger interstellar absorption resulting in harder X-ray counterparts. Current interstellar absorption treatment depends on chosen abundances and scattering cross-sections of the elements as well as on the 3D distribution of the interstellar medium. After a discussion of these factors we use the comprehensive 3D measurements of the Local Bubble by \\cite{Lall03} to construct two simple models of the 3D distribution of the hydrogen column density. We test these models by using a set of soft X-ray sources with known distances. Finally, we discuss possible applications for distance estimations and population synthesis studies. ",
        "introduction": "\\label{intro} The Magnificent Seven (M7), as the ROSAT-discovered X-ray thermal isolated neutron stars are sometimes dubbed, are exceptional because of their soft blackbody-like spectra without non-thermal components; for reviews see e.g. Haberl 2006 (this volume) or \\cite{Haberl2004b}. One of them has the second-best blackbody spectrum after the cosmic background radiation and this lead to one of the best neutron star radius estimates known \\citep{Trumper04}. Thought to be nearby, the M7 represent nearly half of the local young neutron star population. Despite several searches for new candidates of this radio-quiet neutron star class (e.g. \\citealt{Marcel2006,Chieregato2005,Rutledge2003}), no new M7-like object has been confirmed up to now. One aggravating circumstance these searches face is the hardening of the X-ray spectrum due to the interstellar absorption. Caused mainly by photo-electric absorption of photons by heavy elements, the interstellar absorption is usually described by the equivalent hydrogen column density assuming certain chemical abundances. Highly important is the clumpiness of the interstellar medium  (ISM) causing (un)favourable lines of sights to search for new M7-like objects. The knowledge about the local ISM structure is also important in connection with the long-debated issue of  isolated neutron stars accreting from the ISM. \\vspace{-0.7cm} \\subsection{Abundances and cross-sections} \\vspace{-0.3cm} \\label{abundances} Most absorption models today presume element abundances independent of the line of sight. It was outlined by \\cite{Wilms2000} that the total gas plus dust ISM abundances are lower than the local -- Solar -- abundances and the general ISM abundance uncertainties are still in the order of 0.1~dex because the measurements are very difficult. One can choose between several different abundance tables implemented e.g. in $XSPEC$ for analysing X-ray spectra. We note here only two examples commonly used -- the abundance tables by  \\citet{Anders1989} ('angr' in $XSPEC$) and by \\cite{Wilms2000} ('wilm' in $XSPEC$) considering newer measurements by e.g. \\citet{Snow1996,Cardelli1996,Meyer97,Meyer98}.\\\\ The individual abundances and photoionization cross-sections of the elements and their ions as well as dust grain properties (e.g. sizes, shapes) influence the total photo-electric absorption of X-rays. Again, we mention here only two examplary works used very often for analysis of X-ray spectra. \\cite{MM83} did a polynomial fit of the effective absorption cross-section per hydrogen atom based on measured atomic absorption cross-sections by \\cite{Henke82} and abundances mostly from \\cite{Anders82}. With the exception of oxygen, \\cite{MM83} assume the elements to be either entirely in the gas phase or completely depleted into dust grains. The recent improved absorption treatment by \\cite{Wilms2000} takes into account newer element abundances as noted above and more recent photoionization cross-section calculations by e.g. \\cite{Verner93,Verner95}. \\cite{Wilms2000} additionally include an improved molecular cross-section for $H_2$. Besides this updated database \\cite{Wilms2000} considered furthermore a simple spherical composite dust grain model in an MRN distribution \\citep{MathisRN77}. \\\\ Both models are implemented in $XSPEC$ and can be used via the routines '$wabs$' \\citep{MM83} or '$tbabs$' \\citep{Wilms2000}. The various absorption descriptions can lead to different results (see Fig.~\\ref{fig:abun}) or expectations of e.g. observable objects. The differences caused by the diverse treatments show our incomplete knowledge about the interstellar absorption and should be kept in mind especially when interpreting soft X-ray spectra. \\begin{figure} \\centering \\includegraphics[angle=270,width=0.5\\textwidth]{diffWabsTbabs.ps} \\caption{{\\bf{The effect of different abundances and cross-sections on the absorption towards soft X-ray sources}} \\protect\\\\ For the same neutron star (blackbody of 90~eV, 11.67 km radius, and 1.48 Solar masses) the ratio of the X-ray fluxes (as would be expected to be observed by ROSAT) resulting from the two $XSPEC$ routines '$tbabs$' and '$wabs$' are calculated for different distances. The galactic coordinate $x$ varies from 50 to 2500 ~pc (direction: from the Sun towards the galactic center), $y$ is always 0, $z$ is 25 pc north . For the ISM distribution, we applied the same as used by \\cite{Popov2000}, which is described in more detail below. The effect of less absorption, resulting from the use of the newer routine 'tbabs', is stronger at larger column densities, hence distances. } \\label{fig:abun}       % \\vspace{-0.3cm} \\end{figure} \\vspace{-0.6cm} \\subsection{The inhomogeneously distributed ISM} \\label{ISMintro} \\begin{figure} \\centering \\includegraphics[angle=0,width=0.35\\textwidth]{LallNaNHz0log11farb.ps} \\caption{{\\bf{The limits of using the \\cite{Ferlet1985} formula }} \\protect\\\\ The sodium column density as derived from the density data by \\cite{Lall03} is shown in galactic cartesian coordinates at z=0; x is towards the galactic center which is at the right, y towards l=90$^{\\circ}$ is pointing up. The image spans 500~pc times 500~pc. The area within the contour  has  low column densities ($\\log$ N$(Na I) <11$) for which \\cite{Welty1994} doubted whether the \\cite{Ferlet1985} formula is applicable. This region of small sodium column densities  is largest at z=0 and becomes smaller for lower or larger z. } \\label{fig:lal}       % \\vspace{-0.2cm} \\end{figure} In X-ray astronomy with low spectroscopic resolution the chosen abundances and cross-sections are used to derive the hydrogen column density responsible for the instellar absorption. Cold and warm H-atoms (\\HI) as well as hydrogen molecules and ionized hydrogen (\\HII) contribute to the overall amount of hydrogen along a line of sight. These different components of the hydrogen column density are estimated by different measurement methods, e.g. various extinction measurements as summarized by \\cite{Knude2002} or radio 21-cm observations of hydrogen as the extensive study by \\cite{Dickey90} and more recently by \\cite{Kalberla2005}. A wealth of theoretical models has been applied to these observations. However, most of them are only 2D with few more recent exceptions (see e.g.  \\citealt{Drimmel2003}, \\citealt{Amores2005} or \\citealt{Marshall2006astroph} for 3D models). The overall result is a highly inhomogeneously distributed interstellar medium within the Milky Way, present in  bubbles with loops or dense rims, and shaped by molecular clouds, supernova explosions  and stellar activity (for an illustrated overview see \\citealt{Henbest1994}). The ISM distribution is best known in the close Solar neighbourhood, especially in the Local Bubble (e.g. \\citealt{breit1998}). The most detailed view is provided by the study of \\cite{Lall03}, who measured Na {\\sc{i}} absorption towards 1005 sightlines with Hipparcos distances up to 350 pc from the Sun; for more details on the measurements  and the applied density inversion method see \\cite{Lall03}, the precursor paper \\cite{Sfeir99}, and  \\cite{Vergely01}. Na{\\sc{i}} is regarded as good tracer of the total amount of the neutral interstellar gas and the sodium column density is thought to be convertable into a hydrogen column density (e.g. \\cite{Ferlet1985}, but see also discussion in sec.~\\ref{sec:models}). For more distant regions the determination of a 3D distribution of the ISM is hampered mainly by the unknown or imprecise distances of the measured objects and the lack of a well-sampled grid of measurement points. This is true for all extinction-related methods to our knowledge and results in much more  coarse sampling compared to the Solar neighbourhood. The results obtained by the various methods  can differ for the same regions and a combination can be difficult. These difficulties result partly from the different observed astrophysical phenomena like measuring the dust emission or stellar magnitudes. They also result from different sets of stellar absolute magnitude calibrations and the underlying assumptions which can deviate from each other as interstellar extinction determination is an iterative process coupled with the identification of the statistical properties of stars \\citep{Hakkila97}. \\begin{figure} \\centering \\includegraphics[angle=0,width=0.5\\textwidth]{LalSerPisk1kpcGM.ps} \\caption{{\\bf The N(H) predicted by the analytical model at 1~kpc}\\protect\\\\ Positive galactic latitudes are up and $l=0^{\\circ}$ is in the center. The N(H)is in units of $10^{22}$cm$^{-2}$. For description of the model and discussion see text.} \\label{fig:ana}       % \\vspace{-0.2cm} \\end{figure}  \\vspace{-0.2cm} ",
        "conclusions": "An understanding of the interstellar absorption is crucial for the interpretation of low-resolution soft X-ray spectra. Taking into account the 3D distribution of the ISM is important to determine the absorption of known sources as well as for searches of new soft X-ray sources towards a particular direction in the sky. We presented two simple models for 3D ISM distributions together with possible applications: distance estimation and the population synthesis for young isolated thermal, thus soft X-rays emitting neutron stars. Both models are large scale models and have to be used with caution especially considering distance estimations. Concerning population synthesis the two models influence the log$N$--log$S$ curve discriminatively, while the differences for the predicted spatial distribution of observable neutron stars are only small.       \\vspace{-0.4cm}"
    },
    "0609/astro-ph0609796_arXiv.txt": {
        "abstract": "We show that the luminosity of a star forming galaxy is capped by the production and subsequent expulsion of cosmic rays from its interstellar medium.  By defining an Eddington luminosity in cosmic rays, we show that the star formation rate of a given galaxy is limited by its mass content and the cosmic ray mean free path.  When the cosmic ray luminosity and pressure reaches a critical value as a result of vigorous star formation, hydrostatic balance is lost, a galactic scale cosmic ray-driven wind develops, and star formation is choked off.  Cosmic ray pressure-driven winds are likely to produce wind velocities in proportion to and significantly in excess of the galactic escape velocity.  It is possible that cosmic ray feedback results in the Faber-Jackson relation for a plausible set of input parameters that describe cosmic ray production and transport, which are calibrated by observations of the Milky Way's interstellar cosmic rays as well as other galaxies. ",
        "introduction": "In the context of galaxy formation, ``feedback'' processes are commonly invoked to limit the formation of stars. Feedback can be ``intrinsic'' or ``extrinsic,'' where an intrinsic feedback mechanism regulates star formation directly at the molecular cloud level and an extrinsic mechanism controls the supply of stellar fuel on galactic scales.  There is long-standing observational evidence that extrinsic feedback mechanisms operate.  For example, the luminosity of local galaxies correlates with the stellar velocity dispersion, and thus gravitational binding energy, of the entire galaxy itself (Faber \\& Jackson 1976).  Two popular extrinsic feedback mechanism are thermally-driven winds powered by core-collapse SNe (see e.g., Chevalier \\& Clegg 1985; Leitherer et al. 1999) and momentum-driven winds powered by starburst radiation (Scoville 2003, hereafter SCO; Murray et al. 2005, hereafter MQT).  In this work, we consider another possibility: galactic winds driven by cosmic rays.   \\subsection{Basic Idea:  The Eddington Limit in Cosmic Rays }  There is a theoretical upper-limit for the luminosity of gravitationally bound systems, the Eddington limit $L_{\\rm edd}$, given by \\be L_{\\rm edd}=\\frac{4\\pi GM_{\\rm enc}m_pc }{\\sigma_T} \\ee where $M_{\\rm enc}$ is the enclosed mass and $\\sigma_T$ is the Thomson cross section.  Quasars, the most steadily luminous objects in the Universe, are thought to be powered by accretion onto a black hole and thus their maximum luminosity is given by \\be L^{\\rm Q}_{\\rm edd}\\simeq 10^{47}M_9{\\rm erg\\,s^{-1}} \\label{quasarmax} \\ee where $M_9$ is the quasar black hole mass in units of $10^9M_{\\odot}$.  In the case of elliptical galaxies, the mass of the black hole's host galaxy is $\\sim 10^2 -10^3$ times that of the black hole (Magorrian et al. 1998).  However, rather than being the brightest persistent sources in the Universe, typical luminosities of bright star forming galaxies range from $~ 10^{45} -10^{46}\\,{\\rm erg\\, s^{-1}}$ -- at least a factor of ten less powerful than their quasar counterparts.  Photons are not the only type of radiation produced from the act of star formation.  For every core-collapse supernova, a fraction $\\epsilon\\sim 0.1$ of the released energy $\\sim 10^{50} {\\rm erg}$ is converted into cosmic rays (CRs).  Unlike photons, cosmic rays interact many times with the magnetized interstellar medium, and exchange momentum with it, through pitch-angle scatterings with Alfv\\`enic disturbances.  Though the energy injection rate of CRs is insignificant in comparison to other by-products of star formation such as photons, winds, and supernova shocks, the rate at which they inject momentum is not.  In the case of the Milky Way, the CR diffusion time $t_{\\rm diff,CR} \\sim 10^7{\\rm yrs}$ over a characteristic scale of $H\\sim 1{\\rm kpc}$.  Thus the characteristic mean free path for a Galactic CR is given by \\be \\lambdacr\\sim \\frac{3H^2}{c\\,t_{\\rm diff,CR}}\\sim 1\\,{\\rm pc}, \\ee for 1-10 GeV CRs. Whereas the Thomson mean free path $\\lambda_T\\sim 10^{24}{\\rm cm}$. Thus, CRs are much more highly coupled to the Galactic matter in comparison to photons.  Due to the higher level coupling, the threshold for the maximum allowable CR flux and luminosity is lower in comparison to the electron scattering case.  A rough theoretical upper limit for the CR luminosity $L_{\\rm edd, CR}$ is given by \\be L_{\\rm edd, CR}\\simeq L_{\\rm edd}\\frac{\\lambdacr}{\\lambda_T} = 10^{-6}L_{\\rm edd} \\label{CRedd} \\ee Ultimately, the ratio $\\lambdacr/\\lambda_T$ may be thought of as the inverse of the ratio of the effective cross section for a CR to interact with a unit of plasma mass divided by the corresponding value for a photon.  \\subsection{Plan of this work}  In the next section we present the main results of this work.  First, a brief overview of the observed properties of bright starbursting galaxies is given in \\S\\ref{ss: observed}.  The basics of a momentum-driven CR feedback mechanism is given in \\S\\ref{ss: maxL}. In \\S\\ref{ss:comparison}, and to some degree in \\S\\ref{ss: winds} as well, we weigh the advantages and disadvantages of the various galactic feedback mechanisms and then compare them to our CR feedback mechanism.  In \\S\\ref{s:CRinput} we discuss some of the relevant microphysical properties of CRs that are relevant for CR feedback and in \\S\\ref{sss: Faber-Jackson} we speculate as to whether or not CR feedback is responsible for the Faber-Jackson (1976) relation.  We conclude in \\S\\ref{s: conclusion}. ",
        "conclusions": "\\label{s: conclusion}  From the process of star formation, cosmic rays represent only a minor form of energy release.  Yet, their momentum coupling with the interstellar medium is absolute as result of their small scattering length.  As long as the CR mean free path $\\lambdacr$ does not greatly differ from the inferred value in the Milky Way, the collective force that CRs exert upon the interstellar medium as they leak out is great enough to support the gaseous component against gravity and may even unbind it.  If CR feedback is an important mechanism in determining the maximum luminosity of a starbursting galaxy, then the CR optical $\\taucr$ must be large in order to compensate for the relatively small CR luminosity.  Consequently, galactic CRs necessarily impart momentum upon the interstellar medium in the direction opposite to the galaxy's gravitational center.  Thus, CR feedback is highly independent of the spatial distribution of CR injection sites such as SNe and stellar winds.  In analogy with the photon-driven winds of massive and asymptotic giant branch stars, CR-driven galactic winds may reach terminal velocities in excess of the escape velocity by a factor of $\\sim$ a few.  Due to multiple scatterings in the galactic halo, above the main body of the galaxy, the CR beam may indefinitely accelerate a parcel of fluid up to the point where all of the CRs are redshifted away.  The main uncertainty in the CR feedback theory presented here is in determining the manner in which the CR mean free path $\\lambdacr$ (or optical depth $\\taucr$) and radiative efficiency $\\epsilon_{_{\\rm CR}}$ varies from galaxy to galaxy.  The optical depth $\\taucr$ cannot approach arbitrarily high values since adiabatic losses and meson production will eventually remove energy out of the CR beam."
    },
    "0609/hep-ph0609081_arXiv.txt": {
        "abstract": "We identify the range of parameters for which the sterile neutrinos can simultaneously explain the cosmological dark matter and the observed velocities of pulsars. To satisfy all cosmological bounds, the relic sterile neutrinos must be produced sufficiently cold.  This is possible in a class of models with a gauge-singlet Higgs boson coupled to the neutrinos. Sterile dark matter can be detected by the x-ray telescopes. The presence of the singlet in the Higgs sector can be tested at the Large Hadron Collider. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609782_arXiv.txt": {
        "abstract": "{} {Our aim is to constrain the properties of dark matter halos inhabiting high density environments, such as is the case in massive galaxy clusters.} {We use galaxy-galaxy lensing techniques that utilize a maximum likelihood method to constrain the parameters of the lenses. It has been demonstrated that such a technique provides strong constraints on the parameters that characterize a galaxy halo, as well as on the aperture mass of these halos. In this analysis, we only use weak shear data and do not include strong lensing constraints.} {We present the results of a study of galaxy-galaxy lensing in a homogeneous sample of massive \\textsc{x}-ray luminous clusters at $z\\sim$ 0.2. These have been observed in three bands with the \\textsc{cfh12k} instrument.  We find dark matter halos in these clusters to be compact compared to those inferred around isolated field galaxies of equivalent luminosity at this redshift: the half mass radius is found to be smaller than 50\\,kpc, with a mean total mass of order $0.2\\cdot10^{12}$ M$_{\\sun}$. This is in good agreement with previous galaxy-galaxy lensing results and with numerical simulations, in particular with the tidal stripping scenario. We thus provide a strong confirmation of tidal truncation from a homogeneous sample of galaxy clusters. Moreover, it is the first time that cluster galaxies are probed successfully using galaxy-galaxy lensing techniques from ground based data.}{} ",
        "introduction": "Gravitational lensing has become a very powerful tool to measure the mass distribution of structures in the Universe on a large range of scales.  Both the strong lensing and the weak lensing regime are used to map mass distributions. On large scales, the weak distortions detected in the shapes of distant galaxies allows us to study the distribution of matter on cosmic scales \\citep{refregier03}.  On galaxy cluster scales, the strong lensing features observed in the cores of massive clusters can be used to put strong constraints on the inner part of the cluster potential (\\citealt{stronglensing}; \\citealt{smith05}), whereas at larger cluster-centric radius, in the weak lensing limit, the ellipticities of background galaxies give an estimate of the shear field induced by the gravitational potential of the foreground cluster.  On the scale of individual galaxies, much work has been done as well on modeling and understanding multiple quasar systems (\\citealt{fassnacht}; \\citealt{phillips}). These strong galaxy-galaxy lensing analyses provide a test of structure formation in cold dark matter models since lensing provides an interesting way to estimate the inner slopes of density profiles, which can be compared to theoretical expectations.  Weak galaxy-galaxy lensing studies provide constraints on the physical parameters that characterize the dark matter halos of galaxies.  This is accomplished directly using lensing, since the deformation in the shapes of background galaxies produced by the foreground lenses although weak is observationally detected statistically. Galaxy-galaxy lensing has been used to get constraints on field galaxies in different surveys (see Section 2 for a review). This effect has also been used successfully to map substructure in massive galaxy clusters (see Section 2.4), where most previous work has utilized strong and weak lensing features in order to constrain the properties of galaxy halos associated with the locations of bright early-type galaxies in clusters.  In this paper, we report the first detection of galaxy-galaxy lensing in clusters from ground based observations \\emph{without} utilizing strong lensing. We demonstrate that this technique works even without the constraints from multiple images but as a consequence we need more foreground-background pairs. To illustrate this we apply the technique to a well defined homogeneous sample of clusters all at $z \\sim 0.2$. Additionally, using 3 photometric bands, we apply photometric redshift determination techniques to assign redshifts to all the background objects. The details of our method have been presented in a recent theoretical paper (\\citet{limousin05}, hereafter Paper I). It was demonstrated by analysing simulated data that the technique works well and allows us to retrieve the characteristic parameters that describe the dark matter halos and hence to put strong constraints on the aperture mass of these halos. We confirm the fact that galaxy halos in clusters are compact compared to halos of field galaxies of equivalent luminosity. This is a strong confirmation since it relies on a homogeneous sample of galaxy clusters, whereas most early work relied on a heterogeneous sample. The results presented here are averaged on a cluster galaxy population from the centre of the cluster up to $\\sim$ 2 Mpc, which represent a significant fraction of the virial radius (depending on the cluster: for the sample we consider in this work, the virial radius spans from 2.5 Mpc for Abell~2218 to 3.6 Mpc for Abell~1835).  This paper is organized as follows: In Section 2 we review galaxy-galaxy lensing results to date; Section 3 describes the data and reduction techniques, as well as the procedure used to derive catalogues from these images. The determination of photometric redshifts is described in Section 4 and illustrated for the specific case of Abell~1763.  The formalism used for modeling dark matter halos and the maximum likelihood method are outlined in Section 5.  More details on the analysis are given in Paper I. Our results are presented in Section 6 , where we also make comparisons with other galaxy-galaxy lensing results. Discussion of the results and conclusions are presented in Section 7.  All our results in this paper are scaled to the flat, low matter density $\\Lambda$CDM cosmology with $\\Omega_M = 0.3, \\ \\Omega_\\Lambda = 0.7$ and a Hubble constant $H_0 = 65$ km\\,s$^{-1}$ Mpc$^{-1}$. In such a cosmology, at $z=0.2$, $1''$ corresponds to $3.55$ kpc. ",
        "conclusions": "The results presented here come from wide field ground based data. As a consequence, we probe the cluster galaxy population out to a larger radius than previous studies ($\\sim 2$ Mpc). This means that our results are averaged over a galaxy population located all the way from the centre of the cluster to the transition region of the cluster. Our study confirms the fact that galaxy halos in clusters are significantly more compact compared to halos around field galaxies of equivalent luminosity.  The results presented in this work are in very good agreement with numerical simulations. A detailed comparison with realistic numerical simulations will be presented in a forthcoming publication (see also recent work by Natarajan et~al., 2006). The theoretical expectation is that the global tidal field of a massive, dense cluster potential well should be strong enough to truncate the dark matter halos of galaxies that traverse the cluster core. As mentioned before, this expectation has been observationally confirmed by several independent previous studies. Early numerical work (see \\emph{e.g} \\citealt{merritt83}, \\citealt{richstone84}) noticed that a large fraction of the mass initially attached to galaxies in the central megaparsec is stripped. \\citet{bullock} found that halos in dense environments are more truncated than their isolated counterparts of the same virial mass. Avila-Reese et~al. (1999, 2005) found that halos in cluster regions are more concentrated than isolated halos. \\textsc{n}-body simulations of cluster formation and evolution (\\citealt{ghigna98}, \\citealt{ghigna00}) find that the dominant interactions are between the global cluster tidal field and individual galaxies after $z=2$. The cluster tidal field significantly strips galaxy halos. As previously noted, we probe in this paper lens galaxies located all the way from the centre of the cluster to $\\sim$ 2 Mpc. This galaxy population is bound to the cluster and it is reasonable to think that it has experienced the cluster potential at least once: the characteristic time for a galaxy to cross a cluster is about $10^{9}$ years, and the age of a galaxy cluster that formed at $z=1$ is about $10^{10}$ years.  To conclude, we presented the first galaxy-galaxy lensing results to date that probe cluster galaxies from a ground based survey.  This study has confirmed the fact that galaxy halos in clusters are significantly less massive and more compact compared to galaxy halos around field galaxies of equivalent luminosity. Moreover, this confirmation is based on the analysis of 5 massive clusters lenses whose properties are close to each other, hence the confirmation we provide is a strong one since it relies on a homogeneous sample."
    },
    "0609/astro-ph0609261_arXiv.txt": {
        "abstract": "The fine structure and the variations of the optical pulse shape and phase of the Crab pulsar are studied on various time scales. The observations have been carried out on 4-m William Hershel and 6-m BTA telescopes with APD photon counter, photomultiplier based 4-channel photometer and PSD based panoramic spectrophotopolarimeter with 1$\\mu$s time resolution in 1994, 1999, 2003 and 2005-2006 years. The upper limit on the pulsar precession on Dec 2, 1999 is placed in the 10 s - 2 hours time range. The evidence of a varying from set to set fine structure of the main pulse is found in the 1999 and 2003 years data. No such fine structure is detected in the integral pulse shape of 1994, 1999 and 2003 years.  The drastic change of the pulse shape in the 2005-2006 years set is detected along with the pulse shape variability and quasi-periodic phase shifts. ",
        "introduction": "Over the last 30 years the Crab pulsar has been extensively studied. The reasons for it are clear -- it is the brightest pulsar seen in optics, it is nearby and young. However, the post popular groups contemporary theories of the Crab high-energy emission, the ``polar cap'' \\citep{daugherty} and ``outer gap'' \\citep{cheng} ones, can't explain the whole set of observational data.  One of the main properties of the Crab emission is the very high stability of its optical pulse shape despite the secular decrease of the luminosity, related to the spin rate decrease \\citep{pacini,nasuti}.  At the same time the pulsars in general and the Crab itself are unstable. The instabilities manifest itself as a glitches, likely related to the changes of the neutron star crust, timing noise, powered by the collective processes in the superfluid internal parts of it, magnetospheric instabilities, results of the wisps around the pulsar, precession, et al. All these factors may influence the optical pulse structure and change it on various time scales, both in periodic and stochastic way.  However, it has been found early that the variations of the Crab optical light curve, in contrast with the radio ones, are governed by the Poissonian statistics \\citep{Kristian}. A number of observations show the absence of non-stationary effects in the structure, intensity and the duration of the Crab optical pulses, and the restrictions on the regular and stochastic fine structure of its pulse on the time scales from 3$\\mu$s to 500$\\mu$s \\citep{beskin,percival}, the fluctuations of the pulse intensity \\citep{Kristian}.  Along with the increase of the observational time spans and the accuracy of measurements the small changes of the optical pulse intensity, synchronous with the giant radio pulses, have been detected \\citep{shearer}. Also, the evidence for the short time scale precession of the pulsar has been detected by studying its optical light curve \\citep{cadez}.  All this raises the importance of the monitoring of the Crab optical emission with high time resolution. ",
        "conclusions": "We analyzed the data of several sets of optical observations with high temporal resolution of the Crab pulsar performed by our group over the 12 last years.  No evidence for short time scale precession (like 60-sec free precession discovered in \\cite{cadez} ) is detected on the level of $10^{-5}$ - $10^{-7}$ s$^{-1}$ pulsar frequency variation on 10 s - 2 hours time scale on Dec 2, 1999 (see Figure \\ref{fig_shift_1999}), which corresponds to the precession wobble angle to be less than approximately $2\\cdot10^{-3}$. Also, no signatures of short time scale timing noise is seen in this data set.  No significant fine structure is detected in the integral pulse profile of 1994, 1999 and 2003 years data set (see Figure \\ref{fig_pulse_123}), however, each data set alone show the evidence of fine structure on the level of 3-5 sigma, which may be related to its instability on the time scale of years along with the stability of the pulse shape on the same scale.  We discovered the significant change of the time-averaged Crab pulse profile in the Dec 2005 - Jan 2006 set of observations. The pulse profile also shows the variations between the nights. Also, the quasi-periodic phase shifts in respect to the second-order timing solution (up to second frequency derivative) has been detected in the data with amplitude of $\\sim 100\\mu$s and characteristic time scale of 0.7 days. We have not found any hardware or software issue able to mimic such pulsar behaviour. These results may be interpreted as a geometric effects due to the Crab precession suddenly started between our observations of 2003 and 2005-2006 years."
    },
    "0609/gr-qc0609010_arXiv.txt": {
        "abstract": "We demonstrate the use of automatic Bayesian inference for the analysis of LISA data sets. In particular we describe a new automatic Reversible Jump Markov Chain Monte Carlo method to evaluate the posterior probability density functions of the {\\em a priori} unknown number of parameters that describe the gravitational wave signals present in the data. We apply the algorithm to a simulated LISA data set containing overlapping signals from white dwarf binary systems and to a separate data set containing a signal from an extreme mass ratio inspiral. We demonstrate that the approach works well in both cases and can be regarded as a viable approach to tackle LISA data analysis challenges. ",
        "introduction": "The Laser Interferometer Space Antenna (LISA) is a proposed all-sky gravitational wave observatory with an expected launch date of 2015~\\cite{lisa_ppa}. It will consist of a 3-arm laser interferometer in heliocentric orbit trailing the Earth by 20 degrees and counter-rotating with a period of 1 yr in a plane inclined by 60 degrees with respect to the ecliptic. The instrument will be sensitive to gravitational waves (GWs) in the frequency band 0.1 mHz - 0.1 Hz and is expected to observe a great variety of sources, ranging from known galactic binary systems to high redshift binaries harbouring a massive black hole~\\cite{Cutler2002a}.  One of the challenges prior to the launch of the mission is to develop data analysis strategies able to provide accurate estimates of the parameters of the sources that generate the gravitational waves present in the data stream. In fact the LISA data set will contain a large ({\\em a priori} unknown) number of partially overlapping signals. Here we consider an approach to the analysis within the framework of Bayes' inference~\\cite{Brett1988a}. We use techniques known as Reversible Jump Markov Chain Monte Carlo (RJMCMC) methods~\\cite{Green2003a} to evaluate the posterior probability density function (PDF) of the source parameters when the date set contains an unknown number of signals drawn from a large parameter space. Markov Chain Monte Carlo Methods have already been considered in several LISA data analysis case studies, {\\em e.g.}~\\cite{cornish05,Um2005b,cornish06b}. In these proceedings we report the development and application of an ``automatic'' RJMCMC sampler ({\\em e.g.}~\\cite{Green2003a,Hast2004a}), in which the user just provides a model for the signal(s) expected in the data set and the algorithm automatically returns the relevant posterior PDFs.  ``Automatism'' is a challenge for the theory behind RJMCMC sampling. We implement Metropolis-Hastings sampling, with the underlying proposal density for the Markov Chain state transition as the object of an automatic optimisation according to the Adaptive Acceptance Probability (AAP) technique. AAP tries to find the optimal proposal density to achieve a target acceptance rate of 0.25 for Markov Chain state transitions: in this way a quarter of all proposed transitions should be accepted for a new Markov Chain state. AAP suggests bolder transitions if the estimated acceptance rate is larger than 0.25 or stricter proposals if it is smaller than 0.25. This method has proven to be simple in its implementation and computationally effective in a number of applications. In the case of LISA, the number of signals in the data set (and therefore the number of parameters, {\\em i.e., dimensions}, that characterise the problem at hand) is unknown. We implement AAP for fixed dimension updates (jumps within a proposed model)~\\cite{atch2003a}, but also for trans-dimensional updates (jumps between proposed models)~\\cite{Hast2004a}. In addition we support AAP optimisation by finite mixture fits to the joint posterior density function. In order to increase efficiency, the sampler is implemented in parallel on loosely connected CPUs (such as Beowulf-type clusters). This is achieved by running independent Markov chain samples in parallel, which are then combined after the run, as discussed in~\\cite{rosenthal-parallel}. A more detailed discussion of the implementations and techniques, including some expanded examples, will be reported elsewhere.  Here we describe the application of this analysis approach to two case studies: (i) the analysis of a data set containing two highly overlapping monochromatic DWD signals (actual number of signals and noise level in the data set unknown prior to the analysis); (ii) the analysis of a data set containing a signal from an extreme mass ratio inspiral (EMRI) (number of signals known but noise level unknown in the data set prior to the analysis). We demonstrate that the sampler is able to process these two data sets automatically to return posterior density functions that are consistent with the actual parameter values. We discuss the performance of this approach and future work that is necessary to apply this strategy to a realistic LISA data stream. ",
        "conclusions": "We have introduced a new technique for computing the posterior probability density function of the source parameters in the context of LISA data analysis where the number of signals and the noise level in a data set are unknown. In particular, we have described the use of a new unsupervised and automatic Reversible Jump Markov Chain Monte Carlo technique.  We have examined two case studies to demonstrate the use of this sampler: highly overlapping signals from monochromatic white dwarf binary systems, and one complex signal from an EMRI. In the first study, we treated the number of sources as an unknown, but in the latter case we assumed that the number of sources was known to be one. In both cases we were able to successfully identify the signals in the data set automatically. In the case of the overlapping DWDs the recovered source parameters were within one or two sampled standard deviations of the true parameters. The marginalised posterior distributions deviate significantly from the single-source theoretical distribution predicted by the variance-covariance matrix, but this is to be expected due to the high degree of correlation in the system. A full quantitative characterisation of these features will be further investigated in the future. In the case of the EMRI, the sampler can identify and extract the values of the signal's source parameters automatically with high accuracy, with a bias always within one sampled standard deviation of the Markov chain state distribution. This is a promising result which suggests that MCMC may be a useful tool for the challenging task of EMRI detection. However, the waveform model used was too simplified and the observation time too short to allow us to make a generic statement at this stage.  It is interesting to consider why we see, but also expect, bias in our runs. Metropolis-Hastings sampling is an approximate method to evaluate the posterior PDF and as such always introduces bias. This bias can be divided into user-bias and systematical bias. Metropolis-Hastings creates a sample that approximates the real posterior density function asymptotically in time (equally to ongoing Markov chain state sampling). Since we end our program run after $10^7$ ($3 \\times 10^6$) Markov chain states, we introduce bias since at this stage the method has only approximated the real posterior density to a limited degree. Although adaption as performed within AAP speeds up this convergence, this user-bias is still present. The systematical bias comes from two sources --- the noise in the data and, for our overlapping DWD problem, the high degree of overlap between the two binaries in the data. This systematic bias indicates how parameter extraction from a data set may never achieve ideal results, but will be subject to errors that depend on the signal to noise ratio and the complexity and make up of the data stream.  We plan to extend the current work in several ways. Although a parallel code was used to generate the results presented here, the numerical implementation is still too CPU intensive and not suitable for the actual analysis of a realistic LISA data set. In the future we plan to investigate and implement numerical optimisations of various aspects of the sampler to reduce the code run-times. For the EMRI study, such optimisation should allow us to analyse longer data streams, lasting 2 to 5 years. We plan to initially use the toy model to achieve these longer observation times, before switching to the more realistic Barack and Cutler model for further studies (and/or other waveforms that are currently becoming available). For the DWD study, the optimisation will allow us to explore the scenario in which many sources are present in the data set. We will also consider the full set of TDI channels~\\cite{vallis05} and use simulated noise coloured according to the expected spectral density.  \\begin{theacknowledgments} AS wants to thank David Hastie, John Veitch and Graham Woan for helpful discussions. This work was supported in part by St.Catharine's College (JG). \\end{theacknowledgments}"
    },
    "0609/astro-ph0609057_arXiv.txt": {
        "abstract": "{} {A study of the metastable ($J,K$)=(18,18) line of ammonia (NH$_3$) is presented aiming at better defining the physical properties of molecular gas in extreme environments.} {The spectra were collected with the Effelsberg 100-m telescope and are analyzed in combination with other data using a rotation diagram (Boltzmann plot).} {The first astronomical detection of the metastable ($J,K$) = (18,18) line of NH$_3$ is reported. This is the NH$_3$ line with by far the highest energy, 3130\\,K above the ground state, detected in interstellar space. It is observed in absorption toward the galactic center star forming region Sgr~B2. There is a clear detection toward Sgr~B2~(M) and a likely one toward Sgr~B2~(N). An upper limit for emission is determined for Orion~KL. If we combine the (18,18) line results from Sgr~B2~(M) with the previously measured (12,12) absorption line, we find a rotation temperature, $T_{\\rm rot}$, of $>$1300\\,K for the absorbing cloud. This is at least a factor of two higher than previously derived values from less highly excited ammonia lines. $T_{\\rm rot}$ gives a lower limit to the kinetic temperature. There is a hot low density gas component in the envelope of Sgr~B2. It is possible that the (18,18) line arises in this region. The radial velocity of the low density, hot envelope is the same as that of the dense hot cores, so the (18,18) line could also arise in the dense hot cores where non-metastable ($J$$>$$K$) absorption lines from energy levels of up to $\\sim$1350\\,K above the ground state have been observed. A discussion of scenarios is presented.} {} ",
        "introduction": "A number of vibrationally or torsionally excited molecules have been detected in the interstellar medium, among them H$_2$, CS, HCN, NH$_3$, SiO, H$_3C$N, CH$_3$CN, CH$_3$OH, C$_2$H$_3$CN and CH$_3$CH$_2$CN (e.g. Menten et al. 1986; Ziurys \\& Turner 1986; Mauersberger et al. 1988; Hauschildt et al. 1993; Nummelin \\& Bergman 1999; Mehringer et al. 2004). These lines often trace the hottest gas, where molecules can still survive, and thus serve as useful probes of the physical conditions of star forming regions. Particularly attractive are `hot core' environments, with temperatures and densities in excess of 100\\,K and 10$^5$\\,cm$^{-3}$, respectively. Here young massive stars heat the surrounding dust and gas and modify the chemical composition of the latter.  There are, however, also lines in the vibrational ground state, matching or even surpassing the excitation requirements for some of the vibrationally excited transitions seen in hot cores. With the Kuiper Airborne Observatory, Watson et al. (1980) detected the CO $J$=21--20 and 22--21 lines in Orion~KL, corresponding to an energy of $\\sim$1200\\,K. Observations of ammonia (NH$_3$) with the Effelsberg 100-m radio telescope yielded detections to the ($J,K$)=(14,14) inversion transition toward Orion~KL and Sgr~B2. These lines arise from energy levels $\\sim$1950\\,K above the ground state (H{\\\"u}ttemeister et al. 1995).  Ammonia provides an excellent possibility to observe highly excited transitions since inversion transitions ($\\Delta J$=0, $\\Delta K$=0) across a ($J,K$) doublet are located in a band ranging from about 20 to 50\\,GHz. Ammonia is one of the most interesting tracers of the physical properties of dense interstellar gas. It is an excellent diagnostic of kinetic temperature, $T_{\\rm kin}$, over a wide range of densities. To a lesser extent, NH$_3$ is a probe of H$_2$ density, $n$(H$_2$). Inversion lines arise from either non-metastable ($J$$>$$K$) or metastable ($J$=$K$) levels. The non-metastable levels can decay via allowed rotational transitions in the infrared (IR) with ($\\Delta J$=$\\pm$1, $\\Delta K$=0) and have very short lifetimes. Thus, non-metastable lines require exceptional excitation conditions such as very large densities or intense IR radiation fields. In contrast, metastable ($J$=$K$) inversion lines arise from ($J,K$) doublet levels at the bottom of a $K$ ladder, so cannot rapidly decay to other levels. Radiative transitions between metastable levels are first order forbidden, so the populations of these levels are exchanged by collisions. By combining column densities of different metastable levels, one can obtain an estimate of $T_{\\rm rot}$, which is a lower limit to $T_{\\rm kin}$. The value of the excitation temperature, $T_{\\rm ex}$, across an inversion doublet will be one half of $T_{\\rm kin}$ when $n$(H$_2$) $\\sim$ 10$^4$\\,cm$^{-3}$ (see e.g. Rohlfs \\& Wilson 2003 for details).  In this paper, we report the first detection of the (18,18) inversion line, toward Sgr~B2~(M). We also present a likely detection toward Sgr~B2~(N) and an upper limit for Orion~KL. We interpret the data using comparisons with previous results involving lines of lower excitation. ",
        "conclusions": "\\subsection{The Sgr~B2 cloud}  The Sagittarius B2 molecular cloud is one of the most active regions of star formation in our Galaxy. It is located about 40\\arcmin\\ from the galactic center, at a projected distance of $\\sim$100 pc for a Sun-galactic center distance of 8.5~kpc. Sgr~B2 hosts three prominent cores of ionized gas, Sgr~B2~(M), (N) and (S) (see e.g. Martin-Pintado et al. 1999). These are associated with molecular clouds which are in turn embedded in a much more extended, lower density envelope.   \\begin{table} \\caption{Rotational temperatures$^{\\rm a)}$} \\label{masses1} \\begin{flushleft} \\begin{tabular}{ l r r r r r } \\hline \\multicolumn{1}{c}{(1)}             & \\multicolumn{1}{c}{(2)}             & \\multicolumn{1}{c}{(3)}             & \\multicolumn{1}{c}{(4)}             & \\multicolumn{1}{c}{(5)}             & \\multicolumn{1}{c}{(6)}             \\\\ \\multicolumn{1}{c}{Source}          & \\multicolumn{2}{c}{Column}          & \\multicolumn{2}{c}{Normalized}      & \\multicolumn{1}{c}{$T_{\\rm rot}$}   \\\\ \\multicolumn{1}{c}{}                & \\multicolumn{2}{c}{density}         & \\multicolumn{2}{c}{column}          & \\multicolumn{1}{c}{}                \\\\ \\multicolumn{1}{c}{}                & \\multicolumn{2}{c}{}                & \\multicolumn{2}{c}{density}         & \\multicolumn{1}{c}{}                \\\\ \\multicolumn{1}{c}{}                & \\multicolumn{1}{c}{(12,12)}         & \\multicolumn{1}{c}{(18,18)}         & \\multicolumn{1}{c}{(12,12)}         & \\multicolumn{1}{c}{(18,18)}         & \\multicolumn{1}{c}{}                \\\\ \\multicolumn{1}{c}{}                & \\multicolumn{4}{c}{(cm$^{-2}$/K)}   & \\multicolumn{1}{c}{(K)}             \\\\ \\hline Sgr~B2~(M)  & 5.6(12) &    1.0(13) & 1.1(11) &    1.3(11) &$>$1300 \\\\ &         &    2.2(12) &         &    3.0(10) &   1300 \\\\ Sgr~B2~(N)  & 1.5(13) &    6.2(12) & 2.9(11) &    8.4(10) &   1350 \\\\ &         &    1.5(12) &         &    2.0(10) &    600 \\\\ Orion KL    & 2.6(13) & $<$3.0(12) & 5.2(11) & $<$4.1(10) & $<$500 \\\\ \\hline \\end{tabular} \\ \\ \\ \\ \\ \\ \\ \\\\ \\item[a)] The (12,12) data for Sgr~B2 were taken from Table~2 of H{\\\"u}ttemeister et al. (1995). As in Table~1 of this paper, the higher and lower (18,18) column densities refer to the data taken on July 24 and 26, 1998, respectively (see also Sect.\\,3). Errors are not given, because a comparison between the two epochs provides a good measure of the extremes involved. For Orion~KL, the (12,12) column density was taken from Wilson et al. (1993). For the (18,18) column density optically thin emission, $T_{\\rm mb}$=0.1\\,K and a FWHP linewidth of $\\Delta V_{\\rm 1/2}$ = 8\\,km\\,s$^{-1}$ was assumed. The given rotation temperature also accounts for the fact, that Wilson at al. calculated column densities averaged over a 29\\arcsec\\ beam, while the (18,18) upper limit is an average over a 21\\arcsec\\ beam. \\end{flushleft} \\end{table}   For Sgr~B2, the geometry of the molecular clouds and continuum sources is important. There are compact radio continuum sources absorbed by NH$_3$ inversion lines. The telescope beam of the Effelsberg 100-m dish is larger than these continuum sources and the associated hot cores, but smaller than the NH$_3$ containing cloud in the molecular envelope of Sgr~B2 (see e.g. H93 for details of this geometry and 40\\arcsec\\ images of NH$_3$ lines).  For the NH$_3$ ($J,K$) = (1,1), (2,2), (3,3), (4,4), (5,5), (6,6) and (7,7) metastable inversion lines, there is absorption toward Sgr~B2~(M) and (N), and emission one full beamwidth, 40\\arcsec, off the continuum sources. H93 assumed that this is due to a uniform extended layer of NH$_3$. In this case, the main beam brightness temperature of the continuum source, $T_{\\rm C}$, is larger than the excitation temperature, $T_{\\rm ex}$, of the line. Given a uniform NH$_3$ cloud, more extended than the radio telescope beam and calculated collision rates (Danby et al. 1988), one can determine values for $n$(H$_2$). From such models, H93 had estimated $n$(H$_2$)$\\sim$3\\,10$^3$\\,cm$^{-3}$. Since radiative transitions across $K$ ladders are forbidden, the relation of $T_{\\rm rot}$ is close to but less than $T_{\\rm kin}$; this is not strongly dependent on H$_2$ densities. Thus, from measurements of a number of metastable NH$_3$ absorption lines, one estimates that $T_{\\rm rot}$ = 160\\,K in the envelope of Sgr~B2 (Fig.\\,2). H{\\\"u}ttemeister et al. (1995) carried out measurements of the (8,8), (9,9), (10,10), (11,11), (12,12), (13,13), and (14,14) metastable inversion lines toward Sgr~B2~(M) and (N). For SgrB2(M), $T_{\\rm rot}$$\\sim$600\\,K (Fig.\\,2). Toward the continuum sources these lines were found in absorption. Because of the weakness of the lines, there were no measurements of off-source positions, where emission corresponding to the absorption would be expected as with the (1,1) to (7,7) lines.  The rotational transitions of NH$_3$ produce lines in the far infrared (FIR). An extensive study of these FIR lines has been carried out using the ISO (Infrared Space Observatory) LWS (Long Wavelength Spectrometer) with $\\Delta \\lambda$/$\\lambda$ $\\sim$ 8000 and a beam size of $\\sim$80\\arcsec\\ (Ceccarelli et al. 2002). Twenty-one NH$_3$ lines in the wavelength range between 47 and 196$\\mu$m were detected, all of them in absorption. Of these, 13 lines involve metastable levels, while 8 connect non-metastable levels. As Ceccarelli et al. (2002) point out, their lines absorb the spatially extended FIR dust continuum rather than the more compact free-free continuum absorbed by the centimeter wavelength inversion lines. This may account for some of the differences between inversion line results and the FIR line data. Ceccarelli et al. (2002) have presented a model for their data. This involves an absorbing layer 1.15$\\pm$0.15\\,pc in front of the FIR continuum source. This layer has a kinetic temperature of 700$\\pm$100\\,K and an H$_2$ density of $\\le$10$^5$\\,cm$^{-3}$. From the non-detection of FIR CO lines, Ceccarelli et al. (2002) surmise that $n$(H$_2$) $\\le$ 10$^4$\\,cm$^{-3}$. Oka et al.~(2005) deduce H$_2$ densities of order 10$^2$\\,cm$^{-3}$ in the galactic center region, so the envelope of SgrB2 may have even lower densities.  As with the other extended molecular clouds in the galactic center region, dust in the envelope of Sgr~B2 is comparatively cold, with dust temperatures, $T_{\\rm dust}$, of order $\\sim$70\\,K. Thus in the envelope of Sgr~B2, non-metastable levels of NH$_3$ cannot be populated by collisions or IR fields. In addition to the envelope, however, there are a number of hot, dense, compact cores. For SgrB2(M), these have the same radial velocities as the envelope. In the cores, the populations of non-metastable levels can be significant. From maps of the (2,1), (3,2), (4,3), (5,4) and (4,2) non-metastable inversion emission lines, H93 found that the {\\it warm envelope} has a FWHP size of $\\sim$1\\arcmin. These lines arise from energy levels up to 300\\,K above the ground state. Measurements of non-metastable lines from energy levels between 400 and 700\\,K above the ground state show a mixture of absorption and emission. Above 700\\,K, there is only absorption.  H93, de Vicente et al. (1997), Ceccarelli et al. (2002) and Comito et al. (2003) proposed scenarios for Sgr~B2 including (1) compact {\\it hot cores} with $T_{\\rm kin}$$\\sim$200\\,K and $n$(H$_2$)$\\sim$3\\,10$^7$\\,cm$^{-3}$, (2) a more extended {\\it warm envelope} with $T_{\\rm kin}$$\\sim$40\\,K and $n$(H$_2$)$\\sim$3\\,10$^5$\\,cm$^{-3}$, and (3) a very extended {\\it hot envelope} with $T_{\\rm kin}$$\\sim$200\\,K and $n$(H$_2$)$<$10$^4$\\,cm$^{-3}$. In view of the metastable NH$_3$ inversion line data of H{\\\"u}ttemeister et al. (1995) and the FIR line data of Ceccarelli et al. (2002), the $T_{\\rm kin}$ of at least a part of the hot envelope should be raised to 600--700\\,K.   \\begin{figure*}[ht] \\label{boltzmann} {\\includegraphics[scale=0.70]{5591fig2.eps}} \\caption{Normalized column density per K excitation temperature ($T_{\\rm ex}$) versus energy above the ground state. Filled circles with error bars mark the ($J,K$)=(1,1)--(7,7) transitions measured by H{\\\"u}ttemeister et al. (1993), empty circles refer to the (8,8) and (9,9) lines observed by Wilson et al. (1982) and filled circles show the normalized column densities of the (8,8) to (14,14) transitions obtained by H{\\\"u}ttemeister et al. (1995). The inverse slope of a line connecting data points is a direct measure of the rotation temperature, $T_{\\rm rot}$, as long as Local Thermodynamical Equilibrium can be assumed. While the concave form of the diagram suggests the presence of three gas layers with different temperature ($T_{\\rm rot}$ = 160\\,K, 630\\,K and $>$1300\\,K), the gas {\\it might} actually be isothermal with $T_{\\rm rot}$ approaching $T_{\\rm kin}$ with increasing energy above the ground state (see the radiative transfer calculations of Walmsley \\& Ungerechts (1983), Danby et al. (1988), and Flower et al. (1995)). Collision rate coefficients (Danby et al. 1988) were computed for levels as high as ($J,K$) = (6,6) and temperatures as large as 300\\,K. For higher transitions and temperatures, extrapolated rate coefficients have to be used.} \\end{figure*}   \\subsection{Relevance of the ($J,K$)=(18,18) line}  The most definite new result is the detection of the ($J,K$)=(18,18) line in absorption toward Sgr~B2~(M). There are two possible interpretations:  \\noindent (1) As with the (1,1) to (7,7) lines, the (18,18) line arises in the extended envelope and is sub-thermally excited, or  \\noindent (2) the (18, 18) line is thermalized and arises in a region close to the H{\\sc ii} regions in Sgr~B2. After discussing the assumption of LTE used in Table~2 and Fig.\\,3, these scenarios will be analyzed in more detail.  With $T_{\\rm kin}$ $>$ $T_{\\rm rot}$ $>$ 1300\\,K (Table~2), the ($J,K$) = (18,18) line reveals a molecular gas component with a kinetic temperature clearly exceeding the 600--700\\,K estimated by H{\\\"u}ttemeister et al. (1995) and Ceccarelli et al. (2002). Before reaching a conclusion, however, we should ask whether the $T_{\\rm ex}$ values of the various metastable inversion lines are really the same. Does the assumption of LTE hold? One can use the off-peak line and continuum results to show that the value of $T_{\\rm ex}$ for the (1,1) line is $\\sim$6\\,K. A calculation using a two level model for metastable levels with a constant collision rate and the Einstein A coefficients shows that the excitation temperatures across metastable levels decrease from 6\\,K for the (1,1) levels to slightly more than 3\\,K for the (18,18) levels. Our fit to the normalized column densities in Fig.\\,2 is based on the assumption that the value of $T_{\\rm ex}$ is the same for each line. If the value of $T_{\\rm ex}$ does follow the simple relation discussed previously, the normalized column densities will decrease more steeply than shown in Fig.\\,3 with increasing $J$ value. This will lead to a lowering of the $T_{\\rm rot}$ values obtained (Table~2; Fig.\\,3) by about 20\\%; however $T_{\\rm rot}$ will still be $\\ga$1000\\,K.  In Fig.\\,2, there is a clear gradient in the values of $T_{\\rm rot}$. Wilson et al.~(1993) found a similar effect for the high density region Orion~KL. This was interpreted as a gradient in $T_{\\rm kin}$. We note, however, that in a warm enviroment radiative transfer simulations predict higher $T_{\\rm rot}$ values, gradually approaching $T_{\\rm kin}$, for higher metastable levels even if the gas is characterized by a single kinetic temperature (see Walmsley \\& Ungerechts 1983; Danby et al.~1988; Flower et al.~1995). The situation for SgrB2~(M) may be more complex than that for Orion~KL. The lower energy ($J$$\\le$7) metastable absorption lines shown in Fig.\\,2 definitely arise from the envelope, but the higher lying lines {\\it might} arise in hot cores.  There are several arguments suggesting that the ($J,K$)=(18,18) emission arises in a warm, extended low density envelope:  \\noindent (1) Absorption with similar apparent optical depths is seen from the lowest to the highest metastable inversion lines toward both Sgr~B2~(M) and (N). This includes the (18,18) line. Not a single other galactic star formation region shows a similarly excited NH$_3$ absorption component so that statistically it appears unlikely that Sgr~B2 hosts two such clouds. With Sgr~B2~(M) and (N) being separated by only $\\sim$47\\arcsec\\ ($\\sim$2\\,pc), it is not farfetched to assume that the absorption is caused by a single cloud with an extent of at least an arcminute.  \\noindent (2) Maps made in the lower ($J,K$) transitions up to the (7,7) line (see Sect.\\,5.1) reveal emission surrounding the radio continuum sources that appears to arise from the same low density gas component as the gas measured in absorption. Lack of evidence for such emission in $J$$>$7 metastable lines does not argue against this scenario, since such emission is weak and would be difficult to detect.  \\noindent (3) Galactic disk molecular clouds extending over $\\sim$2\\,pc are commonly characterized by a low density. The cloud size-density relation of Larson (1981; his eq.\\,5) suggests $n$(H$_2$)$\\sim$10$^3$\\,cm$^{-3}$. Although an extension of this relation for the galactic disk to galactic center clouds with their larger linewidths may not be justified, this is exactly the density that is required for absorption to dominate in a 30\\arcsec\\ -- 40\\arcsec\\ beam (see Sect.\\,5.1) that characterizes most ammonia measurements obtained at Effelsberg.  These arguments lead one to associate the (18,18) line with warm, extended, low density gas. However a scenario favoring a hot core origin cannot be excluded:  \\noindent (1) $T_{\\rm C}$, averaged over a few arcseconds only, is of order $\\sim$1000\\,K (for detailed continuum maps, see e.g. Gaume \\& Claussen 1990; De Pree et al. 1998). Thus the absorption may arise from compact hot regions of high density ($T_{\\rm ex}$$>$100\\,K) directly associated with the hot cores.   \\noindent (2) Radial velocities of the dense, compact hot cores and less compact lower density envelope are nearly the same and it is not possible to separate their contributions kinematically.  \\noindent (3) While it is sometimes argued that the lack of absorption in the ($J,K$) = (3,2), (4,3) and (7,6) VLA data of Vogel et al. (1987) excludes any compact high density absorbing cloud, the single-dish data of H93 show that absorption dominates the highly excited non-metastable transitions including the (7,6) line toward Sgr~B2~(M). These lines must arise (see Sect.\\,5.1) from small regions that can only be slightly more extended than the radio continuum emission.  \\noindent (4) A comparison of our data from July 24 and 26 shows systematic differences in the Sgr~B2~(M) and (N) absorption profiles (Sect.\\,3), affecting more the high than the low velocity part of the spectra. This suggests that lower ($\\sim$65\\,km\\,s$^{-1}$) and higher ($\\sim$80\\,km\\,s$^{-1}$) velocity components are not smoothly distributed.  A number of heating sources have been proposed for the galactic center region. These include PDRs (Photon Dominated or Photo Dissociation Regions; e.g. Goicochea et al.  2004), shocks (e.g. Flower et al. 1995; Ceccarelli et al. 2002), XDRs (X-ray Dominated Regions; e.g. Mart\\'in-Pintado et al. 2000), cosmic rays (e.g. G{\\\"u}sten at al. 1981) and the dissipation of supersonic cloud motions (e.g. Wilson et al. 1982; Hasegawa et al. 1994). In Sgr~B2, several of these mechanisms are at work. Since kinetic temperatures in excess of 1000\\,K are not widespread ($\\gg$2\\,pc), the latter two mechanisms will not be dominant for the gas traced by the NH$_3$ ($J,K$) = (18,18) line. Due to its low threshhold for photoionization ($\\sim$4.1\\,eV; Suto \\& Lee 1983) NH$_3$ is rapidly destroyed by UV radiation and PDRs can also be excluded (for NH$_3$ in the `PDR galaxy' M\\,82, see Wei{\\ss} et al. 2001; Mauersberger et al. 2003). While detailed models for NH$_3$ being exposed to a strong X-ray radiation field have yet to be reported, shocks are widespread in the galactic center region (e.g. H{\\\"u}ttemeister et al. 1998), can release NH$_3$ from grain mantles and can heat the gas to $T_{\\rm kin}$$\\sim$1000\\,K and more (see Rizzo et al. 2001 for NH$_3$ in a cloud apparently being shocked by the wind of a Wolf-Rayet star). Mart\\'in-Pintado et al. (1999) have found evidence for ring like structures in NH$_3$ in the envelope of Sgr~B2. It may be that the (18,18) emission could arise from such structures, although the $T_{\\rm rot}$ values obtained from lines with $J$$\\leq$4 are less than 200\\,K.  Shocks imply that large amounts of water vapor should also be produced. Simulations using shock models can reproduce the observed NH$_3$ column densities up to the (14,14) line (Flower et al.~1995; their Table~4). These models also predict water vapor column densities of $N$(H$_2$O) $\\sim$ 10$^{18}$\\,cm$^{-2}$ for the envelope. From H$_2^{18}$O absorption line measurements and radiative transfer calculations, Comito et al. (2003) derive $N$(H$_2$O) = 3.5\\,10$^{16}$\\,cm$^{-2}$ in their Hot Layer, which they identify with a part of the Sgr~B2 envelope (for earlier measurements, see Neufeld et al. 2000; Ceccarelli et a. 2002). This is $\\sim$30 times less than the theoretical prediction. Removing the gas seen in the (18,18) absorption line and a part of the gas giving rise to the lower metastable inversion lines from the envelope and associating it instead with the hot cores would reduce the required H$_2$O column density in the envelope and thus might provide a better agreement between theory and observations. Whether associated with the envelope or the hot cores, our data make larger demands on the shock heating model, since the gas temperatures must be considerably higher. An open question is whether larger temperatures alone could reconcile the difference between the model predictions and measurements for the water vapor abundances in the Sgr~B2 envelope.   \\subsection{Orion~KL}  Wilson et al. (1993) identified two NH$_3$ components toward Orion~KL, one with $T_{\\rm rot}$ $\\sim$ 165\\,K, the other with $T_{\\rm rot}$ $\\sim$400\\,K. Our upper line intensity limit to the ($J,K$)=(18,18) transition  is consistent with  $T_{\\rm rot}$$<$500\\,K (Table~2). Thus our data agree with previous results. We conclude that, unlike Sgr~B2, Orion~KL does not host a detectable extremely hot NH$_3$ component."
    },
    "0609/astro-ph0609327_arXiv.txt": {
        "abstract": "{} {To investigate the efficiency of centrifugal acceleration of particles as a possible mechanism for the generation of ultra-high $\\gamma$-ray nonthermal emission from TeV blazars, we study the centrifugal acceleration of electrons by rotating magnetic field lines, for an extended range of inclination angles and determine the maximum Lorentz factors $\\gamma_{max}$ attainable by the electrons via this process.} {Two principal limiting mechanisms for the particle acceleration, inverse Compton scattering and breakdown of the bead-on-the-wire approximation, are examined.} {Particles may be centrifugally accelerated up to $\\gamma_{max} \\simeq 10^8$ and the main limiting mechanism for the $\\gamma_{max}$ is the inverse Compton scattering.} {The energy of centrifugally accelerated particles can be amply sufficient for the generation (via inverse Compton scattering) of the ultra-high energy (up to $20TeV$) gamma emission in TeV blazars.} ",
        "introduction": "One puzzling and interesting problem, related to active galactic nuclei (AGNs) is the origin of the ultra-high $\\gamma$-ray emission from blazars. Some blazars, such as Mrk 421, Mrk 501, PKS 2155-304, 1ES 2344-514, H1426+428, and 1ES 1959+650, emit TeV ($1TeV \\equiv 10^{12}eV$) photons and form a special class of ``TeV blazars\" (\\cite{d02,cw99,h03}).  The standard blazar model implies the presence of a supermassive black hole, surrounded by an accretion disk and ejecting twin relativistic jets, one of which is seen almost end-on. The            % broadband emission spectrum of blazars is made of two components: the low-energy (from radio to optical/UV) part attributed to synchrotron radiation and the high-energy (from X-rays to $\\gamma$-rays) component formed by the inverse Compton scattering (ICS) of softer photons (\\cite{ktk02,ketal98}). The latter part of the spectrum is usually interpreted on the basis of the synchrotron self-Compton (SSC) model. However, the origin of accelerated and/or pre-accelerated electrons and the mechanism that is responsible for their efficient acceleration is still a matter of uncertainty. Proposed mechanisms such as the Fermi-type acceleration process (\\cite{cw99}) and re-acceleration of electron-positron pairs as a feedback mechanism (\\cite{ghb93}) may account for the observed high energy emission up to 20 TeV. However, the Fermi-type acceleration in relativistic jets is efficient when the seed population of \"pre-accelerated\" electrons possesses quite high ($\\gamma_{max}\\ge10^2$) Lorentz factors (\\cite{rm00}). The origin of this \"pre-acceleration\" is still a matter of discussion.  Centrifugally driven outflows (CDOs), centrifugally accelerated particles may acquire quite high energies. This concept (\\cite{g68,g69}) has often been discussed (\\cite{mr94,chetal96,g96,ckf99}) in the context of pulsar emission theory. Regarding the AGNs it was shown that CDOs from accretion disks occurred if the poloidal magnetic field lines are inclined at an angle $\\le 60^0$ to the equatorial plane of the disk (\\cite{bp82}).  In the AGN context, the presence of the CDO would imply that despite the intense UV radiation (when via ICS, soft photons are scattered against accelerated electrons and, as a result, electrons lose energy and photons gain energy), electrons may reach quite high, $\\gamma_{max} \\sim 10^5$, Lorentz factors. If efficient enough, this mechanism could be used not only to justify the pre-acceleration of electrons but it could be considered as an alternative mechanism to generate Blazar high-energy emission (up to $\\sim 20TeV$) (\\cite{oku}).  Using the approach suggested by Gangadhara (\\cite{g96}), where a test charge motion along a rapidly rotating field line was considered in the context of a millisecond pulsar, Gangadhara and Lesch examined the role of the centrifugal force on the dynamics of electrons moving along straight magnetic field lines, fixed in the equatorial plane and co-rotating with the spinning AGN (\\cite{gl97}). They have shown that scattering of low-energy photons against accelerated electrons may lead to the generation of the nonthermal X-ray and $\\gamma$-ray emission.  This problem was re-examined by Rieger and Mannheim (\\cite{rm00}) who tried to specify whether the rotational energy gain of charged particles, moving along straight magnetic field lines, is limited not only by ICS but also by the breakdown of the-bead-on-the-wire approximation (BBW). The latter would happen in the vicinity of the light cylinder when a Coriolis force acting on the particle and trying to `tear it off' the field line  would exceed the Lorentz force binding the particle to the field line. According to their consideration, the maximum value of the Lorentz factor, being the subject of both limitations, is $\\gamma_{max}\\sim1000$, which is not enough to produce the ultra-high energy photons emitted by TeV blazars.  In real, three-dimensional astrophysical situations (for example jets) the magnetic field lines are not localized in the equatorial plane but are inclined with respect to it. In this paper we re-examine the same problem considering the wide range of possible inclinations. We show that, for a wide range of AGN, the mechanism responsible for limiting the attainable maximum Lorentz factors is ICS, which under certain conditions can allow particles to reach quite high Lorentz factors $\\gamma_{max} \\ge 10^{5}$. The BBW becomes important for the low luminosity ($<10^{41}erg/s$) AGN, when $\\gamma_{max}\\sim10^8$. For higher luminosities ($>10^{41}erg/s$) it can be dominant only for relatively small inclinations of the magnetic field lines with respect to the rotation axis ($\\leq10^\\circ$). Therefore, contrary to the conclusions of Rieger and Mannheim (\\cite{rm00}), we argue that CDOs could be efficient enough to account for the TeV blazar emission.  The paper is arranged in the following way: In Sec. 2 we derive basic equations. In Sec. 3 we consider a case when magnetic field lines are located in the equatorial plane and show that the presence of the ICS allows a wide range of AGN to generate ultra-high photon emission. In Sec. 4 inclined magnetic field lines are considered. We show that in this case ICS still remains the dominant limiting mechanism for most of the AGN. In the final section we summarize the obtained results. ",
        "conclusions": "In the present work we focused on the ultra-high emission from AGN, the role of the centrifugal acceleration in this process and the efficiency of different mechanisms limiting this acceleration. We re-examined and generalized the problem considered in Rieger and Mannheim (2000), where the authors argued that together with the ICS, is also BBW an efficient limiting mechanism for electron acceleration. They argued that the BBW limits $\\gamma_{max}$ to a few thousands, which is much less than $\\sim10^5$ (the minimum value of the Lorentz factor needed to produce TeV photons).  Revisiting the same problem  we showed that for AGNs considered by Rieger and Mannheim (2000) as the acceleration limiting mechanism, only ICS is important and consequently the upper limit of the maximum Lorentz factor may be much more than several thousands. In particular, we have shown that when the magnetic field lines are straight and inclined to the axis of rotation by the angle $\\theta=90^\\circ$, for AGN with luminosity $L>8\\times 10^{40}erg/s$, ICS is dominant and BBW becomes significant only for AGN with lower luminosities $L<8\\times 10^{40}erg/s$. Therefore, the example of an AGN with a luminosity power $2\\times 10^{42}erg/s$, when as it was shown only ICS is important (for $\\theta=90^\\circ$), we calculated the approximate value of the maximum Lorentz factor of $10^{5}$ and argued that by virtue of the ICS electrons with such Lorentz factors can produce $1TeV$ energy $\\gamma$-rays.  We also considered the dependence of the validity of the mentioned mechanisms on the inclination angle of the magnetic field lines with respect to the rotation axis and have shown that also in this case, for a wide range of possible inclination angles, ICS is the most important mechanism limiting the Lorentz factor. The latter value may reach up to $\\sim10^8$. BBW becomes dominant mostly for relatively lower luminosities with relatively higher angles, or with higher luminosities and smaller angles. The exception is the case when the luminosity is less than $8\\times 10^{40}erg/s$, when independently of the inclination angle the BBW is dominant. We can estimate by Eq. (14a) that the minimum value of $\\gamma_{max}^{BBW}$ also is very high, $\\sim10^7$.  Letting magnetic field lines have inclination angles in the range $\\sim[0.6^\\circ - 1.2^\\circ]$, we have shown that the corresponding energies of photons due to ICS are of the order of $[1-20]TeV$ which is sufficient to account for the emission of TeV blazars.  For the wide range of AGN with two considered mechanisms, ICS still remains the dominant factor in limiting the maximum Lorentz factor of accelerated particles, letting them produce very high energy photons $[1-20]TeV$. For low luminosity power AGN, at the other hand, BBW may be significant but it also may provide high maximum Lorentz factors of accelerated electrons.  An important restriction of our approach is that we studied only straight magnetic field lines, whereas in realistic astrophysical situations the magnetic field lines are curved. Therefore, it would be interesting to generalize our approach to how the considered problem changes, when the curvature of the magnetic field lines is taken into account. It might be especially important if we consider particle dynamics on a longer time scale, over large length-scales, when the curvature of the field lines cannot be neglected. A mathematical formalism for such a study does exist (\\cite{r03}) and it can be applied to the problem under consideration.  The next limitation is related to the fact that the magnetic field is not influenced by particle motion up to the region  in the immediate vicinity of the light cylinder, where the magnetic field is affected by the motion of the plasma. Our model has been based on the study of the dynamics of a single particle, neglecting complex processes in real relativistic plasmas. We must generalize our approach and examine a more realistic model."
    },
    "0609/astro-ph0609111_arXiv.txt": {
        "abstract": "{} {They are the interpretation of the emission line formation regions in {\\object CI Cygni}} {They involve the examination of radial velocities and fluxes of ultraviolet emission lines at different epochs, deduced from archival IUE and GRHS/HST spectra.} {The line fluxes give electron densities and were in addition used to calculate emission measures, suggesting line formation in regions rather smaller than the binary separation. Examination of the radial velocities led to us to find a systematic redshift of  the high ionization resonance lines with respect to the intercombination, and \\ion{He}{ii} lines. Possible explanations of the redshift and the high resolution GHRS \\ion{C}{iv} profile are discussed. We favour that involving resonance line absorption by a circum-binary region  most  probably in an asymmetric wind interaction shell or in a wind from the accretion disk. } {} ",
        "introduction": "{\\object CI Cyg} is a symbiotic binary, containing a cool giant, of type M5.5 (M\\\"urset \\& Schmid \\cite{ms99}), and a much more compact companion. The infrared spectrum shows no sign of circumstellar dust. It is moreover particularly fruitful to study {\\object CI Cyg}, because it undergoes eclipses. A detailed investigation of this object was performed by Kenyon et al. (\\cite{ken91}). They concluded that the compact object was a main sequence star of $0.5\\, M_{\\sun}$, surrounded by a large accretion disk. Recent work shows however that the compact component of this binary may not be a main sequence star (e.g. Miko{\\l}ajewska \\cite{jm03}).  {\\object CI Cyg} is one of the very few symbiotic binaries in which the cool giant fills or nearly fills its Roche lobe (e.g. Miko{\\l}ajewska \\cite{mik01}, and references therein). {\\object CI Cyg} is also one of the very few symbiotic systems for which a high quality radio spectrum, covering the range between 6 cm and 850 $\\mu$m, has been obtained (Miko{\\l}ajewska \\& Ivison (\\cite{mi01}). These radio data enabled them to determine the free-free turnover frequency of the ionized component, and to critically test the known models for radio emission from symbiotic stars. Unfortunately, they ruled out the two most popular models: ionization of the red giant wind by Lyman continuum photons from its hot companion, and emission resulting from the interaction of winds from the two binary components.  \\begin{table} \\caption[]{Log of the UV observations of {\\object CI Cyg}} \\label{obslog} \\begin{tabular}[bottom]{lcccc} \\hline \\noalign{\\smallskip} Image No. & Date & MJD & Phase$^1$ & T$_{\\rm exp}$ [sec]\\\\ \\noalign{\\smallskip} \\hline SWP\\,06818 &  Dec 11, 1979 & 44157 & 0.718  & ~5400\\\\ SWP\\,09256 & Jun 6, 1980 & 44401 & 0.003 & 10800 \\\\ SWP\\,13795 & Apr 24, 1981 & 44719 & 0.375 & ~4800\\\\ SWP\\,14754 &  Aug 14, 1981 & 44831 & 0.506 & ~5400\\\\ SWP\\,18602 &  Nov 20, 1982 & 45294 & 0.046 & 40800 \\\\ SWP\\,37047 &  Sep 17, 1989 & 47786 & 0.959 & ~6480 \\\\ SWP\\,45102 &  Jul 9, 1992 & 48812 & 0.159 & ~5400 \\\\ SWP\\,50760 &  May 14, 1994 & 49486 & 0.946 & ~6300\\\\ SWP\\,51109 &  Jun 16, 1994 & 49520 & 0.986 & 23640 \\\\ SWP\\,52838 &  Nov 20, 1994 & 49676 & 0.168 & 16200 \\\\ SWP\\,54471 &  Apr 20, 1995 & 49827 & 0.345 & 21600 \\\\ SWP\\,55894 &   Sep 10, 1995& 49970 & 0.513 & 18000 \\\\ \\hline Z1705107M &  Oct 11, 1993 & 49272 & 0.696 & 1795  \\\\ Z1705109T &  Oct 11, 1993 & 49272 & 0.696 & ~~707  \\\\ Z170510BT &  Oct 11, 1993 & 49272 & 0.696 & ~~816  \\\\ Z170510EM &  Oct 11, 1993 & 49272 & 0.696 & ~~816  \\\\ \\noalign{\\smallskip} \\hline \\end{tabular}  \\noindent $^1$  Phase calculated from ${\\rm Min}(m_{\\rm pg}/B)={\\rm JD}\\,2442687.1 + 855.6 \\times E$ (Belczy{\\'n}ski et al. \\cite{bel00}). \\end{table}  A systematic shift in radial velocity between ultraviolet intercombination lines and ultraviolet high ionization permitted resonance emission lines in symbiotic binary spectra, was found by Friedjung et al. (\\cite{fsv83}). The latter are generally redshifted with respect to the former, and Miko{\\l}ajewska \\& Friedjung ({\\cite{mf05}) have recently reported an orbital variation of this redshift in {\\object CI Cyg}. Unlike the intercombination lines, the resonance lines are expected to be optically very thick. Their redshift might be due to the absorption of both the continuum and part of the emission line by P Cygni profile absorpion components, or only to a radiative transfer effect produced by photon scattering in an expanding medium.  In the present work, we have examined in detail the effect for the symbiotic binary, {\\object CI Cyg}, and we find some support for the P Cygni profile interpretation. Our data base is presented in  Sect.2, analyzed in Sect.3, and the results discussed in Sect. 4.  We conclude with a brief summary of these results in Sect. 5.   \\begin{table} \\caption[]{Emission line fluxes and radial velocities from GHRS spectra (MJD=49271/$\\phi=0.696$)} \\label{hst} \\begin{tabular}[bottom]{lcc} \\hline \\noalign{\\smallskip} Line ID & Flux$^1$ & $v_{\\rm rad}$ [km\\,s$^{-1}$] \\cr \\noalign{\\smallskip} \\hline \\ion{Si}{iv}\\, 1393.755 & ~28 &  34 \\cr \\ion{Si}{iv}\\,1402.769 & ~19 & 29 \\cr \\ion{O}{iv}]\\,1397.166 & ~~3 & 35 \\cr \\ion{O}{iv}]\\, 1399.731 & ~~7 & 25 \\cr \\ion{O}{iv}]\\,1401.115 & ~36 & 26 \\cr \\ion{O}{iv}]\\,1404.740 & ~15 & 33\\cr \\ion{O}{iv}]\\,1407.333 & ~~9 & 26 \\cr \\ion{S}{iv}]\\,1406.004 & ~~5& 23 \\cr \\ion{S}{iv}]\\,1416.872 & ~~4 & 20 \\cr \\ion{C}{iv}\\,1548.202 & 350 & 32 \\cr \\ion{C}{iv}\\,1550.774 & 234 & 30 \\cr \\ion{He}{ii}\\,1640.470 & 205 & 12 \\cr \\ion{O}{i}]\\,1641.310 & ~~6 & 14 \\cr \\ion{O}{iii}]\\,1660.800 & ~25 & 20 \\cr \\ion{O}{iii}]\\,1666.150 & ~68 & 18 \\cr \\ion{Mg}{v}\\,2783.03 & 21 & 89 \\cr \\ion{Mg}{ii}\\,2796.352 &~30 & ~39 \\cr \\ion{Mg}{ii}\\,2803.531 & ~22 & ~35 \\cr \\noalign{\\smallskip} \\hline \\end{tabular}  \\noindent $^1$  $10^{-14}\\, \\rm erg\\,s^{-1}cm^{-2}$. \\end{table} ",
        "conclusions": "Archival ultraviolet spectra taken by IUE and at higher resolution at one epoch by the GHRS/HST are studied by us. We examined radial velocities and line fluxes as well as higher resolution HST line profiles, taken at one epoch. The line fluxes were used to determine electron densities and emission measures; line formation in regions rather smaller than the binary separation being indicated.  The relative systemic velocity  shifts of the \\ion{He}{ii} and intercombination lines may be explicable by expansion, possibly in a medium above an optically thick accretion disk, with stratification. In this framework changes in disk properties might be responsible for changes in the intercombination line radial velocity between decline and quiescence.  A systematic redshift between the optically thick resonance lines on the one hand and the optically thin intercombination lines and the usually not optically very thick \\ion{He}{ii}\\,1640\\,\\AA\\  line on the other hand, like that of Friedjung et al. (\\cite{fsv83}), was confirmed by us. Two possible explanations still exist. We favour that involving  a non-classical P Cygni profile due to large circum-binary absorbing region, which is mainly expanding, but of which parts (perhaps inner parts) are contracting towards the binary. This region is most  probably an asymmetric wind interaction shell or a wind from the accretion disk. Such a region could also produce the observed radio emission. The other explanation of the redshift, as due to radiative transfer effects, cannot yet however be completely eliminated. More theoretical work on the production of such a line shift is still required."
    },
    "0609/astro-ph0609394_arXiv.txt": {
        "abstract": "We present infrared spectroscopy of the recurrent nova RS~Ophiuchi, obtained 11.81, 20.75 and 55.71~days following its 2006 eruption. The spectra are dominated by hydrogen recombination lines, together with \\pion{He}{i}, \\pion{O}{i} and \\pion{O}{ii} lines; the electron temperature of $\\sim10^4$K implied by the recombination spectrum suggests that we are seeing primarily the wind of the red giant, ionized by the ultraviolet flash when \\rs\\ erupted. However, strong coronal emission lines (i.e. emission from fine structure transitions in ions having high ionization potential) are present in the last spectrum. These imply a temperature of $930\\,000$~K for the coronal gas; this is in line with x-ray observations of the 2006 eruption. The emission line widths decrease with time in a way that is consistent with the shock model for the x-ray emission. ",
        "introduction": "RS~Ophiuchi is a recurrent nova that has undergone nova eruptions in 1898, 1933, 1958, 1967, 1985, and possibly \\cite{1907} 1907. As in the case of a classical nova, the eruption follows a thermonuclear runaway on the surface of the white dwarf \\cite{tnr}.  The key differences between classical and recurrent novae, in terms of both system properties and outburst behaviour, are reviewed by Anupama \\shortcite{anupama}. The recurrents are a heterogeneous class of objects but the \\rs\\ type is characterized by a semi-detached binary consisting of a roche-lobe-filling M~giant mass donor (M8III in the case of \\rs; \\cite{fekel}) and a massive (close to the Chandrasekhar limit) white dwarf ({\\em classical} novae almost exclusively have cool {\\em dwarf} mass donors).  \\begin{table*} \\begin{center} \\caption{Observing log.} \\begin{tabular}{cccccccccc} \\hline 2006 UTC Date & Day &\\multicolumn{2}{c}{Exp (s)} &  \\multicolumn{4}{c}{Resolution} & Comment  \\\\ &     & $IJ$ & $HK$                     & $I$ & $J$ & $H$ & $K$ &   \\\\ \\hline Feb 24.64    & 11.81 & 80 & 80 &  350 & 350   & 1000 & 1000 &        \\\\ Mar 5.58     & 20.75 & 80 & 80 &  350 & 350   & 1500 & 1500 & Cloud  \\\\ Apr 9.54     & 55.71 & 80 & 80 &  350 & 350   & 1500 & 1500 & Thin cirrus \\\\ \\hline \\end{tabular} \\end{center} \\label{obs} \\end{table*}   \\begin{figure*} \\setlength{\\unitlength}{1cm} \\begin{center} \\leavevmode \\begin{picture}(5.0,5.0) \\put(0.0,4.0){\\special{psfile=data.eps hoffset=-170 voffset=40 hscale=28 vscale=28 angle=-90}} \\put(0.0,4.0){\\special{psfile=data2.eps hoffset=70 voffset=40 hscale=28 vscale=28 angle=-90}} \\end{picture} \\caption[]{(a) 0.87-2.51\\mic\\ spectra for Feb 24.64 (upper spectrum, data shifted upwards by 30\\%, for clarity) and April 9.54 (lower spectrum). (b) detail around 2.4\\mic\\ for Feb 24.64 (upper spectrum) and April 9.54 (lower spectrum). In both figures the H recombination lines are identified by ``$n-m$''.} \\end{center} \\label{data} \\end{figure*}  The 1985 eruption of \\rs\\ was the first to have been observed over the entire electromagnetic spectrum, from the radio to the x-ray (see Bode 1987). What distinguishes the evolution of the eruption in the case of \\rs\\ is the fact that the ejected material runs into the dense red giant wind, which is shocked \\cite{bk}. Observations of the 1985 eruption provided indirect evidence for the shocking of the wind and the ejecta: \\rs\\ was a strong and rapidly evolving x-ray \\cite{mason} and radio source \\cite{padin}, and there was coronal emission over a wide range of wavelengths \\cite{snijders,evans88,shore}.  Infrared (IR) observations of the 1985 eruption are presented in Evans et al. \\shortcite{evans88}, starting on day~23 of the outburst. These authors found that the 1--2.5\\mic\\ spectrum was dominated by hydrogen recombination lines, and \\pion{He}{i}$\\lambda$1.083\\mic. Coronal lines (\\fion{Si}{vi}$\\lambda1.965$\\mic\\ and \\fion{Si}{vii}$\\lambda2.481$\\mic) were present on day 143 of the eruption. They also tentatively noted the presence of first overtone CO emission, although this was based on low resolution circular variable filter data.  \\rs\\ was discovered in eruption \\cite{hirosawa} on 2006 February 12.83, which we take to be day zero for this outburst. The discovery of the 2006 eruption triggered a multi-wavelength campaign of observations \\cite{bode-iauc,bode-b,das,evans06,eyres06a,eyres06b,ness-a,ness-b,ness-c,charo,obrien-iauc,obrien-nature} and yielded for the first time direct evidence, from VLBI imaging, for an expanding shock \\cite{obrien-nature}.  In this paper we present IR spectroscopy of \\rs, obtained in the first 55~days of the 2006 eruption, covering a much earlier phase than that observed in 1985. ",
        "conclusions": "We have reported the early IR spectroscopy of the 2006 eruption of the recurrent nova \\rs, covering the first 55~days. We find a spectrum dominated by hydrogen recombination lines arising from a gas at $\\sim10^4$~K; silicon coronal lines prominent on day~55, implying a temperature for the coronal gas of 930\\,000~K.  IR (and other) observations of this remarkable object are continuing and in subsequent papers we will present contemporaneous observations carried out with UKIRT and the Spitzer Space Observatory."
    },
    "0609/astro-ph0609488_arXiv.txt": {
        "abstract": " ",
        "introduction": "In the last decade, work on Hickson Compact Groups (HCGs) has been characterized by a sense of disappointment. On one hand, HCGs of galaxies, with their high spatial densities and low velocity dispersions appear to be ideal places for interactions and mergers. On the other hand, the interaction-activity paradigm (e.g. Barnes \\& Hernquist 1992) predicts the triggering of star formation and a widespread flux increase at almost every wavelength as a result of these interactions. Although observations have found abundant traces of tidal interactions in HCG galaxies (Mendes de Oliveira \\& Hickson 1994), their star formation levels are surprisingly similar to those found in isolated galaxies (e.g. Zepf, Whitmore \\& Levison 1991; Moles et al. 1994; Allam et al. 1996; Verdes-Montenegro et al. 1998; Iglesias-P\\'aramo \\& V\\'\\i lchez 1999).  There is evidence that this unexpected behavior is not exclusive to the HCGs, but is shared by the broader family of galaxy groups and clusters. The exploitation of large scale redshift surveys has allowed the study of large samples of galaxies in clusters (Hashimoto et al. 1998) and groups (Wilman et al. 2005). One of the most appealing results of these works is the realization that the dense galactic environments, found in groups and clusters, are hostile to star formation. Unfortunately, the mechanism which truncates the star formation in groups remains virtually unknown. Although, there is no shortage of proposed mechanisms to halt star formation in dense environments, they are considered ineffective in group conditions. In addition, the merger of galaxies, favored by the conditions in HCGs, has also been usually excluded because it enhances the star formation instead of stopping it.  This view has been modified based on recent hydrodynamical simulations of galaxy mergers by Di Matteo, Springel \\& Hernquist (2005), which take the feedback from supermassive black holes into account. They show that the star formation enhancement is just a transient phase of a longer process which, ends up depleting the gas and quenching both star formation and nuclear activity. A short time ($\\sim$ 0.5 Gyr) after the merger event, the remnant has evolved to a dead spheroid where the tracers of the transient starburst begin to fade. The short-lived SF tracers, such as the H$\\alpha$ emission, disappear along with the gas, while the color, which slowly reddens after star formation has ceased, reaches u-r $\\sim$ 2.2--2.3. This places the remnant in the red sequence of the bimodal color distribution in about 1-2 Gyr (Springel, Di Matteo \\& Hernquist 2005). From this new perspective, the failure of HCGs to display traces of ongoing or recent star formation would not seem so disappointing.  A more appropriate tool to determine the star formation histories (SFH) of galaxies is the study of their stellar populations. The use of this technique has produced valuable insights into the history of large samples of galaxies in high and low-density environments (Kuntschner et al. 2002; Thomas et al 2005). There also exist a few stellar population studies with smaller samples of early-type galaxies in HCGs (de la Rosa et al. 2001b; Proctor et al. 2004; Mendes de Oliveira et al. 2005). Their common result is that early-type galaxies in HCGs are older and more metal-poor than their counterparts in the field. Opposite results have also been reported (Ferreras et al. 2006). Our goal in this work is to use stellar populations to interpret the relics of the past episodes of galaxy mergers. The present study supersedes a previous work (de la Rosa et al. 2001b), using the same sample but with a less sophisticated analysis.  The study of stellar populations in galaxies has reached maturity. It has overcome several obstacles, like the age-metallicity degeneracy, which was partially lifted with the introduction of suitable Lick spectral indices (Worthey 1994). The non-solar element abundance ratios in the galaxies was taken into account with the inclusion of the abundance ratios effects in the stellar population models (Weiss, Peletier \\& Matteucci 1995; Vazdekis et al. 1997; Tantalo, Chiosi \\& Bressan 1998; Trager et al. 2000a; Thomas, Maraston \\& Bender 2003, hereafter TMB03; Barbuy et al. 2003). Despite these successful achievements, stellar population studies are still not problem-free. Although novel approaches and better models appear at a good pace, the convergence of their results is not accompanying these improvements. Different stellar population approaches applied to the same galaxy do not generally agree on the absolute values of the derived parameters.  Another drawback of stellar population studies arises from the relatively frequent emission contamination of the Balmer absorption indices, which are extensively used to determine galaxy ages. In the present study, the emission fill-in of the Balmer lines is a potential source of systematic errors, due to a possible different incidence of AGNs in the two environments (Coziol, Iovino \\& de Carvalho, 2000). If left uncorrected, the emission fill-in weakens the Balmer absorption lines and leads us to derive systematically larger ages and lower metallicities. Some authors have abandoned emission correction of their galaxies, simply excluding the contaminated cases from their samples.  In this work we use medium-resolution spectra of a homogeneous sample of early type galaxies belonging to both HCGs and the field or loose groups. Our goal is to study the effect of the environment on the stellar population parameters (age, metallicity and abundance ratio) of galaxies within the sample. For this reason, the absolute values of these parameters are much less relevant than their relative values. We have taken special care to double-check the influence of both the choice of a particular stellar population analysis and the emission contamination on the results. We have carried out three parallel analysis of the same data using very different approaches. We first apply the Lick procedure with the low resolution TMB03 models, including the effects of non-solar element abundance ratios. We then use the SSP models presented in Mendes de Oliveira et al. (2005, hereafter MCGB05), which also work at the Lick resolution and include non-solar abundance effects. Because the work by Mendes de Oliveira et al. (2005) contains two kinds of data relevant for this study, i.e. SSP models as well as galaxy parameters, we use the acronym MCGB05 to reference the models, and the full reference to cite the galaxy results, especially in subsection 5.2. Third is the Vazdekis (1999, hereafter Vaz99) approach, with a high resolution Stellar Population (SSP) model and a different stellar population analysis strategy, where the models are adapted to the resolution of the data, instead of the usual converse method. Our aim is not to compare the performances of the three approaches, but to gain insight into the systematic differences to assess the robustness of the results. The emission correction has also been calculated following two distinct strategies. One of them, pioneered by Gonz\\'alez (1993) (hereafter G93), calculates the H$\\beta$ emission contamination from the emission of the nearby [OIII]$\\lambda$5007 line. The second approach, introduced by Caldwell, Rose \\& Concannon (2003), estimates the H$\\beta$ fill-in through several steps, involving the determination of the H$\\alpha$ emission and the use of the ratios between the Balmer line strengths in nebular emission spectra (Osterbrock  1989).  This paper is organized as follows: Section 2 describes the observations and the basic reduction process. In Section 3, the emission correction is discussed. Section 4 is devoted to the extraction of the stellar population parameters with three alternative methods, whose results and correlations are shown in Section 5. In Section 6 a physical scenario is proposed to account for the results, and a summary is presented in Section 7. ",
        "conclusions": ""
    },
    "0609/astro-ph0609731_arXiv.txt": {
        "abstract": "Flat density cores have been obtained for a limited number of clusters of galaxies by strong gravitational lensing. Using a phenomenological equation of state (EOS) describing the full-to-partial degeneracy, we integrate the equation of hydrostatic equilibrium. The EOS is based on an assumption that the local kinetic energy of a classical particle induced by the gravity dissolves the quantum statistical degeneracy. The density profile is uniquely determined by four parameters, the central density, $\\rho(0)$, the properties of a fermion, namely, the mass, $m$, and statistical weight, $g$, and the ratio of the total matter density and fermion density, $\\delta$. As a case study, we model the column density and 2D encircled mass profiles of A1689, whose column density profile has been observationally obtained by Broadhurst et al., using gravitational lensing. The column density and 2D encircled profiles at the core, are reasonably reproduced for models with a limited range of particle properties. In the case that previously unknown fermions with spin 1/2 dominate the dark matter, the acceptable particle mass range is between 2 and 4 eV. In the case that the dark matter consists of the mixture of degenerate relic neutrinos and classical collisionless cold dark matter particles, the mass range of neutrinos is between 1 and 2 eV, if the ratio of the two kinds of dark matter particles is fixed to its cosmic value. Both the pure fermionic dark matter models and neutrino-CDM-mixture models reproduce the observations equally well. ",
        "introduction": "In most of the cosmological studies, dark matter particles are treated as collisionless particles. For instance, to describe the cosmological density evolution, collisionless Boltzmann equation is adopted. In cosmological N-body simulations, a group of collisionless dark matter particles, are represented by a single collisionless particle in the computer from the coarse-grained point of view. In the treatment by the Boltzmann equation, in which the distribution function is defined in a phase space, the particle mass and statistical weight are explicitly dealt with, although the dark matter particles are still regarded as classical particles without interaction except for gravity. In the coarse-grained view adopted by N-body simulations, there is no particle information and only the global mass density distribution is obtained.  However, by adopting the assumption that dark matter particles are collisionless classical particles, we might have lost basic physics in some cases. At high number densities and relatively low temperature, low-mass elementary particles, experience quantum statistical degeneracy due to indistinguishability of identical particles. For instance, it used to be well known that the neutrino black body in the early universe was partially degenerate \\citep{Weinberg1962}. Neutrinos are fermions and the radiation pressure of the neutrino black body can be interpreted as the combination degeneracy pressure and thermal pressure. As we discuss later in this paper, massive relic neutrinos are likely to remain partially degenerate, after decoupling and even after they become nonrelativistic under adiabatic expansion.  After nonlinear evolution of a high-density part of the universe, we might see a high concentration of dark matter particles around the center of a cluster of galaxies. If dark matter particles are composed at least partially of light fermions (e.g. neutrinos), their degeneracy pressure may be large enough to support the density structure near the center of the cluster against gravity. A self-gravitating system supported by degeneracy pressure of fermions, such as a white dwarf or a neutron star, is known to have a flat-top density profile. This is our motivation to explore the possibility that recent results regarding the mass profiles of clusters of galaxies obtained by gravitational lensing \\citep{Tyson98,Sand02,Sand04,TB05A,TB05B} might be explained by the degeneracy pressure of light fermionic dark matter particles.  In this paper, using a phenomenological equation of state (EOS) that describes the physical conditions between fully degenerate fermionic gas and the classical ideal gas, we integrate the equation of hydrostatic equilibrium, under the simple assumption that the local kinetic energy of a classical particle is equal to its gravitational energy determined by the 3D encircled mass. Our model is expected to be valid only near the core of a cluster where dynamical equilibrium is possibly achieved. For pure fermions, the volume density profile is uniquely determined by three parameters, the central density, $\\rho(0)$, and the properties of dark matter particles, namely, the mass, $m$, and statistical weight, $g$. To compare our model with observations, we smoothly connect our model volume density profile describing the inner region to a volume density profile derived from the observed column density profile by assuming spherical symmetry at a radius near the Einstein radius. In this way our model column density profiles can be directly compared with strong lensing results. In reality, there are baryons and possibly other forms of dark matter such as massive cold dark matter particles. These general cases are dealt with by introducing another input parameter, the ratio of the total mass density and light fermion density, $\\delta$.  A1689 studied by \\citet{TB05A,TB05B} is the best studied cluster by both strong and weak lensing, for which column density profile is obtained from the core to radii greater than 1 Mpc. As a case study, we apply our modeling to this cluster and constrain the possible combinations of particle properties. As a natural candidate for light fermions, we consider the case of massive neutrinos.  The paper is organized as follows. We first obtain the volume density of A1689 from the observed column density profile of \\citet{TB05B} and show that eV fermions can be degenerate near the core of this cluster in \\S2. Our modeling procedure is described in \\S3 and the properties of input fermions are discussed in \\S4, and then the results are presented in \\S5. The moderate degeneracy of unbound relic neutrinos and the plausibility that they can fall into the cluster core, are discussed in \\S6. ",
        "conclusions": "In this paper, we have shown that the observed mass profile at the center of the cluster of galaxies, A1689, is reproduced by models assuming the presence of degenerate fermionic dark matter. In the case that previously unknown fermions with spin 1/2 dominate the dark matter, the acceptable particle mass range is between 2 and 4 eV. In the case that the dark matter consists of the mixture of the degenerate relic neutrinos and classical collisionless cold dark matter particles, the mass range of neutrinos is between 1 and 2 eV, if the ratio of the two kinds of dark matter particles is fixed to its cosmic value. Both the pure fermionic dark matter models and neutrino-CDM-mixture models reproduce the observations equally well."
    },
    "0609/astro-ph0609507_arXiv.txt": {
        "abstract": "Supernova experiments to characterize dark energy require a well designed low redshift program; we consider this for both ongoing/near term (e.g.\\ Supernova Legacy Survey) and comprehensive future (e.g.\\ SNAP) experiments. The derived criteria are: a supernova sample centered near $z\\approx0.05$ comprising 150-500 (in the former case) and 300-900 (in the latter case) well measured supernovae. Low redshift Type Ia supernovae play two important roles for cosmological use of the supernova distance-redshift relation: as an anchor for the Hubble diagram and as an indicator of possible systematics.  An innate degeneracy in cosmological distances implies that 300 nearby supernovae nearly saturate their cosmological leverage for the first use, and their optimum central redshift is $z\\approx0.05$.  This conclusion is strengthened upon including velocity flow and magnitude offset systematics.  Limiting cosmological parameter bias due to supernova population drift (evolution) systematics plausibly increases the requirement for the second use to less than about 900 supernovae. ",
        "introduction": "\\label{sec:intro}  Type Ia supernovae (SN) observations discovered the acceleration of the universe \\cite{perl99,riess98} and have proved central in progress elucidating the nature of the dark energy responsible \\cite{snls,knop03,riess04}. SN provide a clear, direct, and mature method for mapping the expansion history of the universe, $a(t)$, with their measured flux giving the distance and hence lookback time $t$ through the cosmological inverse square law and their redshift $z$ giving the scale factor $a$.  Ground-based SN surveys are already underway to obtain hundreds of SN in the range $z=0.2-0.9$, suitable for measuring an averaged, or assumed constant, dark energy equation of state. Plans are well advanced for a comprehensive SN experiment aimed at accurate determination of dark energy properties including dynamics, in the form of a wide field space telescope exquisitely characterizing some 2000 SN over the range $z=0.2-1.7$, with launch planned by 2013.  We consider here what specific low redshift (``local'': $z\\lesssim0.2$) SN program would strengthen either of the higher redshift SN programs, combined with Planck CMB constraints on the distance to last scattering.  (We do not consider adding further probes, since it is useful to obtain a answer through purely geometric probes, as well as comparing results from different techniques.)  Just as \\cite{lh03,fhlt} brought into sharp relief the importance of SN at $z\\gtrsim1.5$ for probing dark energy, due to their ability to break degeneracies, control systematics, and realistically follow dark energy dynamics, here we show similar crucial roles of SN at $z\\lesssim0.1$. Local SN serve two key purposes: leverage in breaking degeneracies by anchoring the low redshift Hubble, or magnitude-redshift, diagram, and providing a well characterized set of SN to search for systematic magnitude effects through subclassification (``like subsets'').  In \\S\\ref{sec:cos} we demonstrate the cosmological need for a low redshift sample, and follow this in \\S\\ref{sec:nz} with detailed analysis of the numbers and redshift distribution for optimum complementarity with the higher redshift SN. Issues of systematics including evolution and intrinsic dispersion are discussed in \\S\\ref{sec:sys}, and a two stage survey program outlined to fulfill the required science criteria. ",
        "conclusions": "\\label{sec:further}  Local supernovae serve critical roles for maximal science return from supernova cosmology experiments -- anchoring the Hubble diagram and identifying possible systematics -- and failure to supply these data can cost a factor of 2 in cosmology parameter determination. General considerations of cosmology in the form of parameter degeneracies, systematics, and biases impose reasonable and somewhat conservative criteria for a low redshift SN program to match and strengthen a comprehensive, next generation higher SN program: 300-900 SN centered at $z\\approx0.05$.  We have tested the robustness of this conclusion by including coherent and random velocity systematics, magnitude offsets, and different intrinsic dispersions.  We propose here a two stage approach to the local sample, with the foundational set composed of 300 well characterized local SN.  This data then determines the necessary size for an extension, or second stage, which the arguments presented here suggest may require 100-600 additional SN to control systematic biases.  The first stage would likely supply the necessary local set to fully match ongoing and near term SN surveys, which we concluded also are optimized by a local data set centered at $z\\approx0.05$.  Though we used straightforward cosmological considerations, for more detail we could also seek future guidance from SN theory on questions such as maximum expected magnitude evolution, population drift predictions, and characteristics for subclasses, and from calibration and measurement error models and simulations.  This detail may be unnecessary and is model dependent.  The empirical foundations -- from well characterized (i.e.\\ spectroscopic time series measured) SN, not just more SN -- will define the relation between subsets and subclasses, and the population distribution (e.g.\\ local percentage of low metallicity SN, differing in the mean magnitude by $X$ mag). The data will be the final arbiter of the sample size required\\footnote{Note that much valuable information will be gained from surveys, even without full spectroscopic followup, including Carnegie Supernova Project \\cite{csp}, CfA Nearby SNe \\cite{cfa}, CTI-II \\cite{cti}, Essence \\cite{essence}, KAIT \\cite{kait}, Nearby Supernova Factory, SDSS II \\cite{sdss2}, SkyMapper \\cite{skymap}, Supernova Legacy Survey, and Texas Supernova Search \\cite{quimby}.}.  To achieve understanding of the physics behind the acceleration of the universe, we need a well designed SN program at $z<0.1$ just as we need a comprehensive SN experiment extending to $z>1.5$. Implementing the criteria derived here should ensure that we reach out into the universe with a firm foundation beneath our feet."
    },
    "0609/astro-ph0609595_arXiv.txt": {
        "abstract": "The multi-wavelength observations of the 2004 December~27 Giant Flare (GF) from \\sgr and its long-lived radio afterglow are briefly reviewed. The GF appears to have been produced by a dramatic reconfiguration of the magnetic field near the surface of the neutron star, possibly accompanied by fractures in the crust.  The explosive release of over $10^{46}\\;$erg (isotropic equivalent) powered a one-sided mildly relativistic outflow. The outflow produced a new expanding radio nebula, that is still visible over a year after the GF. Also considered are the constraints on the total energy in the GF, the energy and mass in the outflow, and on the external density, as well as possible implications for short $\\gamma$-ray bursts and potential signatures in high energy neutrinos, photons, or cosmic rays. Some possible future observations of this and other GFs are briefly discussed. ",
        "introduction": "Magnetars are a small class of young, isolated, neutron stars with extremely large magnetic fields of up to $\\sim 10^{15}\\;$G at the surface \\citep{dun92,pac92,kou98}. These objects give rise to occasional bursts peaking in the hard X-ray to soft $\\gamma$-ray range, thus manifesting themselves as Soft Gamma-ray Repeaters (SGRs) and Anomalous X-ray Pulsars (AXPs) \\citep[for a review see][]{WT04}. They have quiescent X-ray luminosities of $\\sim 10^{34}-10^{35}\\;{\\rm erg\\;s^{-1}}$, and are seen to pulse in the X-rays with periods of $P \\sim 5-12\\;$s (the pulsed fraction of the flux is $\\sim 5-11\\%$ for SGRs and $4-60\\%$ for AXPs). The pulsation period is identified with the rotational period of the star.  They have rapid spin-down rates, $\\dot{P} \\sim 10^{-11}-10^{-10}\\;{\\rm s\\;s^{-1}}$, leading to age estimates of $P/2\\dot{P} \\sim 10^3-10^4\\;$yr.  They have no detectable quiescent radio emission. From the rapid spin-down, and lack of any signs of accretion (e.g., infrared excess, radio emission) {or orbital modulation of their emission}, they are thought to be isolated (and not in binary or multiple star systems).  Only 13 magnetars are known \\citep{har05}, all in our Galaxy or the Large Magellenic Cloud (LMC) satellite galaxy.  The SGRs are more boisterous compared to the AXPs, giving rise to frequent bursts (with energies of $\\lesssim 10^{41}\\;$erg and typical durations of $\\sim 0.2\\;$s), and on rare occasions, also emit giant flares.  SGRs also have somewhat larger period derivatives, and shorter spin-down ages, than AXPs.  SGR-like flares have been detected from a few AXPs \\citep{kas03}, supporting the notion that SGRs and AXPs can be unified within the magnetar model.  Quite rarely, about once every $\\sim 50\\;$yr per source, SGRs produce a giant flare (GF). A GF consists of a very bright initial spike that peaks in soft gamma rays (a few hundred keV) and lasts for about a quarter of a second, followed by a longer and dimmer tail that peaks in the hard X-rays ($\\sim 10\\;$keV) and lasts for a few hundred seconds, with a strong modulation at the rotational period of the neutron star. So far only three GFs have been detected, originating from within our galaxy or the LMC.  On 1979 March 5, SGR~0526$-$66 in the LMC produced the first GF that could be witnessed using satellites in orbit about the Earth.  The peak (isotropic equivalent) luminosity from this flare was estimated at $\\sim 4\\times 10^{44}\\;{\\rm erg\\;s^{-1}}$, which exceeds the luminosity of the entire Galaxy, for the 0.2 seconds of the peak emission. A fading tail lasted 3 minutes and exhibited strong oscillations with a period of 8.1 seconds \\citep{maz79}. The total isotropic equivalent energy release was $\\sim 5\\times 10^{44}\\;$erg. A similar GF event was detected from SGR~1900$+$14 on 1998 August 27, with an isotropic equivalent peak luminosity and total energy in excess of $3\\times 10^{44}\\;{\\rm erg\\;s^{-1}}$ and $10^{44}\\;$erg, respectively. Its initial spike and tail lasted for $\\sim 0.35\\;$s and $\\sim 400\\;$s, respectively. A faint radio afterglow was detected in the days following this GF \\citep{fra99}. In both of these two GFs the (isotropic equivalent) energy in the initial spike was comparable to that in the tail.  There are also some intermediate events between the short, more frequent, bursts and the giant flares, in terms of their total duration and isotropic equivalent peak luminosity and energy \\citep[see][and references therein]{WT04}. They often occur following giant flares. The largest so far of these intermediate events (on 2001 April 18, from SGR~1900+14) also showed several strong flux modulations at the stellar rotational period. All this suggests that there might be a continuum of events, differing mainly in their energy release, and consequently also in their duration and peak luminosity. It might make sense to roughly divide the different types of such events according to their isotropic equivalent energy release, $E_{\\rm iso}$: (i) short, more frequent, bursts ($E_{\\rm iso} \\lesssim 10^{41}\\;$erg), (ii) intermediate events ($10^{41}\\;{\\rm erg} \\lesssim E_{\\rm iso} \\lesssim 10^{44}\\;$erg), and (iii) giant flares ($E_{\\rm iso} \\gtrsim 10^{44}\\;$erg).  The last and most spectacular of the three GFs we have witnessed so far occurred on 2004 December 27 from GSR~1806$-$20.  The GF was preceded by a gradual change in the spectral hardness, photon index, and spin-down rate that peaked several months in advance of the GF \\citep{woods06}.  The sudden energy release of more than $10^{46}\\;$erg in gamma-rays (assuming isotropic emission at a distance of $15\\;$kpc) managed to eject a significant amount of baryons, probably accompanied by some pairs and magnetic fields, from the neutron star \\citep{pal05,gel05,gra06}. As this outflow interacted with the external medium, it powered an expanding radio nebula \\citep{cam05,gae05} at least 500 times more luminous than the only other radio afterglow detected from an SGR GF \\citep{fra99}.  We note that the energetics are reduced by a factor of 6$-$2 if one adopts a lower distance of 6$-$10$\\;$kpc found by \\cite{cam05b} based on \\HI\\ absorption observations.  However, a more detailed analysis by \\cite{MCC05} suggests that the \\HI\\ observations are consistent with the X-ray absorption measurements that give a distance of 14.5 $\\pm$ 1.4 kpc \\citep{cor97}, and with the distance of 15.1$^{+1.8}_{-1.3}$ kpc of the associated stellar cluster in G10.0$-$0.3 \\citep{CE04}. For this review we adopt a distance to \\sgr of $15d_{15}\\;$kpc, so that $1\\;$mas corresponds to $15d_{15}\\;$AU or $2.25\\times 10^{14}d_{15}\\;$cm.  In this brief review we first consider how the GF was powered along with implications for the neutron star (\\S~\\ref{subsec:GF_obs}-\\S~\\ref{subsec:GF_QPOs}), for the possible connection to short gamma-ray bursts (\\S~\\ref{subsec:GF_GRBs}), and for possible signatures in high energy neutrinos, photons or cosmic rays (\\S~\\ref{subsec:GF_HE}). We go on to discuss the radio afterglow that is still being studied over a year after the GF (\\S~\\ref{sec:radio}) and its implications for the properties of the outflow from the GF and its environment. We conclude with a short discussion of possible future work (\\S~\\ref{sec:future}). The primary focus is on the observations and the immediate insights gained from them.  A detailed discussion of the physics of the GF is beyond the scope of this review. ",
        "conclusions": ""
    },
    "0609/astro-ph0609076_arXiv.txt": {
        "abstract": "We present a study of the local $B$- and $K_s$-band Tully--Fisher Relation (TFR) between absolute magnitude and maximum circular speed in S0 galaxies.  To make this study, we have combined kinematic data, including a new high-quality spectral data set from the Fornax Cluster, with homogeneous photometry from the RC3 and 2MASS catalogues, to construct the largest sample of S0 galaxies ever used in a study of the TFR.  Independent of environment, S0 galaxies are found to lie systematically below the TFR for nearby spirals in both optical and infrared bands.  This offset can be crudely interpreted as arising from the luminosity evolution of spiral galaxies that have faded since ceasing star formation.  However, we also find a large scatter in the TFR.  We show that most of this scatter is intrinsic, not due to the observational uncertainties.  The presence of such a large scatter means that the population of S0 galaxies cannot have formed exclusively by the above simple fading mechanism after all transforming at a single epoch. To better understand the complexity of the transformation mechanism, we have searched for correlations between the offset from the TFR and other properties of the galaxies such as their structural properties, central velocity dispersions and ages (as estimated from line indices). For the Fornax Cluster data, the offset from the TFR correlates with the estimated age of the stars in the individual galaxies, in the sense and of the magnitude expected if S0 galaxies had passively faded since being converted from spirals.  This correlation implies that a significant part of the scatter in the TFR arises from the different times at which galaxies began their transformation. ",
        "introduction": "The Tully--Fisher relation (TFR; Tully $\\&$ Fisher 1977) is one of the most important physically-motivated correlations found in spiral galaxies. The correspondence between luminosity and maximum rotational velocity ($V_{\\rm max}$) is usually interpreted as a product of the close relation between the stellar and total masses of galaxies or, in other words, as the presence of a relatively constant mass-to-light ratio in the local spiral galaxy population (e.g. Gavazzi 1993, Zwaan et al. 1995). Such a general property in spirals puts strong constraints on galaxy formation scenarios (e.g. Mao et al. 1998, van den Bosch, 2000) and cosmological models (e.g. Giovanelli et al. 1997, Sakai et al. 2000). Also, the low scatter in the TFR (only $\\sim 0.35$ mag in $I$-band, according to Giovanelli et al. 1997, Sakai et al. 2000, Tully \\& Pierce 2000 and Verheijen 2001) permits us to use this tool as a good distance estimator (e.g. Yasuda et al. 1997).  Attempts to ascertain whether S0 galaxies follow a similar TFR have two main motivations.  First, if there is a TFR of S0s, it would prove useful for estimating distances in the nearby universe, particularly in clusters where S0s are very prevalent (Dressler 1980).  Second, and more related to the present study, a possible scenario where S0 galaxies are the descendants of evolved spirals at higher redshifts (Dressler 1980, Dressler et al. 1997) could leave traces of this evolution in the observed TFR of S0s.  Different mechanisms have been proposed as the channels for such evolution, like small mergers (Schweizer 1986), gas stripping in later types (Gunn \\& Gott 1972), halo stripping (Bekki et al. 2002) and galaxy harassment (Moore et al. 1998). If this kind of picture is correct, it would be expected that S0s retain some memory of their past as spirals, in particular through their TFR, and perhaps even some clues as to which of the channels they evolved down.  Only a few studies of the TFR for S0 galaxies can be found in the literature. The first effort, made by Dressler \\& Sandage (1983), found no evidence for the existence of a TFR for S0 galaxies. However, the limited spatial extent of their rotation curves, the inhomogeneous photographic photometry employed and the large uncertainties in the distances to their sample made it almost certain that any correlation between luminosity and rotational velocity would be lost in the observational uncertainties.  Fifteen years later, Neistein et al. (1999) explored the existence of a TFR for S0s in the $I$-band, using a sample of 18 local S0s from the field. Although some evidence for a TFR was uncovered in this study, they also found a large scatter of $0.7$ magnitudes in the relation, suggesting the presence of more heterogeneous evolutionary histories for these galaxies when compared to spirals. Also, a systematic shift $0.5$ magnitudes was found between their galaxies and the relation for local spirals.  In two papers, Hinz et al. (2001, 2003), explored the $I$- and $H$-band TFRs for 22 S0s in the Coma Cluster and 8 S0s in the Virgo Cluster. By using cluster data, they avoided some of the errors that arise from the uncertainty in absolute distances estimation. The analysis of $I$-band data from the Coma Cluster revealed very similar results to the study by Neistein et al. (1999), implying that the larger scatter of the latter could not be attributed to distance errors or the heterogeneous nature of the data.  In the $H$-band, an even larger scatter of 1.3 magnitudes was found, but with a smaller offset from the corresponding spiral galaxy TFR of only 0.2 magnitudes.  Interestingly, there was no evidence for any systematic difference between the results for the Virgo and Coma Clusters, despite their differences in richness and populations, implying that these factors could not be responsible for the scatter in the TFR.  Given the large scatter and small shift in the $H$-band TFR for S0s compared to spirals, it was concluded that these galaxies' properties are not consistent with what would be expected for spiral galaxies whose star formation had been truncated; instead they suggested that other mechanisms such as minor mergers are responsible for the S0s' TFR.  By contrast, Mathieu, Merrifield \\& Kuijken (2002) found in their detailed dynamical modeling of six disk-dominated field S0s that these galaxies obey a tight $I$-band TFR with a scatter of only 0.3 magnitudes, but offset from the spiral galaxy TFR by a massive 1.8 magnitudes.  The authors therefore concluded that these objects were consistent with being generated by passively fading spirals that had simply stopped producing stars.  This result does not appear to be consistent with the previous studies, although it should be borne in mind that the galaxies in this study were selected to be disk-dominated, so they morphologically resemble spiral galaxies more than those in other work.  In addition, their field locations means that they are less likely to have had their evolution complicated by mergers.  It is therefore possible that these S0s really are just passively-fading spirals where those in clusters have led more complicated lives.  As can be seen, there is no general consensus as to either the scatter or the shift in the TFR for S0 galaxies when compared to spirals, and so no agreement as to their interpretation.  We therefore revisit the TFR of S0 galaxies in this paper, adding our new spectral dataset from the Fornax Cluster to the existing results from literature.  In addition to the new homogeneous spectral data set, we can also take advantage of the 2MASS photometry for these galaxies (Jarrett et al. 2003), to obtain consistent infrared magnitudes and structural parameters for all the galaxies.  With this more systematic study of S0s, it is to be hoped that we can avoid any past problems that may have arisen due to the heterogeneous nature of the available data, to reveal the underlying physics that dictates the TFR in S0 galaxies.  The remainder of the paper is laid out as follows.  In Section~\\ref{sec:data}, we describe the different samples of S0s used in this study, and the data that we have adopted. Section~\\ref{sec:results} presents the main results and then discusses their implications. Finally in Section~\\ref{sec:conc} our conclusions are summarised. ",
        "conclusions": "\\label{sec:conc} This paper presents the largest compilation of data to-date in a study of the Tully--Fisher relation of nearby S0 galaxies in both $B$ and $K_s$-bands, using new and archival results for a sample of 60 galaxies in different environments.  The principal results of this study are:  \\begin{itemize}  \\item The local TFR of S0 galaxies presents a shift with respect to the local TFR of spirals. Interpreted as a shift in luminosity, it amounts to between $-1.3 \\pm 0.1$ and $-1.7 \\pm 0.4$ in the $B$-band and between $-0.8 \\pm 0.4$ and $-1.2 \\pm 0.4 $ in $K_s$-band, with the value depending on the calibration adopted for the spiral galaxy TFR.  \\item In a scenario where S0 galaxies are the descendants of spirals at higher redshifts which have simply passively faded, these offsets are consistent with a population that has been passively fading for somewhere between 1 and 6 Gyr (again depending on the adopted calibration for the spiral galaxy TFR).  If we include a starburst before the truncation of the star formation, the upper end of the possible timescale since transformation increases somewhat. However, the large scatter in the TFR means that such a simple model with a single epoch of formation cannot provide the whole story.  \\item The total scatter in the TFR of S0s is found to be $0.88 \\pm 0.06$ mag in the $B$-band and $0.98 \\pm 0.06$ in the $K_s$-band. Our fitting indicates that only $\\sim 10\\%$ of this scatter can be attributed to the observational errors, with $\\sim 90\\%$ arising from the intrinsic astrophysical spread in the TFR.  Such a scatter could have been ``imprinted'' by the larger scatter in the spiral galaxy TFR at high redshift, or it could have arisen through the subsequent somewhat stochastic evolution of these systems.  \\item To try to distinguish between these possibilities, we have looked for correlations between the offset from the TFR and other properties of these systems.  Morphologically, a number of S0 properties correlate with their absolute magnitudes, as seen in other galaxies, but none seem to correlate strongly with the offset of this magnitude from the spiral galaxy TFR.  \\item By investigating line indices, we have been able to establish a reasonably strong correlation between the central ages of the Fornax Cluster galaxies' stellar populations and their offsets from the TFR.  This offset is in the sense and of the magnitude expected for a population of galaxies that transformed from spirals at various times in the past and has been passively fading ever since. A central starburst before the truncation of star formation would explain why the correlation is stronger with central age than at larger galactocentric distances.  \\end{itemize}  The new data from the Fornax Cluster show that we are finally in a position to use observations of nearby S0 galaxies to obtain important archaeological clues as to the mechanisms by which these systems form. However, the relatively small number of S0s in a single cluster will always limit the significance of the resulting measurements.  Clearly, we need a still-larger homogeneous dataset of the quality now attainable with telescopes like the VLT of S0 galaxies to finish off the work begun here."
    },
    "0609/astro-ph0609289_arXiv.txt": {
        "abstract": "We present high resolution ($R \\approx 60,000$) circular spectropolarimetry of the classical T Tauri star TW Hydrae. We analyze $12$ photospheric absorption lines and measure the net longitudinal magnetic field for 6 consecutive nights. While no net polarization is detected the first five nights, a significant photospheric field of $B_{\\rm{z}} = 149 \\pm 33$ G is found on the sixth night. To rule out spurious instrumental polarization, we apply the same analysis technique to several non-magnetic telluric lines, detecting no significant polarization. We further demonstrate the reality of this field detection by showing that the splitting between right and left polarized components in these 12 photospheric lines shows a linear trend with Land\\'e $g$-factor times wavelength squared, as predicted by the Zeeman effect. However, this longitudinal field detection is still much lower than that which would result if a pure dipole magnetic geometry is responsible for the mean magnetic field strength of $2.6$ kG previously reported for TW Hya. We also detect strong circular polarization in the \\ion{He}{1} $5876$ \\AA\\ and the \\ion{Ca}{2} $8498$ \\AA\\ emission lines, indicating a strong field in the line formation region of these features. The polarization of the \\ion{Ca}{2} line is substantially weaker than that of the \\ion{He}{1} line, which we interpret as due to a larger contribution to the \\ion{Ca}{2} line from chromospheric emission in which the polarization signals cancel.  However, the presence of polarization in the \\ion{Ca}{2} line indicates that accretion shocks on Classical T Tauri stars do produce narrow emission features in the infrared triplet lines of Calcium. ",
        "introduction": "T Tauri stars (TTSs) are newly formed low-mass stars that have recently become visible at optical wavelengths. These young, roughly solar mass stars are still contracting along pre-main sequence evolutionary tracks in the H-R diagram. It is generally believed that classical T Tauri stars (CTTSs) are still surrounded by disks of material that are undergoing accretion onto the central star, producing excess emission in both lines and continuum at multiple wavelengths. Magnetospheric accretion models are the most popular description of the accretion process. Strong, stellar magnetic fields are believed to regulate the accretion and confine disk material to flow onto the stellar surface along the field lines (e.g., Camenzind 1990; K\\\"onigl 1991; Cameron \\& Campbell 1993; Shu et al. 1994). These models generally assume the magnetic structure of TTSs to be dipolar and require magnetic field strengths which vary over a wide range of values, with the field on some stars as high as several kilogauss (see Johns--Krull, Valenti \\& Koresko 1999b). Such high field strengths should be measurable by utilizing the most magnetically sensitive diagnostics.  On the other hand, direct magnetic field measurements are difficult, since T Tauri stars are relatively faint and display various spectral peculiarities. The most successful approach for measuring fields on late--type stars in general has been to measure Zeeman broadening of spectral lines in unpolarized light (e.g., Robinson 1980; Saar 1988; Valenti, Marcy \\& Basri 1995; Johns--Krull \\& Valenti 1996; Johns--Krull et al. 1999b; Johns--Krull, Valenti, \\& Saar 2004). This method is more efficient at infrared wavelengths thanks to the $\\lambda^{2}$ dependence of Zeeman broadening, compared to the $\\lambda^1$ dependence of Doppler broadening. While a sensitive measure of field strength, Zeeman broadening measurements give little information on the magnetic field geometry.  Another direct method for measuring magnetic fields is to detect net circular polarization in Zeeman-sensitive lines. Generally, Zeeman $\\sigma$ components are elliptically polarized and the components of opposite helicity are split to either side of the nominal line wavelength. A net longitudinal component of the magnetic field makes components of Zeeman-sensitive lines distinguishable through a right-circular polarizer (RCP) and left circular polarizer (LCP). The separation between the line observed in RCP and LCP light is $$\\Delta\\lambda = 2{e \\over 4\\pi m_ec^2} \\lambda^2 g_{\\rm{eff}} B_{\\rm{z}} = 9.34 \\times 10^{-7} \\lambda^2 g_{\\rm{eff}} B_{\\rm{z}} \\,\\,\\,\\,\\,\\, \\rm m\\mbox{\\AA}\\,\\eqno(1)$$ where $g_{\\rm{eff}}$ is the effective Land\\'e $g$-factor of the transition, $B_{\\rm{z}}$ is the strength of the mean longitudinal magnetic field in kilogauss, and $\\lambda$ is the wavelength of the transition in Angstroms (see Mathys 1988, 1991). The weights for individual $\\pi$ and $\\sigma$ components in the definition of $g_{\\rm eff}$ assume an optically thin medium, so eq.[1] is only approximately true in our case of moderately strong photospheric lines. Previously, Johns--Krull et al. (1999a) did not detect polarization in the photospheric lines of the CTTS BP Tau, setting a $3\\sigma$ upper limit on \\Bz\\, of $\\pm 200$ G.  Smirnov et al. (2003) report a longitudinal magnetic field of $B_{\\rm{z}} \\sim 150 \\pm 50$ G on the CTTS T Tau, which is very close to their detection limit, while their subsequent observation of T Tau (Smirnov et al. 2004) did not detect a significant field. In an effort to confirm the original Smirnov et al. (2003) detection, Daou et al. (2006) measure a mean longitudinal field of \\Bz$\\, = 12\\,\\pm \\,35$ G on T Tau.  Daou et al. (2006) use upper limits on \\Bz\\, on multiple nights along with the mean field strength detected on T Tau of $\\bar B \\sim 2.4\\ $kG (Guenther et al. 1999; Johns--Krull et al. 2001) to seriously question the assumed dipole field geometry (see also Valenti \\& Johns--Krull 2004).  On the other hand, Johns--Krull et al. (1999a) discovered net polarization in the \\ion{He}{1} $5876$ \\AA\\ emission line on the CTTS BP Tau, indicating a net longitudinal magnetic field of $2.46 \\pm 0.12$ kG in the line formation region. This \\ion{He}{1} emission line is believed to be produced, at least partially, in the shock region formed where disk material accretes onto the stellar surface (Hartmann, Hewett \\& Calvet 1994; Edwards et al. 1994). Circular polarization in the \\ion{He}{1} line has now been observed in several CTTSs (Valenti \\& Johns--Krull 2004; Symington et al. 2005). These observations suggest that accretion onto CTTSs is indeed controlled by a strong stellar magnetic field.   While substantial observational evidence indicates strong fields on the surface of TTSs, the origin of the surface magnetic fields on TTSs are not clear.  Interface dynamo models (e.g., Parker 1993) which are applied to solar-type main-sequence stars probably do not apply to TTSs since their internal structure is significantly different from that of the Sun. Feigelson et al. (2003) summarize several theoretical considerations for the origin of TTS magnetic fields. One possibility is that a distributed dynamo due to turbulent convection could operate in TTSs and generate small-scale magnetic fields. These fields could also be amplified by differential rotation throughout the convective zone (Durney et al. 1993; Kitchatinov \\& R\\\"udiger 1999; K\\\"uker \\& Stix 2001), though surface differential rotation is weak or absent on TTSs (Johns-Krull 1996). It is also possible that T Tauri stars simply maintain a ``fossil\" magnetic field throughout the star formation process and do not have significant field contributions from dynamo processes (Tayler 1987; Mestel 1999; Moss 2003).  More observational studies are needed to put further constraints on theories of the origin and evolution of magnetic fields on pre-main sequence stars.   In order to gain further insight into the magnetic properties of young stars, we present an analysis of high resolution spectropolarimetry of the CTTS TW Hya, a K7Ve(Li) star (Herbig 1978) located at a distance of $56 \\pm 7$ pc (Wichmann et al. 1998). Imaging in multiple wavelengths (e.g., Krist et al. 2000; Weinberger et al. 2002; Qi et al. 2004) indicates a nearly face-on disk around TW Hya that is optically thick in the visible, near-IR, sub-millimeter and millimeter wavelengths. Herczeg et al. (2004) estimate an accretion rate of $2 \\times 10^{-9}$\\ ${\\rm M}_\\odot\\ {\\rm yr}^{-1}$ from its strong optical and near-UV excess emission. Previously, we have modeled magnetic broadening of four K band \\ion{Ti}{1} lines observed in unpolarized light to measure a mean magnetic field strength of $\\bar B = 2.6 \\pm 0.2$ kG on TW Hya (Yang, Johns--Krull \\& Valenti 2005). Here, we use a time series of spectropolarimetric observations to study the geometry of the stellar magnetic field. The remainder of this paper is organized as follows. In \\S\\ 2 we describe our observations and data reduction procedures. In \\S\\ 3 we present the measurements of the longitudinal field in the photosphere as well as in the accretion regions. H$\\alpha$ line profiles are also presented in \\S\\ 3. Finally, a discussion of our results is presented in \\S\\ 4. ",
        "conclusions": "Imaging of the circumstellar disk around TW Hya in the infrared (Krist et al. 2000; Trilling et al. 2001; Weinberger et al. 2002), millimeter (Wilner et al. 2000) and sub-millimeter (Qi et al. 2004) wavelengths all suggest that the inclination of the disk is close to $0^\\circ$. Alencar \\& Batalha (2002) derived an inclination of $18^\\circ \\pm 10^\\circ$ from emission line profile analysis. Lawson and Crause (2005) derive $i \\sim 16^\\circ$ from photometric monitoring. The night-to-night variation in our measurements of the longitudinal field in the \\ion{He}{1} line is small, only about $10\\%$ of the mean value, which is also consistent with a low disk inclination angle of TW Hya.  Such a geometry allows a strong test whether the magnetic field on TW Hya is primarily a dipole field with the magnetic axis aligned with the rotation axis. Yang et al. (2005) find the mean magnetic field strength in the photosphere of TW Hya to be $2.6 \\pm 0.2$ kG from infrared (IR) Zeeman broadening measurements. If we follow Alencar \\& Batalha (2002) and assume an inclination angle between $8^\\circ$ and $28^\\circ$, and if the magnetic dipole axis is aligned with the rotation axis, the $2.6$ kG mean field would predict a mean line of sight field in the photosphere of \\Bz \\ $ = 0.97 - 1.05$ kG. This is much higher than our maximum measured value \\Bz$ = 149 \\pm 33$ G, or our weighted mean from all nights of \\Bz$= 90 \\pm 17$ G. One possibility to explain the low longitudinal magnetic field values yet retain a dipole field geometry is to assume the magnetic axis is highly inclined from the rotation axis (i.e. has high obliquity, $\\beta$). For example, Krist et al. (2000) use the model of Mahdavi \\& Kenyon (1998) and conclude that the photometric variability of TW Hya observed by Mekkaden (1998) can be explained if $i < 10^\\circ$ and the field is a dipole with $\\beta > 55^\\circ$. If we take $i = 10^\\circ$ and $\\beta = 55^\\circ$, the 2.6 kG mean field on TW Hya still implies a photospheric \\Bz$ = 537$ G, again well above our upper limits for \\Bz\\ in the photosphere. If the photospheric field is globally dipolar, then $\\beta$ must be significantly larger than $55^\\circ$. If we ask what is required of a dipole field geometry to match the 2.6 kG mean field and the 149 G longitudinal field for the visible hemisphere of the star, we find that $i + \\beta = 83.5^\\circ$. However, in this case we have an additional constraint from the \\ion{He}{1} polarization.    If the photospheric magnetic field on TW Hya is dipolar, then the magnetic poles must be nearly perpendicular to the line of sight at all rotational phases (see preceding paragraph). For such a magnetic geometry, material accreted from the corotation radius ($6.3\\ R_*$, see Johns--Krull \\& Valenti 2001) will land within $13^\\circ$ of the magnetic pole. In this case, the longitudinal field of $\\sim 1.7$ kG measured in the accretion region using the \\ion{He}{1} line is only the small fraction projected onto the line of sight of the true magnetic field in this region.  The minimum value of the true field is $1.7\\ {\\rm kG} / {\\rm cos(}83.5^\\circ - 13^\\circ{\\rm )} = 5.1$ kG. Such a strong field in the \\ion{He}{1} line has not been observed in any of the previous studies now covering 11 CTTSs (Johns--Krull et al. 1999a; Valenti \\& Johns--Krull 2004; Symington et al. 2005; Daou et al. 2006). We conclude that the surface topology of the magnetic geometry on TW Hya (and likely all CTTSs) is not a $pure$ dipole. It is likely that the field topology at the stellar surface is dominated by small scale structure such as seen on the Sun. However, the dipole component of the field will fall off the least rapidly with distance from the star, so it may well be that the interaction of the stellar field with the disk is governed by a dipole geometry. Such a picture may explain the relatively smooth, sinusoid like modulation of the field traced in the \\ion{He}{1} line of the CTTSs studied by Valenti and Johns--Krull (2004).  Hartmann (1998) gives an expression for the truncation radius for an assumed dipole field geometry (his equation 8.72) derived under the assumption of spherical accretion.  As discussed by Bouvier et al. (2006), this is an upper limit for a disk geometry.  If we assume that the dipole component of the field on TW Hya is responsible for our detection of \\Bz\\ $ = 149$ G on this star, we can use the equation from Hartmann (1998) to estimate a truncation radius.  The estimate again depends on the inclination of the dipole component of the field with respect to our line of sight. If (as is generally the case for the Sun) we assume the dipole component is aligned with the rotation axis, assuming $i = 28^\\circ$ gives the strongest possible dipole component consistent with our \\Bz\\ detection and this correponds to an equatorial field strength of 260 G. Putting this into equation (8.72) from Hartmann (1998) yields a truncation radius of 3 $R_*$ which is significantly less than the co-rotation radius of 6.3 $R_*$.  Eisner et al. (2006) analyze K-band interferometry observations of TW Hya from Keck, along with previous K-band veiling and NIR photometric measurements, to conclude that the inner radius of the optically thin disk is around 0.06 AU, which corresponds to 13 $R_*$.  While there are additional techniques that could shed light on the exact location of this inner truncation radius (e.g. linear polarimetry is a powerful technique which may help in this regard as discussed in Vink et al. 2005), the current data suggests that the system is not in equilibrium as assumed by magnetospheric accretion theories if the dipole component of the field alone is responsible for truncating the accretion disk. There are certainly higher order contributions to the total field at the surface of TW Hya which will contribute some field strength at the co-rotation radius, but just how much depends on the detailed field geometry. We suggest more work is needed to verify whether the fields on TTSs really are strong enough to truncate disks around these stars near the co-rotation radius for realistic magnetic field geometries.  As noted in \\S 3.2, the level of polarization in the \\ion{Ca}{2} $8498$ \\AA\\ line is much less than that detected in the \\ion{He}{1} line. The IRT lines of \\ion{Ca}{2} likely have contributions from different regions in CTTSs.  These lines sometimes show only broad components (BC), narrow components (NC), or a mixture of the two (e.g. Alencar \\& Basri 2000).  It is generally accepted that the BC of the IRT originates in the accretion and/or wind flows associated with CTTSs, as this component is not observed in NTTSs.  However, the origin of the NC of the IRT lines is not completely clear.  The \\ion{Ca}{2} $8498$ \\AA\\ line of TW Hya is dominated by a NC inside a photospheric absorption line.  A chromospheric origin for the NC of the IRT lines in both NTTSs and CTTSs was proposed by Hamann and Persson (1992).  Batalha and Basri (1993) and Batalha et al.  (1996) echo this idea, though they further suggest that accretion activity in CTTSs can enhance the chromospheric emission.  They suggest that this is due to the reprocessing of radiation produced in accretion shocks as the accreting material hits the stellar atmosphere. Another possible scenario results since the accretion occurs along stellar magnetic field lines: the process may launch Alfv\\'en waves in the magnetosphere which deposits energy in the chromosphere, heating it beyond what the star alone would do.  Observationally, there is almost certainly a chromospheric contribution to this line in TW Hya and the NC of other CTTSs given the persistent appearance of NC emission in NTTSs.  For example, Batalha et al. (1996) measured equivalent widths (EWs) of the \\ion{Ca}{2} $8498$ \\AA\\ line for 3 K7 NTTSs, finding values which range from $0.53$ \\AA\\ to $0.59$ \\AA.  These values are weaker than our EW measurements of TW Hya ($0.69 - 0.98$ \\AA), suggesting a contribution to the line in TW Hya from more than just pure chromospheric emission.  This mixture of sources can explain the relatively weak polarization observed in the \\ion{Ca}{2} line relative to the \\ion{He}{1} line.  We generally expect the magnetic field in the stellar chromosphere to show the same behavior as that traced by the photospheric absorption lines. In particular, the weak photospheric polarization suggests there should be comparably weak polarization in the IRT lines, and we further expect the implied field to be of the same polarity.  The field detected in the \\ion{Ca}{2} line is of opposite polarity to that observed in the photosphere, implying that the polarization in the enhanced portion of the IRT emission is more polarized than implied by the reported value of \\Bz\\ in this line. The polarity of the average field in the \\ion{Ca}{2} region is the same as that in the \\ion{He}{1} region on all nights, suggesting that their emission is related. Whether this implies actual formation of (some of) the IRT NC emission in the accretion shock or formation in nearby field regions of the same polarity is not clear.  This dilution of the \\ion{Ca}{2} polarization by chromospheric emission is not expected to be an issue for the \\ion{He}{1} line. The \\ion{He}{1} $5876$ \\AA\\ line of TW Hya is dominated by NC emission as well; however, the contribution to this from a pure stellar chromosphere is likely quite small. For example, Batalha et al. (1996) find the EW of the \\ion{He}{1} $\\lambda 5876$ for the same 3 K7 NTTSs to range between $0.00 - 0.09$ \\AA, while for TW Hya we find values from $2.5 - 4.3$ \\AA. Besides chromospheric activities, both magnetospheric infall and a hot wind could potentially give rise to the \\ion{He}{1} emission. They are believed to be responsible for the blue-shifted or red-shifted broad components of the $5876$ \\AA\\ line profiles in spectra of a number of T Tauri stars (e.g., see Baristain et al. 2001). In the case of TW Hya, an obvious BC is not seen. Thus, the \\ion{He}{1} line in TW Hya, and indeed in most CTTSs, is likely dominated by emission from the accretion shock. In the extreme case of Batalha et al. (1996), the collective non-shock contributions only account for 4\\% (0.09 \\AA / 2.5 \\AA) of the \\ion{He}{1} emission, which makes our estimate that 16\\% of \\ion{Ca}{2} emission is from accretion regions only adjusted down by 0.6\\% at most.  We did not detect any polarization in the H$\\alpha$ line forming region. H$\\alpha$ emission likely results from several zones in the TW Hya system including the accretion shock, magnetospheric accretion flow, and the stellar and/or disk wind flowing away from the star.  Thus, the line forms over a large volume where the mean magnetic field is likely relatively weak and the direction of the field is not constant.  Indeed, there are regions in the magnetosphere (if viewed nearly pole-on) where the field directions are reversed.  This mixture of polarities will reduce the separation of LCP and RCP line profiles.  Depending on where the majority of the line forms, we expect the local field strength to be quite weak. For example, assuming a face-on disk for TW Hya and a dipolar-like magnetospheric accretion flow with the magnetic axis close to the stellar rotation axis, the dipole component of the field at the corotation radius will be equal to the stellar value of that field multiplied by $(R_*/R_{co})^3$.  Taking our estimate above of 260 G for the equatorial value of the dipole component of the field at the stellar surface and $R_{co} = 6.3 R_*$ yields a field of about 1 G at a corotation radius where the magnetospheric flow may originate. This example simply illustrates that the field strength throughout a large portion of the H$\\alpha$ line forming region may be quite weak, so our lack of a detection in this line may well be expected."
    },
    "0609/astro-ph0609130_arXiv.txt": {
        "abstract": "{We present numerical 2D magnetostatic models for sunspot penumbrae consisting of radially aligned field-free gaps in a potential magnetic field, as proposed by Spruit and Scharmer (2006). The shape of the gaps and the field configurations around them are computed consistently from the condition of magnetostatic pressure balance between the gap and the magnetic field. The results show that field-free gaps in the {\\em inner} penumbra are cusp-shaped and bounded by a magnetic field inclined by about $70^\\circ$ from the vertical. Here, the magnetic component has a Wilson depression on the order $200$--$300$~km relative to the top of the field-free gap; the gaps should thus appear as noticeably elevated features. This structure explains the large variations in field strength in the inner penumbra inferred from magnetograms and two-component inversions, and the varying appearance of the inner penumbra with viewing angle. In the outer penumbra, on the other hand, the gaps are flat-topped with a horizontal magnetic field above the middle of the gap. The magnetic field has large inclination variations horizontally, but only small fluctuations in field strength, in agreement with observations. ",
        "introduction": "The spatial resolution achieved with adaptive optics systems, implemented recently on major solar telescopes (Rimmele 2000, Scharmer et al. 2000, Soltau et al. 2002), now approaches the values needed to resolve the intrinsic length scales expected for magnetic, as well as nonmagnetic, structures on the solar surface. The spectacular convergence between numerical simulations (Carlsson et al. 2004, Keller et al. 2004, Steiner 2005) and recent observations of the small scale magnetic field (Lites et al. 2004) provides confidence, both in our theoretical understanding of these structures and in the power of the best current observations to test the theory convincingly.  The more puzzling phenomenology of sunspots still awaits a similar breakthrough. The discovery that penumbral filaments have dark cores flanked by lateral brightenings (Scharmer et al. 2002), however, strongly suggests that the fundamental scales of penumbral filaments are now being resolved as well. These recent observations have also highlighted the evolutionary connection of penumbral filaments with umbral dots and light bridges, as well as their morphological similarities. This hints at the possibility of a common underlying structure.  The center-to-limb variation of  structure in sunspot images provides some geometrical information on vertical structure (e.g. Lites et al. 2004), but  most information about the 3rd dimension is encoded in spectral line profiles and their polarization properties.  Inversions of these data into a vertical structure model, based on techniques developed by e.g. \\ Skumanich and Lites (1987), Ruiz Cobo and del Toro Iniesta (1992) and Frutiger et al. (1999), are very far from unique and must be regularized by assumptions about the vertical structure of the thermodynamics and magnetic field configuration in the atmosphere. Alternatively, forward modeling by radiative transfer in detailed structure models is used (e.g. \\  Solanki \\& Montavon 1993, Martinez Pillet 2000,  M{\\\"u}ller  et al. 2002).  On the theoretical side, 3D MHD simulations including full radiative transfer (Stein  \\& Nordlund 1998, V\\\"ogler et al. 2003) are able to model increasingly large volumes of the solar atmosphere.  Recent progress in understanding of umbral dots with such simulations (Sch\\\"ussler \\& V\\\"ogler 2006) raises hopes that convergence between theory and observation of sunspot structure is a realistic prospect for the near future.  Much of the current thinking about penumbral structure and Evershed flows is based on 1D MHD simulations of thin magnetic flux tubes assumed to move in a background of different magnetic properties (Schlichenmaier et al. 1998a, b). These simulations have led to the view that the magnetic field in the observable layers of the penumbra is intrinsically far removed from its lowest energy state, the potential field. Interpretations of polarized spectra obtained at low spatial resolution made in a similar view are the so-called embedded flux tube or `uncombed' penumbra models (Solanki \\& Montavon 1993).  Spruit and Scharmer (2006, hereafter SS06) have questioned these interpretations of penumbra fine structure. In the layers above the photosphere, the magnetic field is already so dominant that deviations from a potential field configuration is unlikely on the observed time scales of penumbral structure. The Alfv\\'en crossing time on which a magnetic structure changes,  unless it is close to a potential or at least a force free field, would only be of the order of seconds for the embedded tubes of width $\\sim 100$ km assumed in the models. It seems unlikely that such embeddings will be found in realistic 3D MHD simulations, such as may become possible in the near future.  In SS06 an alternative scenario is proposed for understanding penumbra fine structure, its magnetic field configuration and its energy balance. It assumes the existence of field-free, radially aligned gaps below the visible surface, intruding into a nearly potential field above. In this model, the dark penumbral cores outline the centers of the field-free gaps, analogous to the dark lanes running along light bridges in spots (Berger \\& Berdyugina 2003, Lites et al. 2004), and on a smaller scale the `canals' seen in strong field regions outside spots (Scharmer et al. 2002). Remarkably, the dark cores of light bridges have already been reproduced in full 3-D radiative MHD simulations (Nordlund 2005, Heinemann 2006). They are due to the enhanced opacity associated with the higher gas pressure in the field-free gaps combined with an overall drop of temperature with height. Outlining the `cusps' of these gaps, they appear as elevated above and dark relative to their surroundings. This is consistent also with the appearance of the dark cores of penumbral filaments.  Another important motivation for embedded tube models has been the small length scale on which the inclinations of field lines vary. Vertical length scales on the order of a hundred km above the photosphere, as inferred from observations (Sanchez Almeida  \\& Lites 1992, Solanki et al. 1993) led to the idea that the penumbral atmosphere could not be the smooth structure expected from a potential field. The assumption of non-potential structure in the form of small-scale inclusions within the height of formation of a spectral line may have seemed straightforward. In SS06 we have shown that this seemingly obvious conclusion is nevertheless erroneous. Not only are small scale variations consistent {\\em with}, they are actually an inevitable property {\\em of} potential fields. By the nature of the Laplace equation, an inhomogeneity of wavelength $\\lambda$ imposed at the boundary of the domain decays into it with an (e-folding) length scale $\\lambda/2\\pi$. For horizontal structure with $\\lambda\\sim 1\"$ in the penumbra, this predicts a vertical length scale of $\\sim$100 km.  In this way the interpretation of filaments as gaps, solving the penumbral heat flux problem, and the observation of small scale field inclination variations mutually support each other. It does so within the elegance of a potential field model which is expected, on theoretical grounds, to hold approximately in the atmosphere.  The shape of a gap (i.e., its width as a function of depth) is determined by the condition of  pressure balance between the surrounding magnetized fluid and the nonmagnetic stratification in the gap. This is the same mathematical problem as finding the shape of the sunspot flux bundle from the balance of pressures at its outer boundary (e.g., Jahn and Schmidt 1994).  The analytic potential field model used in SS06 took this force balance into account only in the coarsest sense. More quantitative detail is needed when comparing the model with observations. For example, the top of the gap must have a `cuspy' shape in order for the explanation of dark cores (see above) to hold, and the observed ranges of field line inclinations in inner and outer penumbra must be reproduced. In this paper we address this with numerical solutions for the field configuration, assuming  a pressure stratification in the gap approximating that of the normal convection zone, and a similar but reduced pressure distribution inside the magnetic field. These assumptions lead to models that are distinctly different for the inner and outer penumbra and that allow observed properties of dark-cored filaments to be explained. ",
        "conclusions": "We have shown that magnetostatic penumbra models characterized by field-free gaps can be constructed by allowing the shape of the discontinuity to adapt itself to the required force balance between the field-free and magnetic components. This is in contrast to the embedded flux tube model where the geometry, a round flux tube with internal field lines aligned with the flux tube and external field lines wrapping around the flux tube, is assumed given a priori. As pointed out by SS06 and further discussed in Sect. 2, such flux tubes cannot be in magnetostatic equilibrium.  The magnetostatic gappy penumbra model presented here is conceptually simple. We have assumed a potential field interlaced by field-free gaps, a variation of gas pressure with height that is similar to that of the quiet sun for the field-free gap and a gas pressure in the magnetic component that is simply that of the field-free gap scaled by a constant. Further input quantities of the models are the average field strengths and inclinations typical of the outer, mid and inner penumbra. With these constraints, we have calculated boundary shapes and magnetic field configurations.  The calculated models have properties that are distinctly different in the inner and outer penumbra and that agree with observed images and magnetograms. In particular, we find that field-free gaps in the {\\it inner} penumbra are cusp-shaped and associated with a magnetic field that is inclined by about $70^\\circ$ from the vertical for filaments that are separated by $1000$ km. Here, the magnetic component is associated with a Wilson depression on the order $200-300$ km relative to the field-free component that makes field-free gaps appear as elevated, distinct features. This large Wilson depression, in combination with field lines converging with depth, explains the large variations in field strength inferred from magnetograms and two-component inversions. The steep walls of the field-free gaps explain why dark-cored penumbral filaments are seen with lateral brightenings on both sides of the dark core at all azimuth angles only for sunspots that are close to disk center. In the {\\it outer} penumbra, we find that field-free gaps are associated with flat-topped boundaries and a horizontal magnetic field above the center of the gap. Near the surface, this magnetic field shows large inclination variations horizontally, but only small fluctuations in field strength, in agreement with observations.  We associate the atmospheres above our field-free gap and magnetic component with the flux tube and background components respectively in the inversions of Bellot Rubio et al. (2004) and Borrero et al. (2005). Our models are then consistent with these inversions as regards the variation of field strength and inclination with distance from the umbra for the two components. Our  models also show a widening of the cusp and gradually reduced Wilson depression, consistent with the gradual widening and fading of dark cores, towards the outer penumbra. This is  also consistent with a systematic increase of the flux tube filling factor towards the outer penumbra.  Whereas our calculations were made by assuming potential magnetic fields, the differences between the inner and outer penumbra are fundamentally due to magnetic flux conservation, constrained by magnetostatic equilibrium and the average properties of the magnetic field assumed. Flux conservation and a strong {\\it vertical} field in the inner penumbra forces the field-free gap to narrow and the magnetic component to widen, as compared to what is the case in the outer penumbra, and a cusp to form. These {\\it qualitative} differences between the inner and outer penumbra constitute solid results, not likely to change with more accurate models. More realistic magnetic field configurations, intended for detailed comparisons with observations, may however need to take into account the effects of horizontal temperature gradients, convection and flows.  The interpretation of line profiles and polarimetry is frequently formulated in terms of an embedded flux tube paradigm. The practical implementations of such models, however, are generic 2-component inversions, often not particularly consistent with the physics of embedding of a flux tube in a background magnetic field (cf. discussion in section 2). While these inversions therefore cannot be used as support for their embedded tubes, they have allowed general conclusions about magnetic field strength and inclination variations in the penumbra to be drawn from spectropolarimetric data. Since its introduction 13 years ago by Solanki and Montavon (1993),  no magnetostatic embedded flux tube model has yet emerged, and there are good physical reasons why a realistic model is unlikely to materialize. The embedded flux tube model thus remains a conceptual cartoon, useful in a restricted sense as a two-component model for quantifying variations of the magnetic field, temperature, line of sight velocity and other properties within the resolution element.  The moving tube model of Schlichenmaier (1998a, b) leads to a number of predictions that have successfully been tested against observations. In spite of this success, the presence of flux tubes with circular cross sections in the penumbra meets with the same objections as the embedded flux tube models; such flux tubes (if they exist) are more likely to manifest themselves as sheets with an azimuthal thickness much smaller than their vertical extent (Jahn \\& Schmidt 1994). Interchange convection in such flux sheets has been suggested as a  possible heating mechanism for the penumbra (Jahn \\& Schmidt 1994). However, the long measured lifetimes (on the order of $1$~hr) for filaments (e.g. Sobotka \\&  S\\\"utterlin 2001, Langhans et al. 2005) as well as flow channels (Rimmele \\& Marino 2006) leads to the conclusion that this is not a viable heating mechanism (Schlichenmaier \\& Solanki 2003). While the flux tubes simulated by Schlichenmaier (1998a,b) appear to explain Evershed flows, such flux tubes or flux sheets therefore also pose severe problems for explanations of penumbral heating. As shown by Schlichenmaier and Solanki (2003), individual flux tubes cannot heat penumbral filaments extending over more than approximately $1000$--$2000$~km. Such flux tubes must either submerge to give room for new flow channels, or heating of the submerged part of the flux tube must occur. Schlichenmaier and Solanki (2003), relying on simulations by Schlichenmaier (2002), speculate that the submerged part of the flux tube is heated radiatively by hotter gas below the photosphere, causing it to reappear as a hot upflow channel. The efficiency of such radiative heating is restricted to shallow depths below the photosphere. Therefore, this explanation ultimately relies on (efficient) convection to provide the needed heat flux. Evidence to support either of the two scenarios discussed above is absent in the highly resolved images analyzed by Rouppe van der Voort et al. (2004). We also note that Rimmele and Marino (2006) found no evidence for downflows along the observed flow channels.  Our model predicts strong vertical and horizontal gradients in both the magnetic field inclination and azimuth angles. In this paper, we have made no attempts to adjust our models to match such gradients with those inferred from inversion techniques. In future work we intend to use our models to calculate synthetic continuum and narrowband images as well as conventional and polarized spectra, using empirically determined temperatures and velocity fields. Our model predicts not only strong magnetic field gradients, but also a strongly warped $\\tau = 1$ surface that should enable a number of critical tests with observed high-resolution data."
    },
    "0609/astro-ph0609306_arXiv.txt": {
        "abstract": "The detection by the HESS atmospheric Cherenkov telescope of fifteen new sources from the Galactic plane makes it possible to estimate the contribution of unresolved point sources like those detected by HESS to the diffuse Galactic emission measured by EGRET and recently at higher energies by the Milagro Collaboration. Assuming that HESS sources have all the same intrinsic luminosity, the contribution of this new source population can account for most of the Milagro $\\gamma$-ray emission at TeV energies and between 10 and 20 per cent of EGRET diffuse Galactic $\\gamma$-ray emission for energies bigger than 10 GeV. Also, by combining the HESS and the Milagro results, constraints can be put on the distribution and the luminosities of gamma ray emitters in the Galaxy. ",
        "introduction": "The Galactic diffuse $\\gamma$-ray emission is believed to be mostly produced in interactions of cosmic rays with the matter and the radiation fields in the Galaxy, the main production mechanisms being electron non-thermal Bremsstrahlung, Inverse Compton scatterings off the radiation fields and pion decay processes in inelastic collisions of nuclei and matter. Although the standard production mechanisms of $\\gamma$-rays \\citep{Bertsch:1993} explain generally well the spatial and energy distribution of the emission below 1 GeV, the model does not match EGRET observations of the $\\gamma$-ray sky above 1 GeV.\\citep{Hunter:1997we} The discrepancy between the model and the observations is more pronounced towards the direction of the inner Galaxy, where it amounts to 60 per cent of the observed emission. Many possible explanations have been proposed to account for the GeV excess. \\citet{Aharonian:2000iz} have proposed an hardening of the cosmic ray proton spectrum. The average proton spectrum in the Galaxy would have to be harder than locally measured. \\citet{Strong:1999hr,Strong:2004} have suggested a renormalization of the electron spectrum in order to fit EGRET data. \\citet{Kamae:2006} have argued that diffractive processes make the $\\gamma$-ray spectrum harder than the incident proton spectrum by ~0.05 in power-law index, that the scaling violation produces more $\\pi^0$, and, that the inelastic $p-p$ cross section has a logarithmic increase with the incident proton energy, so that the combination of these effects could explain about 20 per cent of the \"GeV Excess\". \\citet{deBoer:2005tm} claimed that the features of the EGRET excess are consistent with the hypothesis of a signal from Dark Matter Annihilation. Recently \\citet{Strong:2006} has argued that the excess in the diffuse emission measured by EGRET and COMPTEL could be explained by a combination of Galactic source populations too dim to be detected by EGRET in addition to the standard production mechanisms.  In any case a major difficulty when studying the diffuse emission is to disantangle the truly diffuse emission from that produced by unresolved sources.  At TeV energies, Milagro has observed diffuse emission from a large region of the Galactic plane from Galactic longitude $l$ of ${40}^{o}<l<{100}^{o}$ \\citep{Atkins:2005}.   The different theories to explain the GeV excess make very different predictions about the TeV flux. Recently \\citet{Prodanovic} have argued that no strong signal of pion decay is seen in the $\\gamma$-ray spectrum. The pion decay mechanism would not be able to explain the excess above 1 GeV. Therefore, \\citet{Prodanovic} claimed that the Milagro measurements of the diffuse flux at higher energies revealed a new excess at TeV energies and concluded that a different mechanism is required to explain the two excesses, either dark matter decay or contribution from unresolved EGRET sources. An extrapolation of the EGRET sources within the range in Galactic longitude of the Milagro observation overestimates the diffuse TeV flux measured. However, there could be a population of sources undetectable by EGRET, but contributing to both the GeV and TeV excess diffuse emission. Recent observations by the TeV observatory, HESS, in fact point to a new class of hard spectrum TeV sources. HESS is an array of telescopes using the imaging atmospheric Cherenkov technique to detect $\\gamma$-rays above 200 GeV. HESS scans the inner part of the Galaxy in the region between $-{30}^{o}$ and ${30}^{o}$ degrees in Galactic longitude ($-{30}^{o}<l<{30}^{o}$) and between $-3^{o}$ and $3^o$ degrees in Galactic latitude ($-3^o<b<3^o$) with an angular resolution better than 0.1 degrees. Its average sensitivity at the energy threshold of 200 GeV is 2 per cent of the Crab flux within $\\pm 2^{o}$ degrees in Galactic latitude and decreases outside this range. (The flux of the Crab nebula is about $2.4 \\times {10}^{-10} \\, {\\mathrm {cm^{-2}}} \\, {\\mathrm {s^{-1}}}$ for $E>200 \\, \\mathrm{GeV}$ \\citep{Aharonian:2004hegra}). HESS has previously reported the detection of 8 unknown sources close to the Galactic Plane. \\citep{Aharonian:2005jn} Recently the total number of discovered sources has been updated to 15. \\citep{Aharonian:2005kn} Three sources have also been detected that are identified with the Galactic Center, the SNR RX J1713.7-3946 and G0.9+01. All HESS sources are located close to the Galactic Plane, within $\\pm 1$ degrees of latitude. \\citep{Aharonian:2005kn} The fluxes above 200 GeV from all sources are fit to power laws, $\\frac{dN}{dE} = F_{0} \\, E^{-\\Gamma}$, whose photon indices $\\Gamma$ range from 1.8 to 2.8, with a relatively hard mean index of 2.32 and RMS of 0.2.  In the inner Galaxy region observed by HESS the existence of 91 SNRs and 389 pulsars was already known at lower energies \\citep{Green:2004gr,Manchester}. Many of these SNRs and pulsars within the inner Galaxy might emit VHE $\\gamma$ rays but only a few $\\gamma$ ray sources were previously known. Although the HESS Collaboration indicates a clear positional coincidence of its sources with a known SNR or pulsar only in a limited number of cases, the distribution of Galactic latitude of the eighteen VHE $\\gamma$ ray sources detected by HESS agrees quite well with the distributions of all SNRs catalogued by \\cite{Green:2004gr} and of all pulsars catalogued by \\cite{Manchester}. Assuming, for example, that SNRs are a single class of counterparts with isotropic luminosity $1.95  \\times {10}^{34}  \\,  {\\mathrm{erg}}  \\,  {\\mathrm{s^{-1}}}$ to the new HESS sources and taking for them a simple radiative model, \\cite{Aharonian:2005kn} found that the location of these sources favours a scale height of less than 100 pc, consistent with the hypothesis that these sources are either SNRs or pulsars in a massive star forming region.  The HESS detections of high energy $\\gamma$ rays from fifteen new sources has improved significantly the knowledge of both the spatial distribution and the spectra and fluxes of VHE $\\gamma$-ray galactic sources. These HESS results make it possible to estimate with unprecedented precision the contribution of unresolved point sources to the Galactic diffuse emission measured up to 100 GeV by EGRET  and recently extended to TeV energies by Milagro.\\citep{Atkins:2005} Based on HESS results, knowing the sensitivity and the field of view of an experiment, we estimate the number of expected sources and their expected VHE $\\gamma$-ray flux. Finally we evaluate the contribution of unresolved point sources to the diffuse $\\gamma$-ray emission for EGRET and Milagro. ",
        "conclusions": "Using HESS new sample of Galactic $\\gamma$-ray emitters we estimated the contribution of unresolved sources to the diffuse emission measured by EGRET and by Milagro. We find that the contribution of unresolved HESS-like sources to the diffuse $\\gamma$ ray emission is large at energies above 100 GeV, whereas it amounts to about 10 per cent of EGRET measurements of the diffuse flux at GeV energies. Since pulsars and SNRs are distributed mostly in the Galactic plane, their unresolved fluxes will mainly contribute to the low latitude diffuse emission. Thus we expect that the total flux due to unresolved sources contributes less at higher latitudes.  The main uncertainties of this calculation consist in assuming that all sources are monoluminous and the distribution of $\\gamma$-ray sources follows the distribution of either pulsars or SNRs. Of the 1509 pulsars catalogued by the Australian National Telescope Facility \\citep{Manchester} only a few are also $\\gamma$-ray emitters. This could be due to strong selection effects as suggested by the existence of pulsars, such as Geminga, which are $\\gamma$-ray loud pulsars, yet not observed in the radio. However, both the SNR and pulsar distribution give similar conclusions, as would any class of sources that follows a similar spatial distribution.  To draw more certain conclusions about the high energy $\\gamma$-ray Galactic sky, new observations are of fundamental importance. VERITAS will be able to survey the same region of the sky as Milagro. Milagro's recently improved sensitivity will better image the diffuse emission. Finally, GLAST will investigate the window of energy between 10 MeV to 300 GeV, covering the energy gap left between EGRET and the ground based gamma-ray observatories."
    },
    "0609/cs0609126_arXiv.txt": {
        "abstract": "%  Are the e-prints (electronic preprints) from the arXiv repository being used instead of the journal articles? In this paper we show that the e-prints have not undermined the usage of journal papers in the astrophysics community. As soon as the journal article is published, the astronomical community prefers to read the journal article and the use of e-prints through the NASA Astrophysics Data System drops to zero. This suggests that the majority of astronomers have access to institutional subscriptions and that they choose to read the journal article when given the choice. Within the NASA Astrophysics Data System they are given this choice, because the e-print and the journal article are treated equally, since both are just one click away. In other words, the e-prints have not undermined journal use in the astrophysics community and thus currently do not pose a financial threat to the publishers. We present readership data for the arXiv category \"astro-ph\" and the 4 core journals in astronomy (Astrophysical Journal, Astronomical Journal, Monthly Notices of the Royal Astronomical Society and Astronomy \\& Astrophysics). Furthermore, we show that the half-life (the point where the use of an article drops to half the use of a newly published article) for an e-print is shorter than for a journal paper.  The ADS is funded by NASA Grant NNG06GG68G. arXiv receives funding from NSF award \\#0404553. ",
        "introduction": "The questions we address in this paper are twofold. The first question is a practical one: are people reading the e-prints from arXiv instead of the journal articles? In other words: are the journals safe? The second question concerns bibliometrics: Are e-prints read in a different way than journal articles?  Besides providing easy access to technical literature, digital repositories like the NASA Astrophysics Data System (ADS;\\citet{kurtz00}) and the arXiv e-print repository (\\citeauthor{ginsparg94} \\citeyear{ginsparg94} and \\citeyear{ginsparg01}) also allow a detailed study of the readership of individual articles. This is the electronic equivalent of studying traditional library circulation records, listing who accessed what book when (\\citet{kurtz05b}). We are interested in trends, however, not the usage of individual users.  The teams of the ADS and arXiv have been collaborating since 1997. In 1997 the ADS created an index for the astro-ph e-prints and made it available through the ADS abstract service with the same searching options as the published literature database. The abstracts of all arXiv categories were included in 2002. The incorporation of the arXiv abstracts resulted in the powerful alerting service {\\it myADS-arXiv} (\\citet{henneken06a}). Since mid-2004, arXiv usage data are incorporated in the ADS \"also-read\" statistics. The collaboration provides the unique opportunity to study differences between the ADS and arXiv readership. In other words: is there a difference between the \"life\" of a journal article and that of an e-print on arXiv? The \"life\" of an article can be characterized in various ways. The number of reads over time is one such way. Because of the stochastic nature of individual article reads, the \"life\" of an article only makes sense in an average sense.  The readership of the technical astronomical literature, available through the ADS, has already been studied extensively (\\citet{kurtz05b},\\citet{kurtz03}). In these studies article readership was measured by looking at how usage declines as a function of article age. For the astronomical literature (going back to the 1890s), a model was introduced describing 4 components (\"modes\") of readership. This model provides a good fit with the data. Each of the \"modes\" is characterized by a time scale, which is the \"half-life\" for that mode: the time where the use of an article drops to half the use of a newly published article. The different \"modes\" correspond with different types of usage. For example, people using the repository to look for historical litature form one group with specific usage characteristics. In general, one could say that articles move through different types of readership over time.  The model described in \\citet{kurtz05b} introduces 4 distinct stages or \"modes\", where an article moves to \"new\" to \"recent\" to \"interesting\" to \"historical\".  Using the arXiv usage logs, we will investigate whether the readership of e-prints follows trends similar to the astronomical literature. Obviously, there will be differences because the articles of the Astrophysical Journal go back to the 1890s and the e-prints from arXiv go back to 1991. So, there is no \"historical\" component in the arXiv readership data. The most interesting time frame will be the first couple of months to a year after the publication of an e-print. The readership data from the arXiv logs will show whether e-prints are read in a different way than journal articles. Besides this bibliometric component, there is also an economic component: it has been said that the presence of e-prints is costing the publishers money. If astronomers read fewer journal articles because of the presence of e-prints, this might be true. At the same time, the opposite might be true as well. Maybe, thanks to the e-prints, journal articles have gained visibility and therefore an increased readership. ",
        "conclusions": "The members of the astronomical community consitute the lion's share of the journal article readers. We have shown (figure~\\ref{fig:histo}) that when given the choice, they prefer to read the article over the e-print. We believe this is because the journal article has been refereed and is accepted as the `official' version. This is also a very important observation with respect to the economic component. The majority of astronomers has access to the online journals through institutional subscriptions. In other words, the e-prints have not undermined journal use in the astrophysical community! Quite the contrary, the e-prints help journal articles to gain more visibility. Additional reasons for choosing the journal article are the cross-linking of references and datasets, the links to astronomical objects (SIMBAD) and the availability of high-resolution and low-resolution versions of plots and plates. In terms of the 4 modes, this means that the \"interesting\" and \"historical\" modes develop with smaller amplitudes in the readership of e-prints, because after the publication of the journal article, only a small subset of the typical users keeps reading the e-prints. The readership study in \\citet{kurtz05a} has shown that the \"new\" mode has a half-life of about 16 days, the \"current\" mode has a half-life of about 1.7 years and the \"interesting\" mode has a half-life of about 10.7 years. The \"historical\" mode is a constant number of reads of about 1.5 reads per paper per month. For the e-prints, the \"historical\" mode seems settle at around 0.5 reads per paper per month.  The e-prints and the journal articles co-exist and show vastly different usage characteristics. Their co-existence is productive, because the prior publication as e-print gives a journal article primacy. This primacy is one of the factors that results in higher citation rates (\\citet{henneken06b} and references herein)."
    },
    "0609/astro-ph0609498.txt": {
        "abstract": "We present echelle spectrophotometry of the blue compact dwarf galaxy (BCDG) NGC~5253. The data have been taken with the Very Large Telescope UVES %%@ echelle spectrograph in the 3100 to 10400 \\AA\\ range. We have measured the intensities of a large number of permitted and forbidden emission lines in %%@ four zones of the central part of the galaxy. In particular, we detect faint \\ion{C}{2} and \\ion{O}{2} recombination lines. This is the first time %%@ that these lines are unambiguously detected in a dwarf starburst galaxy.  The physical conditions of the ionized gas have been derived using a large %%@ number of different line intensity ratios. Chemical abundances of He, N, O, Ne, S, Cl, Ar, and Fe have been determined following the standard %%@ methods. In addition, C$^{++}$ and O$^{++}$ abundances have been derived from pure recombination lines. These abundances are larger than those %%@ obtained from collisionally excited lines (from 0.30 to 0.40 dex for C$^{++}$ and from 0.19 to 0.28 dex for O$^{++}$). This result is consistent with %%@ a temperature fluctuations parameter ($t^2$) between 0.050 and 0.072. We confirm previous results that indicate the presence of a localized N %%@ enrichment in certain zones of the center of the galaxy. Moreover, our results also indicate a possible slight He overabundance in the same zones. %%@ The enrichment pattern agrees with that expected for the pollution by the ejecta of massive stars in the Wolf-Rayet (WR) phase. The amount of %%@ enriched material needed to produce the observed overabundance is consistent with the mass lost by the number of WR stars estimated in the %%@ starbursts. Finally, we discuss the possible origin of the difference between abundances derived from recombination and collisionally excited lines %%@ (the so-called {\\it abundance discrepancy} problem) in \\ion{H}{2} regions, finding that a recent hypothesis based on the delayed enrichment by SNe %%@ ejecta inclusions seems not to explain the observed features. ",
        "introduction": "\\footnotetext{Based on observations collected at the European Southern Observatory, Chile, proposal number ESO 70.C-0008(A)}   Deep cross-dispersed echelle spectrophotometry with large aperture telescopes is a powerful technique for refining our knowledge about the che\\-mical composition of \\ion{H}{2} regions. It permits to observe the whole optical spectral range, to measure crucial faint auroral and recombination lines, to deblend nebular lines, and to decontaminate them for nearby sky spectral features.  Our group has obtained deep high resolution spectra of most of the brightest Galactic \\ion{H}{2} regions \\citep{E98,E99a,E99b,E04,GRE04,GRE05,GRE06} that have led to the precise determination of the abundance of heavy element ions from recombination lines and even to the determination of the carbon radial abundance gradient in the Galactic disk \\citep{E05}. There is still a limited number of similar stu\\-dies devoted to giant extragalactic \\ion{H}{2} regions (hereafter GEHRs), focusing the efforts on the Magellanic Clouds \\citep{P03,Tsamis03} and in a few other nearby irregular and spiral galaxies \\citep{EP02,PPR05}. A common result of both the Galactic and extragalactic works, is the finding of a systematic difference between the abundances determined from recombination lines (hereafter RLs) and collisionally excited lines (hereafter CELs) of the same ions, in the sense that abundances determined from RLs are {\\it always} higher than those determined from CELs. This problem, that has been also reported to occur in some planetary ne\\-bulae (PNe), is known as the {\\it abundance discre\\-pancy}. The origin of this problem is still unknown. \\citet{Pequignot02} and \\citet{Tsamis04} have constructed models with inhomogeneous chemical composition and physical conditions to explain that discre\\-pancy in PNe. A similar idea has been very recently proposed by \\citet{Tsamis05} for explaining the abundance discrepancy in the case of 30 Dor. These authors propose the presence of metal enhanced inclusions that would be responsible for most of the emission in RLs of heavy element ions. However, the abundance discrepancy could be related to another secular problem in nebular astrophysics: the existence of temperature fluctuations in the ionized gas volume (Peimbert 1967; Torres-Peimbert, Peimber \\& Daltabuit 1980), %\\citep{P67,TP80}, whose existence is still controversial although some mechanisms have been proposed (see Esteban 2002 for a review).  One step ahead in our investigation is to obtain abundance determinations from RLs in a sample of bright GEHRs in starburst galaxies and even \\ion{H}{2} galaxies of chemical compositions other than solar. The dwarf irregular galaxy NGC 5253 is one of the closest starbursts and a very appro\\-priate target for our purposes (Figure~\\ref{nombres}). It lies at a heliocentric distance of 3.3 Mpc \\citep{Gibson00} and belongs to the Centaurus Group. Consequently, NGC 5253 has been one of the most studied starbursts, observed at basically all wavelengths from radio to X-rays.  \\begin{figure*}[t!] %\\includegraphics[width=1\\linewidth]{ngc5253_nombres.eps} \\includegraphics[width=1\\linewidth]{fig1.eps} \\caption{\\small{(\\emph{Left}) False color image in filters $B$ (blue), $V$ (green) and $I$ (red) of the starburst galaxy NGC 5253 observed by the 2.5m du Pont telescope, Las Campanas Observatory. (\\emph{Right}) False color image obtained by combining $HST$ \\Ha\\ (red), [O III] (green) and UV (blue) images. The position of our $VLT$ slit is indicated in it. North is up and east at left. %%@ The $HST$ image has been included at the center of the 2.5m du Pont telescope image for comparison.}} \\label{nombres} \\end{figure*}  NGC 5253 is also especially interesting because it is the best candidate for localized chemical enrichment in a GEHR. \\citet{Welch70}, Walsh \\& Roy (1987, 1989), and \\citet{Kobulnicky97} reported the presence of a strong nitrogen overabundance in a particular zone of the starbursting nucleus of the galaxy. \\citet{CTM86} also found a helium enhancement in that zone, but this has not been confirmed in later works. Another peculiarity of the galaxy is that the majority of the gas seems to rotate about the optical major axis of the galaxy \\citep{KS95} but this behaviour is not clear and could be combined with some kind of outflow (Koribalski 2006, priv. communication). However, CO observations by \\citet*{TBH97} and \\citet*{MTB02} suggest that mole\\-cular clouds are infalling into NGC 5253. These two observational facts indicate that the dynamical situation of this galaxy is far from clear. Other authors as \\citet{vdBergh80} and \\citet{CP89} have suggested that the nuclear starburst was triggered as a result of a past interaction (around 1 Gyr ago) with the neighboring galaxy M83. Both galaxies lie at a radial distance of only 500 kpc and the projected separation is 130 kpc \\citep{Thim03}.   Concerning the kinematics of the ionized gas of NGC 5253, it has to be highlighted the presence of two large superbubbles related to the central region of the galaxy, with diameters of the order of 1 arcmin ($\\sim$1 kpc), and expansion velocities of 35 km s$^{-1}$ \\citep{Marlowe95}. These superbubbles coincide with extended X-ray emission as demonstrated by {\\it Chandra} and {\\it XMM-Newton} data analyzed by \\citet{Summers04}. The structure and morphology of the ionized gas in NGC 5253 were deeply analyzed by \\citet{Calzetti04} \\mbox{using} imagery from {\\it Hubble Space %%@ Telescope} (HST), detecting faint arches and filaments in both \\Ha\\ and [\\ion{S}{2}] at $\\geq$1 kpc from the main ionizing cluster (see %%@ Figure~\\ref{nombres}, right panel) that are partially excited by shocks.  NGC 5253 has been considered as one of the youngest starbursts in the local universe \\citep{vdBergh80,MG82,Rieke88}. \\citet*{CTM86} and \\citet{WR87} reported a broad emission feature, indicating the presence of Wolf-Rayet (WR) stars in this galaxy. Seve\\-ral authors \\citep{Schaerer97,Kobulnicky97} have confirmed the presence of late-type WN and early-type WC stars. This fact, and the almost entirely thermal radio spectrum with very little synchrotron emission from supernova remants \\citep{Beck96}, imply the extreme youth of the starburst. Indeed, there is evidence in the literature that multiple starbursts in succession have occurred in the center of this galaxy, leaving behind an evolved, but still young-ish stellar population \\citep{SS00,Tremonti01}. ",
        "conclusions": "We present deep echelle spectrophotometry in the 3100$-$10400 \\AA\\ range of four selected zones of the nucleus of the blue compact dwarf galaxy %%@ NGC~5253. We have identified and measured a large number of emission lines in the spectra. This represents the largest collection of optical emission %%@ lines available for a dwarf starburst galaxy. Electron densities and temperatures have been consistently determined from a large number of emission  %%@ line ratios of different ions. Chemical abundances of He, N, O, Ne, S, Cl, Ar, and Fe have been derived following standard methods. Recombination %%@ lines of \\ion{C}{2} and \\ion{O}{2} have been measured in the spectra of three of the individual zones analysed. The first time these kinds of lines have been unambiguously detected in a dwarf \\ion{H}{2} galaxy. The ionic abundances of C$^{++}$ and O$^{++}$ derived from these lines are from %%@ 0.20 to 0.40 dex larger than those determined from the intensity of collisionally excited lines. This behavior has been found in all the Galactic and %%@ extragalactic \\ion{H}{2} regions where both kinds of lines have been reported for the same ion. The quoted abundance differences can be accounted for %%@ a temperature fluctuations parameter ($t^2$) between 0.050 and 0.072.  The emission line profiles are complex and seem to be the combination of $-$at least$-$ two kinematical components with slightly different radial %%@ velocity and, in some zones, different line width. Apparently, there are no clear differences in the physical and chemical properties of the %%@ kinematical components, except perhaps a somewhat different reddening coefficient and N/O ratio in some particular zones. Position-velocity diagrams %%@ of different emission lines show a sinusoidal shape. However, the position-velocity diagram of stellar absorption spectral features shows a pattern %%@ entirely consistent with rotation. This fact suggests that the radial velocity behavor of the ionized gas should be due to outflows from the central %%@ starbursts and not a product of merging.  We have derived the ages and the massive star population of the starburst. In particular, the detection of %%@ the so-called blue WR bump indicates the presence of a rather small number of WR stars in the starbursts.  Our observations confirm previous results that indicate the presence of localized N enrichment in two of the central starbursts of NGC~5253. %%@ Moreover, we also detect a possible slight He pollution in the same zones. The enrichment pattern is entirely consistent with that expected by the pollution of the winds of masive stars in the WR phase. Moreover, the estimated mass of newly created N and He is also consistent with the number of WR stars determined in the starbursts.  Finally, we conclude that a recent hypothesis that tried to solve the abundance discrepancy problem in \\ion{H}{2} regions in terms of a delayed enrichment by SNe ejecta inclusions \\citep{Tsamis05} seems not to explain the available observational data in a satis\\-factory manner."
    },
    "0609/astro-ph0609183_arXiv.txt": {
        "abstract": "Pulsar \"standard model\", that considers a pulsar as a rotating magnetized conducting sphere surrounded by plasma, is generalized to the case of oscillating star. We developed an algorithm for calculation of the Goldreich-Julian charge density for this case. We consider distortion of the accelerating zone in the polar cap of pulsar by neutron star oscillations.  It is shown that for oscillation modes with high harmonic numbers $(l,m)$ changes in the Goldreich-Julian charge density caused by pulsations of neutron star could lead to significant altering of an accelerating electric field in the polar cap of pulsar.  In the moderately optimistic scenario, that assumes excitation of the neutron star oscillations by glitches, it could be possible to detect altering of the pulsar radioemission due to modulation of the accelerating field. ",
        "introduction": "\\label{sec:introduction}  Neutron stars (NS) are probably the most dense objects in the Universe.  There are extreme physical conditions inside NS, i.e. the magnetic field is close to the quantum limit, the pressure is of the order of the nuclear one and the typical radius of a NS is only about 2-3 times larger than its gravitational radius.  Knowledge of properties of matter under such extreme conditions if very important for fundamental physics.  It is impossible to reconstruct such physical circumstances in terrestrial laboratories, therefore, study of the NS's internal structure would give an unique opportunity for experimental verification of several fundamental physical theories.   To study interiors of a celestial body one have to perform some kind of seismological study, by comparing observed frequencies of eigenmodes with frequencies inferred from theoretical considerations. Eigenfrequencies of oscillations in the crust of NS as well as in its interiors were calculated in several papers \\citep[see e.g.  ][and references there]{McDermott/1988,Chugunov2006}.  However, for seismological study there should exist both i) a mechanism for excitation of oscillations, ii) a mechanism modulating radiation of the celestial object.   There are two types of known NSs: member of binary systems and isolated ones.  The former radiate due to accretion of the matter from the companion.  For these stars there are many possibilities to excite oscillations, for example by instabilities in the accretion flow. However, in this case it would be difficult to distinguish whether a particular feature in the power spectrum of the object is due to oscillations of the NS or it is caused by some processes in the accretion disc/column.  Because of this ambiguity we think that the study of isolated NS should be more promising in regard of the seismology.    The vast majority of known isolated neutron stars are radiopulsars. The glitch (sudden change of the rotational period) is probably the only possible mechanism for excitation of oscillations for isolated pulsars.  Radiation of radiopulsars is produced mostly in the magnetosphere.  In order to judge whether oscillations of the NS could produce detectable changes in pulsar radiation, the impact of the oscillations on the magnetosphere must be considered.  There is a widely accepted model of radiopulsar as a highly magnetized NS surrounded by non-neutral plasma \\citep{GJ}.  Although, there is still no self-consistent theory of radiopulsars, there is a general agreement regarding basic picture for the processes in the magnetosphere.  Oscillations of the star can generate electric field as it happens in the case of rotation.  Generalization of the formalism developed for rotating NS to the case of oscillating star should help to obtain the desired information.    The first attempt to generalize Goldreich-Julian (GJ) formalism to the case of oscillating NS was made in \\citet{TBS2000}.  It was developed a general algorithm for calculation of the GJ charge density in the near zone of an oscillating NS.  Using this algorithm GJ charge density and electromagnetic energy losses were calculated for the case of toroidal oscillations of the NS.  Here we apply this formalism to the case of spheroidal oscillation modes, representing wide class of stellar modes (r-,g-,p- modes).  We consider also impact of stellar oscillations on the acceleration mechanism in the polar cap of pulsar and discuss the possibility of observation on the NS oscillations. ",
        "conclusions": "\\label{sec:discussion}   \\begin{figure} \\centering \\includegraphics[clip,totalheight=6cm]{fig7_1.eps} \\ \\includegraphics[clip,totalheight=6cm]{fig7_2.eps} \\ \\includegraphics[clip,totalheight=6cm]{fig7_3.eps} \\caption{Dependency of the Goldreich-Julian charge density on the harmonic number $m$. Distribution of \\Rgj{} on the NS surface in the polar cap of pulsar is shown for spheroidal oscillation modes (from left to right) $(64,4)$, $(64,8)$, $(64,14)$ in polar coordinates $(\\left|\\Rgj^{lm}(\\RS,\\theta,0)\\right|, \\theta)$. Meaning of the lines is the same as in Fig.\\ref{fig:Rgj_increase_l}.} \\label{fig:Rgj_increase_m} \\end{figure}   \\begin{figure} \\centering \\includegraphics[width=.9\\columnwidth]{fig8.eps} \\caption{$\\Delta\\rho\\equiv\\rho-\\Rgj$ -- difference between the charge density of a space charge limited flow and the local Goldreich-Julian charge density along magnetic field lines in the polar cap of pulsar.  $\\Delta\\rho$ is shown (in arbitrary units) by solid lines for spheroidal modes with different ($l,m$), from top to bottom: (64,14), (54,14),(64,2),(54,2). The same relation for an aligned rotator is shown by the dashed line. Negative values of $\\Delta\\rho$ give rise to an accelerating electric field, negative ones -- to a decelerating field.} \\label{fig:R_m_Rgj} \\end{figure}    We have shown, that oscillations of the NS can induce changes in the accelerating electric field, which are more stronger than a naive estimation $(\\Vosc/c)B$.  Indeed, for high harmonics the induced electric field will be $\\sim{}l$ times stronger.  In order to make definitive predictions about observable parameters is necessary to study the polar cap acceleration zone more detailed.  An accelerating electric potential and the height of the pair formation front should be calculated.   However, we can do simple estimations using equation~(\\ref{eq:l_Eacc_Edec_V}).  Let us estimate the harmonic number of the mode which can cancel the accelerating electric field, assuming the mode is excited by a glitch.  As we pointed in Sec.~\\ref{sec:general-formalism} only modes with non-zero amplitude on the NS surface can produce changes in the magnetosphere.  The distribution of oscillational motion plays crucial role.  If oscillations are trapped in the NS crust a rather small energy will be required to pump the oscillation amplitude to the level high enough for strong disturbance of the accelerating electric field.  If a fraction $\\epsilon$ of the NS mass $\\NS{M}$ is involved in the oscillations the amplitude of the oscillational velocity is of the order \\begin{equation} \\label{eq:Vosc_MNS} \\Vosc \\sim \\sqrt{ \\frac{2\\Osc{W}}{\\epsilon \\NS{M}} }\\:, \\end{equation} where $\\Osc{W}$ is the total energy of the mode.  The energy transferred during the glitch of the amplitude $\\Delta\\Omega$ is \\begin{equation} \\label{eq:Wglitch} W^\\mathrm{glitch} = i\\;\\NS{I} \\Omega \\Delta \\Omega = i\\;\\NS{I} \\left(\\frac{2\\pi}{P}\\right)^2 \\frac{\\Delta\\Omega}{\\Omega} \\:, \\end{equation} where $i$ is the fraction of the total momentum of inertia of the NS $\\NS{I}$ coupled to the crust.  Let us assume that some fraction $\\eta$ of this energy goes into excitation of oscillations.  Using eqs.~(\\ref{eq:Wglitch}),~(\\ref{eq:Vosc_MNS}),~(\\ref{eq:l_Eacc_Edec_V}) we get conditions for harmonic number of modes which would periodically cancel the accelerating electric field, as \\begin{equation} \\label{eq:l_blok_final_SCLF} {\\bf l} > 300 \\; {\\eta_\\%^{-1/2}} \\; {\\epsilon^{1/2}}\\; {i^{-1/2}}\\; \\left( \\frac{\\Delta\\Omega}{\\Omega} \\right)_6^{-1/2} \\;, \\end{equation} for SCLF model, and for \\citet{Ruderman/Sutherland75} model: \\begin{equation} \\label{eq:l_blok_final_RS} {\\bf l} > 2000 \\; {\\eta_\\%^{-1/2}} \\; {\\epsilon^{1/2}}\\; {i^{-1/2}}\\; \\left( \\frac{\\Delta\\Omega}{\\Omega} \\right)_6^{-1/2} \\;. \\end{equation} Here $\\eta_\\%$ is measured in per cents and the relative magnitude of the glitch $\\left(\\Delta\\Omega/\\Omega\\right)_6$ is normalized to $10^{-6}$.  It is widely accepted, that the origin of pulsar glitches is angular momentum transfer from the NS core to the crust. In the frame of this model the fraction of the energy which can go into excitation of NS oscillation $\\eta$ is of the order of $\\Delta\\Omega/\\Omega$, i.e. it is very small, of the order of $\\sim{}10^{-6}$.  We may speculate however, that excited oscillations are trapped in the NS crust, i.e. $\\epsilon$ is also very small. Such global oscillation modes ($l\\sim$ several hundreds) could induce substantial changes in the accelerating electric field.     As we mentioned in Sec.~\\ref{sec:general-formalism} all physical quantities in the solutions obtained here oscillate with the frequency of the star oscillations.  The accelerating electric field close to the local geometrical maxima of the oscillational GJ charge density will be weakened periodically by the decelerating effect due to stellar pulsations.  The field oscillation will influence the particle distribution in the open field line zone of the pulsar magnetosphere and it should produce some observable effects.  Depending on oscillation mode and position of the line of sight a complicated pattern will appear periodically in individual pulse profiles. Although individual pulses are highly variable, the presence of periodical features should be possible to discover in the power spectra of pulsars, provided the oscillations are excited to a high enough level and observations have been made with hight temporal resolution.  If one observes some feature, which appears just after the glitch, then decreases and disappears after some time, and it never appears in the normal pulsar emission, then one can undoubtedly attribute this feature to the NS oscillations."
    },
    "0609/astro-ph0609460_arXiv.txt": {
        "abstract": "The {\\it Microvariability and Oscillations of stars (MOST)} satellite has detected low-amplitude light variations ($\\Delta m\\sim$1 mmag) in the Be star $\\beta$ CMi (B8Ve). The observations lasted 41 days and the variations have typical periods $\\sim 0.3$ days. We demonstrate that the dominant frequencies are consistent with prograde high-order g-modes of $m=-1$ excited by the Fe-bump of opacity in an intermediate-mass ($\\approx 3.5 M_\\odot$) star with a nearly critical rotation period of 0.38 days. This is the first detection of nonradial g-mode pulsations in a Be star later than B6 leading to the possibility that pulsations are excited in all classical Be stars. ",
        "introduction": "Be stars are non-supergiant, rapid rotators, which at times show emission in the Balmer and certain metallic lines. They are often called ``classical Be stars'' to distinguish them from other emission line stars such as Herbig Ae/Be stars and Algol systems \\citep[see][for a recent review on Be stars]{Por03}. The line emission is assumed to arise in a geometrically thin disk consisting of matter ejected from the central star. Such circumstellar disks have actually been imaged by long-baseline interferometry for several Be stars including $\\beta$ CMi \\citep[e.g.][]{Tyc05}. Our target star, $\\beta$ CMi (HD 58715, HR 2845), is a bright but 'quiet' late-type Be star (B8Ve; $V=2.886, B-V=-0.072$); no photometric variations were found by \\citet{Pav97}. The star shows only gradual and small long-term variations in  H$\\alpha$ emission \\citep[e.g.,][]{Han96} arising from a relatively small disk of $\\sim 3$ stellar radii \\citep[e.g.][]{Tyc05}.   Other Be stars, particularly of early-type, often show short-term photometric variations of less than a few days. The causes of these variations are assumed to be pulsation and rotational modulation \\citep[see][for a review]{Por03}. Many early-type Be stars exhibit line-profile variations indicating the presence of non-radial pulsations  \\citep[e.g.][]{Riv03}. Multi-periodic photometric variations are also the signature of non-radial pulsations such as those found recently by {\\it MOST} for $\\zeta$ Oph \\citep[O9.5Ve;][]{Wal05a} and HD 163868 \\citep[B5Ve;][]{Wal05b}. Until now, there has been no clear indication of pulsations for late-type Be stars. \\citet{Baa89} searched for, but failed to detect any, line-profile variations in B8-B9.5 stars. There were very small photometric variations in a few late-type Be stars, but these were judged to be mono-periodic (e.g., $\\mu$ Pic B9Ve, Balona et al 1992; $\\zeta$ Crv B8Vne, Barrera et al. 1991) and attributed to rotational modulation.   It is of critical importance to know if nonradial pulsations are excited even in late-type Be stars because it would indicate that all Be stars may have nonradial pulsations which could play a critical role in mass ejection. Excitation of g-modes in late-type Be stars is likely because the cool boundary of the distribution of Be stars on the HR diagram \\citep{Zor05} roughly coincides with the cool boundary of the Slowly Pulsating B (SPB) stars \\citep{Pam99}, and because prograde g-modes can be excited even in a rapidly rotating star \\citep{Wal05b}. With this motivation we observed $\\beta$ CMi with {\\it MOST} and indeed discovered multi-periodic light variations, albeit of low amplitude. ",
        "conclusions": "We have found nonradial g-mode pulsations excited in the late-type Be star $\\beta$ CMi (B8Ve). This confirms that $\\beta$ CMi belongs to the new class of SPBe stars, like HD163868 \\citep{Wal05b}. We have identified pulsation modes as prograde-sectoral g-modes of $m=-1$, based on the comparison between the observed frequencies and theoretical ones excited by the Fe bump of opacity in rapidly rotating models. The fact that only one group of frequencies is detected by {\\it MOST} is consistent with our prediction of the pulsational stability if $\\beta$ CMi rotates nearly critically. This supports the recent claims \\citep{Tow04,Cra05} that classical Be stars rotate more rapidly than previously thought \\citep[70--80\\% of the critical rate; e.g.][]{Por96,Cha01}. The observed amplitudes are small ($\\la 1$ mmag). Since $\\beta$ CMi lies close to the red edge of the instability region on the HR diagram, the pulsational instability is not strong and hence the pulsation amplitudes are expected to be small.  Our discovery of nonradial pulsations in $\\beta$ CMi suggests the possibility that nonradial pulsations are involved in all rapidly rotating Be stars. In some early-type Be stars mass ejections are observed to occur at certain pulsation phases \\citep[e.g. $\\mu$ Cen:][]{Por03} suggesting that pulsations trigger mass ejections. It is tempting to speculate that nonradial pulsations may play a crucial role in mass ejection in all Be stars. \\citet{Osa86} argued that prograde nonradial pulsations can transport angular momentum toward the surface and spin up the equatorial region critically leading to a mass ejection. \\citet{Owo05} showed numerically that introducing velocity variations corresponding to a prograde ($m=-2$) nonradial pulsation in a nearly critically rotating stellar atmosphere actually yields a circumstellar Keplerian disk.  Our findings in this paper and \\citet{Wal05b}, that {\\it prograde} sectoral g-modes are selectively excited in rapidly rotating Be stars, support the above scenario for mass ejection. Line-profile-variation (LPV) analyses for early type Be stars, however, tend to indicate the presence of retrograde modes \\citep{Riv03} rather than prograde modes, contrary to our model findings.   \\citet{Riv98} obtained four periods around 0.503 days and two period around 0.28 days from LPVs in $\\mu$ Cen, one of the best studied and most active Be stars. [These LPVs of $\\mu$ Cen were modeled as retrograde ($m=2,~{\\rm and}~3$) modes by \\citet{Riv01}.] Frequency spacings among each group are about 0.034 cycles~day$^{-1}$ or less, which are comparable with the frequency difference 0.025 cycles~day$^{-1}$ between the two main frequencies of $\\beta$ CMi (\\S2) and with typical frequency spacings $\\sim 0.03$ cycles~day$^{-1}$ of HD 163868 \\citep{Wal05b}. These frequency spacings are comparable with those of g-modes, although they are close to the limit of the frequency resolution. Longer observations by MOST or COROT would make g-mode seismology of SPBe stars possible. Also, simultaneous photometric and spectroscopic observations of $\\mu$ Cen and other Be stars showing conspicuous LPVs and mass-ejection events would be important to understand the connection between nonradial pulsations and mass ejections in Be stars."
    },
    "0609/astro-ph0609656_arXiv.txt": {
        "abstract": "We present the UV spectroscopy and timing of three nearby pulsars (Vela, B0656+14 and Geminga) recently observed with the Space Telescope Imaging Spectrograph. We also review the optical and X-ray properties of these pulsars and establish their connection with the UV properties. We show that the multiwavelengths properties of neutron stars (NSs) vary significantly within the sample of middle-aged pulsars. Even larger differences are found between the thermal components of Ge-minga and B0656+14 as compared to those of radio-quiet isolated NSs. These differences could be attributed to different properties of the NS surface layers. ",
        "introduction": "\\label{intro}   Optical through X-rays radiation from a typical isolated neutron star (NS) is expected to exhibit two components: thermal radiation from the NS surface and non-thermal radiation from the NS magnetosphere. The spectrum of {\\em non-thermal radiation} can be described by a power-law (PL) model with a spectral index $-1\\lesssim\\alpha\\lesssim 0$ ($F_{\\nu}\\propto \\nu^{\\alpha}$). This radiation is commonly interpreted as synchrotron emission from relativistic electrons/positrons accelerated by the electric fields near the NS surface and from secondary particles produced in the pair cascades. The magnetospheric radiation  is intrinsically anisotropic and, therefore, it shows strong pulsations. The non-thermal emission completely dominates the multiwavelength spectrum in very young pulsars (e.g., Crab and B0540--69).    As the magnetospheric emission becomes fainter with increasing pulsar age, {\\em thermal emission} becomes detect-able at $\\tau\\equiv P/2\\dot{P}\\gtrsim 10$ kyrs, being seen as a ``thermal hump'' on top of the flat non-thermal spectrum. This thermal radiation is emitted from the NS surface which can be non-uniformly heated (e.g., due to anisotropic heat conductivity of the crust and bombardment by relativistic particles). Indeed, in several cases at least two thermal components with different temperatures are need-ed to fit the X-ray spectrum (e.g., Pavlov et al.\\ 2002; Zavlin \\& Pavlov 2004a; De Luca et al.\\ 2005). The cooler {\\em thermal soft} (TS) component is commonly interpreted as emission from the bulk of the NS surface. The hotter {\\em thermal hard} (TH) component is often attributed to the NS polar caps (PCs) which can be additionally heated by the magnetospheric particles. As expected, the TS temperatures generally decrease with pulsar age. However, they do not fall onto a monotonically decreasing cooling curve, which suggests that NSs may have different masses (e.g., Yakovlev \\& Pethick 2004). In X-rays, the TS component usually shows weaker and broader pulsations than the TH component. The TS pulsations can be caused by non-uniformities of the NS surface temperature and by anisotropy of local emissivity in the strong magnetic field.  Spectra of middle-aged pulsars ($10\\lesssim\\tau\\lesssim500$ kyr) are particulary interesting because they often exhibit all the three emission components. The nearby middle-aged pulsars, Vela ($\\tau\\approx11$ kyrs), B0656+14 (hereafter B0656; $\\tau\\approx110$ kyrs) and Geminga ($\\tau\\approx340$ kyrs), are relatively bright and have been extensively studied in both the optical and X-rays. Given their different ages, the three pulsars provide a representative cut through the population of middle-aged pulsars, allowing one to study the evolution of the pulsar properties.  \\begin{table}[t] \\setlength{\\tabcolsep}{0.85\\tabcolsep} \\caption{Basic properties of three middle-aged pulsars and three RQINSs with measured spin-down parameters and/or parallaxes (Kaplan \\& van Kerkwijk 2005ab; see also van Kerkwijk \\& Kaplan, these proceedings).} \\centering \\label{tab:1}       % \\begin{tabular}{lcccccc} \\hline\\noalign{\\smallskip} Name & Dist.\\ & $~~\\tau$ & $P$ & $\\log\\dot{E}$ & B & $r_{\\rm pc}$ \\\\[3pt] & $~~\\rm{pc}$ &  kyr & s & ergs s$^{-1}$ & TG & m \\\\[3pt] \\tableheadseprule\\noalign{\\smallskip}  Geminga  & $\\!\\!\\!\\!\\!\\!\\sim 200$   & ~340 & 0.24 & 34.52 & 1.6 & 440 \\\\ B0656+14 & $288^{+33}_{-27}$ & ~110 & 0.38 & 34.58 & 4.7 & 346\\\\ Vela & $293^{+19}_{-17}$ & ~~11  & 0.09  & 36.84 & 3.8 & 717\\\\ RX~J0720 & $333^{+167}_{-83}$ & 1900 & ~~8.4 & 30.67 & 24 & ~~~7\\\\ RX~J1308 & $\\!\\!\\!\\!\\!\\!\\sim700$ & 1500 & 10.3  & 30.60 & 34 & ~~~7\\\\ RX~J1856 & $161^{+17}_{-14}$ & $\\!\\!\\sim$400 & ...  & ...  & ... & ...\\\\ \\noalign{\\smallskip}\\hline \\end{tabular}\\\\ \\end{table}  The X-ray spectra of these three pulsars can be described by a three-component, TS+TH+PL, model, with the thermal components modeled as BBs. However, such fits result in substantially different model parameters in each case.  The youngest Vela pulsar shows the highest $T_{\\rm TS}\\approx1.2$ MK emitted from a region with $R_{\\rm TS}\\approx3.7$ km. This $R_{\\rm TS}$ is substantially smaller than a typical NS radius, which might indicate that a large fraction of the NS surface is too cold to be seen in X-rays\\footnote{ Alternatively, it may mean that the BB description of the thermal spectrum is inadequate, and a more realistic atmosphere model should be used. For instance, the Vela pulsar spectrum  fits equally well by a two-component, NS atmosphere plus PL, model, which gives $T\\approx0.7$ MK at $R=13$ km (Pavlov et al.\\ 2001; see also Kargaltsev 2004).}. A factor of ten older B0656 has a lower temperature, $T_{\\rm TS}\\approx0.7$ MK, and a much larger emitting area, $R_{\\rm TS}\\approx12$ km, close to that of a typical NS. The oldest of the three pulsars, Geminga, is even colder, $T_{\\rm TS}\\approx0.5$ MK, and has the largest emitting area ($R_{\\rm TS}\\approx13$ km).  The PL component, with $\\alpha_{\\rm X}\\approx-1$, is the strongest in the Vela pulsar, $L_{X,{\\rm PL}}\\approx3\\times10^{31}$ erg s$^{-1}$ (in 0.2$-$10 keV), becoming weaker in B0656 ($\\alpha_{\\rm X}\\approx -0.5$; $L_{X,{\\rm PL}}\\approx1.4\\times10^{31}$ erg s$^{-1}$), and much weaker in Geminga ($\\alpha_{\\rm X}\\sim-0.6$; $L_{X,{\\rm PL}}\\approx2.2\\times10^{30}$ erg s$^{-1}$). This PL component is typically interpreted as synchrotron emission produced by relativistic particles in the pulsar margentosphere.  The strength of the TH component also appears to decrease with pulsar age. Being relatively strong in the Vela pulsar ($L_{{\\rm TH}}\\approx6.3\\times10^{31}$ erg s$^{-1}$), it becomes weaker in B0656 ($L_{{\\rm TH}}\\approx3.5\\times10^{31}$ erg s$^{-1}$) and decreases dramatically in Geminga ($L_{{\\rm TH}}\\approx4.4\\times10^{29}$ erg s$^{-1}$). The sizes of the TH emission regions are close to the conventional PC radii [$r_{\\rm pc} \\equiv (2\\pi R_{NS}^{3}/cP)^{1/2}$; see Table 1] in Vela and B0656, while in Geminga, $R_{\\rm TH}\\approx46$ m, is surprisingly small when compared to its $r_{\\rm pc}=440$ m. This may suggest that the surface of Geminga is more uniformly heated and the ``TH component'' simply mimics a phase-dependent PL component\\footnote{From a phase-resolved analysis of the XMM data, Jackson \\& Halpern (2005) conclude that the X-ray spectrum of Geminga can be alternatively described by a BB+PL model with a phase-dependent PL slope.}. However, regardless of the exact nature of the TH and PL components, the soft X-ray (0.3$-$2 keV) spectra of these three pulsars are clearly dominated by the thermal emission.  Observing the TS radiation over a wide range of frequencies is particulary important. X-ray observations can only probe the Wein tail of the TS spectrum emitted from the NS surface. This information is insufficient to reconstruct the overall shape of the spectrum  which can deviate from the simple black-body due to the presence of an atmosphere or nonuniformity of surface temperature. Chemical composition of the atmosphere can dramatically affect the X-ray spectrum (Romani 1987; Pavlov et al.\\ 1995; Zavlin, Pavlov \\& Shibanov 1996; Zavlin \\& Pavlov 2002), with light element atmospheres leading to a large Wien excess. Any surface temperature inhomogeneities will also complicate the spectrum, with hotter regions being increasingly important at higher frequencies. For these reasons, comparison of thermal X-ray emission with optical-UV emission on the Rayleigh-Jeans (R-J) side of the thermal hump is particularly valuable. The challenge here is that the non-thermal magnetospheric emission becomes increasingly dominant at longer wavelengths.    \\begin{table*}[t] \\caption{STIS observations} \\centering \\label{tab:1}       % \\begin{tabular}{lllcccc} \\hline\\noalign{\\smallskip} Pulsar & Date & Instrument & Exposure, s & $\\Delta\\lambda^{~\\rm a}$, \\AA & ${\\mathcal F}^{~\\rm b}$, cgs & $f_{p}^{~\\rm c}$, \\% \\\\[3pt] \\tableheadseprule\\noalign{\\smallskip} Geminga  & 2002\\,Feb\\,27 & NUV-MAMA/F25SRF2 & 11367 & 1800-3000 & $1.7\\times 10^{-15}$ & 40 \\\\ & 2002\\,Feb\\,26 & FUV-MAMA/G140L & 10674   & 1155-1702 & $3.7\\times 10^{-15}$ & 65 \\\\ B0656+14 & 2001\\,Sep\\,\\,1   & NUV-MAMA/PRISM & \\,\\,\\,6791    & 1790-2950 & $2.6\\times 10^{-15}$ & 67\\\\ & 2001\\,Nov\\,16 & NUV-MAMA/PRISM & 12761   & 1790-2950 & $2.6\\times 10^{-15}$ & 89\\\\ & 2004\\,Jan\\,20  & FUV-MAMA/G140L & \\,\\,\\,4950    & 1153-1700 & $4.2\\times 10^{-15}$ & 64\\\\ Vela & 2002\\,May\\,28 & NUV-MAMA/F25SRF2 & \\,\\,\\,2895  & 1800-3000 & $7.2\\times 10^{-15}$ & 87 \\\\ & 2002\\,May\\,28 & FUV-MAMA/G140L & \\,\\,\\,3060    & 1153-1701 & $8.2\\times 10^{-15}$ & 73 \\\\ \\noalign{\\smallskip}\\hline \\end{tabular}\\\\ \\tablecomments{$^{~\\rm a}$ -- instrument+filter passband; $^{~\\rm b}$ -- observed flux in the corresponding passband (for the NUV observations of B0656, the flux has been measured from the two observations combined); $^{~\\rm c}$ -- intrinsic (corrected for the background) pulsed fraction.} \\end{table*}  \\begin{figure} \\hspace{0.5cm} \\includegraphics[width=3.2in,angle=0]{opt_to_UV_updated.ps} \\caption{Unabsorbed NIR through FUV spectra of the Geminga, B0656 and Vela pulsars (top to bottom). The PL+BB fits are shown by the solid blue lines. The contributions of the individual components are also shown (blue dashed line for PL, blue dash-dotted line for BB). The dashed red line shows the extrapolation of the TS component fitted to the X-ray spectrum (see also Fig.~3). For B0656 we have also plotted the fit to the NUV+FUV spectrum at the pulse minima (black dotted line). } \\label{fig:1}       % \\vspace{-0.5cm} \\end{figure}   The observations with ground-based telescopes and the Hubble Space Telescope (HST) have provided multi-band photometry from near-IR (NIR) to near-UV (NUV) and showed that the NIR-optical spectra of the three pulsars are non-thermal, with a hint of a R-J component seen in Geminga and B0656+14 at $\\lambda\\lesssim3000$ \\AA. The observations in UBVRIHJ filters show that the optical spectrum of the Vela pulsar remains very flat from NIR to UV, with $\\alpha_{\\rm O}\\approx-0.1$ (Shibanov et al.\\ 2003). For the other two pulsars, the NIR-optical spectra are steeper, $\\alpha_{\\rm O}\\approx-0.4$ (Koptsevich et al.\\ 2001; Shibanov et al.\\ 2006).  For all the three pulsars optical pulse profile measurements have been carried out. The Vela and B0656, show strong non-sinusoidal pulsations (Gouiff\\'es 1997; Kern et al.\\ 2003). Optical pulsations of Geminga were only marginally detected in the B band (Shearer et al.\\ 1998).  Despite the extensive coverage in the optical and X-rays, very few NUV observations had been carried out until recently, and the FUV parts of the spectra have remained completely unexplored. To  fill this gap, we have undertaken an observational campaign with the MAMA-NUV and MAMA-FUV detectors of the Space Telescope Imaging Spectrograph (STIS) aboard the  HST. ",
        "conclusions": "\\label{sec:concl}  We have observed three pulsars (Geminga, B0656+14, and Vela) at the far-UV and near-UV wavelengths with the STIS MAMA detectors and measured the pulsar spectra and pulse profiles. We have also analyzed the X-ray data and found that optical through X-ray spectra consist of thermal and non-thermal components in all the three pulsars. In particular, we found the following.  {\\bf 1.}  Thermal contribution to the FUV passband increases with pulsar age. The thermal component is not seen in the younger Vela pulsar ($\\tau\\approx11$ kyrs), but it dominates the FUV spectrum in the Geminga pulsar ($\\tau\\approx340$ kyrs). As the FUV spectrum grows more thermal, the ``thermal hump'', whose maximum is located in soft X-rays/EUV, shifts toward lower frequencies.  {\\bf 2.} The X-ray spectra of the two older pulsars, Gemin-ga and B0656, resemble those of RQINSs. However, their UV properties are noticeably different. In sharp contrast to RQINSs, whose spectra show large {\\em UV/optical excess}, Geminga shows {\\em UV deficit} while the continuation of the X-ray TS component approximately matches the R-J component in B0656. Likely, this reflects the differences in the magnetic field and phase state of the NS surface layers.    {\\bf 3.} While no UV pulsations have been detected in RQINSs, Geminga shows strong, non-sinusoidal pulsations in the FUV range, where the spectrum is dominated by the thermal component. To explain the strong FUV pulsations, one may need to invoke magnetospheric absorption at certain rotation phases."
    },
    "0609/astro-ph0609526_arXiv.txt": {
        "abstract": "We present data from mid-infrared Keck Telescope imaging of 18 radio-selected ultra-compact \\hii region candidates at diffraction-limited resolution\\footnote{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.}. The goal of these observations is to determine the sizes, luminosities, and morphologies of the mid-infrared emitting dust surrounding the stellar sources. All 18 sources were imaged at 11.7 $\\mu$m and at 17.65 $\\mu$m, and 10 of them were imaged also at 24.5 $\\mu$m. All the sources were  resolved. We have generated dust temperature and optical depth maps and combine them with radial velocity measurements and radio data (1.4 and 5 GHz) to constrain the properties of these star-forming regions. Half of our objects are excited by B-stars, and all our objects have derived types that are later than an O6 star. We find a significant correlation between infrared and radio flux densities, and a weaker one between infrared diameters and the central source ionizing photon rates. This latter correlation suggests that the more compact sources result from later spectral types rather than young age. Our new data may suggest a revision to infrared color selection criteria of ultra-compact \\hii regions at resolutions $\\ltorder 1''$. These 18 sources are part of a sample of 687 sources dominated by ultra-compact \\hii regions selected by matching radio and infrared maps of the first Galactic quadrant by Giveon and coworkers. The new mid-infrared images constitute a significant improvement in resolving sub-structure at these wavelengths. If applied to all of this sample our analysis will improve our understanding of embedded star-formation in the Galaxy. ",
        "introduction": "\\label{intro}  Embedded star-forming regions constitute an important contribution to the overall Galactic star-formation (Wood \\& Churchwell 1989a; Churchwell 2002). The energetic radiation fields of massive stars ionize the surrounding interstellar medium, producing  radio free-free radiation, and heat their natal dust cocoons, generating mid-infrared (MIR) radiation. Measuring the radio and IR luminosities of \\hii regions can therefore constrain the properties of the stellar objects embedded in them. The physical basis for the radio and IR emissions determines the spectral range most suited for these observations. The shape of the thermal free-free spectrum in the radio is such that the optical depth decreases at higher frequencies, enabling more sources to be detected at these frequencies. In the IR, the longer the wavelength, the cooler the objects that can be observed, making the observations more sensitive to more embedded sources in the mid and far-IR (FIR). Understanding the processes and the various evolutionary phases of embedded high-mass stars (e.g., protostars, young stellar objects, and ultra-compact [UC] \\hii regions) requires the study of large, complete samples. Statistically complete samples can establish the distribution of properties such as IR colors (e.g., Wood \\& Churchwell 1989b), radio spectral indices, spatial distributions (e.g., Zoonematkermani et al. 1990; Becker et al. 1994; Giveon et al. 2005a; White, Becker, \\& Helfand 2005), and ultimately, the spectral types and the initial mass-function and formation-rate of massive stars (e.g., Wood \\& Churchwell 1989b).  Numerous studies have addressed the issue of determining the stellar content, the evolutionary phase, and the dust properties of individual embedded sources or small to medium ($<100$ sources) samples of them, by combining IR and radio observations: embedded sources and molecular clouds (Brand \\& Wouterloot 1991; Mookerjea et al. 1999; Saito et al. 2006), protostars and star-forming regions (Homeier 2003; Kraemer et al. 2003; Klein et al. 2005; van der Tak, Tuthill, \\& Danchi 2005; Williams, Fuller, \\& Sridharan 2005), UC and compact \\hii regions (Chini, Kruegel, \\& Wargau 1987; Wood \\& Churchwell 1989a,b; Ball et al. 1996; Mart\\'{i}n-Hern\\'{a}ndez, van der Hulst, \\& Tielens 2003; Crowther \\& Conti 2003; Alvarez et al. 2004), and masers (Testi et al. 1994; De Buizer, Pi\\~{n}a, \\& Telesco 2000; De Buizer 2003a, 2005a; Ellington 2006; Wu et al. 2006). Relatively few studies (e.g., Zoonematkermani et al. 1990; Becker et al. 1994; Giveon et al. 2005a; White, Becker, \\& Helfand 2005; Beltr\\'{a}n et al. 2006) combined IR and radio observations to determine the characteristics of massive star-formation on a Galactic scale by analyzing large samples ($>100$ sources). One of the major drawbacks of these larger samples was the poor angular resolution the IR surveys used -- IRAS with $\\sim 1'$, and the Midcourse Space Experiment (MSX) with $18.3''$, compared to the resolution of the radio surveys (typically $1''$ to a few arcseconds). Not only does the resolution mismatch severely limit comparison between the IR and radio data, it also increases the risk of false positives in IR-radio samples.  This work presents high-resolution diffraction-limited MIR observations of 18 UC \\hii region candidates. These objects are part of a sample of 687 candidates selected by matching IR and radio maps of the first Galactic quadrant (Giveon et al. 2005a). The 18 sources sample almost triples the number of long wavelength ($\\ge 18 \\mu$m) MIR observations of UC \\hii regions observed at high-resolution (De Buizer et al. 2002a,b, 2003b, 2005b, and Kraemer et al. 2003). The IR maps for finding these 687 matches were taken from the MSX archive (Egan, Price, \\& Kraemer 2003), presenting a relative improvement in angular resolution compared to IRAS. The radio maps were generated from VLA observations (White, Becker, \\& Helfand 2005) at an unprecedented sensitivity ($>90\\%$ complete for sources with $F_{5 \\rm GHz}\\ge 3$ mJy) and angular resolution ($6''$). This sample is the most complete of its kind so far, due to the quality and spatial extent of the radio maps. Giveon et al. (2005a) showed that the radio survey is $>90\\%$ complete in detecting all embedded O stars across the Galaxy in the survey's area, assuming ionization-bounded nebulae. The still relatively poor resolution of MSX can lead to confusion regarding IR-radio associations and prohibits the study of source morphologies. With higher IR resolution, it becomes feasible to study the relationships between the central stellar sources, the ionized gas, and the surrounding dust. The goal of the present observations is to verify our selection of UC \\hii regions by determining their sizes, luminosities, morphologies, and stellar types of the ionizing stars.  The SPITZER program, Galactic Legacy Infrared Mid-Plane Survey Extraordinaire (GLIMPSE and GLIMPSE {\\small II}), will provide maps of large parts of the Galactic plane ($|l|\\le 65^{\\circ}$, $|b|\\le 1^{\\circ}$; Benjamin et al. 2003) at wavelengths $<8\\ \\mu$m. The major advantages of this survey compared to the MSX survey is its higher angular resolution ($\\sim 2''$), and its much higher sensitivity (0.4 mJy compared to 0.1 Jy at 8$\\mu$m). While GLIMPSE will be ideal in confirming the radio and IR coincidences, and will help improve our understanding of the morphologies of our full sample of candidates, observations at longer wavelengths are needed for evaluating the dust properties.  The new Keck observations present a factor of 50 improvement in angular resolution compared to MSX. While we have begun our study in a limited sub-sample, especially for our longest wavelength (24.5 $\\mu$m), in the future, we hope to include more sources in the analysis presented in this paper. In \\S \\ref{obs} we describe the observations and present the sample. In \\S \\ref{results} we present flux density maps, dust temperature maps, and dust optical depth maps of our sample, and we derive MIR luminosities and spectral types. In \\S \\ref{discuss} we discuss correlations between the measured and calculated properties. In \\S \\ref{conc} we present our conclusions. ",
        "conclusions": "\\label{conc}  We have presented high-resolution observations of a sub-sample of UC \\hii regions candidates drawn from the sample of Giveon et al. (2005a) in the first Galactic quadrant. We have presented sub-arcsecond resolution flux density maps, dust temperature and emission optical depth maps of these sources, and have studied the relations between their central stellar objects and the dust properties. The improved angular resolution of our observations ($\\times 50$ compared to the MSX; $\\times 150$ compared to IRAS) lead to better pinpointing of the UC \\hii regions within the MIR sources, and thus, to better flux density and morphological determination. Some of the sources are still compact, but the peaks of all sources are resolved.  Our main findings are: \\begin{itemize} \\item[1.] Half of our objects are excited by B-stars and not by O-stars, and all our objects have derived types that are later than an O6 star, even for radio-derived luminosities and assuming far kinematic distances when ambiguous. Only two sources (10\\%) show a significant inconsistecy between the MIR and radio-derived spectral type, which may be explained by the MIR luminosities ($L_{25/18}$) strongly underestimating the bolometric luminosities, or by foreground extinction of the MIR wavelengths (equivalent to $A_V=30$). In general, approximating the bolometric luminosities using the longer wavelengths (17.65 and 24.5 $\\mu$m) is better compared to the shorter wavelengths (11.7 and 17.65 $\\mu$m). The average spectral type derived from $L_{25/18}$ is B0 ($\\pm 1$, i.e., in the range B1--O9) -- the same as the radio-derived average spectral type.  \\item[2.] The majority of our sources are optically thin at MIR wavelengths, with a significantly optically thinner ($\\tau_{11.7}^{max}<0.07$ for all sources) hot dust (150--250 K), compared to an optically thicker ($\\tau_{24.5}^{max}>0.08$ for all sources and $>1$ in two cases) component of warm dust (60--110 K). Most of the optical photons are absorbed by the warm dust.  \\item[3.] The 24.5 $\\mu$m flux densities of our sources are correlated with the 6cm radio flux densities. This relation is fitted by $F_{\\rm 24.5\\mu m}[{\\rm Jy}]=(0.38\\pm 0.05)\\cdot F_{\\rm 5GHz}[{\\rm mJy}] + (45\\pm 8)$. The correlations of the 11.7 and 17.65 $\\mu$m bands are weaker, further supporting our result that the 24.5 $\\mu$m flux density is a better estimator of the total flux density reprocessed by the dust.  \\item[4.] The MIR sizes are correlated with the source ionizing photon rates, confirming the picture of dust cocoons enveloping Str\\\"omgren spheres of ionized gas. We find the relation $\\log N_{lyc}\\ [s^{-1}]=(1.98\\pm 0.09)\\log D_{11.7}\\ [pc]-(50.1\\pm 0.1)$. This correlation suggests that the more compact sources are actually due to later spectral types rather than young age, which may lead to alleviation of the age problem of UC \\hii regions. A possible explanation for missing the earliest O-stars might be a bias imposed by selecting the most compact radio sources (5 GHz diameters $\\le 5''$).  \\item[5.] The new flux density estimations lead on average to redder MIR colors or cooler temperature for our sources compared to IRAS observations. For at least two of our sources, the reddening effect is too big to be exclusively explained by the different bandpass response of IRAS and LWS. We conclude that in these cases the significantly smaller beam of LWS focuses on larger dust columns towards the more accurate positions of the UC \\hii regions, while the larger beam of IRAS averaged out many lines-of-sight. Our sources are thus shifted to redder colors in the Wood \\& Churchwell (1989b) color diagram that is used for selecting UC \\hii regions based on their MIR colors. This may account for some contaminations in selecting UC \\hii regions using the IRAS criteria, and may lead to a change in the estimated star-formation when larger samples are observed at sub-arcsecond resolution.  \\item[6.] The MIR maps show a good overall match in shape when compared to existing radio images with comparable resolution. \\end{itemize}  The significant improvement in angular resolution and the ability to pinpoint the UC \\hii regions lead us to conclude that such an analysis applied to a larger sample would improve our understanding of star-formation in the Milky Way: what are the effects of including a possible large population of B-stars, and the reddening effect we discovered on color-selection criteria of UC \\hii regions in both the Galactic and extragalactic objects? Our observations had a 25\\% rate of non-detections with LWS of sources with high MIR-radio matching probabilities based on the MSX and the VLA data. We suspect the reason is the large difference in angular resolution between LWS and MSX, combined with the small field of view of LWS. This either caused a false match between the IR source and a radio source, or even if the match is real, the IR source might be a low surface brightness source, below the detection limit of LWS. The GLIMPSE survey is most suitable for confirming the MIR-radio associations of sources in our original sample (Giveon et al. 2005a), since it has higher resolution than the MSX. However, it is limited in wavelength. Additional observations which will compensate for that deficiency are thus needed."
    },
    "0609/astro-ph0609710_arXiv.txt": {
        "abstract": "{We present an analysis of high-resolution FLAMES spectra of approximately 50 early B-type stars in three young clusters at different metallicities, NGC\\,6611 in the Galaxy, N\\,11 in the Large Magellanic Cloud (LMC) and NGC\\,346 in the Small Magellanic Cloud (SMC). Using the {\\sc tlusty} non-LTE model atmospheres code, atmospheric parameters and photospheric abundances (C, N, O, Mg and Si) of each star have been determined. These results represent a significant improvement on the number of Magellanic Cloud B-type stars with detailed and homogeneous estimates of their atmospheric parameters and chemical compositions. The relationships between effective temperature and spectral type are discussed for all three metallicity regimes, with the effective temperature for a given spectral type increasing as one moves to a lower metallicity regime. Additionally the difficulties in estimating the microturbulent velocity and the anomalous values obtained, particularly in the lowest metallicity regime, are discussed. Our chemical composition estimates are compared with previous studies, both stellar and interstellar with, in general, encouraging agreement being found. Abundances in the Magellanic Clouds relative to the Galaxy are discussed and we also present our best estimates of the base-line chemical composition of the LMC and SMC as derived from B-type stars. Additionally we discuss the use of nitrogen as a probe of the evolutionary history of stars, investigating the roles of rotational mixing, mass-loss, blue loops and binarity on the observed nitrogen abundances and making comparisons with stellar evolutionary models where possible. ",
        "introduction": "\\label{s_intro}  The differing environments of the Magellanic Clouds, compared with our own Galaxy and with each other, makes them ideal laboratories for the study of both stellar and galactic evolution. Their proximity, combined with their relatively low extinction has contributed to the intensity with which they have been studied in recent years (see, for example, Bouret et al. \\cite{bou03}; Korn et al. \\cite{kor02}; Garnett \\cite{gar99}; Westerlund \\cite{wes97}). For stellar evolutionary theories it is necessary that the present-day chemical composition of the host material in which star-birth occurs is well understood. The interstellar material can be directly sampled via, for example, \\ion{H}{ii} regions (Dufour \\cite{duf84}; Russell \\& Dopita \\cite{rus90}; Peimbert et al. \\cite{pei00}; Kurt \\& Dufour \\cite{kur98}) but such studies may be complicated by depletion onto dust grains.  Young massive stars allow for an alternative means to determine chemical compositions. However, there is evidence (Gies \\& Lambert \\cite{gie92}; Maeder \\cite{mae87}; Langer et al. \\cite{lan98}; Heger \\& Langer \\cite{heg00}) that mixing of CN-cycled material into the stellar photosphere can occur giving rise to N enrichment and to smaller C and O underabundances. Nevertheless by selecting targets that do not appear to have had material mixed into their atmospheres, massive stars can provide reliable indicators of the pristine chemical composition of a recent or ongoing star formation region.  An important challenge for theoretical evolutionary predictions is to reliably model the wide range of N/H abundances seen in observations of A-type supergiants (Venn et al. \\cite{ven99}) as well as in B-type stars, both supergiant and main-sequence (see, for example, Korn et al. \\cite{kor02}, Lennon et al. \\cite{len03} and Trundle et al. \\cite{tru04}). The role of rotational mixing has become increasingly important particularly for main-sequence objects (Maeder \\& Meynet \\cite{mae01}, Meynet \\& Maeder \\cite{mey02} and Lamers et al. \\cite{lam01}). Futhermore, observational evidence suggests that stellar rotational velocity distributions may depend on metallicity whilst cluster stars may have greater rotational velocities than those in the field (Keller \\cite{kel04}, Strom et al. \\cite{str05}). As the amount of mixing depends on the initial stellar rotational velocity, it is important that these effects are fully understood.  Metallicity plays an important role in all aspects of massive star evolution. For example, at low metallicity the stars are more compact, have lower mass-loss rates (Massey \\cite{mas03}) and therefore they may lose less angular momentum and hence may rotate faster (Meynet \\& Maeder \\cite{mey02}). As discussed above, this scenario then leads to enhanced mixing between the stellar interior and surface. Additionally photospheric mixing is more evident at lower metallicity as small absolute changes in the abundances will be easier to observe and hence the Magellanic Clouds provide ideal laboratories within which metallicity effects upon steller evolution can be tested.  Binarity effects have been less studied, both because these effects do not globally affect a given population of stars and also due to the wide range of free parameters that are involved in modelling a binary stystem. However mass-transfer in binary systems may be an important contributor to the abundance variations found between stars within a given population. Additionally, the implications of binarity for other processes such as rotational mixing is not fully understood.  We have undertaken a high resolution spectroscopic survey of young clusters in the Galaxy, LMC and SMC (see Evans et al. \\cite{eva05a}, \\cite{eva05b}) and observations have been obtained for a total of 750 stars, mainly with O and B spectral types. Mokiem et al. (\\cite{mok05a}, \\cite{mok05b}) have derived atmospheric parameters and rotational velocities of the O-type objects in this survey. In this paper we discuss B-type stars with narrow absorption line spectra in the young Magellanic Cloud clusters, N\\,11 and NGC\\,346 along with NGC\\,6611, a Galactic cluster with a similar age.  We determine atmospheric parameters and abundances (C, N, O, Mg and Si) of approximately 50 early B-type stars in these three clusters with the aim of determining the base-line chemical abundances of each region from a large sample of objects, taking into account any chemical enrichments or depletions which may have occured during the stellar lifetime. Additionally we examine our results in terms of the stellar evolution of these objects, through both internal mixing and binary interaction.  In Sect.~\\ref{s_obs} we discuss our target selection. In Sect.~\\ref{s_atmos} we discuss our methodology and present the photospheric abundances derived for each star in our sample along with estimates of the uncertainites in these abundances. In Sect.~\\ref{s_chem} we discuss the elemental abundances and in Sect.~\\ref{s_comparisons} we compare our results to previous stellar and interstellar analyses of the three regions. In Sect.~\\ref{s_discussion} we discuss both abundance differences between the three clusters and between stars within the same region, comparing where possible our observations with current evolutionary theory. Finally in Sect.~\\ref{s_conclusions} we present our best estimates for the present-day chemical composition of the LMC and SMC and summarize the evolutionary effects observed in the sample. ",
        "conclusions": "\\label{s_conclusions}   \\subsection{Chemical composition of the Magellanic Clouds}  Atmospheric parameters and abundances have been estimated for approximately 50 stars in three clusters, NGC\\,6611 in the Galaxy, N\\,11 in the LMC and NGC\\,346 in the SMC. In Table~\\ref{t_finalabund} we present our best estimates for the present-day chemical composition of the LMC and SMC. Our carbon (corrected values), oxygen, magnesium and silicon estimates are taken directly from Table~\\ref{t_averages} rounded to the nearest 0.05~dex. Given the agreement between our lowest observed nitrogen abundance in N\\,11 and NGC\\,346 and the nitrogen abundance estimated from these \\ion{H}{ii} regions we adopt 6.90\\,dex and 6.50\\,dex in Table~\\ref{t_finalabund} as the pristine nitrogen abundance of the LMC and SMC respectively.  Given the large number of stars which have been analysed and the use of a single methodology throughout, with sophisticated non-LTE calculations in determining the abundances given in Table~\\ref{t_finalabund}, we believe that these are the best estimates currently available for the present-day chemical composition of the Magellanic Clouds derived from B-type stars and are suitable for use as the base-line abundances in stellar evolution models.  \\begin{table}[htbp] \\caption{Present-day chemical composition of the LMC and SMC. Abundances are presented on the scale 12+log[X/H]. For comparison the solar abundances of Asplund et al. (\\cite{asp05}) are also given.} \\label{t_finalabund} \\centering \\begin{tabular}{lccc}\\hline \\hline & \\multicolumn{1}{c}{LMC} & \\multicolumn{1}{c}{SMC} & \\multicolumn{1}{c}{Solar}\\\\ \\hline \\\\ C  &                    7.75 & 7.35 & 8.40 \\\\ N  &                    6.90 & 6.50 & 7.80 \\\\ O  &                    8.35 & 8.05 & 8.65 \\\\ Mg &                    7.05 & 6.75 & 7.55 \\\\ Si &                    7.20 & 6.80 & 7.50 \\\\ \\\\ \\hline \\end{tabular} \\end{table}  \\subsection{Chemical evolution of massive stars}  Given the large number of stars which we have analysed and the wide range of observed nitrogen abundances, the evolutionary effects of rotational mixing, mass-loss, blue loops and binarity have been discussed and compared with theory where possible. Variable rates of rotational mixing and mass-loss can be used to explain the observed photospheric nitrogen enhancements of the main-sequence objects. However for objects beyond the main-sequence, particularly low mass objects, these effects do not account for the observed abundances in many cases. At LMC metallicity Lamers et al. (\\cite{lam01}) predictions of nitrogen abundances by the end of the main-sequence are greater than those observed in almost all cases. At SMC metallicity, the models of Maeder \\& Meynet (\\cite{mae01}) do not predict the high nitrogen abundances observed in many B-type giant and supergiant objects. The nitrogen abundances predicted by the end of the blue loops at SMC metallicity are in good agreement with those observed in many of the non-main-sequence objects, but these loops do not extend far enough into the blue or to high enough masses in current stellar evolutionary calculations to make this a viable explanation. Tidal forces in binary systems may supress rotational mixing and all our main-sequence binary objects have close to baseline nitrogen abundances. Several evolved binary objects have high nitrogen enhancements, which may be explained via mass-transfer events. However these abundances are in agreement with those observed from apparently single stars and hence binary evolution may not be necessary to account for the observed abundances.  Although this analysis represents one of the largest systematic abundance analyses of B-type stars in the Magellanic Clouds, much still remains unclear about the chemical evolution of these objects. In particular, as only narrow lined stars have been analysed, our main-sequence sample may be biased towards slowly rotating stars. Additionally this analysis highlights the need for evolutionary models that adopt the correct initial chemical composition of these regions. In a future paper we intend to estimate nitrogen abundances (or limits) for all the B-type stars (including fast rotators) observed by Evans et al. (\\cite{eva05b}), some 300 objects from two LMC and two SMC clusters, in order to place better constraints on the evolutionary processes that occur in massive stars."
    },
    "0609/astro-ph0609532_arXiv.txt": {
        "abstract": "We present the results of a 9.3 square degree infrared ($ZYJHK$) survey in the Upper Scorpius association extracted from the UKIRT Infrared Deep Sky Survey (UKIDSS) Galactic Cluster Survey Early Data Release. We have selected a total of 112 candidates from the ($Z-J$,$Z$) colour-magnitude diagram over the $Z$ = 12.5--20.5 magnitude range, corresponding to M = 0.25--0.01 M$_{\\odot}$ at an age of 5 Myr and a distance of 145 pc. Additional photometry in $J$ and $K$ filters revealed most of them as reddened stars, leaving 32 possible members. Among them, 15 have proper motion consistent with higher mass members from Hipparcos and optical spectra with strong H$\\alpha$ in emission and weak gravity features. We have also extracted two lower mass candidate members for which no optical spectra are in hand. Three members exhibit strong H$\\alpha$ equivalent widths ($>$20\\AA{}), suggesting that they could still undergo accretion whereas two other dwarfs show signs of chromospheric activity. The likelihood of the binarity of a couple of new stellar and substellar members is discussed as well. ",
        "introduction": "\\label{USco:intro}  The Scorpius Centaurus is the nearest OB association, located at a mean distance of 145$\\pm$2 pc from the Sun \\citep{deBruijne97}. The association covers 120 square degrees and is composed of 3 subgroups: Upper Scorpius, Upper Centaurus Lupus, and Lower Centaurus Crux \\citep{blaauw64,deZeeuw99}. The age of the Upper Sco association is about 5 Myr with little scatter \\citep{preibisch02}. The region is free of extinction with Av $\\leq$ 2 mag, suggesting that star formation has already ended \\citep{walter94}.  The first survey conducted in Upper Sco was performed with the Einstein X-ray satellite by \\citet{walter94}, yielding 28 low-mass pre-main-sequence stars in 7 square degrees. Later on, \\citet{deBruijne97} extracted 115 members using kinematic information and trigonometric parallaxes for 1215 Hipparcos stars. A larger scale X-ray survey carried out by \\citet{preibisch98} revealed several hundreds of X-rays sources confirmed spectroscopically as young pre-main-sequence members of the Scorpius Centaurus association. Deeper optical surveys complemented by near-infrared photometry \\citep{ardila00,martin04,slesnick06} crossed the hydrogen-burning limit with the discovery of members with spectral types later than M6 \\citep{martin96,luhman98}. The current census of spectroscopic brown dwarfs in Upper Sco includes 46 M6, 10 M7, 6 M8, and 2 M9 dwarfs, all confirmed as members on the basis of their chromospheric activity and weak gravity features \\citep{martin04,slesnick06}. Recently, we extended the cluster sequence down to 0.008 M$_{\\odot}$ using UKIDSS Galactic Cluster Survey science verification data. A dozen new brown dwarf candidates with spectral types later than M9 and masses below 0.02 M$_{\\odot}$ were extracted over 6.5 square degrees (Lodieu et al., subm.\\ to MNRAS).  \\begin{figure*} \\centering \\includegraphics[width=1.00\\linewidth]{0882fig1.ps} \\caption{The UKIDSS GCS coverage in the Upper Sco association from the Early Data Release is depicted with boxes. Early-type stars from Hipparcos are displayed as 5-pointed star symbols \\citep{deZeeuw99}. We have included known members: triangles are from \\citet{preibisch02}, diamonds from \\citet{ardila00}, squares from \\citet{martin04}, and plusses from \\citet{slesnick06}. The new spectroscopic members and non-members from this study are shown as filled circles and star symbols, respectively. Diamonds located in the area surveyed by the EDR are potential reddened members. Filled circles outside the boxes represent new members of the association extracted from the UKIDSS GCS science verification (Lodieu et al.\\ 2006). Spectroscopic non-members are located in an area affected by reddening at (RA,dec)$\\sim$(250$^{\\circ}$,$-$21.8$^{\\circ}$). } \\label{fig_USco:coverage} \\end{figure*}  \\begin{table} \\centering \\caption{Log of the observations. Central coordinates (in J2000) of each WFCAM tile is provided along with the observing date. Each tile was observed in all five passbands ($ZYJHK$). } \\label{tab_USco:log_obs} \\begin{tabular}{c c c c c c} \\hline Tile  &  R.A.\\  & Dec.\\ & Date & \\\\ \\hline 1    &    15    42    &    --21    36    &    2005-06-05     \\\\ 2    &    15    46    &    --21    38    &    2005-06-05     \\\\ 3    &    15    50    &    --21    38    &    2005-06-05     \\\\ 4    &    15    53    &    --21    38    &    2005-06-05     \\\\ 5    &    16    13    &    --25    16    &    2005-06-11     \\\\ 6    &    16    25    &    --25    03    &    2005-06-14    \\\\ 7    &    16    28    &    --21    36    &    2005-06-05     \\\\ 8    &    16    31    &    --21    38    &    2005-06-05     \\\\ 9    &    16    35    &    --21    38    &    2005-06-05     \\\\ 10   &    16    36    &    --25    16    &    2005-06-06     \\\\ 11   &    16    39    &    --21    38    &    2005-06-05     \\\\ 12   &    16    40    &    --25    09    &    2005-06-06    \\\\ \\hline \\end{tabular} \\end{table}  The UKIRT Infrared Deep Sky Survey \\citep[UKIDSS;][]{lawrence06} constitutes the new generation of deep large-scale infrared surveys. The project is subdivided into 5 surveys: The Large Area Survey, the Galactic Cluster Survey (GCS), the Galactic Plane Survey, the Deep Extragalactic Survey, and the Ultra-Deep Survey. The GCS will survey over thousand square degrees in 10 star-forming regions and open clusters down to $K$ = 18.4 mag at two epochs. The main goal of the GCS is to study the Initial Mass Function well down into the substellar regime in an homogeneous manner in 10 star-forming regions and open clusters.  In this paper we concentrate on the selection of new cluster member candidates in a 9.3 square degree area in the Upper Scorpius association from the UKIDSS Early Data Release \\citep[EDR;][]{dye06}. The photometric, proper motion, and spectroscopic results presented in this paper represent the first step towards the ultimate goal of the GCS\\@. In Sect.\\ \\ref{USco:ukidss_GCS} we briefly present the infrared photometric observations conducted in Upper Sco. In Sect.\\ \\ref{USco:cand} we describe the selection of new candidates from various colour-magnitude diagrams complemented by proper motion. In Sect.\\ \\ref{USco:spectro} we present the optical spectroscopic follow-up carried out with the ESO 3.6-m telescope. We infer spectral types and measure equivalent widths for the H$\\alpha$ emission line and Na{\\small{I}} doublet to assess the membership of the new candidates. Finally, we summarise our work in Sect.\\ \\ref{USco:summary}. ",
        "conclusions": "\\label{USco:summary}  We have presented the analysis of a 9.3 square degree survey in the Upper Scorpius association conducted for the UKIDSS Galactic Cluster Survey and extracted from the Early Data Release. Combining infrared photometry, proper motion, and optical spectroscopy, the membership of the original 112 candidates can be summarised as follows: \\begin{enumerate} \\item 80 photometric candidates selected from the ($Z-J$,$Z$) colour-magnitude diagram were rejected on the basis their $J-K$ affected by reddening. Among them, 15 sources could be reddened members as they exhibit proper motion consistent with the association \\item 13 photometric candidates have proper motions inconsistent with the Upper Sco association and optical spectra of early-type stars affected by reddening \\item 3 objects have optical spectroscopy inconsistent with membership despite their photometry and proper motion suggesting that they are possible members \\item 15 candidates are definite members confirmed from photometry, proper motion, and optical spectroscopy, including 4 new brown dwarfs (spectral types later than M6). They span masses between 0.2 and 0.01 M$_{\\odot}$, assuming an age of 5 Myr and a distance of 145 pc for Upper Sco \\item 2 faint photometric candidates with estimated masses below 20 M$_{\\rm Jup}$ according to the DUSTY models \\citep{chabrier00c}. One object has proper motion consistent with Upper Sco and no spectroscopy whereas the other one has neither proper motion nor spectroscopy currently available \\end{enumerate} We detected variation in the H$\\alpha$ equivalent widths for one Upper Sco members and one field dwarf, hence showing signs of chromospheric activity. In addition, 3 new members exhibit H$\\alpha$ emission lines stronger than the lower limit of the empirical relation between accreting and non accreting low-mass stars and brown dwarfs. Finally, a handful of members present a high probability of being binary systems from their location in the colour-magnitude diagrams."
    },
    "0609/astro-ph0609018_arXiv.txt": {
        "abstract": "% We study the luminosity function (LF) of the star clusters in M51. Comparing the observed LF with the LF resulting from artificial cluster populations suggests that there exists an upper mass limit for clusters and that this limit and/or the cluster disruption varies with galactocentric distance. ",
        "introduction": "In cluster populations with a power law cluster initial mass function (CIMF) the highest appearing mass in a sample is determined by the number of clusters (besides statistical fluctuations). This is no longer true if there exists a physical upper mass limit for clusters.  Information about the CIMF can be derived from the present day luminosity function (LF). Although there is not a one-to-one relation between the CIMF and the LF, modeling the LF with an artificial cluster population with varying CIMF parameters can aid the interpretation of the LF. According to the models of \\cite{gieles06}, a bend in the LF (i.e. a double power law distribution function) may be the signature of a CIMF, truncated at the high mass end. The (filter dependent) location of the bend is intimately connected to the value of the upper mass limit (brighter bends correspond to higher possible cluster masses).  Bends in the LF are observed in NGC 6946, M51 \\cite{gieles06, gieles06a} and NGC 4038/4039 \\cite{whitmore99}. Here we again examine the cluster population of M51, dividing it in subsamples at different galactocentric radii. By comparing to the artificial populations we draw conclusions on the CIMF and cluster disruption parameters across the disk. ",
        "conclusions": "Although more detailed modeling is necessary for quantitative results, the conclusions following the models of \\cite{gieles06} for the galactocentric dependence of the upper mass limit and disruption parameters are: \\begin{enumerate} \\item Subsets closer to the center have their bends at brighter magnitudes, suggesting higher possible cluster masses. \\item Subsets closer to the center have shallower faint-end slopes, indicating faster disruption. \\end{enumerate} The two effects are not to be seen separately though. Faster disruption alters the location of the bend in the same direction as a higher upper mass limit. In a follow up study the statistical significance of the results will be further investigated and a more quantitative analysis will be given \\cite{haas06}."
    },
    "0609/astro-ph0609642_arXiv.txt": {
        "abstract": "{% There is still no consensus about progenitor masses of Type IIP supernovae. }{% We study a normal Type IIP SN 1999em in detail and compare it to a peculiar Type IIP SN 1987A. }{% We computed the hydrodynamic and time-dependent atmosphere models interpreting simultaneously both the photometric and spectroscopic observations. }{% The bolometric light curve of SN 1999em and the spectral evolution of its H$\\alpha$ line are consistent with a presupernova radius of $500\\pm200 R_{\\sun}$, an ejecta mass of $19.0\\pm1.2 M_{\\sun}$, an explosion energy of $(1.3\\pm0.1)\\times10^{51}$ erg, and a radioactive $^{56}$Ni mass of $0.036\\pm0.009 M_{\\sun}$. A mutual mixing of hydrogen-rich and helium-rich matter in the inner layers of the ejecta guarantees a good fit of the calculated light curve to that observed. Based on the hydrodynamic models in the vicinity of the optimal model, we derive the approximate relationships between the basic physical and observed parameters. The hydrodynamic and atmosphere models of SN 1999em are inconsistent with the short distance of 7.85 Mpc to the host galaxy. }{% We find that the hydrogen recombination in the atmosphere of a normal Type IIP SN 1999em, as well as most likely other Type IIP supernovae at the photospheric epoch, is essentially a time-dependent phenomenon. It is also shown that in normal Type IIP supernovae the homologous expansion of the ejecta in its atmosphere takes place starting from nearly the third day after the supernova explosion. A comparison of SN 1999em with SN 1987A reveals two very important results for supernova theory. First, the comparability of the helium core masses and the explosion energies implies a unique explosion mechanism for these core collapse supernovae. Second, the optimal model for SN 1999em is characterized by a weaker $^{56}$Ni mixing up to $\\approx 660$ km\\,s$^{-1}$ compared to a moderate $^{56}$Ni mixing up to $\\sim 3000$ km\\,s$^{-1}$ in SN 1987A, hydrogen being mixed deeply downward to $\\sim 650$ km\\,s$^{-1}$. } ",
        "introduction": "The supernova (SN) 1999em was discovered by the Lick Observatory Supernova Search on October 29.44 UT in the nearly face-on SBc galaxy NGC 1637 (Li \\cite{li99}). Detected shortly after the explosion at an unfiltered magnitude of $\\sim 13.5$, SN 1999em was bright enough to be observed well both photometrically and spectroscopically for more than 500 days (Hamuy et al. \\cite{hpm01}; Leonard et al. \\cite{lfg02}; Elmhamdi et al. \\cite{edc03}). SN 1999em was the first Type II-plateau supernova (SN IIP) detected at both X-ray and radio wavelengths, being the least radio luminous and one of the least X-ray luminous SNe (Pooley et al. \\cite{plf02}). The X-ray data indicated the interaction between SN ejecta and a pre-SN wind with a low mass-loss rate of $\\sim 2 \\times 10^{-6} M_{\\sun}\\,yr^{-1}$. Leonard et al. (\\cite{lfab01}) presented the first spectropolarimetry of SN IIP based on the optical observations of SN 1999em during $\\sim 160$ days after SN discovery. The weak continuum polarization increasing from $p \\approx 0.2\\%$ on day 7 to $p \\approx 0.5\\%$ in the final observations was detected with an unchanging polarization angle $\\theta \\approx 160^{\\circ}$. To date this event has become the most studied normal SN IIP.  In order to study properly any individual object, an accurate distance is required in addition to high quality observations. It is this point that is not clear for SN 1999em. There was a reasonable agreement in determining the distance to SN 1999em using the expanding photosphere method (EPM; Kirshner \\& Kwan \\cite{kk74}): $7.5 \\pm 0.5$ Mpc (Hamuy et al. \\cite{hpm01}), $8.2 \\pm 0.6$ Mpc (Leonard et al. \\cite{lfg02}), and $7.838 \\pm 0.331$ Mpc (Elmhamdi et al. \\cite{edc03}). However, Leonard et al. (\\cite{lknt03}) identified 41 Cepheid variable stars in the galaxy NGC 1637 and found that the Cepheid distance to the host galaxy was $11.7 \\pm 1.0$ Mpc. This distance to SN 1999em is $\\sim 50\\%$ larger than the earlier EPM distance estimates. Assuming a correlation between the total $^{56}$Ni mass and the explosion energy for SNe IIP, Nadyozhin (\\cite{n03}) evaluated the distance to the galaxy NGC 1637 of 11.08 Mpc close to the Cepheid distance. Using the spectral-fitting expanding atmosphere model (SEAM), Baron et al. (\\cite{bnbh04}) obtained the distance to SN 1999em of $12.5 \\pm 1.8$ Mpc in a good agreement with the Cepheid distance scale. Finally, studying various ingredients entering the original EPM, Dessart \\& Hillier (\\cite{dh06}) improved this method and also achieved better agreement to the Cepheid distance with an estimate of $11.5 \\pm 1.0$ Mpc.  Starting the investigation of SN 1999em, a normal SN IIP, it is impossible to pass by the well-known and studied SN 1987A in the Large Magellanic Cloud (LMC), a peculiar SN IIP. Recent progress in reproducing the bolometric light curve observed in SN 1987A with a modern hydrodynamic model (Utrobin \\cite{utr04}) and in modelling the H$\\alpha$ profile and the Ba II 6142 \\AA\\ line in SN 1987A at the photospheric phase with a time-dependent approach for the SN atmosphere (Utrobin \\& Chugai \\cite{uc02}, \\cite{uc05}) makes this experience very instructive for other SNe IIP. First, it turned out that a good agreement between the hydrodynamic models and the photometric observations of SN 1987A did not guarantee a correct description of this phenomenon as a whole. Second, the strength of the H$\\alpha$ line and its profile provided hard constraints on the hydrodynamic models. The most important lesson from this study of SN 1987A is that we have to take \\emph{both} the photometric \\emph{and} spectroscopic observations into account to obtain an adequate hydrodynamic model (Utrobin \\cite{utr05}). Now such an approach should be applied to SN 1999em.  Here we present the comprehensive hydrodynamic study of SN 1999em complemented by the atmosphere model with the time-dependent kinetics and energy balance. A brief description of the hydrodynamic model and the atmosphere model based on it is given in Sect.~\\ref{sec:snmodel}. The study of SN 1999em we begin with a construction of the optimal hydrodynamic model (Sect.~\\ref{sec:optmod}). Then we continue with a question: at what distance to SN 1999em, short or long, the hydrodynamic model and the proper atmosphere model are consistent with the photometric and H$\\alpha$ line observations (Sect.~\\ref{sec:evidis}). The time development of the optimal hydrodynamic model is presented in Sect.~\\ref{sec:devmod}, and its general regularities in Sect.~\\ref{sec:genpro}. The basic relationships between the physical and observed parameters for SNe IIP similar to the SN 1999em event are obtained in Sect.~\\ref{sec:phyobs}, while in Sect.~\\ref{sec:sn87a} we address the comparison of SN 1999em with SN 1987A. In Sect.~\\ref{sec:disc} we discuss our results from the theoretical and observational points of view. Finally in Sect.~\\ref{sec:concl} we summarize the results obtained.  We adopt here a distance of 11.7 Mpc (Leonard et al. \\cite{lknt03}), a recession velocity to SN 1999em of 800 km\\,s$^{-1}$ (Leonard et al. \\cite{lfg02}), an explosion date of JD 2451476.77, and a total extinction $A_V=0.31$ (Baron et al. \\cite{bbh00}; Hamuy et al. \\cite{hpm01}; Elmhamdi et al. \\cite{edc03}). ",
        "conclusions": "\\label{sec:concl} We have presented a comprehensive study of a normal Type IIP SN 1999em. Our main aim in this paper was to show the absolute necessity of the simultaneous interpretation of \\emph{both} the photometric \\emph{and} spectroscopic observations. To do this, we calculated the hydrodynamic models of SN 1999em complemented by the atmosphere models with the time-dependent kinetics and energy balance. We focused mainly on the optimal hydrodynamic model of SN 1999em and subsequent comparison to a peculiar Type IIP SN 1987A. Our results can be summarized in the following conclusions.  \\begin{itemize}  \\item[--] The bolometric light curve of SN 1999em and its spectral evolution of the H$\\alpha$ line are consistent with a radius of the pre-SN of $500\\pm200 R_{\\sun}$, a mass of the ejected envelope of $19.0\\pm1.2 M_{\\sun}$, an explosion energy of $(1.3\\pm0.1)\\times10^{51}$ erg, and a mass of radioactive $^{56}$Ni of $0.036\\pm0.009 M_{\\sun}$, the bulk of which is confined to layers ejected with velocities less than $\\approx 660$ km\\,s$^{-1}$. The optimal hydrodynamic model with the helium core mass in the range of 5.6--8.1 $M_{\\sun}$ matches the observed bolometric light curve.  \\item[--] The adequate hydrodynamic and atmosphere models of SN 1999em distinguish between the short distance of 7.85 Mpc, the average value of the EPM distance estimates to the host galaxy, and the Cepheid distance of 11.7 Mpc. They are inconsistent with the short distance of 7.85 Mpc, which should be discarded.  \\item[--] It is shown that the hydrogen recombination in the atmosphere of a normal Type IIP SN 1999em, as well as most likely other SNe~IIP, at the photospheric epoch is essentially a time-dependent phenomenon in accordance with the analysis of Utrobin \\& Chugai (\\cite{uc02}, \\cite{uc05}).  \\item[--] This study reveals the time development of a normal SN~IIP with the following stages: a shock breakout, an adiabatic cooling phase, a phase of cooling and recombination wave, a phase of radiative diffusion cooling, an exhaustion of radiation energy, a plateau tail, and a radioactive tail.  \\item[--] The plateau tail of the theoretical bolometric light curve coinciding in time with the second plateau observed in SN 1999em (Elmhamdi et al. \\cite{edc03}) runs slightly above a line of the radioactive tail in contrast to the latter. It seems that the real bolometric luminosity of SN 1999em is greater than the ``UBVRI'' luminosity during the second plateau. The duration of the plateau tail, as well as the second plateau, is correlated with the amount of ejected $^{56}$Ni.  \\item[--] The bolometric light curve depends weakly on the chemical composition in the outer layers of the ejecta beyond the helium core, while hydrogen is abundant there and controls the opacity of matter. A mutual mixing of hydrogen-rich and helium-rich matter in the inner layers of the ejecta guarantees a good fit of the calculated bolometric light curve to what is observed. Note that the sharp boundary between the hydrogen-rich and helium-rich layers at the edge of the helium core --- a characteristic of the evolutionary models of pre-SNe --- changes the light curve producing an unobserved bump at the end of the plateau.  \\item[--] A contribution of numerous metal lines to opacity plays a fundamental role in reproducing the observed light curve of SN 1999em. The decrease in opacity due to disregarding the line contribution speeds up the radiation diffusion and causes the bolometric luminosity to increase significantly compared to what is observed during nearly the whole outburst and, as a consequence, causes the characteristic duration of the light curve to shorten by about 10 days.  \\item[--] The comparison of a normal Type IIP SN 1999em with a peculiar Type IIP SN 1987A reveals two very important results for SN theory. First, the masses of helium cores in the pre-SNe of SN 1999em and SN 1987A are close enough to suppose that in the final stages of stellar evolution nearly the same iron cores form within the pre-SNe. This fact and roughly the same explosion energies of SN 1999em and SN 1987A together imply a unique explosion mechanism for these core collapse SNe. Second, the optimal model for SN 1999em is characterized by a weaker $^{56}$Ni mixing in velocity space up to $\\approx 660$ km\\,s$^{-1}$ compared to a moderate $^{56}$Ni mixing up to $\\sim 3000$ km\\,s$^{-1}$ in the case of SN 1987A, hydrogen being mixed deeply downward to $\\approx 700$ km\\,s$^{-1}$ and $\\sim 600$ km\\,s$^{-1}$.  \\item[--] A significant excess of the gamma-ray luminosity over the bolometric one, shown by the optimal model of SN 1987A at the radioactive tail, is of great interest and should be taken into account in measuring the $^{56}$Ni amount from the observed bolometric luminosity at this stage when the SN envelope still remains optically thick for gamma-rays. The higher the rate of work done by radiation pressure in an SN envelope, the larger an excess of the gamma-ray luminosity over the bolometric one at the radioactive tail.  \\item[--] It is shown that the approximation of homologous expansion may be used in the EPM and SEAM for determining the distances to normal SNe IIP starting from nearly day 3 after the SN explosion.  \\item[--] Based on the hydrodynamic models in the vicinity of the optimal model of SN 1999em, we derive the approximate relationships between the basic physical and observed parameters. The basic physical parameters are the pre-SN radius, the ejecta mass, the explosion energy, and the total $^{56}$Ni mass. The observed properties of SNe IIP are described by the three parameters: the characteristic duration of the light curve, the bolometric luminosity and the photospheric velocity in the middle of the plateau. To cover a wider range of the parameters, beyond the vicinity of the optimal model, we have to construct a global approximation that should be based on a set of hydrodynamic models of \\emph{many real} SNe IIP.  \\end{itemize}  Finally, we emphasize that the simultaneous analysis of the photometric and spectroscopic observations is in fact crucial for correctly interpreting core collapse SNe and express the hope that the obtained basic parameters and properties of SN 1999em determine its evolutionary state and will be a reliable reference for the proper mechanism of SN explosions."
    },
    "0609/astro-ph0609368_arXiv.txt": {
        "abstract": "% The topic of wind-clumping has been the subject of much activity in recent years, due to the impact that it can have on derived mass-loss rates. Here we present an alternative method of investigating wind-clumping, that of polarimetry. We present simulations of the polarization produced by a clumpy wind, and argue that the observations may be reproduced just by statistical deviations from spherical symmetry when the outflow is only slightly fragmented. Here, the polarization scales with $\\dot{M}$, which is consistent with observations of LBVs, WRs and O supergiants. Finally, we find clumping factors in the inner 2$R_{\\star}$ of $\\sim 2-3$, and speculate as to the clumping stratification of hot stars. ",
        "introduction": "Many massive stars show evidence of variable intrinsic polarization, from OB supergiants \\citep{L-N87}, to Luminous Blue Variables \\citep[LBVs, ][ \\& refs therein]{Davies05}, and Wolf-Rayet stars \\citep[WRs, e.g. ][]{StL87}. This is typically attributed to scattering off density-inhomogeneities, or `clumps', at the base of the wind.  Clumping is an active research area at present due to the significant effect it has on derived mass-loss rates when incorporated into model stellar atmospheres (see e.g. Puls et al., this volume). Typically, quantitative studies of clumping involve comparison of the synthetic spectra produced by these model atmopsheres with observed data. Here we present an alternative avenue of investigation into clumping, that of polarimetry. Below, the basics of the technique are described. The results of a recent spectropolarimetric survey of LBVs are reviewed, and a quantitative investigation into the data is described. ",
        "conclusions": "We show that, through observational monitoring and quantitative modelling, polarimetry can be used to independently investigate the phenomenon of wind-clumping in hot stars. From modelling the polarimetric variability, it is found that the observations are reproduced when the wind consists of a large number of low-density clumps. Here the polarization can arise from statistical deviations from spherical symmetry in only a slightly `fragmented' wind.  It is predicted that polarization scales linearly with mass-loss rate, and so is consistent with result of higher polarization in LBVs than WRs and O supergiants. Short-timescale variability is also predicted, and this has been detected in a polarimetric-monitoring pilot-study. Further observations are planned.  In an investigation of the wind clumping-factor $f_{cl}$, we find that the wind is already significantly clumped in the inner 2$R_{\\star}$, in agreement with recent combined H$\\alpha$/IR/radio observations. The model could be further tested by extending these observations to B/A supergiants."
    },
    "0609/astro-ph0609197_arXiv.txt": {
        "abstract": "We investigated the frequency distributions of flares with and without coronal mass ejections (CMEs) as a function of flare parameters (peak flux, fluence, and duration of soft X-ray flares). We used CMEs observed by the Large Angle and Spectrometric Coronagraph (LASCO) on board the Solar and Heliospheric Observatory (SOHO) mission and soft X-ray flares (C3.2 and above) observed by the GOES satellites during 1996 to 2005.  We found that the distributions obey a power-law of the form: $dN/dX \\propto X^{-\\alpha}$, where $X$ is a flare parameter and $dN$ is the number of events recorded within the interval [$X$, $X+dX$]. For the flares with (without) CMEs, we obtained the power-law index $\\alpha=1.98\\pm0.05$ ($\\alpha=2.52\\pm0.03$) for the peak flux, $\\alpha=1.79\\pm0.05$ ($\\alpha= 2.47 \\pm 0.11$) for the fluence, and $\\alpha=2.49\\pm0.11$ ($\\alpha=3.22\\pm0.15$) for the duration. The power-law indices for flares without CMEs are steeper than those for flares with CMEs. The larger power-law index for flares without CMEs supports the possibility that nanoflares contribute to coronal heating. ",
        "introduction": "Heating of the solar corona is one of the fundamental problems in solar physics. Solar flares have been proposed as a heat source, but the observed flares do not supply enough energy to keep the coronal temperature at million degrees. However, tiny flares known as nanoflares, whose intensity is below the observational limits may be able to heat the corona \\citep{parke88}. Since the nanoflares cannot be detected as discreet events with the current observational capability, their occurrence frequency distribution is often extrapolated from the observed flares. The flare frequency distributions can be represented by a power-law of the form: $dN/dE~\\propto~E^{-\\alpha}$, where $E$ is flare energy and $dN$ is the number of events recorded within the interval [$E$,~$E+dE$]. When $\\alpha < 2$, only larger flares dominantly contribute to coronal heating \\citep{hudso91}, meaning that nanoflares cannot contribute. Flare peak flux or peak count rate have been used to obtain the power-law index since it is difficult to measure the total flare energy. Many authors have examined $\\alpha$ for various parameters of flares and flare-related phenomena, and found to be smaller than 2 \\citep[e.g.,][and references therein]{crosb93,aschw98}. The only exception was for quiet-region flares observed in EUV \\citep[$\\alpha=2.3-2.6$;][]{kruck98}.  After the discovery of CMEs in 1971, the relation between flares and CMEs have been studied extensively (see \\citealt{kahle92} for review). A close relation is also indicated from the similarity between the derivative of the X-ray light curve and CME acceleration profile \\citep{zhang01,vrsna04}. However, not all flares are associated with CMEs. Even X-class flares (about 10\\% of them) lack CME association \\citep{yashi05}. Since a large, uniform and extended data base on CMEs has become available for the first time from SOHO, we can perform an extensive statistical analysis for a detailed examination of flares with and without CMEs. In this paper, we show the frequency distributions for flares with and without CMEs, and discuss their implications for the problems of coronal heating.  \\begin{figure*} \\epsscale{0.9} \\plotone{f1.eps} \\caption{CME association rate as a function of (a) X-ray peak flux, (b) fluence, and (c) duration. The gray straight line is the least-squares fit to the data points.} \\end{figure*} ",
        "conclusions": "Since small flares are unlikely to be associated with CMEs, a power-law index obtained from small flares should be similar to that from flares without CMEs. \\citet{kruck98} examined the distribution of small flares in the quiet regions observed by SOHO/EIT, and found the power-law index to be $2.3 - 2.6$, which is consistent with our result.  However, from X-ray data, \\citet{shimi95} found the index to be in the range $1.5 - 1.6$. He examined the distribution of the transient brightenings in active regions observed by Yohkoh/SXT. A similar power-law index ($1.7 - 1.8$) was found for transient brightenings in the \\ion{Fe}{19} line observed by SOHO/SUMER \\citep{wang06}. The transient brightenings correspond the GOES B-class flares and below, but the obtained power-law indices were very different from ours. The different temperature response of the three instruments above might have resulted in the different power-law indices. Since the power-law index for the smaller flares is not observationally determined yet, more studies are needed before reaching firm conclusions.  \\citet{hudso91} showed that smaller flares are able to contribute dominantly to coronal heating when the power-law index~$\\alpha$ is larger than 2. By separating flares with and without CMEs, we showed that the flare frequency distribution may obey a double power-law distribution.  Flares without CMEs dominate at small flare sizes and with $\\alpha = 2.47 \\pm 0.11$ for fluence, indicating that nanoflares contribute to coronal heating if the frequency distribution keeps the same power-law below the observational limit.  Flares without CMEs are thus a potential source for heating the corona since they do not have energy loss due to CMEs. The CME kinetic energy ranges from $10^{28}$ to $> 10^{32}$ erg \\citep{gopal04a}, which is generally higher than the flare energy. In flares with CMEs, more than half of the released energy is used by CMEs to escape from the Sun. On the other hand, lack of CMEs allows the entire released energy to go into flare thermal energy. This is consistent with the observational result that, for a given flare class, flares without CMEs tend to have a higher temperature than those with CMEs \\citep{kay03}.  CME observations by SOHO/LASCO over the past 10 years enabled us to perfom an extensive statistical analysis of flares with and without CMEs. We examined the CME associations of flares from 1996 to 2005 and found that the CME association rate clearly increases with flare's peak flux, fluence, and duration. These results have been known from the SMM and Solwind era, but the large sample in our study has shown these relations clearer. The primary result of this paper is that the power-law index for the distributions of flares without CMEs is much steeper than that for distributions of flares with CMEs. This result supports the possibility that flares without CMEs is a likely source of coronal heating and is consistent with the observation that flares without CMEs have a higher temperature."
    },
    "0609/gr-qc0609093.txt": {
        "abstract": " ",
        "introduction": "The possibility of extra spatial dimensions, in addition to the three that we perceive in everyday life, is a crucial aspect of a number of constructions in modern particle physics, from string theory to methods to address the hierarchy problem~\\cite{Kaluza:tu,Klein:tv,Rubakov:1983bb,Antoniadis:1990ew,Lykken:1996fj, Arkani-Hamed:1998rs,Antoniadis:1998ig,Randall:1999ee,Randall:1999vf,Lykken:1999nb, Arkani-Hamed:1999hk,Antoniadis:1993jp,Dienes:1998vg,Kaloper:2000jb}. In many of these models, the matter content of our three-dimensional universe is confined to a submanifold, or 3-brane, while gravity, via the equivalence principle, propagates in the entire space-time, or bulk.  Large extra dimension models with matter localized on a 4d brane have been re-popularized recently by their application to the electroweak hierarchy problem \\cite{Arkani-Hamed:1998rs,Antoniadis:1998ig}. In these models the bulk metric, typically a flat hypertorus, is assumed to be unaffected by the energy-momentum content of the brane, which allows for a formulation independently of the total dimension. At the background level, this is justified by the absence of a brane tension. The effect of the brane matter is less trivial in general. For energies small compared to the typical inverse radius $R^{-1}$ of the bulk, the Kaluza-Klein (KK) zero-mode truncation applies as a valid approximation for the bulk metric. The non-trivial curvature induced on the bulk metric may be studied above these energies by including the tower of non-zero KK gravitons as effective 4d particles. Brane-world cosmology with flat bulk has received much careful study, although previous work has mostly focused on the low-energy regime (although far above $R^{-1}$), under a ``normalcy temperature\" $T_*$ evaluated in \\cite{Arkani-Hamed:1998nn, Macesanu:2004gf}, where the bulk size is assumed to have been already stabilized by an unspecified high-energy mechanism and the non-zero KK gravitons do not have a significant effect on the evolution of the brane metric. The effect of brane matter on the bulk metric evolution at earlier times, and the consequent alteration of the brane metric evolution, has not yet been studied systematically in an arbitrary number of dimensions.  The problem extends to unwarped brane models with non-flat but homogeneous transverse space. The conclusions of cosmological studies where the brane is a point in a compact hyperbolic transverse space \\cite{Starkman:2001xu,Starkman:2000dy}, for instance, are valid only as long as the brane content does not significantly alter the bulk homogeneity condition.  Another very interesting class of brane models considers the cosmological evolution resulting from an induced kinetic term for gravity localized on the brane in addition to the bulk one \\cite{Deffayet:2001xs,Deffayet:2001pu,Dvali:2000hr, Deffayet:2002sp,Deffayet:2000uy,Dvali:2003rk,Lue:2004rj,Lue:2002sw} ; we focus here on the more simple case of models with pure bulk gravity.  Compact and non-compact models with extra dimensions warped by the presence of a large brane tension provide an elementary example in which the backreaction effect is taken into account. In five dimensions, the brane is codimension one and leads to an exponential warping useful for explaining the Standard Model hierarchies as a gravitational effect \\cite{Randall:1999ee}, or for localizing gravity with non-compact bulks \\cite{Randall:1999vf}. However, here again it is a non-trivial task to treat the backreaction of a more general cosmological brane source beyond the zero-mode approximation.  In five and six dimensions, exact and perturbative results have been obtained for such a source. In the simplest case, the cosmological evolution of a brane with tension $\\lambda$ in an infinite $D=5$ bulk leads to exactly calculable $O(\\rho^2/\\lambda^2)$ corrections to the Friedmann equation \\cite{Binetruy:1999hy}. The compact case is more subtle as one must also allow for the evolution of the radion field~\\cite{Binetruy:2001tc}. The aberrant Hubble rate, quadratic in $\\rho$ for $D=5$ flat compact bulks \\cite{Binetruy:1999ut,Chung:1999zs} for instance, has been shown \\cite{Csaki:1999mp} to result from the unnatural requirement of a static radion without introducing a stabilization mechanism, which leads to an overconstraint on the extra matter on a second brane. Stabilized models treated perturbatively in $\\rho$ again lead to a correct Friedmann equation, linear in $\\rho$ with quadratic corrections \\cite{Cline:2002ht}. We leave the review and discussion of more general five-dimensional configurations, including non-$\\mathbb{Z}_2$ symmetric bulks and ``dark radiation\" terms, for section \\ref{bulkspher}.  The problem is qualitatively different for a codimension two brane in $D=6$, a configuration reminiscent of the much studied cosmic string solutions in four dimensions. Indeed, for any brane source differing from a pure tension, the curvature diverges on the brane on which observers are supposed to live. This problem has been solved by considering regularized or ``fat\" branes with finite width. In the simple case of the compact ``football\" shaped bulk solution whose size is stabilized by a bulk magnetic flux \\cite{Carroll:2003db,Navarro:2003vw}, the Friedmann equation once again exhibits perturbative $\\rho^2$ corrections at high energies \\cite{Vinet:2004bk,Vinet:2005dg}. Backreaction effects in cosmological brane models for $D>6$ remain to be investigated.  Since it is expected that higher codimension infinitely thin branes will be singular, regularization is necessary, by an homogeneous spherical core for example. Alternatively, if the brane matter is taken to be localized at the boundary of this core (see figure \\ref{figure with non-symmetric bulk}), the codimension (D-4) brane is regularized as an infinitely thin codimension one brane shell with a (D-1)-dimensional worldvolume. The extra (D-5) spatial dimensions on the worldvolume in addition to the three observable ones may then be eliminated by compactification with a sufficiently small radius\\footnote{After completion of this work, reference \\cite{Cuadros-Melgar:2005ex} was pointed to our attention, following a similar approach for the particular $D=6$ case.}.  In this paper, we derive the effective cosmological equations for a non-$\\mathbb{Z}_2$ symmetric codimension one brane embedded in an arbitrary D-dimensional bulk spacetime. Our results generalize the $D=5,6$ cases which have been exhaustively studied previously by other authors, exactly in the case of infinitely thin branes for $D=5$, and perturbatively with regularized core branes for $D=6$. As we shall see, in the case of higher codimension, there is significantly more choice for the bulk geometry and hence, in the equivalent 4-dimensional effective theory, for the nature of the extra terms induced in the Friedmann equation.  The paper is organized as follows. In the next section we review the derivation of the Einstein equations induced on the brane in an arbitrary number of dimensions. In section~\\ref{bulkspher} we consider the case of a general unspecified bulk with spherical symmetry around the brane, including non-$\\mathbb{Z}_2$ symmetric contributions. We then derive a non-conservation law for the dark radiation specific to $D>5$ brane models, and we apply our findings to a specific non-compact orbifold bulk. In section~\\ref{adsbulk} we take the example of an AdS-Schwarzschild bulk in which spherical symmetry is broken by a toroidal compactification of the transverse space on the brane. In section~\\ref{conclusions} we provide a discussion of our results and conclude. ",
        "conclusions": "\\label{conclusions}  In this paper, we have derived the effective Einstein equations for the cosmological evolution of a codimension one brane embedded in a general D-dimensional bulk with spherical symmetry around the brane. The geometrical side of the effective Friedmann equation corresponds to the product of a 4d cosmological spacetime with an internal sphere with time-dependent radius. When the brane tension $\\lambda$ enjoys a full (D-2)-dimensional isotropy, the source side of the Friedmann equation contains the standard term linear in the brane energy plus quadratic corrections of order $O(\\rho^2/\\lambda^2)$, in addition to terms linear in the bulk energy-momentum. When it does not, the effective gravity does not have the correct tensor structure for the part linear in the brane source. We have included contributions to the effective source arising from the breaking of the $\\mathbb{Z}_2$ symmetry across the brane. We have shown that the bulk radiation manifests itself as an extra conformally invariant ``dark radiation\" source, which is in general not conserved, contrary to the $D=5$ case. By taking the example of an AdS-Schwarzschild bulk, we have illustrated the problem that any static bulk necessarily leads to an overconstraint on the brane matter equation of state.  These simple analytic examples provide a useful guide for the explicit construction of realistic cosmological metrics in the bulk in $D\\geq 7$. As an alternative approach to codimension one branes, it may be interesting to derive the effective Friedmann equations in a regularization-independent way for fat branes of codimension (D-4), as a perturbation expansion in the brane energy density integrated over the core. \\\\ \\ \\\\   {\\large \\bf Acknowledgments} \\\\  NC and MT are supported by the National Science Foundation under grant number PHY-0354990, by Research Corporation, and by funds provided by Syracuse University. NC and CM are supported by the Department of Energy, under grant number DE-FG02-85ER40231."
    },
    "0609/nucl-th0609006_arXiv.txt": {
        "abstract": "I review progress that has been made in nuclear astrophysics over the past few years and summarize some of the questions that remain.  Topics selected include solar neutrinos, supernovae (the explosion and associated nucleosynthesis), laboratory astrophysics, and neutron star structure. ",
        "introduction": "Perhaps no problem in nuclear astrophysics has drawn more attention than solar neutrinos and the associated discoveries of neutrino mass and mixing, the first experimental evidence for physics beyond the minimal standard model. These results grew out of careful solar modeling, patient laboratory measurements of pp-chain nuclear cross sections, and heroic efforts on new experiments.  The detectors that have come on line in the past decade and resolved the solar neutrino puzzle, SNO \\cite{SNO} and Super-Kamiokande \\cite{SuperK}, are remarkable instruments.  But much remains to be done.  Direct event-by-event measurements are so far limited to the high-energy $^8$B neutrinos, a 0.01\\% branch in the sun.  There are important reasons for extending such measurements to lower energies: \\\\ $\\bullet$ No phenomenon directly connected with matter-enhanced oscillations, such as spectral distortions or day-night effects due to the earth's matter, has been observed.   While indirect evidence of spectral distortions comes from the comparison of  GALLEX/GNO/SAGE \\cite{Gallex} integrated rates with SNO and Super-Kamiokande results, a direct measurement of changes in  $P_{\\nu_e}(E_\\nu)$ as a function of $E_\\nu$ is lacking. The parameters deduced from SNO, Super-Kamiokande, and earlier radiochemical experiments predict that low-energy pp and $^7$Be neutrinos experience only vacuum oscillations, as the matter contribution to the effective neutrino mass \\begin{equation} m_{\\nu_e}^{2~eff} = 4 E_\\nu \\sqrt{2} G_F \\rho_e(x) \\end{equation} is insufficient to cause a level crossing for low-energy $\\nu$s.  The survival probability for these sources is almost twice that for the high-energy portion of the $^8$B $\\nu_e$ flux. There is some chance that current experiments will be able to see variations in $P_{\\nu_e}$. SNO-III (use of $^3$He neutron counters) may be able to operate at a lower threshold due to improvements in the charge-current(CC)/neutral current(NC) event separation. KamLAND \\cite{KamLAND} will be attempting a $^7$Be measurement (though this is a difficult task at Kamioka depths), and Borexino \\cite{Borexino} may soon turn on at Gran Sasso.  \\\\ $\\bullet$ Measurement of the principal component of the solar flux, pp neutrinos, would isolate one of nature's best calibrated sources of $\\nu_e$s. In part because SNO and Super-Kamiokande measurements of $^8$B neutrinos have constrained the temperature of the solar core, the standard solar model (SSM) uncertainty in the pp flux is $\\lsim$ 1\\%.  For this reason pp experiments could improve the accuracy of mixing angle ($\\theta_{12}$) determinations.  New detectors, such as CLEAN \\cite{CLEAN} (NC scattering in liquid neon) and LENS \\cite{LENS} (CC scattering in In), are being developed with this goal in mind.\\\\ $\\bullet$  But perhaps the most intriguing solar neutrinos are those from H-burning through the CNO cycle, which contributes to solar energy generation at about the 1\\% level. The CNO cycle is the dominant mechanism for stellar energy generation in heavier main-sequence stars, including the very massive popIII stars that formed and produced \"first light\" (and which turned on the CNO cycle by first synthesizing carbon in the triple-alpha process) \\cite{firststars}.  Indeed, nuclear astrophysics arguably began with Bethe's search for the mechanism that could account for energy generation in massive main-sequence stars.  Because the CNO $\\nu$ flux is proportional to core metallicity Z, a measurement would test this important parameter. It is believed that the sun was highly convective as it entered the main sequence, and thus thoroughly mixed.  A plausible consequence -- that the zero-age solar core metallicity Z can be equated to today's surface abundance -- is a key postulate of the SSM \\cite{bahcallSSM}. SSM predictions are sensitive to Z because of the importance of free-bound atomic transitions to the solar opacity.  Furthermore, metals can have interesting evolutionary consequences by driving out-of-equilibrium burning of CNO elements in the early sun.  Such burning is thought to have powered an early convective solar core for $\\sim$ 10$^8$ years.  The sun's metallicity also controls the location of the radiative/convective boundary: a shift to lower Z makes radiative transport in the interior more efficient, and thus leads to a shallower convective zone.  Recently improved 3D models of the solar atmosphere have altered the interpretation of photospheric absorption lines, lowering deduced metal abundances by 30\\% \\cite{3Datmos}.  This has destroyed the previous agreement between SSM predictions and helioseismology: the acoustic modes are quite sensitive to the depth of the convective zone.  This is yet another reason for attempting an independent measurement of the sun's interior Z. ",
        "conclusions": "I believe at no time has the field of nuclear astrophysics been more exciting than today. The discoveries recently made in neutrino astrophysics have given us our first glimpses of physics beyond the standard model.  The neutrino parameters that have emerged from these discoveries are now inspiring a new generation of accelerator neutrino experiments to resolve questions about the neutrino mass hierarchy, $\\theta_{13}$, and leptonic CP violation.  The results will be important to some of the most interesting open questions in astrophysics: the origin of matter, the contribution of neutrinos to hot dark matter, the mechanism for core-collapse SN.  Laboratory astrophysics will remain important to this field.  LUNA's measurement of the controlling reaction of the CNO cycle underscores how uncertain our knowledge is of key nuclear reactions governing stellar evolution. The field's most important successes -- identification and resolution of the solar neutrino problem, the BBN prediction of $\\eta$ -- came from a combination of observation, theory, and laboratory measurements.  Today we need laboratory data not only for cross sections, but also to constrain the new phenomena associated with neutrino masses and mixing.  To understand nucleosynthesis in stars, laboratory measurements have to include both stable and unstable targets.  Large-scale computing has become the third leg of nuclear astrophysics, complementing theory and experiments/observations.   It now appears likely that the Type II SN mechanism is intrinsically multi-dimensional, and that successful models of the collapse will require the coupling of magneto-hydrodynamics, shock-wave propagation, neutrino transport of energy and lepton number, and rotation.   Success in building a ``standard model'' of SN could produce benefits comparable from those derived from the SSM. The supernova environment tests unique aspects of neutrino physics ($\\theta_{13}$, the MSW potential of neutrinos), nucleosynthesis (the explosive conditions necessary for the r-process, the neutrino fluences required for the $\\nu$-process), and gravity (gravitational waves tagged by the SN's optical signal).  Observations of millsecond pulsars like Terzan5 and the double-pulsar system PSRJ0737-3039 suggest that stringent constraints on the properties of dense, neutron-rich nuclear matter will soon be available.   NSs could provide crucial constraints on models of the nuclear EOS at high density.  This year is also one for remembering three colleagues who did so much to build this field, Hans Bethe, John Bahcall, and Vijay Pandharipande.   Almost every topic in my talk today has roots in the work of these pioneers.   Hans is truly the father of nuclear astrophysics, as it was his efforts to understand hydrogen burning as the mechanism for main-sequence stellar evolution that built the bridge from the stars to the laboratory.  Because of John's tireless efforts, the solar neutrino problem was finally recognized as a triumph for quantitative nuclear astrophysics and as the first low-energy manifestation of the physics that lies beyond the standard model.   Vijay pioneered techniques to derive the properties of nuclei and nuclear matter directly from NN and NNN forces, including properties of nature's largest nucleus, the neutron star.   We will miss them all.  This work was supported in part by the Office of Nuclear Physics, U.S. Department of Energy, and by the DOE's SCIDAC program."
    },
    "0609/math-ph0609038_arXiv.txt": {
        "abstract": "This paper is related to our previous works \\cite{uno} \\cite{due} on the error estimate of the averaging technique, for systems with one fast angular variable. In the cited references, a general method (of mixed analytical and numerical type) has been introduced to obtain precise, fully quantitative estimates on the averaging error. Here, this procedure is applied to the motion of a satellite in a polar orbit around an oblate planet, retaining only the $J_2$ term in the multipole expansion of the gravitational potential. To exemplify the method, the averaging errors are estimated for the data corresponding to two Earth satellites; for a very large number of orbits, computation of our estimators is much less expensive than the direct numerical solution of the equations of motion. ",
        "introduction": "\\label{intro} Celestial mechanics often requires approximation techniques to compute the motion of bodies over very long times. Giving reliable error estimates on these techniques is a very practical problem, of not simple solution. \\par In this paper we apply to a typical astronomical problem the general scheme of \\cite{uno} \\cite{due} to estimate the error of the averaging principle, in the case of one fast angular variable $\\te$ (in the one-dimensional torus) and many slow variables $I = (I^i)$ (the actions). The application we consider refers to the motion of a satellite around an oblate planet: of course, oblateness gives a perturbation of the Keplerian, purely radial gravitational potential. There is a classical way to express the perturbed potential as a series of spherical harmonics, where the multipole moments of the planetary mass distribution appear as coefficients. In the case of axial symmetry, this series only involves zonal harmonics and the $\\ell$-th term contains a multipole coefficient $J_{\\ell}$ (for a reminder, see subsection \\sate \\, and references therein). For a slightly oblate planet, all these coefficients are small if they are defined in a convenient, dimensionless fashion; often, $J_1 \\simeq 0$ and one retains only the $\\ell=2$ term of the previous expansion; the corresponding equations of motion for the satellite are the classical \"$J_2$ problem\". \\par In this paper the $J_2$ problem is considered in the restricted case of polar orbits; this problem involves a fast angular variable $\\Te$ (the angle between the polar axis and the radius vector of the satellite) and three slow variables $\\I = (\\P,\\E,\\Y)$ representing the parameter, the eccentricity and the argument of the pericenter for the osculating Keplerian ellipse. In principle, all these variables are unknown functions of the \"physical\" time $t$; however, it is customary to consider as a ``time'' variable the angle itself or, rather, the ratio $\\ti$ of the angle to $2 \\pi$. With this choice of the time variable, by definition the angle evolves with the law $\\Te(\\ti) = [2 \\pi \\ti]$ (where $[~]$ indicates equivalence mod.$2 \\pi$); for obvious reasons, $\\ti$ will be called the \\textsl{orbit counter}. The evolution of $\\I = (\\P,\\E,\\Y)$ is described by a set of equations of the form \\beq {d \\I^i \\over d \\ti} = \\ep f^i(\\I,[2 \\pi \\ti])~, \\qquad \\I^i(0) = I^i_0~, \\label{per} \\feq where $\\ep$ is a small parameter proportional to $J_2$ (see again subsection \\sate). Independently of this specific application, a system of evolution equations like \\rref{per}, with any number of unknown functions $\\I = (\\I^i)_{i=1,...,d}$ and a small parameter $\\ep > 0$, will be called in the sequel a \\textsl{perturbed periodic system}. To any such system one can associate an averaged system, replacing the functions $f^i$ in \\rref{per} with their averages over the angle $\\Te = [2 \\pi \\ti]$; the solution $\\J = (\\J^i)$ of the latter is in a function of $\\tau := \\ep \\ti$. \\par The difference between the solutions $\\I$ of \\rref{per} and $\\J$ of the averaged system, on a time scale of order $1/\\ep$, can be estimated in a fully quantitative way with the general method of \\cite{uno} \\cite{due}; this is what we are going to do in the present work, for the periodic system describing the polar $J_2$ problem. As in the cited papers, the expression \"fully quantitative\" means that our final estimate will have the form \\beq | \\I^i(\\ti) - \\J^i(\\ep \\ti) | \\leq \\ep \\en^i(\\ep \\ti)~\\qquad \\mbox{for $\\ti \\in [0,U/\\ep)$}~, \\label{formm} \\feq where $\\en^i : [0,U) \\vain [0,+\\infty)$ are computable functions, and $U$ is a constant, whose choice determines quantitatively the interval of observation; we note that $U/\\ep$ is the total number of orbits. \\par The availability of an algorithm to construct the estimators $\\en^i$ is the main difference between the approach of \\cite{due} and the classical, qualitative result $\\I^i(\\ti) - \\J^i(\\ep \\ti) = O(\\ep)$, see, e.g., \\cite{Arn} \\cite{Loc} \\cite{Ver}. We are aware of the existence of higher order versions of the averaging method, in which the error behaves like $O(\\ep^p)$ for $t \\in [0, O(1/\\ep))$ ($p=2,3,...$; see again \\cite{Loc} \\cite{Ver}). However, the classical theory of higher order averaging gives little more than the qualitative error estimate $O(\\ep^p)$; perhaps an extension of our methods could give quantitative estimators of the form $\\ep^p \\, \\en^i_{(p)}(\\ep t)$, but this problem is left to future work. \\par Let us return to the functions $\\en^i$ in Eq. \\rref{formm}. Our method to compute them for the $J_2$ problem is mainly analytical, but, in its final steps, it requires the numerical solution of an ODE on the interval $[0,U)$; however, this numerical computation is much faster than the direct numerical solution of \\rref{per} (a fact already appearing in other applications, different from the $J_2$ problem, considered in our previous papers). \\par The general treatment proposed in this paper for the polar $J_2$ problem is subsequently exemplified, choosing for $\\ep$ and for the initial conditions $I_0 = (P_0, E_0, Y_0)$ the values corresponding to the Earth and to the Polar and Cos-B satellites. In this case, the estimators $\\en^i$ have been computed up to $60000$ orbits (i.e., one or two centuries), spending few seconds of CPU time on a PC. To test their reliability, we have computed by direct numerical integration the differences $\\I^i(\\ti) - \\J^i(\\ep \\ti)$; this is a more expensive operation, that we have been able to perform up to $3000$ orbits with the same PC; our  estimators are quite satisfactory on this interval, since the functions $\\ti \\mapsto \\ep \\en^i(\\ep \\ti)$ are close to the envelopes of the rapidly oscillating functions $| \\I^i(\\ti) - \\J^i(\\ep \\ti) |$. \\par Of course, the treatment of real satellites of the Earth should include, besides the $J_2$ gravitational term, other minor perturbations such as: gravitational forces corresponding to higher order moments of the Earth, atmospheric dragging, solar wind, tidal effects due to the Moon gravity. All these effects could be treated, giving rise again to a system of the form \\rref{per} with more complicated perturbation components $\\ep f^i$ (even for non polar orbits: in general, if the analysis is not restricted to the orbits in a fixed plane, the actions are not three but five). Presumably, the corresponding averaged system and our estimators $\\en^i$ for its error could be computed as well; this is left to future work. In spite of the limitation to the polar $J_2$ problem, we think that the results of this paper have some interest, because they show that a general method for error estimates works over very long times on a non-trivial problem. \\par To conclude, we describe in few words the organization of the paper. In Section \\ref{secav}, we show how to compute the error estimators for the averaging of any periodic system \\rref{per}; this illustration specializes to the periodic case the slightly more general setting of \\cite{uno} \\cite{due} for one-frequency systems. In Section \\ref{satell}, the basic facts on the motion of a satellite in a fixed plane $\\Pi$ are reviewed; the outcomes are the equations of motion for $(\\I^i) = (\\P,\\E,\\Y)$, for any perturbation of the Keplerian potential keeping the satellite on $\\Pi$. In Section \\ref{polmot}, we specialize the previous equations to the polar $J_2$ problem; the averaged system is solved, and we construct the algorithm to compute the error estimators $\\en^i$ $(i=\\p,\\e,\\y)$. In Section \\ref{seces}, we perform the computations for the Polar and Cos-B satellites of the Earth. Some details on the constructions of Sections \\ref{satell}-\\ref{polmot} are presented in Appendices \\ref{appkepl}-\\ref{appinv}. \\salto ",
        "conclusions": ""
    },
    "0609/astro-ph0609638_arXiv.txt": {
        "abstract": "The bright K1\\,III--IV star \\gamcep\\ has been reported previously to have a companion in a $\\sim$2.5-yr orbit that is possibly substellar, and also has a stellar companion at a larger separation that has never been seen. Here we determine for the first time the three-dimensional orbit of the stellar companion accounting also for the perturbation from the closer object. We combine new and existing radial velocity measurements (of both classical precision and high precision) with intermediate astrometric data from the {\\it Hipparcos\\/} mission (abscissa residuals) as well as ground-based positional observations going back more than a century. The orbit of the secondary star is eccentric ($e = 0.4085 \\pm 0.0065$) and has a period $P = 66.8 \\pm 1.4$~yr and a semimajor axis of $19.02 \\pm 0.64$~AU. We establish the primary star to be on the first ascent of the giant branch, and to have a mass of $1.18 \\pm 0.11$~M$_{\\sun}$, an effective temperature of $4800 \\pm 100$~K, and an age around 6.6~Gyr (for an assumed metallicity [Fe/H] = $+0.01 \\pm 0.05$). The unseen secondary star is found to be an M4 dwarf with a mass of $0.362 \\pm 0.022$~M$_{\\sun}$, and is expected to be $\\sim$8.4~mag fainter than the primary in $V$ and $\\sim$6.4~mag fainter in $K$. The minimum mass of the putative planetary companion is $M_p \\sin i = 1.43 \\pm 0.13$~M$_{\\rm Jup}$, the inclination angle of its orbit being unknown.  Taking advantage again of the high-precision {\\it Hipparcos\\/} observations we are able to place a dynamical upper limit on this mass of 13.3~M$_{\\rm Jup}$ at the 95\\% confidence level, and 16.9~M$_{\\rm Jup}$ at the 99.73\\% (3$\\sigma$) confidence level, thus confirming that it is indeed substellar in nature. The orbit of this object (semimajor axis $1.94 \\pm 0.06$~AU) is only 9.8 times smaller than the orbit of the secondary star (the smallest ratio among exoplanet host stars in multiple systems), but it is stable if coplanar with the binary. ",
        "introduction": "\\label{sec:introduction}  The bright evolved star $\\gamma$~Cephei ($V = 3.21$, SpT = K1~III--IV, $\\alpha = 23^{\\rm h} 39^{\\rm m} 21\\fs01$, $\\delta = +77\\arcdeg 37\\arcmin 55\\farcs2$, J2000; also known as HD~222404, HR~8974, HIP~116727) is among the first objects to be subjected to high-precision radial-velocity measurements in an effort to discover substellar-mass companions around nearby stars \\citep{Campbell:88}. This group of investigators \\citep{Campbell:79} pioneered the use of a hydrogen fluoride gas absorption cell on the Canada-France-Hawaii telescope (CFHT) and achieved internal errors around 13~\\ms\\ for bright stars, inaugurating the era of Doppler searches that has been so successful in finding extrasolar planets in the last 10 years.  Small radial-velocity variations in \\gamcep\\ were indeed seen by \\cite{Campbell:88} suggesting the presence of a Jupiter-mass object in a $\\sim$2.5-yr orbit. Those variations with a semi-amplitude of only about 25~\\ms\\ were superimposed on a much larger variation caused by a previously unnoticed stellar companion with a period of decades. However, the interpretation of the residual 2.5-yr variation as due to a planetary object was subsequently put in doubt by the same group \\citep[see][]{Irwin:89, Walker:89, Walker:92}. They argued that changes with a similar period were observed in a chromospheric activity indicator in \\gamcep\\ (the \\ion{Ca}{2} $\\lambda$8662 emission line index), and thus that the velocity variations were spurious and probably due only to changes in the spectral line profiles caused by surface inhomogeneities (spots) driven by stellar rotation. More recently the planetary interpretation was reinstated by \\cite{Hatzes:03} on the basis of new high-precision velocity observations at the McDonald Observatory.  They showed convincingly that the 2.5-yr variation is coherent in phase and amplitude throughout the entire 20-yr interval covered by the merged CFHT and McDonald data sets, as would be expected for Keplerian motion, and that no changes were observed in the spectral line bisectors. On the other hand, a careful re-analysis of the changes in the activity indicator reported by \\cite{Walker:92} revealed that the periodicity of the \\ion{Ca}{2} $\\lambda$8662 measurements (2.14 yr) is not only slightly different from that in the velocities, but it is transitory in nature, thus ruling out a connection.  \\gamcep\\ also carries the distinction of being among the first planet host stars to be found in a binary system, which raises interesting issues related to the dynamical stability of such configurations. A recent study by \\cite{Raghavan:06} points out that among the known planet host stars \\gamcep\\ happens to be the system with the smallest ratio ($\\sim$11) between the size of the binary orbit and the planetary orbit. However, the outer orbit is at present poorly known, and the secondary star is presumably very faint and has never been seen. Reported values for the binary period have ranged between 29.9~yr \\citep{Walker:92} and 66~yr \\citep{Griffin:02}, and have been based on only part of the data available, in some cases spanning much less than a full cycle. A number of authors have carried out numerical investigations of the gravitational influence of the secondary star on the orbit of the planet \\citep[e.g.,][]{Dvorak:03, Thebault:04, Haghighipour:06}, but have used rather different parameters for the binary or have pointed out the uncertainty in those elements as a limiting factor.  The motivation for this paper is thus threefold: $i$) To improve the determination of the orbit of the secondary star (including for the first time an estimate of the inclination angle, and of the mass of the secondary) in order to allow more definitive dynamical studies of the stability and evolution of the system. We do this by using all available radial velocity data for \\gamcep\\ including new measurements reported here and other historical observations not previously used. We incorporate also astrometric measurements from the {\\it Hipparcos\\/} mission \\citep[``abscissa residuals'';][]{ESA:97} as well as transit circle and other positional information spanning more than a century. $ii$) To carry out a critical review of previous studies of the physical properties of the primary star and use all available information to estimate its absolute mass, a key parameter influencing the mass of the substellar companion. $iii$) To place firm dynamical upper limits on the mass of this companion by taking advantage of the high-precision {\\it Hipparcos\\/} intermediate data and modeling the reflex motion of the primary star on the plane of the sky. We show that this modeling allows us to confirm the substellar nature of the companion, although it is not yet possible to rule out a mass in the brown dwarf regime. ",
        "conclusions": "\\label{sec:discussion}  In their paper \\cite{Hatzes:03} attempted to place limits on the mass of the substellar companion in a different way by using their measured projected rotational velocity for \\gamcep\\ ($v \\sin i = 1.5 \\pm 1.0$~\\kms) along with the period they determined for the variation of the \\ion{Ca}{2}~$\\lambda$8662 emission-line index ($781 \\pm 116$~days) and the estimated radius of the star \\citep[$R=4.66$~R$_{\\sun}$, adopted from][]{Fuhrmann:04}. They relied on two assumptions: that the spin axis of the star is parallel to the axis of the planetary orbit, and that the period of variation of the \\ion{Ca}{2} index represents the true rotation period of the star. The comparison between the measured $v \\sin i$ and the expected equatorial rotational velocity ($v_{\\rm eq}$) then gives limits on $\\sin i$ (or $\\sin i_{\\rm A}$ in our notation).  As it turns out, however, there is a mathematical error in their calculation of $v_{\\rm eq}$: they reported $v_{\\rm eq} = 4.9$~\\kms, while the correct value is 0.3~\\kms\\ \\citep[see also][]{Walker:92}. Since this is \\emph{smaller} than their $v \\sin i$, no limits can be placed on $\\sin i_{\\rm A}$ in this way. Their statement on the probable mass range of the planetary companion is therefore not valid.  Either the measured $v \\sin i$ is overestimated, or the period of the \\ion{Ca}{2} variations is not the true rotation period of the star. The former explanation is perhaps supported by a measurement by \\cite{Gray:85}, who gave $v \\sin i = 0.0 \\pm 0.8$~\\kms\\ (along with a sizeable radial-tangential macroturbulence of $\\zeta_{\\rm RT} = 4.2 \\pm 0.6$~\\kms, which can effect rotational velocity measurements if not properly accounted for). The study by \\cite{deMedeiros:99} reported $v \\sin i < 1.0$~\\kms.  Alternatively, the rotation period would have to be considerably shorter than 781 days ($\\approx$100--500 days), and another explanation would have to be found for the variations in the emission-line index.  Our dynamical constraint on $M_p$ thus shows for the first time that the companion is substellar in nature, although a mass in the brown dwarf regime (as opposed to the planetary regime) cannot be completely ruled out with the present observations.  \\gamcep\\ is one of more than two dozen examples of substellar companions found in stellar binaries \\citep[see, e.g.,][]{Raghavan:06}. Such systems have attracted considerable interest in recent years, and numerical studies have been carried out specifically for the case of \\gamcep\\ to assess not only the dynamical stability of the orbit of the substellar companion \\citep[e.g.,][]{Dvorak:03, Solovaya:04, Haghighipour:06}, but also the stability of the orbits of other (possibly Earth-like) planets that might be present in the habitable zone of the primary star. \\cite{Thebault:04} have also investigated the conditions under which the substellar companion may form through core accretion in the binary environment.  With a relative semimajor axis for the planet orbit of $a_{{\\rm Aa}-p} = 1.94 \\pm 0.06$~AU (for an adopted primary mass $M_{\\rm Aa} = 1.18 \\pm 0.11$~M$_{\\sun}$), the size of that orbit is only 9.8 times smaller than the size of the binary orbit (19.02~AU; see Table~\\ref{tab:elements}), currently the lowest value among the known exoplanets in binaries\\footnote{The slightly smaller orbit size ratio compared to the value of $\\sim$11 given by \\cite{Raghavan:06} is largely due to the significant improvement in the elements of the binary orbit in the present work.}. Orbit stability depends quite strongly on the parameters of the binary system, in particular the semimajor axis and eccentricity, as well as on the masses of the components. The dynamical studies mentioned above have all had to make do with the rather poorly determined binary properties and also often inconsistent results from various authors.  \\cite{Holman:99} have derived a simple empirical formula for computing the maximum value of the semimajor axis of a stable planetary orbit (``critical'' semimajor axis, $a_{\\rm crit}$) in a coplanar S-type planet-binary system.  \\cite{Haghighipour:06} pointed out in his study that the uncertainty in the binary orbital elements made for a very large parameter space to be explored numerically for \\gamcep. Furthermore, the inclination of the binary orbit was unknown at the time and therefore so was the mass of the secondary star. As a result, he was only able to provide a rather wide range of critical semimajor axes as a function of the adopted binary eccentricity (see his Figure~1). With the present study that situation has changed, and the critical semimajor axis can now be computed directly with a relatively small formal uncertainty. We obtain $a_{\\rm crit} = 3.61 \\pm 0.36$~AU, which is considerably larger than the semimajor axis of the planet orbit, implying the latter is stable if coplanar with the binary.  The combination of classical as well as high-precision radial velocity measurements of \\gamcep\\ with ground- and space-based astrometry have allowed a significant improvement in the binary orbital elements (and a first determination of the inclination angle) as well as a better knowledge of the stellar masses. Nevertheless, the secondary star remains unseen. Even though the predicted angular separation of \\gamcep~B (0\\farcs84 for 2007.0; 0\\farcs99 for 2009.0) is not particularly challenging, the 8-magnitude brightness difference in the visual band relative to the glaringly bright primary explains all negative results (e.g., the speckle interferometry attempts by Mason et al.\\ 2001, as well as the imaging by Hatzes et al.\\ 2003). We expect the contrast to be much more favorable in the near infrared ($\\Delta m \\sim 6.4$ in $K$), and that this detection should not be very difficult at those wavelengths with adaptive optics on a large telescope. Such measurements of the relative position would allow a dynamical determination of the mass of both stars, free from assumptions."
    },
    "0609/astro-ph0609312_arXiv.txt": {
        "abstract": "{ High spectral resolution studies of cool Ap stars reveal conspicuous anomalies of the shape and  strength of many absorption lines. This is a signature of large atmospheric chemical gradients (chemical stratification) produced by the selective radiative levitation and gravitational settling of chemical species. } { Previous observational studies of the chemical stratification in Ap stars were limited to fitting simple parametrized chemical profiles. Here we present a new approach to mapping the vertical chemical structures in stellar atmospheres. } { We have developed a regularized chemical inversion procedure that uses all information available in high-resolution stellar spectra. The new technique for the first time allowed us to recover chemical profiles without making \\textit{a priori} assumptions about the shape of chemical distributions. We have derived average abundances and applied the vertical inversion procedure to the high-resolution VLT UVES spectra of the weakly magnetic, cool Ap star \\hd. } { Our spectroscopic analysis yielded improved estimates of the atmospheric parameters of \\hd. We show that this star has negligible \\vs\\ and the mean magnetic field modulus \\bs\\,=\\,$1.1\\pm0.1$~kG. We have derived average abundances for 43 ions and obtained vertical distributions of Ca, Si, Mg, Fe, Cr, and Sr. All these elements except Mg show high overabundance in the deep layers and solar or sub-solar composition in the upper atmosphere of \\hd. In contrast, the Mg abundance increases with height. } { We find that transition from the metal-enhanced to metal-depleted zones typically occurs in a rather narrow range of depths in the atmosphere of \\hd. Based on the derived photospheric abundances, we conclude that \\hd\\ belongs to the rare group of evolved cool Ap stars, which possesses very large Fe-peak enhancement, but lacks a prominent overabundance of the rare-earth elements. } ",
        "introduction": "\\label{intro}  Michaud (\\cite{M70}) considered in detail a process of element separation in stellar atmospheres under the mutual action of radiative acceleration and gravitational settling. He showed that in atmospheres with suppressed convection and turbulence this process could lead to chemical anomalies in the form of under- or overabundances. This \\textit{radiative diffusion} hypothesis was first applied to interpret the abundance anomalies inferred for the atmospheres of magnetic chemically peculiar Ap stars (Michaud et al. \\cite{MRC74}), which were known to be slow rotators and whose atmospheres are stabilized by the strong magnetic field.  Diffusion processes are not only responsible for the observed average atmospheric abundance anomalies but also lead to an inhomogeneous abundance distribution through the stellar atmosphere depending on the balance between radiative acceleration and surface gravity. Borsenberger et al. (\\cite{BPM81}) calculated Ca and Sr abundance profiles for atmospheres with \\te$\\ge$10000~K and demonstrated the influence of chemical stratification on the profile of the resonance  \\ion{Ca}{ii} 3933~\\AA\\ line. These interesting theoretical predictions notwithstanding, the absence of high-resolution, high signal-to-noise spectroscopic observations at that time did not permit a direct comparison between observations and diffusion calculations. Later Babel (\\cite{Babel92}) carried out detailed diffusion calculations for Ca, Ti, Cr, Mn, and Sr in \\te=8500~K stellar atmosphere and applied his model to explain an unusual shape of the \\ion{Ca}{ii} 3933~\\AA\\ line, which often shows extremely wide wings and a very sharp narrow core in the spectra of Ap stars. Babel's calculations inspired the use of a simple approximation of the vertical element distribution in the stellar atmosphere in the form of a step function. This parametrization is now widely employed in observational stratification studies, for example Mn in HgMn stars (Sigut \\cite{sigut01}), Si, Ca, Cr, Fe in Ap stars (Wade et al. \\cite{WLRK03}; Ryabchikova et al. \\cite{RPK02,RLK05}).  Despite the encouraging success of the previous attempts to model chemical gradients using the prescribed parametrized shape, real element distributions in the atmospheres of peculiar stars may differ considerably from the simple step function, at least for some chemical elements. Some theoretical studies also hinted at the possibility of complex vertical abundance distributions (e.g. fig.~6 in Borsenberger et al. \\cite{BPM81}). Moreover, recent self-consistent model atmosphere calculations including diffusion (LeBlanc \\& Monin \\cite{LM04}) have identified a number of additional effects, like mass loss and weak turbulent mixing, which may significantly affect the shape of the vertical chemical profiles but are difficult to predict \\textit{ab initio}. These interesting hydrodynamical effects, treated as free parameters in the current theoretical diffusion modelling, could be potentially constrained by comparison with the observed vertical abundance distributions.  Meanwhile modern spectrographs at large telescopes, such UVES at the ESO VLT, have reached remarkable precision and spectral coverage, and now allow us to record very high signal-to-noise ratio data covering the whole optical region spectra of moderately bright Ap stars. This dramatic improvement in the quality and quantity of the available observational material suggests that unprecedented details of  the stellar atmospheric properties, and, in particular, the vertical abundance distributions, can be deduced with the help of detailed modelling of the spectral line profiles. These facts stimulated a recent surge of interest in the observational analysis of chemical stratification and emphasized the need to improve stratification modelling techniques.  In this paper we present a new procedure to study vertical chemical inhomogeneities in Ap-star atmospheres. We have chosen \\hd\\ (HR~5623, HIP~74181) for our modelling. This star was classified as an A0p Sr-Cr-Eu object by Jaschek \\& Jaschek (\\cite{JJ59}). Mathys (\\cite{M90}) noted that \\hd\\ exhibits remarkably sharp lines and no evidence of a strong magnetic field. Martinez \\& Kurtz (\\cite{MK94}) detected no rapid pulsational variability in \\hd. Despite its brightness and extremely low \\vs, facilitating detailed spectroscopic studies, the star was never analysed with the model atmosphere method.  The paper is structured as follows. The principles and numerical details of the vertical inversion procedure are given in Sect.~\\ref{vip}. Observations of \\hd\\ and spectra reduction are described in Sect.~\\ref{observ}. Fundamental stellar parameters are determined in Sect.~\\ref{params} and the abundance analysis is presented in Sect.~\\ref{abund}. We derive the vertical distribution of chemical elements in Sect.~\\ref{strat}. Sect.~\\ref{discus} concludes the paper with the summary and discussion.  \\begin{figure}[!th] \\figps{5607f1.eps} \\caption{Depth variation of the iron abundance and opacities for a chemically homogeneous (\\textit{left column}) and stratified (\\textit{right column}) stellar atmosphere. {\\bf a)} Fe abundance as a function of the optical depth in continuum at $\\lambda$ 4924~\\AA; {\\bf b)} depth dependence of the line centre (\\textit{filled symbols}) and continuous (\\textit{open symbols}) opacities. The bottom panels ({\\bf c}) show the maximum error in the normalized intensity profile of the \\ion{Fe}{ii} 4923.93~\\AA\\ line computed with different radiative transfer algorithms and different number of layers in the model atmosphere.} \\label{fig1} \\end{figure}  \\begin{figure*}[!th] \\figps{5607f2.eps} \\caption{Normalized disk-centre intensity profile of the \\ion{Fe}{ii} 4923.93~\\AA\\ line computed for the stratified iron abundance distribution shown in Fig.~\\ref{fig1} using different radiative transfer algorithms: {\\bf a)} DELOQ, {\\bf b)} DELOB, {\\bf c)} Feautrier, {\\bf d)} MW, {\\bf e)} IS. In each plot the lower panel shows the line profiles computed for the model atmosphere with 49 (\\textit{dotted line}), 118 (\\textit{dashed line}), and 332 (\\textit{solid line}) layers. The upper panels show the difference with respect to the computation for the model atmosphere discretized into 628 layers.} \\label{fig2} \\end{figure*} ",
        "conclusions": "\\label{discus}  Thanks to recent advances in the instrumentation at large telescopes, high-resolution, very high $S/N$ spectra have been obtained for a large number of Ap stars. This dramatic improvement of the data quality stimulated a renewal of interest in the atmospheric properties of magnetic CP stars. Modern investigations of the spectral line intensities and shapes show that many peculiar features in the Ap-star spectra originate from an inhomogeneous vertical distribution of chemical elements in stellar atmospheres. Thus, magnetic chemically peculiar stars represent the only type of non-degenerate stellar objects where direct observations and diagnosis of the chemical diffusion signatures becomes possible. Observational study of the chemical stratification in Ap stars is able to provide crucial and unique constraints for the theoretical modelling of chemical diffusion in stellar atmospheres and also helps to understand the relation between diffusion and hydrodynamic mixing processes.  Until now observational analyses of stratification in peculiar stars have been limited to fitting simple parametrized vertical chemical profiles to a small number of diagnostic lines (Wade et al. \\cite{WLRK03}; Ryabchikova et al. \\cite{RPK02,RLK05}). Most often, chemical distributions were approximated with a step function, whose shape is described by two abundance values and the position of the transition zone.  In the present study we have developed and applied the first assumption-free method of reconstruction of the chemical stratification in stellar atmospheres. In our technique, the individual chemical composition is derived for all atmospheric layers contributing to the line absorption. Uniqueness and stability of the chemical inversion is achieved by applying the Tikhonov regularization function. This means that our code finds the simplest elemental distribution sufficient to fit observations.  The vertical inversion procedure is successfully applied to the weakly magnetic, evolved Ap star \\hd. We have reconstructed vertical distributions of 6 elements. Magnesium is found to have an abundance close to the solar one in the lower atmosphere and shows an increase of concentration with height. All other elements have the opposite vertical distribution: high abundance deep in the photosphere and a lower abundance in the upper layers. The transition region is narrow for Si, Ca and Sr, and is located close to $\\log\\tau_{\\rm 5000}=0$. Fe and Cr are overabundant over the whole atmosphere of \\hd, and the respective change of the concentration of these elements is smaller than for other species and occurs between $\\log\\tau_{\\rm 5000}=-1$ and $-2$.  Our investigation of the diffusion signatures in \\hd\\ confirms the validity of the step-like, parametrized stratification models applied in previous studies of Ap stars. However, we find that for some species the transition region can be rather extended (Mg), or the vertical distribution is more complex than a simple one-step function (Sr). Thus, the applicability of the latter approximation is wide, but not universal. In this context, our automatic regularized inversion approach appears to be more robust and should be preferred whenever possible.  Stratification analysis of \\hd\\ is supplemented by a detailed chemical abundance study assuming homogeneous vertical distribution of elements. We have measured the concentration of 43 ions of 32 chemical elements. The outstanding characteristic of the chemical composition of \\hd\\ is a very large overabundance of the iron-peak elements in its atmosphere, combined with a fairly small enhancement of heavy elements, in particular REE. Strontium and, possibly, palladium are the only heavy elements in \\hd\\ showing a very large overabundance with respect to the solar chemical composition.  \\hd\\ is the second star with a detailed abundance analysis, including stratification effects, whose position on the H-R diagram clearly shows that the star is close to finishing its main sequence life. The first such star was HD~204411 (Ryabchikova et al. \\cite{RLK05}). Both stars possess weak magnetic fields and  have a small overabundance of the REEs compared with many other Ap stars in the 8500--9500~K effective temperature region. The most remarkable abundance characteristic of both stars is an extremely high overabundance of Cr and Fe, in particular in HD~133792.  Cowley \\& Henry (\\cite{CH79}) have defined a group of 5 stars characterized by strong lines of the Fe-peak elements and REE lines weaker than in ``normal'' Ap stars. HD~204411 is a member of this group. The H-R diagram position of these objects indicates that all of them are close to the end of their main sequence life (Kochukhov \\& Bagnulo \\cite{KB06}).  We were able to extract high-resolution spectra from the ELODIE  archive\\footnote{{\\tt http://atlas.obs-hp.fr/elodie/}} for the three other stars of this group: 73 Dra (HD~196502), HD~8441, and HD~216533. Photometric indicators suggest \\te\\,=\\,8900--9200~K for these three objects. A preliminary determination of the Pr and Nd abundance based on the strongest lines of the  second ions shows that in all three stars abundances of these rare-earth elements are comparable to those found in \\hd. In contrast, REE abundances in the young star HD~66318 of similar effective temperature, much stronger magnetic field (Bagnulo et al. \\cite{bagn03}) and the same large Cr and Fe anomalies are  by an order of magnitude higher. For all stars these abundance estimates were based on the same lines of \\ion{Pr}{iii} and \\ion{Nd}{iii}.  Evidently, on the basis of our detailed abundance analysis we may extend the definition of the group of stars with weak REE lines as \\textit{a group of evolved Ap stars with small to moderate REE anomalies, and very large Cr and Fe overabundances}. All stars of this group are lie near maximum in plots of Cr and Fe abundance versus \\te\\ (Ryabchikova et al. \\cite{abun04}). Comparison between Cr and Fe abundances in evolved Ap stars with weak magnetic fields and in the young star HD~66318 with an extremely large magnetic field shows that neither evolutionary status nor magnetic field play an important role in creating the Fe-peak abundance anomalies. On the other hand, the REE overabundance is clearly less pronounced in the group of evolved Ap stars. Further detailed chemical and evolutionary analysis of Ap stars in the effective temperature range 8500--10000~K is required to understand the role of the magnetic field and/or evolution in creating and maintaining REE atmospheric anomalies in magnetic peculiar stars."
    },
    "0609/astro-ph0609124_arXiv.txt": {
        "abstract": "Galaxy cluster surveys will be a powerful probe of dark energy. At the same time, cluster abundance is sensitive to any non-Gaussianity of the primordial density field. It is therefore possible that non-Gaussian initial conditions might be misinterpreted as a sign of dark energy or at least degrade the expected constraints on dark energy parameters. To address this issue, we perform a likelihood analysis of an ideal cluster survey similar in size and depth to the upcoming South Pole Telescope/Dark Energy Survey (SPT-DES). We analyze a model in which the strength of the non-Gaussianity is parameterized by the constant $\\fNL$; this model has been used extensively to derive Cosmic Microwave Background (CMB) anisotropy constraints on non-Gaussianity, allowing us to make contact with those works. We find that the constraining power of the cluster survey on dark energy observables is not significantly diminished by non-Gaussianity provided that cluster redshift information is included in the analysis.  We also find that even an ideal cluster survey is unlikely to improve significantly current and future CMB constraints on non-Gaussianity. However, when all systematics are under control, it could constitute a valuable cross check to CMB observations. ",
        "introduction": "Of the many fascinating discoveries in cosmology over the last decade, perhaps none have aroused more interest than the discovery of the accelerated expansion of the Universe \\citep{Riess:1998cb, Perlmutter:1998np}. Probing the nature of the dark energy thought to be driving this acceleration has become a top priority for the community, and among the promising tools under consideration are surveys of galaxy clusters. Since the number of clusters as a function of redshift and mass depends on both the growth of structure and on the volume of space, the cluster abundance is sensitive to the matter density, the density fluctuation amplitude, and the expansion history of the Universe. For this reason, upcoming cluster surveys will be powerful probes of cosmology \\citep[e.g.][]{Haiman:2000bw, Holder:2001db,Battye:2003bm,Molnar:2003eq,Wang:2004pk,Rapetti:2004aa,Marian:2006zp}.  Although constraining dark energy is a leading motivator for much of the interest in cluster surveys, it is worth noting that the cluster abundance is potentially sensitive to various cosmological parameters beyond those associated with dark energy. For example, it has been recognized for some time that slight deviations from Gaussianity in the primordial matter distribution would cause a significant change in the high mass tail of the halo distribution \\citep{Lucchin:1987yv, Colafrancesco:1989px,Chiu:1997xb,Robinson:1998dx,Robinson:1999se,Koyama:1999fc,Robinson:1999wh, Matarrese:2000iz}. In this paper, we use a maximum likelihood analysis to investigate the extent to which dark energy constraints from cluster surveys are degraded by including the possibility of non-Gaussian initial conditions, in particular when considered within the limits allowed by present and future CMB observations.  The specific form of non-Gaussian initial conditions we consider here is of the local type, described in position space by a primordial curvature perturbation of the form \\citep{Verde:1999ij,Komatsu:2001rj} \\beq\\label{fNLmodel} \\Phi(\\xv)=\\phi(\\xv)+\\fNL[\\phi^2(\\xv)-\\lan\\phi^2(\\xv)\\ran] \\eeq where $\\phi(\\xv)$ is a Gaussian random field and the degree of non-Gaussianity is parameterized in terms of the constant $\\fNL$. For this model, tight constraints of the order of $\\dfNL\\sim 40$ are provided by CMB observations \\citep{Komatsu:2003fd,Creminelli:2005hu,Spergel:2006hy,Chen:2006ew}, while constraints that are somewhat weaker but that are closer in physical scale to that of clusters are expected from higher-order galaxy correlations \\citep{Scoccimarro:2003wn}. From a theoretical point of view, the non-Gaussian model of Eq.~(\\ref{fNLmodel}) is motivated in part by studies of the generation of density perturbations in inflationary scenarios; while single-field inflation models typically predict an unobservably small value for $\\fNL$ \\citep[e.g.][]{Acquaviva:2002ud,Maldacena:2002vr}, multi-field inflation models can lead to much higher values \\citep[e.g.][]{Lyth:2002my,Dvali:2003em,Zaldarriaga:2003my, Creminelli:2003iq,Arkani-Hamed:2003uz,Alishahiha:2004eh,Kolb:2005ux,Sasaki:2006kq}. For a review, see \\citet{Bartolo:2004if}.  While we believe it is worthwhile to keep an open mind to other forms of non-Gaussianity which may not be properly described by the simple expression in Eq.~(\\ref{fNLmodel}), and which might make the extrapolation of current CMB constraints to cluster scales less straightforward than we assume here \\citep[see, e.g.,][]{Mathis:2004vi}, we note that the physical scale probed by clusters differs from that of the Planck survey by roughly a factor of two, so that the two probes are likely to be affected more or less equally by deviations from Eq.~(\\ref{fNLmodel}).  As we discuss below in greater detail, the parameters in our likelihood analysis include $\\fNL$, the matter density $\\Omega_m$ and the matter fluctuation amplitude $\\sigma_8$, while we consider both a constant and time-varying dark energy equations of state described in terms of one ($w$) and two \\citep[$w_0$ and $w_a$,][]{Chevallier:2000qy,Linder:2002et} parameters respectively.  For definiteness, we assume a fiducial ideal survey similar in size and depth to that of the upcoming South Pole Telescope/Dark Energy Survey \\citep[SPT-DES,][]{Ruhl:2004kv,Abbott:2005bi}. We assume a $\\Lambda$CDM fiducial cosmology, for two values of $\\sigma_8$, since cluster number counts are extremely sensitive to this parameter.  This paper is organized as follows. In section~\\ref{secModel} we introduce our model for the non-Gaussian mass function and describe our analysis of the dependence of the expected errors on cosmological parameters on the non-Gaussian component. In section~\\ref{secResults} we present our results and we conclude in section~\\ref{secConclusions}. ",
        "conclusions": "\\label{secConclusions}  The success of the $\\Lambda$CDM standard cosmological model in recent years has been nothing short of spectacular.  Upcoming surveys will either continue to confirm this model and constrain its parameters with unprecedented accuracy, or they will uncover discrepancies which will point the way toward improvements in our understanding of fundamental physics.  Two questions addressing cosmology beyond the standard model that have been the subject of substantial attention in recent years are: what is the nature of the dark energy which is driving the accelerated expansion of the universe? and second, are fluctuations in the primordial matter distribution Gaussian, and therefore consistent with the predictions of the simplest inflationary models? In this paper, we obtain a rough estimate of the success that one of the most promising cosmological probes, galaxy cluster counts, is likely to have in answering these fundamental questions.  We have assumed an ideal cluster survey with survey parameters expected for the upcoming SPT and DES projects. Our fiducial cosmological model includes both dark energy and primordial non-Gaussianity using popular parameterized models, $w$ and $w_a$ for the former, and $f_{NL}$ for the latter. Cluster number counts as a function of mass, redshift, and cosmology were estimated using a standard fit to simulations \\citep{Jenkins:2000bv}, which we adjusted to allow for mild non-Gaussian initial conditions, and all clusters above a threshold mass were considered to be \"found\" by our fiducial survey. We then performed a simple likelihood analysis on the cluster counts using priors from current WMAP and expected Planck and SDSS constraints on non-Gaussianity as well as approximate priors on the two other relevant cosmological parameters from other present and future data sets.  Our principal conclusion is that dark energy constraints are in all cases \\emph{not} substantially degraded by primordial non-Gaussianity when the model parameterized by the constant $\\fNL$ and current limits from CMB observations are assumed. This is true despite the fact that variations in $\\fNL$ close to current uncertainties induce differences in the mass function comparable in magnitude to variations of $10\\%$ in the dark energy parameter $w$. A stronger degeneracy is observed instead between $\\fNL$ and $\\sigma_8$; in this case, the expected errors on $\\sigma_8$ from future cluster surveys can be noticeably affected when non-Gaussianity is included in the analysis.  \\squeezetable \\begin{table}[t] \\caption{\\label{Table_2DE_nb10} Same as Table~\\ref{Table_2DE} but using 10 cluster mass bins.} \\begin{ruledtabular} \\begin{tabular}{l|cccc} prior:            & $\\dfNL=0$ & $\\dfNL=5$ & $\\dfNL=37$ & $\\dfNL=145$                          \\\\ \\hline\\hline \\multicolumn{5}{l}{No priors on $\\Omega_m$ and $\\sigma_8$.}                                    \\\\ \\hline $\\Delta w_0$      & $0.180$   & $0.180$~($0\\%$)   & $0.186$~($3\\%$)    & $0.197$~($9\\%$)       \\\\ $\\Delta w_a$      & $0.68$    & $0.68$~($0\\%$)    & $0.71$~($4\\%$)     & $0.78$~($15\\%$)       \\\\ $\\Delta \\Omega_m$ & $0.0141$  & $0.0141$~($0\\%$)  & $0.0151$~($7\\%$)   & $0.0170$~($21\\%$)     \\\\ $\\Delta \\sigma_8$ & $0.0078$  & $0.0079$~($1\\%$)  & $0.0113$~($45\\%$)  & $0.0158$~($100\\%$)    \\\\ $\\dfNL$           & -         & $5.0$             & $31$               & $51$                  \\\\ \\hline\\hline \\multicolumn{5}{l}{Gaussian priors: $\\Omega_m = 0.27 \\pm 0.035$, $\\sigma_8 = 0.9 \\pm 0.05$}    \\\\ \\hline $\\Delta w_0$      & $0.167$   & $0.168$~($1\\%$)   & $0.170$~($2\\%$)    & $0.175$~($5\\%$)       \\\\ $\\Delta w_a$      & $0.64$    & $0.64$~($0\\%$)    & $0.66$~($3\\%$)     & $0.70$~($9\\%$)        \\\\ $\\Delta \\Omega_m$ & $0.0129$  & $0.0130$~($1\\%$)  & $0.0136$~($5\\%$)   & $0.0146$~($13\\%$)     \\\\ $\\Delta \\sigma_8$ & $0.0072$  & $0.0073$~($1\\%$)  & $0.0103$~($43\\%$)  & $0.0138$~($92\\%$)     \\\\ $\\dfNL$           & -         & $5.0$             & $30$               & $48$                  \\\\ \\hline\\hline \\multicolumn{5}{l}{Gaussian priors: $\\Omega_m = 0.27 \\pm 0.0035$, $\\sigma_8 = 0.9 \\pm 0.01$}   \\\\ \\hline $\\Delta w_0$      & $0.081$   & $0.081$~($0\\%$)   & $0.83$~($2\\%$)     & $0.84$~($4\\%$)        \\\\ $\\Delta w_a$      & $0.42$    & $0.42$~($0\\%$)    & $0.42$~($0\\%$)     & $0.42$~($0\\%$)        \\\\ $\\Delta \\Omega_m$ & $0.0033$  & $0.0033$~($0\\%$)  & $0.0033$~($0\\%$)   & $0.033$~($0\\%$)       \\\\ $\\Delta \\sigma_8$ & $0.0021$  & $0.0023$~($10\\%$) & $0.0049$~($130\\%$) & $0.0063$~($200\\%$)    \\\\ $\\dfNL$           & -         & $5.0$             & $26$               & $35$                  \\\\ \\hline\\hline \\multicolumn{5}{l}{Fixed $\\Omega_m=0.27$ and $\\sigma_8=0.9$}                                   \\\\ \\hline $\\Delta w_0$      & $0.067$   & $0.068$~($1\\%$)   & $0.070$~($4\\%$)   & $0.070$~($4\\%$)        \\\\ $\\Delta w_a$      & $0.37$    & $0.038$~($3\\%$)   & $0.40$~($8\\%$)    & $0.40$~($8\\%$)         \\\\ $\\dfNL$           & -         & $3.8$             & $5.9$             & $6.0$                  \\\\ \\end{tabular} \\end{ruledtabular} \\end{table}  A secondary conclusion is that the cluster survey itself might have sufficient statistical power to provide a valuable cross check on any detection or non-detection of primordial non-Gaussianity in CMB experiments, particularly when information on the cluster distribution as a function of the mass is taken into account.  However, we must emphasize that we \\emph{have not} attempted to include in our analysis any of the systematic and statistical errors in the clusters mass determination, which are likely to cause trouble for real surveys, as well as uncertainties on the predictions for the mass function, and our results must be interpreted with this in mind. We believe that our principal result should be quite robust, since any significant increase in the error budget will reduce constraining power on dark energy parameters and de-emphasize the confusion caused by any non-Gaussian initial conditions. On the other hand, the effectiveness of clusters as a cross check of primordial non-Gaussianity estimates from the CMB could be dramatically worsened, and should therefore be the subject of future work."
    },
    "0609/astro-ph0609254_arXiv.txt": {
        "abstract": "We study the long-term evolution of relativistic jets in collapsars and examine the effects of viewing angle on the subsequent gamma ray bursts.  We carry out a series of high-resolution simulations of a jet propagating through a stellar envelope in 2D cylindrical coordinates using the FLASH relativistic hydrodynamics module.  For the first time, simulations are carried out using an adaptive mesh that allows for a large dynamic range inside the star while still being efficient enough to follow the evolution of the jet long after it breaks out from the star.  Our simulations allow us to single out three phases in the jet evolution: a precursor phase in which relativistic material turbulently shed from the head of the jet first emerges from the star, a shocked jet phase where a fully shocked jet of material is emerging, and an unshocked jet phase where the jet consists of a free-streaming, unshocked core surrounded by a thin boundary layer of shocked jet material.  The appearance of these phases will be different to observers at different angles.  The precursor has a wide opening angle and would be visible far off axis. The shocked phase has a relatively narrow opening angle that is constant in time.  During the unshocked jet phase the opening angle increases logarithmically with time. As a consequence, some observers see prolonged dead times of emission even for constant properties of the jet injected in the stellar core. We also present an analytic model that is able to reproduce the overall properties of the jet and its evolution.  We finally discuss the observational implications of our results, emphasizing the possible ways to test progenitor models through the effects of jet propagation in the star.  In an appendix, we present 1D and 2D tests of the FLASH relativistic hydrodynamics module. ",
        "introduction": "The association of gamma-ray bursts (GRBs) with supernova explosions creates an apparent contradiction. On the one hand, the observed gamma-ray spectra require the emitting material to be outflowing at highly relativistic speed (Lithwick \\& Sari 2001), limiting its rest mass to a fraction of a solar mass. On the other hand, the jet expands through a dense stellar core, potentially picking up several solar masses in its way. Pioneering work by Woosley and collaborators (MacFadyen \\& Woosley 1999; MacFadyen et al. 2001) and by Aloy and collaborators (Aloy et al. 2000) showed that a light jet can reach the surface of a star virtually unpolluted if its head travels sub-relativistically inside the star. In this way, the cold and dense stellar material can be pushed aside. Once the stellar surface is reached, the light jet accelerates and reaches the high Lorentz factors required by $\\gamma$-ray observations (see also Matzner 2003 for a simplified analytic treatment).  Once this basic question is answered, two main issues remain open. First, what is the jet launching mechanism at the base of the star and what are the properties of the injected outflow? Second, how does the propagation of the jet influence its properties in the radiative phase, which takes place when the jet has traveled far from its birth site? The first question has been addressed with MHD numerical simulations (Proga et al. 2003; Mizuno et al. 2004; De Villiers et al. 2006; Nagataki et al. 2006). Due to the complexity of dealing with 3D magnetic phenomena in the general relativistic regime, these studies are still in their infancy. This paper deals with the second of these questions. Zhang et al. (2003) performed fixed grid 2D special relativistic simulations of jets in the cores of massive stars. They observed a diverse phenomenology of time dependent properties in the emerging jets, including variations in the opening angle, but their resolution was not large enough to draw robust conclusions at large radii, after the jet breaks out of the star. Lazzati \\& Begelman (2005, hereafter LB05) showed, with an analytic model, that the star-jet interaction can be strong and can create observable consequences. They argued that even if a steady jet is injected in the stellar core, its properties at the surface of the star will be strongly time-dependent, so that even the concept of a jet opening angle becomes hard to define in a general sense. They concluded that the time-integrated jet structure would match that of the ``Universal Jet'' proposed by Rossi et al. (2002). Even though - as we will show in this paper - their approximations were inaccurate in places, and the jet structure they predicted is not observed in the simulations, we confirm the basic concept that the propagation of the jet through the star will strongly influence both, generating strong temporal evolution in their properties. The influence of the jet's initial conditions has been studied recently with a set of 2D simulations (Mizuta et al. 2006). Their jets, however, are cylindrical at the base, different from any previous work and from our simulations presented here, where the outflow is set up to be conical in the core of the star.  In this paper, we present the results of a set of high-resolution 2-dimensional simulations of a jet propagating through a stellar progenitor. The jet is assumed to be already developed and mildly relativistic at $10^9$~cm inside the star, even though it is still hot and entropy-rich. The energy release is constant and the engine is left to run for a time longer than the breakout time, i.e., the time at which the jet breaks through the stellar surface. We show that the jet-star interaction creates three well-defined phases. The first one is characterized by turbulence. During this phase the jet propagates through and eventually breaks out of the star. The second phase is a transition phase, during which a heavily shocked jet flows through the opened channel. Finally, a more stable configuration is attained, with a freely flowing jet surrounded by a thin boundary layer that progressively widens its opening angle. Some of these phases can be dealt with analytically or semi-analytically with sufficient accuracy, allowing us to understand the origin of the observed behavior.  The observational implications of such non-steady flows are numerous, including long dead times during the prompt emission and bumpy, complex afterglows that do not obey the rules of the simple external shock model.  This paper is organized as follows. In \\S~2 we describe the numerical code used. In \\S~3 we describe the initial conditions for both the stellar progenitor and the jet in its center while in \\S~4 we detail our results. In \\S~5 we develop an analytic modeling of the jet and its propagation, in \\S~6 we present afterglow calculations based on the numerical results of \\S~4. We finally summarize our results in \\S~7. ",
        "conclusions": "We have presented high resolution 2D simulations of the propagation of light relativistic jets inside the cores of massive stars. We use an adaptive mesh code (FLASH) that allows us to study the behavior of the jet-star interaction over a long timescale and a wide spatial range. Different stellar progenitors as well as initial conditions of the jet are explored. Thanks to the high resolution and large spatial and temporal domain of our simulations we can confirm and study in more detail jet features discussed in previous work (MacFadyen \\& Woosley 1999; Aloy 2000; MacFadyen et al. 2001; Zhang et al. 2003, 2004; Mizuta et al. 2006), as well as identify new phenomenologies.  The main conclusion of this paper is that, even if the central engine is stationary, the jet that propagates out of the progenitor star is characterized by three phases, all of which display significant variability. The first phase is a wide angle release of mildly to moderately relativistic material, which we call the precursor. This phase is due to the release of the turbulent shocked material that accumulates around the jet due to vortex shedding before it breaks out of the star (Ramirez-Ruiz et al. 2002; Zhang et al. 2003). This initial phase is followed by a shocked jet phase. In this second phase the jet material that flows out of the star has been heavily shocked. The energy flow is highly variable in this phase, and no temporal trend in the properties can be identified. During this phase the jet is most highly collimated. Finally, as the pressure of the cocoon decreases, the jet settles into a stable configuration with a freely expanding core surrounded by a shocked shear layer at the boundary with the cocoon material. In this phase the energy flow is almost constant, and the jet opening angle increases logarithmically with time.  This temporal evolution of the jet is also associated with the angular distribution of energy, and therefore determines what different observers see from different directions. The precursor is characterized by a wide opening angle and can be seen from most directions. The shocked jet phase is the most concentrated and can be seen only by observers within several degrees of the axis. As a consequence, such observers will see a very bright event. Observers who lie a few degrees outside the shocked jet cone will see the precursor, followed by a dead time of several tens of seconds. Eventually, when the jet opens to contain their line of sight, they will see a second phase of emission. These observers will therefore measure dead-times much longer than any timescale of the inner engine, as found in several BATSE light curves (Lazzati 2005).  The overall angular distribution of energy is complex and does not seem to follow any simple correlation with the jet or progenitor properties. In some cases, the jet is characterized by a flat core with a sudden cutoff, very similar to a top-hat jet (16TIg5), while in other cases the distribution is more centrally peaked. It seems, however, that our high resolution simulations were not able to reproduce previous analytic or numerical results like the $\\theta^{-2}$ universal jet of LB05 or the $\\theta^{-3}$ distribution found by Zhang et al. (2003). Such behaviors are observed only in a limited range of angles from the jet axis. We computed afterglow light curves from our energy distributions showing that differences in the afterglow phase can result from the different properties of the progenitor and/or of the jet in the core of the star. We also developed further the analytic model of LB05, refining some approximations to obtain a model that can reproduce such basic features of the simulations as the propagation of the bow shock inside the star and the precursor energy, as well as the qualitative evolution of the jet opening angle.  It is easy to find paths along which this work can be developed further. Higher resolution and dimensionality will certainly be worth exploring (Zhang et al. 2004), as well as the consequences of a more accurate equation of state capable of describing both relativistic and non-relativistic material. Including a non-relativistic wind from the accretion disk is also worthwhile, as this seems to be a ubiquitous outcome of jet-launching simulations (Proga et al. 2003; de Villiers et al. 2006). This wind component may alter the cocoon properties, which are so important in defining the behavior of the precursor and shocked jet phases. Simulations also suggest that the jet luminosity should be highly variable if not intermittent. Preliminary studies (Aloy et al. 2000) suggest that this may enhance the propagation of the jet through the star, but higher resolution simulations are required. An interesting hypothesis is that the three phases of jet propagation will respond differently to an intermittent engine, with the precursor virtually unaffected while the unshocked jet should retain most of the engine variability.   \\bigskip  The software used in this work was in part developed by the DOE-supported ASC/Alliance Center for Astrophysical Thermonuclear Flashes at the University of Chicago.  We thank Alex Heger for providing us with the tabulated properties of his stellar models, Andrew MacFadyen for useful discussions and advice on the testing of the relativistic FLASH code and Miguel Aloy for useful discussions. This work was supported by NSF grant AST-0307502, NASA Astrophysical Theory Grant NNG06GI06G, and Swift Guest Investigator Program NNX06AB69G."
    },
    "0609/astro-ph0609548_arXiv.txt": {
        "abstract": "We analyze very deep HST, VLT and Spitzer photometry of galaxies at $2<z<3.5$ in the Hubble Deep Field South.  The sample is selected from the deepest public K-band imaging currently available.  We show that the rest-frame $U-V$ vs $V-J$ color-color diagram is a powerful diagnostic of the stellar populations of distant galaxies.  Galaxies with red rest-frame $U-V$ colors are generally red in rest-frame $V-J$ as well.  However, at a given $U-V$ color a range in $V-J$ colors exists, and we show that this allows us to distinguish young, dusty galaxies from old, passively evolving galaxies.  We quantify the effects of IRAC photometry on estimates of masses, ages, and the dust content of $z>2$ galaxies.  The estimated distributions of these properties do not change significantly when adding IRAC data to the UBVIJHK photometry.  However, for individual galaxies the addition of IRAC can improve the constraints on the stellar populations, especially for red galaxies: uncertainties in stellar mass decrease by a factor of 2.7 for red ($\\rfUV > 1$) galaxies, but only by a factor of 1.3 for blue ($\\rfUV < 1$) galaxies.  We find a similar color-dependence of the improvement for estimates of age and dust extinction.  In addition, the improvement from adding IRAC depends on the availability of full near-infrared JHK coverage; if only K-band were available, the mass uncertainties of blue galaxies would decrease by a more substantial factor 1.9.  Finally, we find that a trend of galaxy color with stellar mass is already present at $z>2$.  The most massive galaxies at high redshift have red rest-frame $U-V$ colors compared to lower mass galaxies even when allowing for complex star formation histories. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0609/astro-ph0609062_arXiv.txt": {
        "abstract": "% Current observational constraints on the dynamical evolution of star clusters are reviewed. Theory and observations now agree nicely on the mass dependency and time scales for disruption of young star clusters in galactic disks, but many problems still await resolution.  The origin of the mass function of old globular clusters, and its (near) invariance with respect to host galaxy properties and location within the host galaxy remain prominent puzzles. Most current models fail to reproduce the globular cluster mass function as a result of dynamical evolution from an initial power-law, except under very specific conditions which are not generally consistent with observations. How well do we actually know the proper initial conditions? The cluster initial mass function (CIMF) seems to be consistent with a power-law with exponent $\\alpha\\approx-2$ in most present-day star forming galaxies, but the limits of the mass range over which this approximation is valid remain poorly constrained both observationally and theoretically. Furthermore, there are hints that some dwarf galaxies may have CIMFs which deviate from a power-law. ",
        "introduction": "The idea that star clusters evolve dynamically has been around for nearly as long as the very concept of star clusters itself. The first to discuss the general properties of clusters in some detail was William Herschel, who attributed their various degrees of ``central concentration'' to an evolutionary sequence in which the globular clusters represented the final stage \\citep{wh1789}. At that time, however, a real theory of stellar dynamics was still far in the future and the subject of cluster evolution remained dormant for well over a century.  In the early 20th century, the attention was mostly focussed on external perturbations.  \\cite{jeans1916} remarked that ``A cluster of stars will suffer continual disintegration from its encounters with other stars or clusters''.  In his monograph on ``Star Clusters'', \\citet[][p.\\ 208]{shap1930} devoted only a brief statement to dynamical evolution, noting that ``some of the globular clusters may be affected in freedom, form, and eventual survival by contacts with galaxies or other clusters.'' The stability of ``moving clusters'' was investigated by Jeans and later by \\citet{bok1934}.  The important effect of internal two-body encounters was first studied by \\citet{amba1938} who noted that stars in the high-velocity tail of the Maxwellian velocity distribution will gradually escape, leading to disruption of typical open clusters on time scales of a few Gyrs, depending on cluster mass. With remarkable foresight, Ambartsumian also realised that low-mass stars will escape preferentially, and suggested that more evolved clusters might be recognisable by their lack of low-mass stars. Similar conclusions were reached (in the context of old globular clusters) by \\citet{spitzer1940}.  These ideas have been continuously refined ever since and it would lead too far to review the relevant literature on models for the dynamical evolution of star clusters here. Excellent reviews on the dynamical evolution of open and globular clusters are in \\cite{king1980}, \\cite{spitzer1987} and \\cite{mh1997} and the field remains very active as illustrated by several contributions in this volume.  The remainder of this review concentrates on observational constraints on cluster evolution.  In order to address dynamical evolution in a meaningful way observationally, a method to age-date individual star clusters becomes a necessity.  The catalogue of open clusters compiled by \\citet{trump1930} listed the spectral types of the brightest stars in 100 Milky Way open clusters, but the use of this information as a suitable chronometer had to await the emergence of the first detailed sets of stellar model calculations in the 1950s.  Once it became possible to derive cluster ages from colour-magnitude diagrams, it quickly became clear that the number of old open clusters is lower than one would naively predict from a continuous formation rate, even taken into account the fact that younger clusters are more easily detected than older ones because they contain brighter stars.  The relative paucity of old open clusters was noted almost simultaneously in at least three papers \\citep{vdb1957,oort1958,vonh1958}, all of which based on Trumpler's catalogue.  It appears that they all arrived at the same conclusion quite independently (van den Bergh, priv.\\ comm.).  All three papers also singled out disruption as the most likely cause of the decline in number of clusters with ages greater than $\\sim10^8-10^9$ years.  Considering the effects of stellar mass loss, a tidal field and relaxation due to stellar encounters, \\citet{vonh1958} estimated an average cluster lifetime of $10^9$ years, in reasonable agreement with his best observational estimate of 500 Myrs.  \\citet{spitzer1958} calculated cluster disruption times of $10^8-10^9$ years for typical open cluster densities, in fair agreement with the observational estimates, and also noted that encounters between clusters and interstellar clouds might play an important role for cluster disruption. Thus, one might argue that theoretical and observational estimates of cluster disruption times were in agreement to within a factor of 2 already 50 years ago! In the context of this meeting, it is also interesting to note that the connection between stellar mass loss and the evolution of star clusters was recognised early on. Of course, many of the details still remained (and remain!) to be worked out.  Using more recent compilations of open cluster data, \\citet{wielen1971} estimated that about 50\\% of open clusters in the Milky Way disintegrate within about 200 Myrs. He also noted a significant scatter in the disruption times, with 2\\% of the clusters surviving for longer than a Gyr, and suggested that the disruption time probably depends on the size and mass of the cluster (as anticipated by Ambartsumian).  Observational support for the importance of giant molecular cloud (GMC) encounters was provided by \\citet{vm1980} who pointed out that the oldest open clusters ($>1$ Gyr) are strongly concentrated towards the galactic anti-centre direction, where the density of GMCs is lower. This observation, as well as the fact that old open clusters tend to be found further from the Galactic plane than young ones, has since been confirmed by several studies \\citep[][and references therein]{friel1995}. Theoretically, further support for the importance of GMC encounters came from N-body simulations by \\citet{terle1987}.  The first observational evidence that cluster disruption times depend on mass came from \\citet{ja1982}.  However, it has also been clear for some time that the disruption times may not be the same in all galaxies.  Both the LMC and SMC show a less pronounced lack of old clusters than the Milky Way \\citep{ef1985,hodge1987}, and already \\citet{wielen1985} suggested that different GMC densities in the Milky Way, SMC and LMC might play a role. However, a direct comparison with the Milky Way remained somewhat hampered by different detection limits and the bursty star formation history of the LMC. Prior to the launch of the \\emph{Hubble Space Telescope} in 1990, little was known about the detailed properties of (young) cluster populations in other galaxies, and even in the Local Group spirals M33 and M31 the data were insufficient to strongly constrain cluster lifetimes.  Little has been said explicitly thus far about the \\emph{globular} clusters (GCs), which present a number of interesting problems.  In the Milky Way, they are associated with the spheroidal component(s) and thus considered typical ``Population II'' objects \\citep{baade1944}. In fact, every major galaxy is surrounded by ancient GCs akin to those found in the Milky Way.  The mass distribution of GCs is markedly different from that of Milky Way open clusters, with a relative deficit of low-mass GCs \\citep[\\S\\ref{sec:gclf},][]{vl1984}. This difference has sometimes lead to the notion that GCs may be fundamentally different from star clusters forming today, and might have formed by quite different processes which were unique in the early Universe \\citep{pd1968}. However, the young ``populous'' star clusters in the LMC, some of which have GC-like masses, have remained a puzzle since the early 20th century \\citep{shap1930}, and continue to serve as a reminder that the distinction between open and globular clusters is not necessarily as clear-cut in other galaxies as it may seem in our own.  The idea that GC formation may not have been restricted to the early Universe received renewed interest when \\citet{schweizer1987} suggested that new GCs form in galaxy mergers. This might solve the ``specific frequency problem'' (i.e., the fact that elliptical galaxies contain many more GCs per unit host galaxy luminosity than do spirals), which had been raised as a major objection against the idea that elliptical galaxies form from spiral-spiral mergers \\citep{vdb1982} as suggested by \\citet{tt1972}. The subsequent discovery by HST observations of young ``super'' star clusters in a number of interacting and merging galaxies in the 1990s \\citep[as reviewed by][]{whitmore2003} further stimulated interest in this idea, but the problem remains that the ``cluster initial mass function'' (the CIMF) seems very different from that of the ancient GCs. Besides other objections to the major merger picture \\citep*[e.g.][]{fbg1997} this still represents a major challenge, as I will discuss shortly. ",
        "conclusions": "For several decades it has been clear that a star cluster, even if left in isolation, will gradually disrupt as the velocity distribution approaches a Maxwellian form and stars in the high-velocity tail escape. This process can be greatly accelerated in the presence of external perturbations such as the tidal field of the host galaxy, or tidal shocks from encounters with bulges, disks, spiral arms and/or GMCs. In the context of this meeting, it is also worth emphasising that stellar mass loss may play an important part for the dynamical evolution of star clusters, especially in the earliest stages.  There is now excellent agreement between the observed mass- and age distributions of open clusters in the Milky Way and theoretical predictions, assuming evolution from an initial power-law mass function. It also appears that we are well under way to understanding differences in disruption time-scales in different environments. However, explaining the mass distribution of ancient globular clusters remains an unsolved problem, and this is probably where most work lies in the future. Once again, the effect of stellar mass loss may turn out to be important, possibly in combination with a more realistic treatment of the structural properties of star clusters. An additional complication is that ancient GCs are subject to destruction mechanisms which are unimportant for star clusters forming today and evolving in disk galaxies (bulge, disk shocks), while in contrast the role of other mechanisms which may now be important (spiral arm passages, GMC encounters) is difficult to assess for GCs."
    }
}